{
    "0311/astro-ph0311385_arXiv.txt": {
        "abstract": "Magnetic fields in galaxy clusters can be investigated using a variety of techniques. Recent studies including radio halos, Inverse Compton hard X-ray emissions and Faraday rotation measure, are briefly outlined. A numerical approach for investigating cluster magnetic fields strength and structure is presented. It consists of producing simulated rotation measure, radio halo images, and radio halo polarization, obtained from 3-dimensional multi-scale cluster magnetic field models, and comparing with observations. ",
        "introduction": "A complete description of the astrophysical processes in cluster of galaxies requires knowledge of magnetic fields. The most detailed evidence for this component comes from the radio observations:\\\\ $\\bullet$ Some clusters of galaxies exhibit diffuse non-thermal synchrotron radio halos, associated with the intra-cluster medium, which extend up to Mega-parsec scales. Using minimum energy assumptions, it is possible to estimate an equipartition magnetic field strength averaged over the entire halo volume. These estimates give equipartition magnetic field strengths of $\\simeq$0.1 to 1 $\\mu$G (e.g. Bacchi et al. 2003).\\\\ $\\bullet$ In a few cases, clusters containing a radio halo show an hard X-ray excess emission. This emission could be interpreted in terms of Inverse Compton scattering of the cosmic microwave background photons with the relativistic electrons responsible for the radio halo emission. In this case, the measurements of the magnetic field strength (e.g. Fusco-Femiano et al. 1999, Rephaeli et al. 1999) inferred from the ratio of the radio to X-ray luminosities are consistent with the equipartition estimates.\\\\ $\\bullet$ Indirect measurements of the magnetic field intensity can also be determined in conjunction with X-ray observations of the hot gas, through the study of the Faraday Rotation Measure (RM) of radio sources located inside or behind clusters. By using a simple analytical approach, magnetic fields of $\\sim$$5-30$ $\\mu$G, have been found in cooling flow clusters (e.g. Allen et al. 2001; Taylor \\& Perley 1993) where extremely high RMs have been revealed. On the other hand, significant magnetic fields have also been detected in clusters without cooling flows: the RM measurements of polarized radio sources through the hot intra-cluster medium leads to a magnetic field of $2-8$ $\\mu$G which fluctuate on scales as small as $2-15$ kpc. (e.g. Feretti et al. 1995, Feretti et al. 1999, Clarke et al. 2001, Govoni et al. 2001, Taylor et al. 2001, Eilek \\& Owen 2002).  The magnetic field strength obtained by RM studies is therefore higher than the value derived from the radio halo data and from Inverse-Compton X-ray studies. However, as pointed out by Carilli \\& Taylor (2002) and references therein, all the aforementioned techniques are based on several assumptions. For example, the observed RMs have been interpreted, until now, in terms of simple analytical models which consider single-scale magnetic fields. On the other hand, magneto-hydrodynamic cosmological simulations (Dolag et al. 2002) suggest that cluster magnetic fields may span a wide range of spatial scales with a strength that decreases with distance from the cluster center.  We developed a numerical approach for investigating the strength and structure of cluster magnetic fields. It consists of comparing simulated rotation measure, radio halo images, and radio halo polarization, obtained from 3-dimensional multi-scale cluster magnetic field models, with observations (Murgia et al. 2003, submitted). ",
        "conclusions": "The numerical approach presented here demonstrates how the dispersion and mean of the RM measured in radio galaxies embedded in a cluster of galaxies can be used to constrain not only the strength but also the power spectrum slope of the intra-cluster magnetic fields. Moreover, the study of the polarization properties of a large scale radio halo, if it is present in a cluster, can be used to improve the estimates based on the RM analysis.\\\\  \\noindent {\\it Acknowledgment.}~ We thank L.Feretti, G.Giovannini, D.Dallacasa, R.Fanti, G.B. Taylor and K.Dolag for their collaboration."
    },
    "0311/astro-ph0311499_arXiv.txt": {
        "abstract": "We present a general approach to the modelling of the brightness and polarization structures of adiabatic, decelerating relativistic jets, based on the formalism of \\citet{MS}.  We compare the predictions of adiabatic jet models with deep, high-resolution observations of the radio jets in the FR\\,I radio galaxy 3C\\,31.  Adiabatic models require coupling between the variations of velocity, magnetic field and particle density. They are therefore more tightly constrained than the models previously presented for 3C\\,31 by \\citet{LB02a}. We show that adiabatic models provide a poorer description of the data in two crucial respects: they cannot reproduce the observed magnetic-field structures in detail, and they also predict too steep a brightness decline along the jets for plausible variations of the jet velocity.  We find that the innermost regions of the jets show the strongest evidence for non-adiabatic behaviour, and that the adiabatic models provide progressively better descriptions of the jet emission at larger distances from the galactic nucleus.  We briefly discuss physical processes which might contribute to this non-adiabatic behaviour. In particular, we develop a parameterized description of distributed particle injection, which we fit to the observed total intensities.  We show that particles are preferentially injected where bright X-ray emission is observed, and where we infer that the jets are over-pressured. ",
        "introduction": "\\label{Introduction}  This paper is one of a series whose aim is to develop quantitative models of jets in low-luminosity (FR\\,I; \\citealt{FR74}) radio galaxies on the assumption that they are decelerating relativistic flows. Our hypothesis is that the jets are close enough to being intrinsically symmetrical, axisymmetric and antiparallel that the observed differences between them are dominated by the effects of relativistic aberration.  This hypothesis is motivated by the results of \\citet{LPdRF}, who presented a statistical study of jets in the B2 sample of radio galaxies and showed that the observed correlations between fractional core flux density and side-to-side asymmetries in intensity and width are consistent with jet deceleration and the presence of transverse velocity gradients. They inferred that the jets slow from $\\approx$0.9$c$ where they start to expand rapidly to $\\approx$0.1$c$ over distances $\\sim$1 to 10\\,kpc and that the deceleration scale is an increasing function of jet power.  In \\citet[][hereafter LB]{LB02a}, we demonstrated that an intrinsically symmetrical, relativistic jet model provides an excellent description of the total intensity and linear polarization observed at 8.4\\,GHz from the jets of the FR\\,I radio galaxy 3C\\,31.  We were able to estimate the angle to the line of sight and the three-dimensional distributions of velocity, emissivity and magnetic-field structure.  By combining this kinematic model with a description of the external gas density and pressure derived from {\\em Chandra} observations \\citep{Hard02} and using conservation of particles, energy and momentum, we demonstrated that the jet deceleration could be produced by entrainment of thermal matter and we derived the spatial variations of pressure, density and entrainment rate \\citep{LB02b}.  The {\\em free models} developed by LB for 3C\\,31 were designed to fit the observed images without embodying specific preconceptions about the (poorly known) internal physics. In those models, we adopted simple and arbitrary functional forms for the spatial variations of velocity, synchrotron emissivity and field ordering and allowed the emissivity and field ordering to vary separately as smooth functions of position.  To proceed beyond such purely empirical descriptions of the jets, we must make further assumptions about processes that affect the components of the modelled emissivity -- the relativistic particle energy spectrum and the strength of the magnetic field. The separation of the emissivity into its components is ill-determined unless inverse-Compton emission can be detected from the synchrotron-emitting electrons, in which case the particle density and field strength can be determined independently. Unfortunately, inverse-Compton emission from the inner jets in 3C\\,31 is too weak to detect \\citep{Hard02}.  There are, as yet, no prescriptive theories for dissipative processes such as particle acceleration and field amplification or reconnection in conditions appropriate to FR\\,I radio jets.  The energy-loss processes for the radiating particles can be quantified, however.  It is inevitable that the particles will suffer adiabatic losses as the jets expand.  We argue that synchrotron and inverse-Compton losses are negligible by comparison for electrons radiating at 8.4\\,GHz in 3C\\,31, as the jets have accurately power-law spectra with indices close to 0.5 \\citep{LPBFGMP04} and show no evidence of spectral curvature until much higher frequencies \\citep{Hard02}.  It is therefore worthwhile to compare the observations with models in which the radiating particles are accelerated before entering the region of interest and then lose energy only by the adiabatic mechanism while the magnetic field is frozen into and convected passively with the flow, which is assumed to be laminar.  Following conventional usage, we refer to such models as {\\em adiabatic}.  There are analytical solutions for the adiabatic evolution of emissivity and surface brightness with jet radius and velocity if the magnetic field is exactly parallel or perpendicular to the flow direction.  Models of this kind were first considered by \\citet{Bur79}, who showed that the decrease of brightness with distance from the nucleus would be much steeper than he observed in 3C\\,31 if the jets have constant velocity. \\citet{Fan82} pointed out that the brightness would decline more slowly with distance in a decelerating jet. This idea underlies the turbulent jet models of \\citet{ Bic84,Bic86}.  A relativistic generalisation was developed by \\citet{Bau97} to model the jets in 3C\\,264 and was applied by \\citet {Fer99} and \\citet{ Bondi00} to other FR\\,I jets.  These treatments all assumed that there is no velocity gradient across the jet, and that the magnetic field is exactly parallel or perpendicular to the flow. While this approach is self-consistent, it cannot be applied if there is velocity shear in a direction perpendicular to any component of the field.  In LB, we developed a numerical approach to modelling a relativistic jet with a more general magnetic field structure.  This approach allowed us to use the variations of total intensity and linear polarization as independent constraints on the jet velocity field.  We concluded that the magnetic structure and velocity field in 3C\\,31 are indeed significantly more complex than those assumed by \\citet{Bau97}. Nevertheless, we compared our models and the data for 3C\\,31 with their analytical solutions.  We showed that:  \\begin{enumerate} \\item the adiabatic approximation is qualitatively inconsistent with the variations of brightness and polarization along the first $\\approx$3 kpc of the jets, but \\item further from the nucleus, the observed variations are closer to those expected if the adiabatic approximation holds. \\end{enumerate} This comparison motivated us to develop a more general approach to modelling of adiabatic, relativistic jets, which we present in this paper.  Our approach allows us to calculate how the brightness and polarization structure evolve along an adiabatic jet from {\\it prescribed initial conditions} (specified as profiles across the jet), given the jet geometry and more complex magnetic structures and velocity fields of the type inferred for 3C\\,31 by LB.  Detailed comparison of the new adiabatic models with the free models of LB, which fit the observations better, can then diagnose whether and where other physical processes, such as particle acceleration, may be significant in the jets.  Section~\\ref{ad-approx} reviews our assumptions and the previous analytical solutions, and then outlines our calculation of synchrotron emission from adiabatic flows. Section~\\ref{Obs-mod} describes our approach to modelling of adiabatic jets; it briefly recapitulates material from LB before discussing new aspects specific to the present study. Section~\\ref{Outer-region} applies our adiabatic models to the outer regions of the jets in 3C\\,31 and shows that they can give a fair description of the VLA observations of these regions. Section~\\ref{Whole-jet} confirms that the adiabatic models fail to describe the inner jet regions and critiques the adiabatic hypothesis in the light of this result; it also discusses the extent to which distributed particle injection can bring the adiabatic models into better agreement with the data. Section~\\ref{Conclusions} summarizes our conclusions. ",
        "conclusions": "\\label{Conclusions}  We have investigated the hypothesis that the radio jets in 3C\\,31 can be modelled as adiabatic, decelerating, relativistic flows. Our technique has several advantages over previous work in: modelling linear polarization as well as total intensity; including the effects of velocity shear; taking account of anisotropic emission in the rest frame and fitting to two-dimensional images rather than to longitudinal profiles.  We have found that optimized adiabatic models give a fair description of the observed brightness and polarization distributions in the outer parts of the modelled region.  Their fit to the data is inferior to that of the free models of LB, but is obtained with many fewer free parameters.  The adiabatic models cannot describe the inner or flaring regions of the jets in 3C\\,31, however: the predicted distributions of total intensity and linear polarization are inconsistent with those observed.  In the innermost region, the jets are clearly non-adiabatic unless the emission comes primarily from a part of the jet volume which is not expanding with distance from the nucleus (we do not resolve the region transversely). In the flaring region, a much higher initial velocity is required than is allowed by our measurements of jet/counter-jet sidedness ratio. The emission in this region is well resolved, and comes from the entire jet volume.  We have shown that a modified adiabatic model can still be fitted to the total intensity in this region if we add distributed injection of relativistic particles which then evolve adiabatically; the region where these particles must be injected is also one where there is independent evidence for recent particle acceleration from the detection of X-ray synchrotron radiation \\citep{Hard02} and of a local over-pressure in the jet from dynamical arguments \\citep{LB02b}.  While the total intensity distributions in the jet and counter-jet of 3C\\,31 can be reproduced satisfactorily by decelerating adiabatic jet models with particle injection in the flaring region, the polarization data (degree of linear polarization and apparent magnetic field direction) cannot.  The apparent magnetic field configuration contains several features that are qualitatively incompatible with adiabatic evolution of all the field components in laminar-flow models, even in the presence of a velocity shear.  We infer that the departures from adiabatic conditions in 3C\\,31 must also include deviations either from the flux-freezing description of the magnetic fields or from a laminar, axisymmetric velocity field.  The free models of LB achieved better consistency with the observed polarization structure by allowing the field to become roughly isotropic at the edges of the jet in the flaring region.  This could be achieved by a turbulent velocity component even if the field is frozen into the flow, in which case the resulting shear would also contribute to the enhancement in emissivity required to fit the adiabatic models in the flaring region. Processes such as dynamo action or field-line reconnection might also be significant.  More detailed analysis of the particle injection process inferred for the flaring region would benefit from more sensitive X-ray, optical and infra-red data and higher-resolution radio observations of this region, i.e.\\ from deeper exposures with {\\em Chandra} and {\\em HST} and from deep higher-resolution polarimetry with the EVLA.  Improved knowledge of the intensity and apparent magnetic field structures within this region might also assist development of models for the magnetic field microphysics in this region of the 3C\\,31 jet, which all of our analysis suggests is crucial in determining the deceleration dynamics of the jet, and its brightness and polarization properties further from the nucleus."
    },
    "0311/astro-ph0311516_arXiv.txt": {
        "abstract": "Numerical calculations and linear theory of radiation magnetohydrodynamic flows indicate that the photon bubble and magnetorotational instability (MRI) may produce large density inhomogeneities in radiation pressure supported media. We study the effects of the photon bubble instability on accretion disk spectra using  2-D Monte Carlo (MC) and 1-D Feautrier radiative transfer calculations on a snapshot of a 2-D numerical simulation domain. We find an enhancement in the thermalization of the MC spectra over that of the Feautrier calculation.  In the inner-most regions of these disks, the turbulent magnetic pressure may greatly exceed that of the gas.  It is then possible for bulk turbulent Alfv\\'enic motions driven by the MRI to exceed the thermal velocity making turbulent Comptonization the dominant radiative process. We estimate the spectral distortion due to turbulent Comptonization utilizing a 1-D MC calculation. ",
        "introduction": "The standard assumption central to most accretion disk models is that the turbulence responsible for angular momentum transport and gravitational energy release can be parameterized by an ``$\\alpha$-viscosity'' prescription (Shakura \\& Sunyaev 1973). The hypothesis that turbulence drives accretion has been borne out by the  identification of the magnetorotational instability (MRI, Balbus \\& Hawley 1991) as a generic source for magnetohydrodynamic turbulence in differentially rotating flows. However, it is not yet known how accurately $\\alpha$-disk models describe real accretion disks. It is possible that both the structure and emitted spectrum of a disk in which the MRI is active may deviate significantly from the predictions of ad hoc  $\\alpha$-disk models.  Numerical simulations of the radiation magnetohydrodynamic (RMHD) equations (Turner et al. 2003) show that MRI turbulence can produce a factor of 20 variation in density. In addition to the MRI, local linear analyses of the RMHD equations demonstrate that a radiatively driven photon bubble instability exists in radiation dominated accretion disks (Gammie 1998). This instability is present as long as the magnetic field is finite, but the growth rate for the instability is greatest when the magnetic pressure exceeds the gas pressure (Blaes \\& Socrates 2003). Numerical RMHD calculations of radiation pressure supported disks which neglect shear (Turner et al., in preparation) show that the non-linear development of the photon bubble instability does indeed produce large density variations.  In section 2 we discuss the effects of these density inhomogeneities on the emergent spectrum.  Fluid motions associated with the MRI may also influence the emission. MRI turbulence is fundamentally magnetic in nature with a characteristic velocity $v_w$ given by the Alfv\\'en speed $v_A$ which is less than the local sound speed $c_s$.  Close to the hole, where the radiation pressure is dominant, $c_s$ is given by the radiation sound speed which is larger than the sound speed of the gas.  Thus, it is possible for the MRI turbulence to be {\\it highly supersonic relative to the gas} and Compton scattering off of the turbulent eddies may produce a larger spectral modification than thermal Comptonization (Socrates, Davis, \\& Blaes 2003, hereafter SDB03).  In section 3, we summarize these results. ",
        "conclusions": "Large density fluctuations due to the photon bubbles instability may alter the spectrum emitted in radiation pressure dominated accretion disks. These results are based on calculations which neglect shear.  If shear is not neglected the MRI may affect the photon bubble instability, possibly altering the disk structure. However, for sufficiently large $v_A$, the MRI may be compressible in radiation dominated regions and may produce significant inhomogeneities on its own (Turner et al. 2003).  We also analyze the interactions of photons with turbulent eddies and show that for sufficiently large turbulent stresses, a prominent spectral component in the far UV and X-ray bands emerges. Observations of accreting systems may provide a means to directly probe the mechanism responsible for angular momentum transport in these sources."
    },
    "0311/astro-ph0311270_arXiv.txt": {
        "abstract": "We investigate self-similar dynamical processes in an isothermal self-gravitational fluid with spherical symmetry. In reference to earlier complementary solution results of Larson, Penston, Shu, Hunter and Whitworth \\& Summers, we further explore the `semi-complete solution space' from an initial instant $t\\rightarrow 0^{+}$ to a final stage $t\\rightarrow +\\infty$. These similarity solutions can describe and accommodate physical processes of radial inflow, core collapse, oscillations and envelope expansion (namely, outflow or wind) or contraction as well as shocks. In particular, we present new classes of self-similar solutions, referred to as `envelope expansion with core collapse' (EECC) solutions, that are featured by an interior core collapse and an exterior envelope expansion concurrently. The interior collapse towards the central core approaches a free-fall state as the radius $r\\rightarrow 0$, while the exterior envelope expansion gradually approaches a constant radial flow speed as $r\\rightarrow +\\infty$. There exists at least one spherical stagnation surface of zero flow speed that separates the core collapse and the envelope outflow and that travels outward at constant speed, either subsonically or supersonically, in a self-similar manner. Without crossing the sonic critical line where the travel speed of nonlinear disturbances relative to the radial flow is equal to the sound speed, there exist a continuous band of infinitely many EECC solutions with only one supersonic stagnation point as well as a continuous band of infinitely many similarity solutions for `envelope contraction with core collpase' (ECCC) without stagnation point. Crossing the sonic critical line twice analytically, there exist infinitely many discrete EECC solutions with one or more subsonic stagnation points. Such discrete EECC similarity solutions generally allow radial oscillations in the subsonic region between the central core collapse and the outer envelope expansion. In addition, we obtained complementary discrete ECCC similarity solutions that cross the sonic critical line twice with subsonic oscillations. In all these discrete solutions, subsonic spherical stagnation surfaces resulting from similarity oscillations travel outward at constant yet different speeds in a self-similar manner. With specified initial boundary or shock conditions, it is possible to construct an infinite number of such EECC similarity solutions, which are conceptually applicable to various astrophysical problems involving gravitational collapses and outflows. We mention potential applications of EECC similarity solutions to the formation process of proto-planetary nebulae connecting the AGB phase and the planetary nebula phase, to H{\\sevenrm II} clouds surrounding star formation regions, and to a certain evolution phase of galaxy clusters. ",
        "introduction": "Hydrodynamical processes of a self-gravitational spherical gas have been investigated from different perspectives and in various contexts of astrophysical or cosmological problems on entirely different temporal and spatial scales. As an important example, solutions to the model problem of spherical gravitational collapse provide physical insights for understanding processes of star formation and outflows (Ebert 1955; Bonner 1956; Hayashi 1966; Hunter 1967; Spitzer 1968; Larson 1973; Mestel 1974). Another important field is the formation and evolution of galaxy clusters that contain massive dark matter haloes and hot gases with temperatures of $\\sim 10^7-10^8$K (e.g. Fabian 1994; Sarazin 1988; Gunn \\& Gott 1972; Bertschinger 1985; Fillmore \\& Goldreich 1984; Navarro, Frenk \\& White 1996). Detailed numerical computations to solve the fully coupled nonlinear partial differential hydrodynamic equations show that solutions, sufficiently away from initial and boundary conditions, appear to evolve into self-similar behaviours (e.g. Sedov 1959; Landau \\& Lifshtiz 1959). Over past several decades, researchers have studied self-similar solutions under spherical symmetry and isothermal conditions. Simultaneously and independently, Penston (1969a, b) and Larson (1969a, b) found one discrete self-similar infall solution without involving a central core mass, the so-called LP-solution referred to in the literature, on the basis of their numerical calculations of spherically symmetric gravitational collapse for star formation. In addition to this LP-solution, Shu (1977) derived other types of self-similar solutions, including the so-called `expansion-wave collapse solution' (with a discontinuous derivative at the sonic critical point), which leads to the physical scenario of `inside-out collapse' in the context of molecular cloud and/or star formation (e.g. Shu, Adams \\& Lizano 1987). Shu (1977) showed clearly the existence of the sonic critical line across which two types of eigensolutions may be possibly specified.  Following Shu's analysis, Hunter (1977) found a class of discrete `complete solutions' by extending the time domain way back to the pre-catastrophic period starting from $t\\rightarrow -\\infty$ and by applying a zero-speed boundary condition there (see Fig. 6). One  may regard the LP-solution as a special member of Hunter's class and Shu's expansion-wave collapse solution turns out to be a limiting case of Hunter's solution class [see solution parameters (11) and figs. 1 and 2 of Hunter (1977)]. In terms of the `semi-complete solution' structure (see Fig. 6), the `complete' LP-solution and each of `complete' solutions $a$, $b$, $c$, $d$ and so forth of Hunter's (1977) can be split into two branches (i.e. one inflow and one outflow) at $x\\equiv r/(at)\\rightarrow +\\infty$. In the `semi-complete' perspective, Hunter's solutions $a$, $b$, $c$, $d$ and so forth involve radial oscillations with increasing number of nodes or stagnation points in the subsonic region with either envelope expansion or envelope contraction but without initial onset of central core collapse (see Fig. 6).  Within the same overall framework of Shu (1977) and Hunter (1977), Whitworth \\& Summers (1985, hereafter WS) introduced more elaborate mathematical techniques to construct similarity solutions across the sonic critical line with weak discontinuities (Lazarus 1981; Hunter 1986; Ori \\& Piran 1988; Boily \\& Lynden-Bell 1995). Whitworth \\& Summers noted the existence of two types of sonic points (namely, saddle and nodal points defined later) in the problem (see Jordan \\& Smith 1977 for details) and carefully examined stability properties of numerical integration directions. They realized the distinctly different properties of Shu's two types of eigensolutions at the sonic critical line and referred to the two types as `primary' and `secondary' directions\\footnote{For the two eigensolutions crossing the nodal sonic critical point and in terms of absolute values, the one of larger flow speed gradient is along the primary direction and the one of smaller flow speed gradient is along the secondary direction.} in the case of nodal points (Jordan \\& Smith 1977). By allowing one solution to pass the nodal sonic point along the secondary direction and an infinite number of solutions to pass it tangentially along the primary direction as well as locally linear combinations of the two eigensolutions, WS constructed two-parameter continuum bands of solutions in addition to the previous discrete solutions. Regarding the analysis of WS, Hunter (1986) promptly pointed out that the solutions of WS involve weak discontinuities across the sonic critical line in that physical variables and all their first derivatives in space and time are continuous\\footnote{This does not include those cases when a solution on one side is along the primary direction and that on the other side is along the secondary direction.} and derivatives of some higher orders may also be continuous. Moreover, Hunter (1986) suggested that the solutions of WS may be unstable against small perturbations as numerical computations did not show tendency towards any of such solutions except for the LP-solution and therefore might not be acceptable physically. Ori \\& Piran (1988) proceeded to verify Hunter's conjecture and derived a necessary stability criterion for `complete solutions'.  Lynden-Bell \\& Lemos (1988) and Lemos \\& Lynden-Bell (1989a, 1989b) studied self-similar solutions for a cold fluid, not only in the Newtonian regime but also in the regime of general relativity. Foster \\& Chevalier (1993) studied the gravitational collapse of an isothermal sphere by hydrodynamic simulations. They recovered the LP-solution in the central region where a core forms and the self-similar solutions of Shu (1977) when the ratio of initial outer cloud radius to core radius is $\\gsim 20$. Boily \\& Lynden-Bell (1995) studied the self-similar solutions for radial collapse and accretion of radiative gas in the polytropic approximation. Hanawa \\& Nakayama (1997) also investigated the stability problem of the complete self-similar solutions in a normal mode analysis and concluded that only the LP-solution is stable. The most recent work of Harada, Maeda \\& Semelin (2003) studied spherical collapse of an isothermal gas fluid through numerical simulations and showed a critical behaviour for the Newtonian collapse that turns out to follow the self-similar form of the first member (i.e. solution $a$) of Hunter's solutions (Hunter 1977).  Lai \\& Goldreich (2000) studied the growth of nonspherical perturbations in the collapse of a self-gravitating spherical gas cloud. They found through numerical computations that nonspherical perturbations damp in the subsonic region but grow in the supersonic region with asymptotic scaling relations for their growths. Lai \\& Goldreich mentioned potential applications to core-collapse of supernova explosions (Goldreich \\& Weber 1980; Yahil 1983), where the asymmetric density perturbation may lead to asymmetric shock propagation and breakout, giving rise to asymmetry in the explosion and a kick to the new-borne neutron star. Subsequently, Lai (2000) presented a global stability analysis for a self-similar gravitational collapse of a polytropic gas.  In addition to the shock-free self-similar solutions, Tsai \\& Hsu (1995) found a similarity shock solution in reference to Shu's expansion-wave collapse solution through both numerical and analytical analyses for the formation of low-mass stars. Their results were recently expanded by Shu, Lizano, Galli, Cant\\'o \\& Laughlin (2002) in the context of self-similar `champagne flows' driven by a shock in H{\\sevenrm II} regions. One shock solution found by Tsai \\& Hsu (1995) turns out to be the limiting case of many self-similar shocked LP-solutions found by Shu et al. (2002).  The perspectives of complete (Hunter 1977, 1986) and semi-complete (Shu 1977; WS) similarity solutions are both valid with proper physical interpretations and with the corresponding identification of an initial moment. In view of the invariance under the time-reversal transformation of self-gravitational fluid equations $(1)-(3)$ below, that is, $t\\rightarrow -t$, $\\rho\\rightarrow\\rho$, $u\\rightarrow -u$, the complete and semi-complete similarity solutions are closely related to each other. For example, the complete LP-solution (fig. 1 of Hunter 1977) describes a self-similar collapse process starting from the pre-catastrophic instant $t\\rightarrow-\\infty$ with a zero flow speed and a density distribution, passing through the sonic critical point once, gaining a finite radial infall speed at $t=0$, and approaching eventually a free-fall state involving a central core collapse as $t\\rightarrow +\\infty$. However, the complete LP-solution (Hunter 1977) can be broken into two branches of similarity solution in the semi-complete space. One branch is only partially shown in the lower-left corner of Shu's fig. 2 and is described by equations $(A7)$ and $(A8)$ in Shu's appendix (see also Fig. 5). This branch corresponds to a self-similar outflow without core mass but with a finite central density decreasing with time $t$ in a scaling of $t^{-2}$. The other branch corresponds a constant radial inflow speed at $x\\rightarrow +\\infty$ and a core collapse at $x\\rightarrow 0^{+}$ without crossing the sonic critical line. To some extent, Shu's criticism on the utility of the semi-complete LP-solution may be compromised by joining the other branch of the semi-complete LP-solution of infall-core-collapse solution (Hunter 1977) or by inserting a shock to match with an outward `breeze' for `champagne flows' in H{\\sevenrm II} regions (Tsai \\& Hsu 1995; Shu et al. 2002).  With the complementarity of complete and semi-complete perspectives for self-similar solutions, we re-visit this classic problem in the semi-complete solution space with the strong motivation of finding similarity solutions for concurrent processes of central core collapse and envelope expansion or contraction that may or may not involve subsonic radial oscillations and/or shocks across the sonic critical line. Specifically, we have learned Shu's expansion-wave collapse solution as the limit of a class of core collapse solutions with zero infall speed at $x\\rightarrow +\\infty$. We have also noted that Hunter's complete collapse solutions, including the LP-solution, can be used to describe radial outflows with or without subsonic radial oscillations in the semi-complete perspective. It is then physically plausible to derive self-similar solutions for concurrent envelope expansion with core collapse (EECC) that may or may not involve subsonic radial oscillations and/or shocks. If one allows solutions that pass through the sonic critical line with weak discontinuities (Lazarus 1981), then the mathematical solution space can be further expanded enormously (WS). This paper gives a positive account of constructing EECC similarity solutions that are distinctly different from known similarity solutions. With proper selections and adaptations, we believe that these EECC similarity solutions, due to their generality and simplicity, provide an important basic conceptual framework for a wide range of astrophysical problems involving gravitational collapses, outflows, contractions and subsonic radial oscillations etc. We shall mention a few potential applications, including the dynamical process of forming a proto-planetary nebula that evolves from the asymptotic giant branch (AGB) phase to the planetary nebula phase with a central hot proto white dwarf. Specifics will be described in forthcoming papers.   Several authors have approached the same set of basic isothermal fluid equations with different notations, choice of variables and conventions of presenting their results. It is sometimes mind-boggling to find correspondences among their solutions and results of analysis. Without further complications, we shall more or less adopt Shu's notations and present our results in a similar manner. The basic equations and the self-similar transformation are summarized in Section 2. In contrast to Shu's focus on the first quadrant of $-v$ versus $x$, we pay equal attention to both first and fourth quadrants for $-v$ versus $x$ presentation as we are interested in both outflows and collapses. In order to place our work in a proper perspective in reference to known results, we have explored the entire solution structure, including the counterparts of those solutions found by previous authors, in Section 3. Physical interpretations for the obtained solutions are described in Section 4. We summarize and discuss our results in Section 5. Some technical details and procedures can be found in the Appendix. ",
        "conclusions": "In this paper, we investigated the spherical self-similar solutions for self-gravitating isothermal flows in the `semi-complete solution space' extending from the initial instant $x\\rightarrow +\\infty$ to the final instant $x\\rightarrow 0^{+}$. The relevant similarity solutions are obtained and classified to compare with previous solutions and analyses [Larson 1969a; Penston 1969a; Shu 1977; Hunter 1977; Whitworth \\& Summers 1985 (WS)].  The novel class of infinitely many similarity solutions obtained in this paper, i.e., the EECC self-similar solutions and the envelope contraction with core collapse (ECCC) self-similar solutions (e.g. solution labelled 2 in Figs. 10 and 11 that passes the sonic critical line twice with specific parameters $V=-0.766$, $A=1.209$ and $m_0=5.171\\times 10^{-4}$), is constructed in the `semi-complete solution space'. This is to be compared with the similarity solutions derived in the `complete solution space' first introduced by Hunter (1977). By the invariance property under time reversal in this classical problem, these new self-similar solutions may be formally extended to the `complete solution space' by properly joining two separate branches in the first and fourth quadrants of the `semi-complete solution space' but with diverging asymptotic behaviours of $v(x)$ and $\\alpha(x)$ at $t\\rightarrow -\\infty$. Similarity solutions without crossing the sonic critical line as shown in Figs. 2 and 3 also carry ECCC and EECC features, respectively, and are much easier to construct technically in the isothermal approximation; in a polytropic gas with various laws of radiative losses, some qualitatively similar features were mentioned in subsection 8.2 of Boily \\& Lynden-Bell (1995).  Regarding the isothermal approximation, the core collapsing (CC) portion may be justified qualitatively on the ground that the gas would heat up by adiabatic compressions. In the presence of a sufficient amount of radiative agents in a relatively high density region involving more frequent collisions, such radiative cooling processes may sustain an approximate condition of quasi-constant temperature. To sustain an isothermal condition for the envelope expansion (EE) portion, some heating mechanisms are necessary to compensate the inevitable cooling resulting from adiabatic expansions. Naturally, radiations from gas in the CC portion can be absorbed by gas in the EE portion to heat up the latter to some extent (depending on the escape efficiency of photons). In the astrophysical context of star formation, a gas cloud can be submerged in an intense electromagnetic radiation field when the nuclear reaction has been ignited in the central protostar. In such a situation, the large-scale gas cloud dynamics of contraction and expansion may be regarded as quasi-isothermal. Similar conditions may be applicable in the emergence of a planetary nebula system as well as in a certain evolution phase for hot gas in a galaxy cluster. Although not fully understood in details, there are other possible heating processes that may happen in the EE portion such as micro-turbulence, flow instabilities, shocks, wave dampings and magnetic fields etc.  With these considerations in mind, it is worthwile to note that the isothermal assumption is a special case of the more general polytropic approximation (Cheng 1978; Goldreich \\& Weber 1980; Bouquet et al. 1985; Yahil 1983; Suto \\& Silk 1988; Maeda et al. 2002; Harada et al. 2003). Asymptotic solutions parallel to those of $(16)-(18)$ as well as eigensolutions across the sonic critical line can be readily derived. Therefore, EECC self-similar solutions in the polytropic approximation can be obtained with combined numerical and analytical analyses. One major difference is that, in the polytropic approximation, the asymptotic solution takes the form of \\begin{equation} v\\rightarrow Vx^{(1-\\Gamma)/(2-\\Gamma)}\\ , \\end{equation} \\begin{equation} \\alpha\\rightarrow Ax^{-2/(2-\\Gamma)}\\ , \\end{equation} \\begin{equation} m\\rightarrow Dx^{(4-3\\Gamma)/(2-\\Gamma)}\\ , \\end{equation} where $x\\propto rt^{(\\Gamma -2)}$ is the independent similarity variable, $\\Gamma$ is the polytropic index, and $V$, $A$ and $D$ are three constant coefficients to the leading order. For $1<\\Gamma<2$, the radial flow speed (either outflows or inflows) vanishes as $x\\rightarrow +\\infty$. For a finite radial range, the basic concept of EECC similarity solutions advanced in this paper remains valid. Additional similarity solution features such as shocks, subsonic radial oscillations, central core collapse etc. are also expected. Polytropic EECC similarity solutions should have a wider range of astrophysical applications with, at least, one more degree of flexibility (i.e. $\\Gamma$ parameter).  Conceptually, the EECC similarity solutions provide a simple model scenario for the concurrence of inner infalls (including central core collapses) and outer envelope outflows (including asymptotic winds) with the separation surface (or stagnation point) travelling outward at a constant speed. In any self-gravitating spherical fluid system, the emergence of an EECC similarity phase depends upon allowed physical conditions that can be reached via a certain set of initial/boundary conditions of density and velocity at some time ago through other processes that were usually not self-similar.  For applications of the EECC similarity solutions, as well as other similarity solutions, we have noted several potential areas. The first one is the dynamical evolution phase connecting the AGB or post-AGB phase and the proto planetary nebula (PN) phase, where the central core collapse to form a proto white dwarf and the envelope wind to create various PN morphologies are expected to happen simultaneously in a timescale of $\\sim 10^{3}$ yrs (e.g. Balick \\& Frank 2002). By this process, a late-type star can continue to lose its envelope mass while forming a compact object at the collapsing core. After the self-similar phase, the system continues to evolve in a detached manner with the core and envelope separated. Shocks can be introduced either at the end of the AGB phase (Kwok 1982, 1993) when the faster envelope expansion during the EECC phase catches up with the slower pre-existing wind or when infalling materials bounced back from the central compact object. Much as Tsai \\& Hsu (1995) and Shu et al. (2002) introduced shocks to the similarity solutions in star formation and `champagne flows' of H{\\sevenrm II} regions, one can construct EECC similarity solutions with shocks in the context of planetary nebulae. The solutions other than the EECC ones (including their time-reversal counterparts) have been used in star-formation as well as other gravitational collapse and flow problems.  For astrophysical systems of even larger scales such as clusters of galaxies (e.g. Sarazin 1988; Fabian 1994), the similarity solutions can be valuable for understanding a certain phase of their evolution. Specifically, if their evolution involves a self-similar phase of central collapse (e.g. Gunn \\& Gott 1972; Fillmore \\& Goldreich 1984; Bertschinger 1985; Navarro et al. 1997), then isothermal similarity solutions (17) and (18) seem to suggest two possible classes of galaxy clusters that emit X-rays through hot gases virialized in gravitational potential wells, namely, those with steep gravitational potential wells and thus extremely high X-ray core luminosities, and those with relatively smooth and shallow gravitational potential wells and thus normal X-ray core luminosities.  Finally, we note that the stability of similarity solutions is of interest in various contexts. As several authors have declared (Ori \\& Piran 1988; Hanawa \\& Nakayama 1997), some self-similar solutions may not be stable when one performs a normal mode analysis. The instability growth rates are different for various types of similarity solutions in general. At this moment, the stability problem for the EECC similarity solution remains open for future research. In some cases, depending on the growth rates, an unstable EECC similarity solution might be of interest to disrupt a self-similar phase within proper timescales. For example, we do not expect a self-similar phase to last forever, especially for the ephemeral phase of $\\sim 10^3$ yrs between the end of AGB phase and the PN phase (Kwok 1982, 1993; Balick \\& Frank 2002)."
    },
    "0311/astro-ph0311046_arXiv.txt": {
        "abstract": "The $\\simeq$ 100 kg highly radiopure NaI(Tl) set-up of the DAMA project (DAMA/NaI) took data over seven annual cycles up to July 2002 and has achieved results on various rare processes. Its main aim has actually been the exploitation of the model independent WIMP annual modulation signature. After this conference the total exposure, collected during the seven annual cycles, was released. This cumulative exposure (107731 kg $\\times$ day) has given a model independent evidence for the presence of a Dark Matter particle component in the galactic halo at 6.3 $\\sigma$ C.L.; this main result is summarised here. Some of the many possible corollary model dependent quests for the candidate particle are mentioned. At present, after about five years of new developments, a second generation low background set-up (DAMA/LIBRA with a mass of $\\simeq$ 250 kg NaI(Tl)) was built and is taking data since March 2003. New R\\&D efforts toward a possible NaI(Tl) ton set-up, we proposed in 1996, have been funded and started in 2003. ",
        "introduction": "The DAMA project has been proposed by the italian group to INFN and firstly funded in 1990 \\cite{1}; in 1992 the chinese colleagues joined the project. DAMA is an observatory for rare processes based on the development and use of various kinds of radiopure scintillators. Several low background set-ups have been realised; the main ones are: i) DAMA/NaI ($\\simeq$ 100 kg of radiopure NaI(Tl)), which took data underground over seven annual cycles and was put out of operation in July 2002; ii) DAMA/LXe ($\\simeq$ 6.5 kg liquid Xenon), which has been upgraded various times and is in operation; iii) DAMA/R\\&D, which is a set-up devoted to measurements on prototypes/small--scale--experiments and is in operation; iv) the new second generation set-up DAMA/LIBRA ($\\simeq$ 250 kg; more radiopure NaI(Tl)) in operation since March 2003. Moreover, in the framework of devoted R\\&D for higher radiopure detectors and PMTs, sample measurements are regularly carried out by means of the low background DAMA/Ge detector, installed deep underground since $\\gsim$ 10 years and, in some cases, by means of Ispra facilities. Recent results from DAMA/LXe are presented in these proceedings elsewhere \\cite{2}, while in ref. \\cite{3} some recent results achieved with DAMA/R\\&D can be found.  The DAMA/NaI set-up and its performances have been described in details in ref. \\cite{Nim98}. Since then some upgrading has been carried out; in particular, in summer 2000 the electronic chain and data acquisition system were completely substituted as reported in ref. \\cite{RNC}.  As mentioned, DAMA/NaI has taken data over seven annual cycles up to July 2002 \\cite{RNC}. The first data release and publication of DAMA/NaI data occurred in 1996 \\cite{Psd96} for an exposure of 4123 kg $\\times$ day (DAMA/NaI-0) analysed by using the pulse shape discrimination (PSD) technique in NaI(Tl) in order to discriminate between electromagnetic background and recoils\\footnote{The NaI(Tl) Dark Matter experiments have discrimination capability as largely stated in literature by several groups and at present exploited as main goal by the NAIAD experiment. In Zeplin I the pulse shape discrimination technique has been employed in LXe as well (sometimes this rejection procedure has been indicated as rejection by {\\it timing}), similarly as done and published by DAMA/LXe in 1998. Moreover, we remind that the WIMP annual modulation signature acts (as pointed out since the 80's by \\cite{Freese}) itself as a powerful background rejection procedure and that it can be effectively exploited only in large mass experiments such as so far DAMA/NaI and DAMA/LIBRA.}. Investigation of possible diurnal effects has been carried out as well \\cite{Diu}. The main aim of the experiment was actually the investigation of the presence of a Dark Matter particle component in the galactic halo by means of the model independent WIMP annual modulation signature, which has been deeply investigated over seven annual cycles (107731 kg $\\times$ day total exposure).  DAMA/NaI has also obtained results on the investigation of some exotic Dark Matter candidates (such as neutral SIMPs, neutral nuclearities and Q-balls) \\cite{exotic}, has searched for solar axions \\cite{Dax} and has given results on several other rare processes such as e.g. spontaneous emission of protons in $^{23}$Na and $^{127}$I with violation of the Pauli exclusion principle \\cite{35}, nuclear level excitation of $^{23}$Na and $^{127}$I during charge-non-conserving processes \\cite{cnc}, electron stability and non-paulian transitions in Iodine atoms by L-shell\\cite{Ela99}.  In the following, the result obtained by DAMA/NaI in the investigation of the WIMP annual modulation signature over the seven annual cycles will be briefly summarised. This cumulative 107731 kg $\\times$ day exposure has given a model independent evidence for the presence of a Dark Matter particle component in the galactic halo at 6.3 $\\sigma$ C.L.; this main result and corollary interpretations in some of the possible different model scenarios have been discussed in experimental and theoretical details in ref. \\cite{RNC} and have firstly been presented after this Conference in a seminar at Gran Sasso National Laboratory (LNGS) at end of July 2003.  The new second generation larger mass and higher radiopure NaI(Tl) set-up, named DAMA/LIBRA, is now in operation and will be briefly introduced. ",
        "conclusions": "DAMA/NaI has been a pioneer experiment running at LNGS for about a decade and investigating as first the WIMP annual modulation signature with suitable exposed mass, sensitivity and control of the running parameters. During seven independent experiments of one year each one, it has pointed out the presence of a modulation satisfying the many peculiarities of a WIMP induced effect, reaching an evidence at 6.3 $\\sigma$ C.L.. No other experiment has so far been in position to give a result directly comparable in a model independent way with that of DAMA/NaI. As a corollary, it has also pointed out the complexity of the quest for a WIMP candidate because of the present poor knowledge on the many astrophysical, nuclear and particle physics aspects.  After the completion, on July 2002, of the DAMA/NaI data taking, the second generation DAMA/LIBRA set-up ($\\simeq$ 250 kg mass) has been installed and preliminarily put in operation, as a result of continuous efforts by the DAMA collaboration toward the creation of ultimate radiopure NaI(Tl). DAMA/LIBRA, having a larger exposed mass and an higher overall radiopurity, will significantly contribute in the incoming years to the further understanding of the field. Presently, new R\\&D efforts toward a possible NaI(Tl) ton set-up have been funded and related works have already been started."
    },
    "0311/astro-ph0311100_arXiv.txt": {
        "abstract": "We present a simple technique to estimate mass-to-light ($M/L$) ratios of stellar populations based on two broadband photometry measurements, i.e.\\ a color-$M/L$ relation. We apply the color-$M/L$ relation to galaxy rotation curves, using a large set of galaxies that span a great range in Hubble type, luminosity and scale size and that have accurately measured H{\\small I} and/or H$\\alpha$ rotation curves. Using the color-$M/L$ relation, we construct stellar mass models of the galaxies and derive the dark matter contribution to the rotation curves.  We compare our dark matter rotation curves with adiabatically contracted Navarro, Frenk, \\& White (1997, NFW hereafter) dark matter halos. We find that before adiabatic contraction most high surface brightness galaxies and some low surface brightness galaxies are well fit by a NFW dark matter profile. However, after adiabatic contraction, most galaxies are poorly fit in the central few kpc. The observed angular momentum distribution in the baryonic component is poorly matched by $\\Lambda$CDM model predictions, indicating that the angular momentum distribution is not conserved during the galaxy assembly process.  We find that in most galaxies the dark matter distribution can be derived by scaling up the H{\\small I} gas contribution. However, we find no consistent value for the scaling factor among all the galaxies. ",
        "introduction": "To constrain the stellar mass distribution in our sample galaxies, we use the color-$M/L$ relationships for mixed stellar populations of Bell \\& de Jong (2001). In this paper we showed that for almost any reasonable star formation history and consistent chemical evolution of a galaxy, the stellar population models predict a strong correlation between the optical color of the population and its $M/L$. This was shown to be robust against gas inflow and outflow, different star formation prescriptions and mild starbursts (Fig.\\,1a).   \\begin{figure} \\mbox{ \\epsfysize=5.8cm \\epsfbox[74 355 216 513]{dejong_r_fig1a.eps} \\epsfysize=5.8cm \\hspace{1.cm} \\epsfbox[188 370 410 597]{dejong_r_fig1b.eps} } \\caption{ {\\bf Left (a):} Bell \\& de Jong (2001) model stellar population $M/L$ ratios in the $B$, $I$ and $K$-band versus $B$--$R$ color (points) and the fitted slopes (lines). {\\bf Right (b):} The maximum disk $M/L_K$ ratios for the Verheijen (1997) galaxies (upper limit symbols) versus $B$--$R$ color. The solid line shows the stellar population model predictions for a Salpeter IMF (solid line) and a renormalized Salpeter IMF (dotted line).} \\end{figure}  The one major assumption we had to make was that all galaxies have similar stellar Initial Mass Functions (IMFs), independent of their other properties. The IMF used in our color-$M/L$ analysis determines the slope of the relation, but not the zero-point. This reflects the fact that the high mass end of the IMF determines the luminosity and the color of a population, but the low mass end sets its mass, shifting the relation in Fig.\\,1a up and down.   We used the maximum disk $M/L_K$ limits of the Ursa Major galaxies observed by Verheijen (1997) to constrain the IMF. In Fig.\\,1b we show the maximum allowed $M/L_K$ ratios for these galaxies (hence the upper limit symbols) versus their $B$--$R$ color. We compare these to the predictions of our models for a standard Salpeter IMF (solid diagonal line). If our models are correct, all $M/L$ ratios of galaxies should scatter around this line. Clearly, this is in conflict with observations, as some galaxy upper limits are below the line, and thus using the $M/L$ of our model would over-predict the rotation curve of these galaxies. Hence, we use a normalization that is lower by 0.3 dex (dotted line). In principle, the normalization could be much lower, but there are many indications that in particular high surface brightness galaxies are close to maximum disk (see e.g.\\ Weiner, these proceedings) and therefore we chose a normalization that has these galaxies close to maximum disk. If our color-$M/L$ relation is correct and all true $M/L$ values are near this line, clearly many galaxies are substantially sub-maximal, with some galaxies in the top part of the diagram being sub-maximal by a factor of 10. We use here the updated color-$M/L$ relations of Bell et al.~(2003), which were calibrated with SDSS and 2MASS data. ",
        "conclusions": "Stellar population synthesis models predict a strong color-$M/L$ relation. We use this correlation to model rotation curves of galaxies and determine their DM distribution. Some of the main conclusions of our investigation are:\\\\ - There are large degeneracies in the fitted NFW parameters, even with the use of stellar $M/L$ constraints.\\\\ - With the chosen color-$M/L$ normalization, adiabatically contracted halos not only over-predict central rotation speeds of dwarf galaxies, but also those of high surface brightness galaxies.\\\\ - The angular momentum distribution in the observed baryons differs from $\\Lambda$CDM DM halo predictions, suggesting that detailed angular momentum is not conserved during the galaxy assembly process.\\\\ - H{\\small I} rotation curves cannot consistently be scaled to produce the DM signature."
    },
    "0311/astro-ph0311322_arXiv.txt": {
        "abstract": "The radial velocities of the 13 globular clusters in the Large Magellanic Cloud have a dispersion of 28 km s$^{-1}$ relative to the HI rotation curve of the LMC, compared to a dispersion of 30 km s$^{-1}$ with regard to the mean globular cluster velocity. This shows that, contrary to a suggestion by Schommer et al. (1992), one cannot yet rule out the possibility that the LMC globular clusters formed in a pressure supported halo, rather than in rotating disk. The globular clusters in the LMC may therefore, after all, exhibit a relationship between age and kinematics that is similar to that of the clusters in M33. ",
        "introduction": "In a pioneering investigation of the clusters in the Large Magellanic Cloud, Freeman, Illingworfth \\& Oemler (1983) surmised that ``[T]here is no old, kinematical halo population among the clusters of the LMC.'' At that time radial velocities were only available for nine old clusters in the Large Cloud. Presently 13 globular clusters are known to be associated with the Large Magellanic Cloud (Schommer 1991, Schommer et al. 1992, Suntzeff 1992, van den Bergh 2000, p. 104). From their discussion of the kinematics of old clusters in the LMC Schommer et al. (1992) concluded that: ``The oldest clusters still present an enigma; they do not have the kinematics of an isothermal, or slowly rotating, pressure-supported halo. These objects rotate with an amplitude comparable to that of the younger disk''. Taken at face value this conclusion would appear to indicate that the oldest clusters in the Large Cloud formed quite differently from those in M33 (Schommer et al. 1991, Chandar et al. 2002), in which the oldest (globular) clusters are observed to have halo kinematics, whereas younger (open) clusters exhibit disk-like motions. The view that late-type galaxies such as M33 (M V = -18.9) and the LMC (M V = -18.5), which have similar uminosities, had very different evolutionary histories would not fit comfortably with most current views of galaxy evolution. This prompts one to ask how secure is the conclusion that the oldest massive clusters in the Large Cloud formed in a disk, rather than in the LMC halo? ",
        "conclusions": "By lumping old and intermediate-age LMC clusters together Schommer et al. (1992) found that these clusters exhibited significant rotation. However, the data plotted in Figure 1 show that the oldest (SWB class VII) clusters do not exhibit a significant rotational component. In this respect the Large Cloud cluster system resembles M33 in which Chandar et al. find that 85 $\\pm$ 5 percent of the oldest clusters have halo kinematics. It is therefore concluded that worries about the apparent difference between the kinematics of the oldest population components in M33 and in the LMC (van den Bergh 2000, p. 83) were probably premature.  I thank Ken Freeman for pointing out to me that the number of old LMC globular clusters might not be large enough to prove the conclusion that these clusters have disk kinematics. Thanks are also due to Peter Stetson for statistical advice.  \\clearpage"
    },
    "0311/astro-ph0311608_arXiv.txt": {
        "abstract": "We report infrared spectroscopy of the unusual eclipsing pre-main sequence object KH~15D, obtained using NIRSPEC on Keck II.  During eclipse, observations using low spectral resolution ($\\lambda/\\delta\\lambda \\sim 1000$) reveal the presence of prominent molecular hydrogen emission in 5 lines near $2~\\mu$m.  The relative line strengths are consistent with thermal excitation at $T \\sim 2800 \\pm300K$.  Observations out of eclipse, at both low and high spectral resolution ($\\lambda/\\delta\\lambda \\sim 2 \\times 10^{4}$), show reduced contrast with the stellar continuum. The change in contrast for the strongest line, 1-0~S(1), is consistent with an approximately constant emission line superposed on a variable stellar continuum. Emission in the 1-0~S(1) line is observed to extend by $\\gtrsim 4$ arc-sec both east and west of the stellar point spread function ($\\gtrsim 3000$ AU).  Observed at high spectral resolution, the velocity and intensity structure of the 1-0~S(1) profile are both asymmetric.  East of the stellar PSF (by $1.1 - 2.3$ arc-sec) the emission is blueshifted ($-63$ km sec$^{-1}$), and has significantly greater intensity than the marginally redshifted component ($+2$ km sec$^{-1}$, $\\sim$ consistent with zero) which dominates west of the stellar PSF. The spatial extent of the emission, and the excitation temperature, suggest shock-excitation of ambient gas by a bipolar outflow from the star and/or disk.  However, it is difficult to account for the observed radial velocity unless the outflow axis is inclined significantly to the plane of the sky. ",
        "introduction": "KH~15D is a weak-lined T Tauri star (K7V) in NGC 2264 ($d=760$ pc), which shows deep (3.5 mag) periodic eclipses, each lasting for about 20 days, a large fraction of the 48.36 day period \\citep{her02,hh03}. The eclipses occur abruptly, indistinguishable from obscuration by a `knife edge' \\citep{her02}.  During eclipse, the lack of reddening in the spectrum \\citep{ham01}, and polarization which varies only weakly with wavelength \\citep{agol03}, imply that the obscuring matter must be comprised of large particles. A consensus view to date is that the eclipses are caused by structure in a protoplanetary disk, fortuitously edge-on to our line of sight. Candidate structures include a density wave \\citep{her02}, a vortex \\citep{bv03}, or a warp in the disk induced by planet-disk interactions \\citep{agol03,winn03}.  The considerable recent interest in KH~15D derives from the diagnostic potential inherent in the eclipse geometry, and the serendipity of finding a protoplanetary system whose structure suggests that planetary formation is actively occurring.  Although the eclipses in KH~15D are evidently caused by large solid particles, observations of the gaseous component may help to deduce the nature of the system.  \\citet{ham03} observed double-peaked atomic hydrogen emission lines, indicative of a bipolar outflow.  In this {\\it Letter} we report the spectroscopic discovery of $2~\\mu$m molecular hydrogen lines in emission, also showing a double-peaked structure indicative of outflow.  This H$_2$ emission has been independently discovered by \\citet{tok03}, by imaging in a narrow-band filter. Molecular hydrogen emission is commonly observed from pre-main sequence objects \\citep{bach96}, where it can be excited by fluorescence \\citep{black87} or by shocks \\citep{sb82}, and has been well-studied in many cases.  H$_2$ observations should therefore help to relate KH~15D to other pre-main sequence systems. ",
        "conclusions": "\\citet{tok03} have imaged KH~15D in a narrow band filter centered on the 1-0 S(1) line.  They find a prominent jet starting in close coincidence with the star, and extending many arc-sec to the north. They have considered whether the jet is physically associated with KH~15D or is a chance superposition (H$_2$ emission is common in star-forming regions like NGC 2264).  Independent of our spectroscopy, they conclude that the association of the jet with KH~15D is probably physical.  We find that: 1) our low resolution spectrum during eclipse shows bright S(1) emission centered on the stellar position to $\\sim 0.2$ arc-sec (Sec. 3.1), and 2) our high resolution spectrum resolves the line into two velocity components, whose relative intensities change aburptly at the stellar position (Sec 3.2).  Therefore our results strongly support the conclusion that S(1) emission is physically associated with KH~15D.  Recently, quiescent H$_2$ 1-0~S(1) emission has been detected near several T-Tauri stars \\citep{wein00, bary03}. This quiescent emission shows narrow line profiles, just as we see in KH~15D.  We considered the possibility that the KH~15D emission originates from quiescent gas, with Doppler shifts resulting from Keplerian orbital motion.  We cannot exclude that some fraction of the emission, especially in the red component, arises from quiescent emission.  Nevertheless, because the KH~15D emission is spatially extended and shows evidence of outflow, it seems more closely related to the shock-heated H$_2$ which has been observed in hundreds of young stellar objects, including classical T-Tauri stars \\citep{bach96, her97}.  Interpretation of our KH~15D observations requires knowing the stellar radial velocity, and the heliocentric radial velocity we adopt for the system ($+10$ km sec$^{-1}$) may be in error if the star is a binary. Nevertheless, the double-peaked profiles we observe in the close vicinity of the star suggest outflow.  However, a simple picture of a symmetric bipolar jet emergent normal to a single edge-on disk is hard to reconcile with our observations, and with the \\citet{tok03} imaging.  In such a simple interpretation, the line-of-sight shock velocity in the 1-0 S(1) line would have to be at least half of the difference between the red- and blueshifted components, i.e. $\\gtrsim 30$ km sec$^{-1}$.  For a disk inclination of $\\sim 84^{\\circ}$ \\citep{ham03}, the flow would be inclined by only $\\sim 6^{\\circ}$ to the plane of the sky, so the total shock velocity would have to exceed $\\sim 250$ km sec$^{-1}$.  This is well above the velocity where shocked H$_2$ is observed in other objects (H$_2$ becomes dissociated at high shock velocities, \\citet{sb82}).  Therefore the outflow revealed in molecular hydrogen is likely to be inclined significantly to the plane of the sky.  Similar arguments would not apply to lines formed very close to the star, where a bipolar outflow may not be collimated.  However, the emission we observe is spatially resolved from the star, occurring at distances where bipolar outflows are often highly collimated to a single axis.  \\citet{tok03} find only one bright jet of line emission extending to the north.  Since our slit was oriented primarily east-west (PA $=66^{\\circ}$), we have observed approximately (but not precisely) perpendicular to the jet axis.  Our spectra will reveal the flow pattern in the jet to the extent that the total angle between our slit and the velocity vector differs from ${\\pi}/2$.  We believe that our brightest emission component (blueshifted `east' spectrum in Figure~2) samples the \\citet{tok03} jet, whereas the red components are consistent with either a possible fainter counter-jet (not yet detected by imaging), or quiescent emission near the stellar velocity.  The extent and complexity of H$_2$ emission from KH~15D is reminiscent of T~Tauri itself, a known binary system.  Consistent with its binary nature, T~Tauri shows two nearly perpendicular outflow systems \\citep{her97}.  Our observations are not yet sufficient, in terms of spatial sampling, to say whether KH~15D shows a single, or multiple, velocity systems.  Mapping spectroscopy of KH~15D in the 1-0~S(1) line is needed to clarify the kinematics of this interesting object."
    },
    "0311/astro-ph0311052_arXiv.txt": {
        "abstract": "We describe the first results from two observational projects aimed at measuring the amount and spatial distribution of dark matter in distant early-type galaxies (E/S0s) and clusters of galaxies. At the galaxy scale, the Lenses Structure and Dynamics (LSD) Survey is gathering kinematic data for distant (up to $z\\sim1$) E/S0s that are gravitational lenses. A joint lensing and dynamical analysis constrains the fraction of dark matter within the Einstein radius, the mass-to-light ratio of the stellar component, and the total slope of the mass density profile. These properties and their evolution with redshift are briefly discussed in terms of the formation and evolution of E/S0 galaxies and measurement of the Hubble Constant from gravitational time delay systems. At the cluster scale -- after careful removal of the stellar component with a joint lensing and dynamical analysis -- systems with giant radial arcs can be used to measure precisely the inner slope of the dark matter halo. An HST search for radial arcs and the analysis of a first sample are briefly discussed in terms of the universal dark matter halos predicted by CDM simulations. ",
        "introduction": "Decades after the discovery of dark matter around spiral galaxies little is known about dark matter in early-type galaxies. Mass tracers at large radii (such as stellar kinematics, kinematics of globular clusters and of planetary nebulae, and X-ray halos) generally indicate that a constant mass-to-light ratio cannot reproduce the observations, although there are typically large uncertainties and a wide variety of behavior is seen. The main source of uncertainty in interpreting kinematic measures is that the derived mass profile depends on the assumed orbital structure. This is commonly referred to as the mass-anisotropy degeneracy.  The situation is even more uncertain outside the local Universe, because traditional kinematic tracers at large radii are generally not feasible at cosmological distances. However, additional constraints can be gathered by looking at early-type galaxies that are gravitational lenses. The configuration of multiple images provides information on the mass distribution of the lens; first and foremost a robust measurement of the mass enclosed within the Einstein Radius (typically larger than the effective radius), {\\it independent} of its kinematic status.  The Lenses Structure and Dynamics (LSD) Survey (Koopmans \\& Treu 2002, 2003; Treu \\& Koopmans 2002a,2003; hereafter KT) is obtaining internal kinematics for a sample of of 11 relatively isolated E/S0 lens galaxies at $z=0.04-1.01$. A joint lensing and dynamical analysis is used to break the mass-anisotropy degeneracy and determine the amount and distribution of luminous and dark matter. The first results from this project are briefly summarized and discussed in Section~2.  At the cluster scale, the presence of dark matter has been known for a long time (Zwicky 1937). The distribution of dark matter has been probed over a wide range of scales using X-ray analysis, dynamical studies and gravitational lensing, showing that dark matter dominates in mass. The wealth of mass tracers in clusters -- in particular the existence of giant radial arcs -- makes them an ideal laboratory to test the existence of universal mass density profiles predicted by cold dark matter (CDM) simulations. Combining a lensing analysis of clusters with giant arcs with a dynamical analysis of the kinematics of the brightest cluster galaxy (BCG), we can measure very accurately the mass density profile within $\\sim100$ kpc, disentangle the stellar and dark matter, and pinpoint the inner logarithmic slope of the dark matter halo. This has motivated a search of suitable systems with a central dominant BCG and giant arcs through the entire HST-WFPC2 archive (Sand et al.\\ 2004, in preparation), and the spectroscopic follow-up at Keck (Sand, Treu \\& Ellis 2002; Sand et al.\\ 2003). The first results from this project are briefly summarized and discussed in Section 3. ",
        "conclusions": ""
    },
    "0311/astro-ph0311578_arXiv.txt": {
        "abstract": "We present the high resolution X-ray spectrum of the BL Lac object PKS~2155-304 taken with the RGS units onboard {\\em XMM-Newton} in  November 2000. We detect a  O~{\\sc vii} K$\\alpha$  resonant absorption line from warm/hot local gas at 21.59 \\AA \\/ ($\\sim 4.5$~$\\sigma$ detection). The line profile is possibly double peaked. We  do not confirm the strong 20.02 \\AA \\/ absorption  line seen with {\\em Chandra} and interpreted as $z \\sim 0.05$ O{\\sc viii} K$\\alpha$. A $3$~$\\sigma$ upper limit of 14 m\\AA \\/ on the equivalent width is set. We also detect the $\\sim 23.5$~\\AA \\/  interstellar O{\\sc i} 1s$\\rightarrow$2p  line and derive a factor $\\leq 1.5$ subsolar O/H ratio in the ISM along \\pks \\/ line of sight. ",
        "introduction": "The observed baryon density at $z > 2$ \\citep[e.g.][]{rau98,wei97} agrees well with Standard Big-Bang nucleosynthesis predictions, when combined with observed light-element ratios \\citep{bur98}. At lower redshift, however, less than $1/3$ of the baryons observed at $z > 2$ have been detected so far \\citep{fuk98}. According to simulations for the formation of structures in the Universe, most of such baryons would be located, at the present epoch, in low density intergalactic gas, which has been shock-heated to temperatures of $\\sim 10^{5} - 10^{7}$~K \\citep[i.e. the warm-hot intergalactic medium WHIM, e.g.][]{hel98}. The most efficient way to detect the presence of WHIM is through resonant absorption lines from highly ionized metals (e.g. O{\\sc vi}, O{\\sc vii}, O{\\sc viii}, Ne~{\\sc ix}) imprinted in the far UV (FUV) and soft X-ray spectra of background sources \\citep{ald94, mul96, hel98, per98, fan00}. However, severe instrumental limitations have hampered so far the detection of conspicuous amount of WHIM. The Far Ultraviolet Spectroscopic Explorer ({\\em FUSE}) allows the OVI doublet ($\\lambda = 1031.93$ and $\\lambda=1037.62$ \\AA ) to be detected only up to z$\\sim 0.2$. More importantly OVI dominates the relative abundance distribution of O in shock-heated gas in pure collisional ionization equilibrium, only at relatively low temperatures, $T\\sim (1-5) \\times 10^{5}$~K, and so it tracks just the low-temperature tail of the WHIM distribution \\citep[e.g.][]{fan00}. Soft X-rays are far more promising. The strongest resonant transitions from O{\\sc vii}, Ne{\\sc ix} (both helium-like ions, and therefore highly stable) and O{\\sc viii}, all fall in the soft X-ray band. These ions dominate the O and Ne relative abundance distributions of both collisionaly ionized gas, and mixed, photoionized and collisionaly ionized gas, over a broad range of temperatures, between $T \\sim 5\\times 10^5$ K and $T \\sim 10^7$ K, where the WHIM distribution peaks \\citep{dav01,fan00}. The major fraction of the  WHIM should then be visible in the X-rays. However, despite the large relative abundances of ``X-ray'' ions in the WHIM, the current sensitivity and resolution of X-ray spectrometers has allowed so far only the strongest (EW$\\geq 10$ m\\AA) of these systems to be detected, and only against spectra of very bright background sources (Nicastro et al. 2002; Mathur, Weimberg \\& Chen 2003; Fang et al. 2002; Fang, Weimberg \\& Canizares 2003; with {\\em Chandra} and Rasmussen, Kahn \\& Paerels 2003 with {\\em XMM-Newton}; for a recent review see Paerels \\& Kahn 2003).  Three out of these six detections have been made against the spectrum of the bright ($F_{\\rm 2-10 keV} \\simeq 2 \\times 10^{-11} - 5 \\times 10^{-10}$ erg cm$^{-2}$ s$^{-1}$), nearby \\citep[z=0.116]{fal93}, blazar PKS~2155-304 \\citep{nic02, fan02}. A high quality ($\\geq 700$ counts per resolution element at the relevant wavelengths) {\\em Chandra} (HRCS/LETG) observation of PKS~2155-304, revealed the existence of O{\\sc vii}, O{\\sc viii} and Ne{\\sc ix} absorption lines at $z \\simeq 0$, identified with a  WHIM system, pervading our Local Group \\citep{nic02}. Lower quality {\\em Chandra} (ACIS/LETG) spectra of PKS~2155-304 confirmed the above findings \\citep{fan02}, and also show a line at 20.02 \\AA, that \\citet{fan02} identify as the O{\\sc viii} WHIM counterpart of a known Ly$\\alpha$ and O{\\sc vi} system at $z \\sim 0.05$, where a concentration of galaxies is seen \\citep[][and references therein]{mar88,shu98,shu03}. However it is not clear why such a feature is not detected in the higher quality HRCS/LETG spectrum published by \\citet{nic02}. Additional data are needed to clarify this issue.  In this paper we present the analysis of the high resolution {\\em XMM-Newton} Reflection Grating Spectrometers (RGS hereinafter) spectra of \\pks .\\\\ The structure of the paper is as follows: in \\S~\\ref{sec_obs}  we report on the  RGS data reduction and analysis. The spectral fitting and comparison with previous measurements is given in \\S~\\ref{sec_fits}. The discussion of the detected features is presented in \\S~\\ref{sec_disc}. ",
        "conclusions": "\\label{sec_disc}  At $z \\sim 0$ we confirm the existence of a O{\\sc vii} K$\\alpha$ absorption feature, but we do not detect features of higher ionization elements  (O{\\sc viii}, Ne{\\sc ix} and Ne{\\sc x}) seen by {\\em Chandra} \\citep{nic02}. This lack of detection is likely to be related to the limited number of photons (e.g. $\\sim 15300$ photons between 19.5 and 20.5 \\AA \\/, i.e. $\\sim 100$ photons per 0.0065 \\AA \\/ resolution element) collected in the {\\em XMM-Newton} observation, when the source was $\\sim 2-3$ times fainter than in \\citet{nic02}. The  O{\\sc vii} K$\\alpha$ appears to be unsaturated and the curve of growth presented in  \\citep{nic02} suggests an O{\\sc vii} column density of $\\sim 10^{16}$ cm$^{-2}$. If the weak evidence of a double peak in the profile of the O{\\sc vii} K$\\alpha$ line  is considered real, such shape may indicate that we are starting to detect the different sheets of the local filament.  The $3$~$\\sigma$ upper limit on the O{\\sc viii} Ly$\\alpha$ line EW, computed fixing the line FWHM to the  O{\\sc vii} K$\\alpha$ FWHM value, is $12.61$ m\\AA , consistent with the value of $9.0^{+2.6}_{-2.7}$ m\\AA \\/ found by \\citet{ras03} in the RGS spectrum of \\pks . Under the assumption of unsaturated lines the EW ratios between different ions of the same element depend on the gas temperature and density. Using  the ratio between the upper limit on the EW of O~{\\sc viii} Ly$\\alpha$ to the  O{\\sc vii} K$\\alpha$ (best fit value $-2$~$\\sigma$), we obtain an upper limit on the gas temperature of $\\sim 2.5-3.5 \\times 10^6$~K, for a  gas density of $10^{-6}$- 1 atom cm$^{-3}$ \\citep[see Fig.~5 in][]{nic02}. These values are consistent with the temperature range predicted for the WHIM.  It is not possible to discriminate with the present data whether the O{\\sc vii} features are produced by a WHIM, as proposed by \\citet{nic02} or by radiatively cooling gas within our Galaxy, as proposes by \\citet{hec02}. In the {\\em FUSE} spectrum  of \\pks \\/ \\citet{nic02} found two unsaturated O{\\sc vi} absorption lines: a narrow component (FWHM = $106 \\pm 9$ km s$^{-1}$) at $cz = 36 \\pm 6$ km s$^{-1}$ with EW = $(2.1 \\pm 0.2) \\times 10^{-3}$~eV, probably related to a cloud in the  Galactic disk, and a broad component (FWHM = $158 \\pm 26$ km s$^{-1}$) at $cz = -135 \\pm 14$ km s$^{-1}$ with EW = $(1.6 \\pm 0.4) \\times 10^{-3}$~eV possibly associated with a WHIM filament. Using our O{\\sc vii} K$\\alpha$ EW and the EW of the O{\\sc vi} broad component from \\citet{nic02} we derive a temperature range of $\\sim (4.5 -200) \\times 10^{5}$~K for a gas density of $\\sim 10^{-6}$ cm$^{-3}$, which restricts to $\\sim (4.5 - 25) \\times 10^{5}$~K once the upper limit from the O{\\sc viii} is considered. We note, however,  that the  O{\\sc vii} absorption features at local redshift have been found in the spectra of other 4 AGNs (3C~273 \\citet{fan03}; MRK~421 \\citet{nic01,cag03}; NGC~4593 \\citet{mck03} and NGC~3783 \\citet{kas02}) and might be associated to  the high velocity O{\\sc vi} lines (sampling a gas at $T\\sim 10^5-10^6$ K) recently seen in the spectra of bright AGNs by FUSE \\citep{sem02, nic02b}. If this is the case, the ``local''  O{\\sc vii} K$\\alpha$ would indicate the presence of WHIM within the Local Group  \\citep{nic02b}.    The line at 23.5 \\AA \\/ is  the interstellar O{\\sc i} 1s$\\rightarrow$2p transition  \\citep[e.g.][and references therein]{mcl98}. Assuming a neutral ISM gas, no consistent velocity broadening and unsaturated lines, the comparison of the measured  EW with the curve of growth derived using \\citet{sto97} cross sections implies $\\sim 2 \\times 10^{16}$ O atoms cm$^{-2}$. The H absorbing column density in \\pks \\/ direction is $1.36 \\times 10^{20}$ cm$^{-2}$ \\citep{loc95}, which brings to a O/H ratio of $1.4 \\times 10^{-4}$. This value is  $\\sim 6$ times smaller than in the Solar system \\citep[$8.5 \\times 10^{-4}$][]{and89}, but consistent with the ISM elemental abundances in the Magellanic Clouds \\citep[e.g.][]{rus90}. Assuming that the line EW is enhanced by velocity broadening would further decrease the O/H ratio. If we adopt the theoretical cross-sections of \\citep{mcl98}, as in \\citet[][(see their Fig.~3)]{dev03}, we derive $\\sim 7 \\times 10^{16}$ O atoms cm$^{-2}$, i.e. a O/H ratio of $\\sim 5 \\times 10^{-4}$, still $\\sim 1.5$ times lower than the Solar value. We note that \\citet{dev03}  O{\\sc i} 1s$\\rightarrow$2p EW in \\pks \\/ direction, derived from a coaddition of RGS observations for a total of 346~ks, is $15 \\pm 3$ m\\AA \\/ consistent with our measurement ($17.57^{+1.21}_{-5.81}$ m\\AA ) and the line curve of growth presented in the paper takes saturation effects into account, but still there is indication of a subsolar O/H ratio.    We do not confirm the O{\\sc viii} K$\\alpha$ line by \\citet{fan02}. Note that the lack of detection of WHIM outside our Local Group is consistent with the expectations of theoretical models. In fact only one O{\\sc viii} absorption line with EW$>3$ m\\AA \\/ (i.e. detectable by the present X-ray satellites) in the spectrum of a  random background source at $z\\sim 0.3$ is predicted \\citep[e.g.][]{hel98}. The non-detection of WHIM up to $z =0.116$  therefore does not pose any problem.  Our {\\em XMM-Newton} \\pks \\/ spectrum shows a hint ($1.4$~$\\sigma$ detection; EW $\\sim 5.56$~m\\AA ) of a 22.7 \\AA \\/ feature, seen by \\citet{dev03} in the  {\\em XMM-Newton} spectra of \\pks \\/ and MRK~421, and in the {\\em Chandra} spectrum of \\pks . Since they do not find evidence of such line in the {\\em XMM-Newton} spectra of  Sco~X-1 and 4U~0614$+$091, they  exclude a possible instrumental origin of the line and tentatively identify it with O{\\sc iv} absorption from a local WHIM filament. Higher statistics and other time spaced observations are needed to investigate the reality of the 22.7 \\AA \\/ absorption feature."
    },
    "0311/astro-ph0311264_arXiv.txt": {
        "abstract": "We present the results of a photometric and spectroscopic survey of the globular cluster system of \\cena\\ (Centaurus~A), a galaxy whose proximity makes it an important target for early-type galaxy studies. We imaged three fields in {\\it UBVRI} that extend 50 and 30~kpc along the major and minor axes, respectively.  We used both color and size information to develop efficient selection criteria for differentiating between star clusters and foreground stars.  In total, we obtained new velocities for 138 globular clusters, nearly tripling the number of known clusters, and bringing the confirmed total in \\cena\\ to 215.  We present a full catalog of all known GCs, with their positions, photometry, and velocities.  In addition, we present catalogs of other objects observed, such as foreground stars, background galaxies, three Galactic white dwarfs, seven background QSOs, and 52 optical counterparts to known X-ray point sources. We also report an observation of the cluster G169, in which we confirm the existence of a bright emission line object. This object, however, is unlikely to be a planetary nebula, but may be a supernova remnant. ",
        "introduction": "Systems of globular clusters are a nearly ubiquitous feature of all nearby galaxies.  Globular clusters (GCs), which are typically old with sub-solar metallicities, are the most visible remnants of intense star formation that occurred in a galaxy's distant past.  As single-age, single-metallicity stellar populations, GCs are ideal for studying the fossil remains of a galaxy's star formation and metal-enrichment history.  The past decade has seen a rapid growth in the study of extragalactic GC systems.  The availability of the Hubble Space Telescope (HST) in particular has enabled the study of GC systems in galaxies well beyond the Local Group.  One of the more striking results of these studies is the frequency with which the metallicity distributions of these GC systems are bimodal (e.g. Larsen \\etal 2001; Kundu \\& Whitmore 2001). Different scenarios of galaxy formation (or at least of globular cluster system formation) have been proposed or adapted to explain the observed metallicity distributions and other properties of GC systems: mergers of spiral galaxies (Ashman \\& Zepf 1992), multiple {\\it in situ} star formation epochs (Forbes, Brodie, \\& Grillmair 1997), dissipationless hierarchical merging of protogalactic clumps (C\\^{o}t\\'{e}, Marzke, \\& West 1998), and more generalized hierarchical merging (Beasley \\etal 2002).  The question that all of these scenarios address largely concerns the nature and time frame of merging or gas dissipation. Elements of all of these scenarios are supported by various studies of different facets of galaxy formation and evolution.  For example, the supernova winds that some use to explain the suppression of star formation in dwarf galaxies at high redshift (Dekel \\& Silk 1986) has also been proposed as a possible mechanism for the truncation of metal-poor GC formation at early times (Beasley \\etal 2002).  Similarly, hierarchical merging models are often used to explain the mass assembly of present-day galaxies.  While most spheroid formation may have occurred in the past (at redshifts $z > 2$), present-day examples of recent merger remnants may give us a window onto this distant epoch. Locally, there is compelling evidence that some ellipticals have recently interacted or merged with another galaxy. Many ellipticals possess large-scale disks of gas and dust.  Faint structure in the form of loops, shells, ripples, and tails are especially visible in the outer regions of ellipticals (Malin \\& Carter 1983), and are presumably the aftermaths of a recent interaction. One such galaxy, and perhaps the best candidate for study, is the nearby elliptical, \\cena.  We have chosen to examine of the stellar content of \\cena\\ in order to further elucidate its formation history and gain insight on the formation of other ellipticals.  Our survey of the planetary nebula (PN) system in \\cena's outer halo represents the kinematics of the field star population, and is presented in a separate paper (Peng, Ford, \\& Freeman 2004).  In this paper, we describe our photometric and spectroscopic survey for globular clusters, and present the resulting catalogs of objects.  \\vspace{0.3cm} ",
        "conclusions": "We conducted an optical imaging and spectroscopic survey for globular clusters across $\\sim 1\\degr$ of sky around \\cena.  Using \\ubvri\\ photometry, size, and morphological information, we developed an efficient algorithm for selecting likely GCs candidates for spectroscopic follow-up---a necessity given the large number of foreground stars.  We obtained radial velocities for over 400 objects.  Of these, we identified 102 previously unknown GCs, confirmed the nature of 24 GCs from Rejkuba (2001), 6 GCs from HCH99, 4 GCs from MAGJM96, and 2 from HHH86, providing 138 new GC velocities.  We also obtained new velocities for 25 previously confirmed GCs from HGHH92 and HHH86.  The total number of confirmed GCs in \\cena\\ is now 215.  We present a spectrum of HGHH-G169, which shows an interesting split emission line feature.  It is unlikely to be a planetary nebula, as MR02 suggest, but it could be a supernova remnant, and certainly merits a deeper spectrum."
    },
    "0311/astro-ph0311028_arXiv.txt": {
        "abstract": "In response to the claim by Dziembowski {\\em et al.}\\ (2001) that the solar radius decreases with magnetic activity at the rate of 1.5 km yr$^{-1}$, we consider the theoretical question whether a radius variation is expected with the solar cycle. If the radius variation is caused by the magnetic pressure of toroidal flux tubes at the bottom of the convection zone, then the dynamo model of Nandy and Choudhuri (2002) would suggest a radius decrease with magnetic activity, in contrast to other dynamo models which would suggest a radius increase.  However, the radius decrease is estimated to be only of the order of hundreds of metres. ",
        "introduction": "By analyzing the f-mode frequency data from MDI for the period 1996--2001, Dziembowski {\\em et al.}\\ (2001) claimed that the Sun shrinks in radius from the solar minimum to the solar maximum, at the rate of about 1.5 km per year.  If correct, then this will certainly be a most dramatic discovery in solar physics. However, this claim has not yet been supported by other researchers, who often present conflicting results.  Antia (2003) provides a review of the subject and concludes that there is no compelling evidence for solar radius change.  It may be noted that some earlier studies (N\\\"oel, 1997; Emilio {\\em et al.},\\ 2000) reported much larger radius variation (up to 700 km) based on other (non-helioseismic) methods.  It is clear that more work is needed before a definitive conclusion can be reached. The aim of this paper is to discuss whether on theoretical grounds we expect a solar radius variation with solar cycle.  Other than some early work by Spruit (1982, 1994), we have not been able to find any other theoretical analysis of this problem.  In the light of recent developments of solar physics, this theoretical question certainly needs revisiting.   There are two ways in which a magnetic field present in the solar interior may lead to an increase of solar radius. \\begin{enumerate} \\item The pressure of the magnetic field embedded in the gas may cause the gas to expand, thereby inflating the size of the Sun. \\item The inhibition of convection by the magnetic field may cause a piling up of heat below the magnetic field.  This would result in heating of the surrounding regions, leading to an expansion. \\end{enumerate} From both of these arguments, it appears that the radius of the Sun should increase when more magnetic fields are present.  Since we normally think of the solar maximum as being the time when the Sun has more magnetic fields in the interior, we would expect the Sun to be larger in size during the solar maximum.  Shrinking of the Sun with increasing magnetic activity, therefore, would seem improbable at the first sight.  We shall argue that the recent solar dynamo model of Nandy and Choudhuri (2002), in principle, allows for the possibility of shrinking of the Sun during the solar maximum.  The important theoretical question is whether the theoretically expected radius decrease agrees with the claim of Dziembowski {\\em et al.}\\ (2001). Spruit (1982, 1994) studied the effect of heat blocking by magnetic fields and concluded that the total radius variation cannot be more than 0.14 km. We mainly look at the other effect of magnetic pressure and present order-of-magnitude estimates, suggesting that the radius variation, for reasonable assumptions, cannot be as large as what Dziembowski {\\em et al.}\\ (2001).  We have also made a mathematical formulation of the problem, to go beyond order-of-magnitude estimates. To our surprise, we found that the mathematical problem is much more difficult to solve than what we anticipated.  Although our mathematical analysis so far has not yielded very definite results, we present a brief account of it to help future researchers who may wish to analyze this problem in greater depth.  Section 2 addresses the question whether we expect the radius to increase or decrease during the solar maximum.  An order-of-magnitude estimate is presented in Section~3, showing that a much smaller variation of radius is expected than what is reported by Dziembowski {\\em et al.} \\ (2001). Our incomplete mathematical analysis is given in Section~4.  Finally, our conclusions are summarized in Section~5. ",
        "conclusions": "Within the currently understood framework of solar MHD, we have not been able to provide a theoretical explanation of the controversial claim made by Dziembowski {\\em et al.}\\ (2001) that the Sun shrinks with increasing magnetic activity at the rate of 1.5 km yr$^{-1}$.  Spruit (1982, 1994) estimated the radius change due to heat blocking by sunspots near the solar surface.  Not only the radius variation was found to be much smaller, the radius is expected to increase with the solar maximum.  We considered the possibility if the excess magnetic pressure of flux tubes near the base of the SCZ can give us a different result.  Even though traditional dynamo models regard the solar maximum as the time when the magnetic flux in the solar interior peaks and would suggest a radius increase at the time of the solar maximum, we pointed out the model of Nandy and Choudhuri (2002) provides a different scenario.  Although this model would predict a decrease of the solar radius with increasing magnetic activity (i.e.\\ with increasing loss of magnetic flux from the storage region), various reasonable assumptions give a radius decrease rate about one order of magnitude smaller than what is claimed by Dziembowski {\\em et al}.\\ (2001).  It is true that our estimate is based on a very rough calculation.  However, we believe that it is still an over-estimate rather than an under-estimate. Only if the magnetic field at the base of the SCZ was concentrated to values as high as $10^6$ G, the radius shrinkage would have been as large as reported by Dziembowski {\\em et al.}\\ (2001).  There is no theoretical reason to expect a magnetic field as strong as $10^6$ G at the base of the SCZ.  In fact, it is not easy to generate even a $10^5$ G magnetic field there (Choudhuri, 2003).  We conclude that either the claim of Dziembowski {\\em et al.}\\ (2001) is incorrect, or else we do not understand some basic physics of the solar cycle."
    },
    "0311/astro-ph0311502_arXiv.txt": {
        "abstract": "We report further results from a 191~ks \\emph{Chandra} observation of the core of the Perseus cluster, Abell 426. The emission-weighted temperature and abundance structure is mapped detail. There are temperature variations down to $\\sim 1 \\kpc$ in the brightest regions.  Globally, the strongest X-ray surface brightness features appear to be caused by temperature changes. Density and temperature changes conspire to give approximate azimuthal balance in pressure showing that the gas is in hydrostatic equilibrium.  Si, S, Ar, Ca, Fe and Ni abundance profiles rise inward from about 100~kpc, peaking at about 30--40~kpc. Most of these abundances drop inwards of the peak, but Ne shows a central peak, all of which may be explained by resonance scattering. There is no evidence for a widespread additional cooler temperature component in the cluster with a temperature greater than a factor of two from the local temperature. There is however evidence for a widespread hard component which may be nonthermal. The temperature and abundance of gas in the cluster is observed to be correlated in a manner similar to that found between clusters. ",
        "introduction": "The Perseus cluster, Abell 426, is the brightest cluster of galaxies in the sky in the X-ray band.  The most prominent features in the core of the cluster are the cavities in the X-ray emitting gas (B\\\"ohringer et al 1993), associated with the bubbles blown by the central source 3C~84 (Pedlar et al 1990). The coolest X-ray gas in the cluster lies on the rims of these lobes (Fabian et al 2000, 2003a). A swirl of cool gas occupies the inner 70 kpc radius of the cluster (Churazov et al 2000; Fabian et al 2000; Schmidt, Fabian \\& Sanders 2002; Churazov et al 2003a; Fabian et al 2003a), winding around the core in an anti-clockwise direction.  The X-ray emission contains a number of quasi-periodic fluctuations in emission, which are interpreted as sound waves propagating in the intracluster medium (ICM), driven by the active nucleus (Fabian et al 2003a). In addition there appears to be a weak shock to the NE of the core, about 10 kpc from the bubble rim. The abundance in the cluster rises towards the centre, peaking around $0.7\\Zsun$ about 40-50 kpc from the core, and dropping back down in the centre (Schmidt et al 2002; Churazov et al 2003a).  We previously observed the Perseus cluster with \\emph{Chandra} (Fabian et al 2000; Schmidt et al 2002) for about~25 ks. The initial results from a much longer 191.2~ks \\emph{Chandra} observation were reported in Fabian et al (2003a), showing evidence for the presence of shocks and ripples in the core of the cluster. Here we present further results from this observation, in particular the temperature and abundance structures in this object, and tests for the presence of multiphase gas. As this is the deepest \\emph{Chandra} observation of a bright nearby diffuse object, our results are affected by systematic uncertainties which are not visible in shorter observations.  We do not discuss the high-velocity system (HVS) in depth here (see Gillmon, Sanders \\& Fabian 2003), nor the relationship between the H$\\alpha$ nebulosity and the X-ray emission (see Fabian et al 2003b).  The Perseus cluster is at a redshift of 0.0183. We assume that $H_0 = 70 \\kmpspMpc$; therefore 1 kpc corresponds to about 2.7 arcsec. ",
        "conclusions": "We have studied a 191~ks \\emph{Chandra} image of the core of the Perseus cluster for small-scale temperature and abundance structures. The obvious X-ray surface brightness features are seen to be due to temperature variations, with corresponding density variations ensuring approximate azimuthal pressure balance. Such variations are seen on scales down to above 2.5~arcsec ($\\sim1$~kpc) in the brightest regions.  We have made extensive searches for multi-temperature gas components. There is little evidence in the deprojected spectra for any widespread components at a temperature more than a factor of two away from the local temperature. There is some evidence for a widespread hot component (at 16~keV) in the projected data. This emission could alternatively be due to an extended nonthermal component. The solution to the cooling flow problem in the Perseus cluster must produce a fairly uniform mass distribution of gas from 2.5 to 8~keV, with little gas below~2 keV.  The abundances of Si, S, Ar, Ca, Fe and Ni rise inward from about 100~kpc, peaking at about 30--40~kpc. Most of these abundances level out or drop again inward of the peak. The extent of the drop is unclear, but it is plausibly explained by resonance scattering. O and Mg are more uniform across the core and Ne shows a central peak. The overall abundance pattern is similar to that found by others using XMM-Newton data of M87 in the Virgo cluster, except that S/Si exceeds unity in the Perseus cluster core. The abundance peaks are likely due to Type Ia supernovae in the central galaxy, NGC\\,1275."
    },
    "0311/astro-ph0311444_arXiv.txt": {
        "abstract": "The Extended Self-Similarity (ESS) of cosmic microwave background (CMB) radiation has been studied using recent data obtained by the space-craft based Wilkinson Microwave Anisotropy Probe. Using the ESS and the high angular scale resolution (arcminutes) of the data it is shown that the CMB temperature space {\\it increments} exhibit considerable and systematic declination from Gaussianity for high order moments at the small angular scales. Moreover, the CMB space increment ESS exponents have remarkably close values to the ESS exponents observed in turbulence (in magnetohydrodynamic turbulence). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311391_arXiv.txt": {
        "abstract": "At the end of their lives low mass stars such as our Sun lose most of their mass. The resulting planetary nebulae show a wide variety of shapes, from spherical to highly bipolar. According to the generalized interacting stellar winds model, these shapes are due to an interaction between a very fast tenuous outflow, and a denser environment left over from an earlier slow phase of mass loss. Previous analytical and numerical work shows that this mechanism can explain cylindrically symmetric nebulae very well. However, many circumstellar nebulae have a multipolar or point-symmetric shape. With two-dimensional calculations, Icke showed that these seemingly enigmatic forms can be easily reproduced by a two-wind model in which the confining disk is warped, as is expected to occur in irradiated disks. Here, we present the extension to fully three-dimensional adaptive mesh refinement simulations of such an interaction. ",
        "introduction": "In the final phases of stellar evolution, low mass stars, such as our Sun, first swell up and shed a dense, cool wind in the asymptotic giant branch (AGB) phase. This episode is followed by a fast, tenuous wind that is driven by the exposed stellar core, the future white dwarf. The planetary nebulae (PNe) resulting from this expulsion phase come in a wide variety of shapes, from spherical to highly bipolar. Some even have a multipolar or point-symmetric shape. Balick \\citep{1987AJ.....94..671B} proposed that such forms arise due to an interaction between a slow disk-shaped inner AGB nebula and the fast `last gasp' of the star. Analytical \\citep{1988A&A...202..177I, 1989AJ.....97..462I} and numerical \\citep{1989ApJ...339..268S, 1991A&A...251..369I, 1991A&A...252..718M} work shows that this {\\it generalized interacting stellar winds} (GISW) mechanism works very well (for an up-to-date review, see \\citet{2002ARA&A..40..439B}). Several scenarios for obtaining a disk around a PN exist, and it is in general assumed in the models that the shape of the dense gas around the star is a disk or a toroid.  Icke \\citep{2003A&A...405L..11I} proposed that the point-symmetric shapes observed in a number of PNe are formed in an interaction between a spherical stellar wind and a {\\it warped} disk. It is prossible to produce such a warp around a single star, through the combined effects of irradiation and cooling \\citep[e.g.][]{1996MNRAS.281..357P, 1996ApJ...472..582M}. Whereas Icke's computations were restricted to a two-dimensional proof-of-principle, we now present a first series of fully three-dimensional hydrodynamic computations of such a wind-disk interaction using the technique of adaptive mesh refinement (AMR). ",
        "conclusions": "Our computations show that the wind-disk interaction model in which the confining disk is warped results in a wide variety of point-symmetric shapes. Nebulae that show 'punched holes', such as NGC~7027 \\citep{2002A&A...384..603C}, are readily accommodated in this model. Other candidates for our model are probably NGC~6537 and, to a lesser extent, NGC~7026, in which the inner disk is still visible, while the outer nebula shows clear point-symmetric structures. Also, the shapes of the (proto-)PNe He~2-155 and M~4-18 (and maybe He~2-47) can be explained with our model. As a further application, large-scale explosions in non-planar disks, such as might occur in active galaxies, are expected to show similar patterns, provided that the disk material can cool rapidly enough. Movies of these simulations can be found at \\url{http://www.strw.leidenuniv.nl/AstroHydro3D/}. \\subparagraph{Acknowledgements} V.I. expresses his gratitude to Raghvendra Sahai and Hugo Schwarz for lively discussions that were the primary cause for taking up this subject.  The research of G.M. has been made possible by a fellowship of the Royal Netherlands Academy of Arts and Sciences.  The software used in this work was in part developed by the DOE-supported ASCI/Alliance Center for Astrophysical Thermonuclear Flashes at the University of Chicago.  Our work was sponsored by the National Computing Foundation (NCF) for the use of supercomputer facilities, with financial support from the Netherlands Organization for Scientific Research (NWO), under grant number $614.021.016$."
    },
    "0311/hep-ph0311171_arXiv.txt": {
        "abstract": "{ Within the Quantum Field Theory context the idea of a ``cosmological constant'' (CC) evolving with time looks quite natural as it just reflects the change of the vacuum energy with the typical energy of the universe. In the particular frame of Ref. \\cite{Letter}, a ``running CC'' at low energies may arise from generic quantum effects near the Planck scale,  $M_P$, provided there is a smooth decoupling of all massive particles below $M_P$. In this work we further develop the cosmological consequences of a ``running CC'' by addressing the accelerated evolution of the universe within that model. The rate of change of the CC stays slow, without fine-tuning, and is comparable to $H^2\\,M_P^2$. It can be described by a single parameter, $\\nu$, that can be determined from already planned experiments using SNe Ia at high z. The range of allowed values for $\\nu$ follows mainly from nucleosynthesis restrictions. Present samples of SNe Ia can not yet distinguish between a ``constant'' CC or a ``running'' one. The numerical simulations presented in this work show that SNAP can probe the predicted variation of the CC either ruling out this idea or confirming the evolution hereafter expected.} ",
        "introduction": "The Standard Cosmological Model fits our universe, in the large, into an homogeneous and isotropic Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) cosmological type\\, \\cite{Peebles}. Its 4-curvature is determined from the various contributions to its total energy density, namely in the form of matter, radiation and cosmological constant.  Evidence for a dominant content of energy in the form of cosmological constant was found by tracing the rate of expansion of the universe along z with high--z Type Ia supernovae \\cite{p99, Riess98}. This measurement combined with the measurements of the total energy density $\\OT^0$ from the CMB anisotropies \\cite{CMBR,WMAP}, indicates that $\\OL^0\\sim 70\\%$ of the critical energy density of the universe is cosmological constant (CC) or a dark energy candidate with  a similar dynamical impact on the evolution of the expansion of the universe. The matter content, on the other hand is dominated by the dark matter, whose existence is detected by dynamical means \\cite{Peebles}, and amounts to $\\OM^0\\sim 30\\%$ of the critical density.  The CC value found from Type Ia supernovae at high z \\cite{p99,Riess98} is: \\begin{equation}\\label{CCvalue} \\CC_0=\\OL^0\\,\\rc^0\\simeq 6\\,h_0^2\\times 10^{-47}\\,GeV^{4}\\,. \\end{equation} {Here $\\rc^0 \\simeq \\left( 3.0\\,\\sqrt{h_{0}}\\times 10^{-12}\\,GeV\\right) ^{4}$ is the present value of the critical density, and $h_{0}\\sim 0.7\\pm 0.1$\\ sets the typical range for today's value of Hubble's constant $H_{0}\\equiv 100\\;(Km/sec\\, Mpc)\\;h_{0}$}. In the context of the Standard Model (SM) of electroweak interactions, this measured CC should be the sum of the original vacuum CC in Einstein's equations, $\\CC _{vac}$, and the induced contribution from the vacuum expectation value of the Higgs effective potential, $\\CC_{ind}=\\langle V_{\\rm eff} \\rangle$: \\begin{equation} \\CC=\\CC _{vac}+\\CC _{ind}\\,. \\label{lambdaphys} \\end{equation} It is only this combined parameter that makes physical sense, whereas both $\\CC _{vac}$ and $\\CC _{ind}$ remain individually unobservable\\footnote{{In general the induced term may also get contributions from strong interactions, the so-called quark and gluon vacuum condensates. These are also huge as compared to (\\ref{CCvalue}), but are much smaller than the electroweak contribution $V_{\\rm eff}$.}}. From the current LEP 200 numerical bound on the Higgs boson mass, $\\,M_{H}> 114.1\\,GeV$ \\cite{LEPHWG}, one finds $\\,\\,\\,\\left| \\CC _{ind}\\right|=(1/8)M_H^2\\,v^2> 1.0\\times 10^{8}\\,GeV^{4}$, where $v\\simeq 246\\,GeV$ is the vacuum expectation value of the Higgs field. Clearly, $\\left| \\CC _{ind}\\right|$ is $55$ orders of magnitude larger than the observed CC value (\\ref{CCvalue}). Such discrepancy, the so-called ``old'' cosmological constant problem \\cite{weinRMP,CCRev}, manifests itself in the necessity of enforcing an unnaturally exact fine tuning of the original cosmological term $\\CC _{vac}$ in the vacuum action that has to cancel the induced counterpart $\\,\\CC _{ind}$ within a precision (in the SM) of one part in $\\,10^{55}$.  The measured CC remains very small as compared to the huge CC value predicted in the SM of Particle Physics. {Actually, if the physical value of the CC would conform with that one predicted in the SM, the curvature of our universe would be so high that the Special Theory of Relativity could not be a solution to Einstein equations to any reasonable degree of approximation. Therefore, the SM prediction of the CC violently contradicts our experience, whereas the small measured value (\\ref{CCvalue}) is perfectly compatible with it.  The Cosmological Constant Problem (CCP) is a fundamental problem. It is most likely related to the delicate interplay between Gravity and Particle Physics, and it has become one of the main poles of attention \\cite{weinRMP,CCRev}. All attempts to deduce the small value of the cosmological constant from a sound theoretical idea ended up with the necessity of introducing severe fine-tuning. This concerns also, unfortunately, the use of supersymmetry and string theory (see e.g. \\cite{wittenDM,Carroll2}). In this respect we recall that, for a realistic implementation of the existing versions of M-Theory, one would like to have a negative (or at least vanishing) cosmological constant in the remote future, such that it does not prevent the construction of the asymptotic S-matrix states in accelerated universes\\,\\cite{MTheory}. {Since the presently observed value of the CC is positive, there is the hope that a variable cosmological term may solve this problem.}  {There is a permanently growing flux of proposals concerning the CCP.} On the first place there is the longstanding idea of identifying the dark energy component with a dynamical scalar field \\cite{Dolgov,PSW}. More recently this approach took the popular form of a ``quintessence'' field slow--rolling down its potential\\,\\cite{Caldwell}. This proposal has, on its own, given rise to a wide variety of models \\cite{PeebRat,Moreq}. Extended models of this kind (``k--essence'') are also based on scalar fields but with a non-canonical kinetic energy \\cite{kessence}. {The main advantage of the quintessence models is that they could explain the possibility of a variable vacuum energy.} This may become important in case such variation will be someday detected in the observations. Recently other approaches have appeared in which the dark energy is mimicked by new gravitational physics \\cite{Carroll}. From the point of view of the CCP, all these approaches lead {to the introduction of either a very small parameter} or a very high degree of fine-tuning. In another, very different, vein the possibility to accept the observed value of the CC within the context of a many world pool is offered by the anthropic proposal \\cite{antrop}. Let us finally mention the intriguing proposal of non-point-like gravitons at sub-millimeter distances suggested in \\cite{Sundrum2}, or the possibility of having multiply degenerate vacua\\,\\cite{Yokoyama}.  When assessing the possibility to have variable dark energy, other no less respectable possibilities should be taken into account. In a series of recent papers \\cite{JHEPCC1,cosm}, the idea has been put forward that already in standard Quantum Field Theory (QFT) one would not expect the CC to be constant, because the Renormalization Group (RG) effects may shift away the prescribed value, in particular if the latter is assumed to be zero.  Thus, in the RG approach one takes a point of view very different from e.g. the quintessence proposal, as we deal all the time with a ``true'' cosmological term. It is however a variable one, and therefore a time-evolving, or redshift dependent: $\\CC=\\CC (z)$. Although we do not have a QFT of gravity where the running of the gravitational and cosmological constants could ultimately be substantiated, a semiclassical description within the well established formalism of QFT in curved space-time (see e.g. \\cite{birdav,book}) should be a good starting point.  From the RG point of view, the CC becomes a scaling parameter whose value should be sensitive to the entire energy history of the universe -- in a manner not essentially different to, say, the electromagnetic coupling constant. One of the main distinctions between our approach and all kinds of quintessence models is that these models imply the introduction of a phenomenological equation of state $p_{\\chi}=w\\,\\rho_{\\chi}$ for the scalar field $\\chi$ mimicking the CC, where $w$ is a negative index (smaller than $-1/3$). Whether constant or variable,  a ``true'' cosmological parameter has, instead, no other equation of state associated to it apart from the exact $w=-1$ one.  Attempts to apply the RG for solving the CC problem have been made earlier \\cite{Polyakov,lam}. The {canonical form} of renormalization group equation (RGE) for the $\\CC$ term at high energy is well known -- see e.g. \\cite{book,Nelson}. {However, at low energy decoupling effects of the massive particles may change significantly the structure of this RGE, with important phenomenological consequences.} This idea has been retaken recently by several authors from various interesting points of view \\cite{JHEPCC1,cosm,Babic,Reuter03a}. However, it is not easy to achieve a RG model where the CC runs smoothly without fine tuning at the present epoch. {In Ref.\\,\\cite{Letter,IRGA03} a successful attempt in this direction has been made, which is based on possible quantum effects near the Planck scale.} At the same time, the approximate coincidence of the observed $\\CC$ and the matter density, $\\ \\OL^0\\sim \\OM^0\\,$, i.e.  the ``new'' CC problem, or ``time coincidence problem'' \\cite{weinRMP,CCRev} {can be alleviated in this framework if we assume the standard (viz. Appelquist-Carazzone\\,\\cite{AC}) form of the low-energy decoupling for the massive fields.}  In the present paper we elaborate on this idea further. We develop a semiclassical FLRW model whose running CC is driven smoothly, without fine tuning, due to generic quantum effects near the Planck scale. We show that, due to the decoupling phenomenon, the low-energy effects (in particular the physics from the SM scale) are irrelevant for the CC running, and so the approximate coincidence between $\\OM^0$ and $\\OL^0$ is not tied to any particular epoch in the history of the universe. Furthermore, the new effects imply deviations from the standard cosmological equations due to quantum effects. Our ``renormalized'' FLRW model provides a testable framework that can be thoroughly checked from SNAP data on Type Ia supernovae at very high z -- see \\cite{SNAP, goods} and references in \\cite{rev}. If these experiments  detect a $z$-dependence of the CC similar to that predicted in our work, we may suspect that some relevant physics is going on just below the Planck scale. Alternatively, if they find a static CC, this might imply the existence of a desert in the particle spectrum near the Planck scale.  The structure of the paper is as follows. In the next section we compare constant versus variable CC models. In Section \\ref{sect:RCC} we present our variable CC model based on the Renormalization Group. In Section \\ref{sect:solving} we solve the FLRW cosmologies with running CC. In Section \\ref{sect:numanalysis} we study the numerical behaviour of these cosmologies, and the predicted deviations from the standard FLRW expectations. In Section \\ref{sect:newmodel}, we introduce the magnitude-redshift relation for the analysis of the  SNe Ia. In Section \\ref{sect:snap} we perform the simulations on the SNAP data in order to test the sensitivity with which the features of the new model can be determined. In the last section we draw our conclusions. {Two appendices are included at the end: one to discuss some technical issues inherent to our QFT framework, the other providing some background related to the statistical analysis.} ",
        "conclusions": "\\label{sect:conclusions}  To summarize, we have constructed a semiclassical FLRW model with variable CC at the present cosmic scale. The variation of the vacuum energy is provided without introducing special scalar (quintessence-like) fields and is completely caused by quantum effects of vacuum. The evolution of the CC is due to the Renormalization Group running triggered by the smooth decoupling of the massive fields at low energies, while the RG scale $\\mu$ being associated to the Hubble parameter, $H$, at the corresponding epoch. Although the $\\bCC$-function itself is proportional to the fourth power of the masses, the decoupling does still introduce an inverse power suppression by the heavy masses, and thus one is left with a residual quadratic law $\\bCC\\sim H^2 M^2$. The effective scale $M$ summarizes the presence of the heaviest degrees of freedom available. This peculiar form of decoupling stems directly from: \\,i) the Appelquist-Carazzone (AC)-decoupling theorem, \\,ii) general covariance, and also from\\, iii) the \\textit{non}-fine-tuning hypothesis between the $n=1$ (``soft decoupling'') terms of $\\bCC$ (Cf. Eq. (\\ref{newRG1})) whereby the overall coefficient of the quadratic contribution $H^2 M^2$ does not vanish. This particular form of decoupling is a specific characteristic of the CC because it is of dimension four. There is no other parameter either in the SM or in the GUT models with such property.  In constructing the cosmological model we have explored the possibility that the heaviest d.o.f. may be associated to particles having the masses just below the Planck scale. This assumption is essential to implement the soft decoupling hypothesis within the $\\mu=H$ setting. Indeed,  the present value of $\\,H_0^2 M^2\\,$ is just of the order of the CC, and therefore it insures a smooth running of the cosmological term around the present time. In this setting the $\\bCC$-function has only one arbitrary parameter $\\,\\nu\\,$ (\\ref{nu}), and as a result the model has an essential predictive power. In general we expect $\\,|\\nu|\\ll 1\\,$ from phenomenological considerations, mainly based on the most conservative hypotheses on the nucleosynthesis framework. Furthermore, for $\\,|\\nu|\\ll 1$, we insure the absence of the trans-Planckian energies. It is not completely clear how much the cosmological index $\\,\\nu\\,$ can approach the value $1$ from below. Despite the values $\\nu\\lesssim 1$ are not completely ruled out, it is not clear that the conditions for the nucleosynthesis could be safe. However, if accepted, it opens up the possibility that the CC could be in transit from a $\\,\\CC>0\\,$ regime into a future $\\,\\CC\\leq 0\\,$ one, which would be a welcome feature for string/M-theory.  Our RG model offers the possibility to explore the existence of the sub-Planck physics in direct cosmological experiments, such as SNAP (and the very high--z SNe Ia data to be obtained with HST). For example, for the flat FLRW case and a moderate and positive value $\\nu \\sim 10^{-2}$, we predict an increase of $10-20\\%$ in the value of $\\Omega_{\\Lambda}$ at redshifts $z\\gtrsim 1-1.5$ perfectly reachable by SNAP. For similar, but negative, values of $\\nu$ we predict that the CC should become negative beyond redshifts $z\\gtrsim 2$. For $\\nu\\sim 0.1$ corrections to some FLRW cosmological parameters become as large as $50\\%$ or more. In general, this model has a wide spectrum of implications that could be tested by SNAP, even for fairly moderate values of $\\nu$ compatible with the most conservative bounds from nucleosynthesis. The simulations of the SNAP data in Section 7 show that the single parameter $\\nu$ of our model could be pinned down, for $\\nu\\lesssim 0.1$, with a precision of around $20-60\\%$. Although this accuracy is not very high, the sign of the parameter could be determined and the effects would be manifest. Undoubtedly, this would be a good starting point to identify the presence of quantum corrections to the FLRW classical cosmology.  The semiclassical FLRW model that we have proposed here explains the variation of the CC due to the ``relic'' quantum effects associated to the decoupling of the heaviest degrees of freedom below the Planck scale, and suggest that a time dependence of the CC may be achieved without resorting to scalar fields mimicking the cosmological term or to modifications of the structure of the SM of the strong and electroweak interactions and/or of the gravitational interactions. At the same time, our model provides a phenomenological parametrization for possible correlated deviations of the cosmological equations for matter density and dark energy.  From a more fundamental point of view, we have shown that the cosmological constant and the matter density may evolve in a correlated way due to quantum effects without resorting to exceedingly exotic frameworks. Using the RG as a basic QFT tool and extrapolating the standard AC law of decoupling to the CC case, we have straightforwardly predicted a cubic redshift evolution law for the CC and the matter density. The larger the redshift that we can eventually explore the larger the effects that we predict. Most important, we have shown that this cubic law can be thoroughly tested by the next generation of cosmological measurements, which will be able to reach depths up to $z\\sim 2$."
    },
    "0311/astro-ph0311629_arXiv.txt": {
        "abstract": "When a main sequence star evolves into a red giant and its Kuiper Belt Object's (KBO's) reach a temperature of ${\\sim}$170 K, the dust released during the rapid ice-sublimation of these cometary bodies may lead to a detectable infrared excess at 25 ${\\mu}$m, depending upon the  mass of the KBO's. Analysis of IRAS data for 66 first ascent red giants with 200 L$_{\\odot}$ $<$ $L$ $<$ 300 L$_{\\odot}$ within 150 pc of the Sun provides an   upper limit to the   mass in KBO's at 45 AU orbital radius that  is usually less than ${\\sim}$0.1 M$_{\\oplus}$. With improved infrared data, we may detect systems of KBO's around first ascent red giants that are analogs to our Solar System's KBO's. ",
        "introduction": "The formation and evolution of large solids such as comets, asteroids and planets in astrophysical environments is of great interest.  Here, we investigate whether other stars possess Kuiper Belt Objects (KBO's) similar to those found in the Solar System.  We assume that the hypothetical KBO's around other stars resemble those in our Solar System and have  a composition of ice and dust similar to that of comets (see Luu \\& Jewitt 2002).  During the main sequence phase of the star's evolution, the KBO's quiescently remain in stable orbits.  However, when the star becomes a red giant, the KBO's  become sufficiently hot that ice sublimates and  previously embedded dust particles are ejected into the surroundings.   In $\\S$2, we sketch our argument,  which is presented in full detail in $\\S$3, that sufficient dust may be  released during this red giant phase that the star can display a detectable infrared excess.  Stern, Shull \\& Brandt (1990) and Ford \\& Neufeld (2001) computed the fate of ice sublimated from comets that have orbital radii greater  than 100 AU when the host star evolves onto the Asymptotic Giant Branch and attains a luminosity near 10$^{4}$ L$_{\\odot}$.  Now that the Kuiper Belt in the Solar System  with most objects lying near  45 AU from the Sun  is becoming better understood (Luu \\& Jewitt 2002), it is possible to imagine the response of such a system to a star's first ascent up the red giant branch when its luminosity exceeds 100 L$_{\\odot}$.  However, although progress is being made, there are still significant uncertainties about the Kuiper Belt in the Solar System.  Recent estimates of its  mass range  from 0.01 M$_{\\oplus}$ (Bernstein et al. 2003) to 0.1 M$_{\\oplus}$ (Luu \\& Jewitt 2002).  Therefore, although we describe a procedure to investigate Kuiper Belt-like systems around other stars,  our own outer Solar System is so poorly understood that an exact comparison is not yet possible.  Two  current descriptions of the KBO's differ mainly but not exclusively in estimates of the numbers of objects with radii smaller than ${\\sim}$200 km.    With the expectation that short period comets originate in  the Kuiper Belt,  Luu \\& Jewitt (2002) propose a rapid rise in the numbers of such smaller KBO's.  In contrast to this expectation,  Bernstein et al. (2003) used the Hubble Space Telescope to search for KBO's as faint as 28.3 mag, corresponding to radii of ${\\sim}$15 km, and found many fewer objects than expected from the  model given by Luu \\& Jewitt (2002).   However, with this observational result of Bernstein et al. (2002), the origin of the short period comets becomes a mystery. Perhaps there is a bimodal distribution of KBO's.  In this paper, given the uncertainties in current knowledge of the outer Solar System, we consider models both with and without large numbers of KBO's smaller than 200 km in radius.   Most first ascent red giants do not show any evidence for circumstellar dust.  The minority of  first ascent red giants with infrared excesses (Judge, Jordan \\& Rowan-Robinson 1987, Jura 1990, 1999, Kim, Zuckerman \\& Silverstone 2001, Smith 1998, Zuckerman, Kim \\& Liu 1995, Plets et al. 1997) do not seem to have  sublimating KBO's, since the characteristic inferred dust temperature is less than 70 K instead of 170 K, as we predict in $\\S$3.   Here, we focus on the possibility of detecting relatively warm dust from disintegrating KBO's. ",
        "conclusions": "We have computed the evolution of a system of KBO's when a star becomes a red giant. For  KBO's at 45 AU from the host star, we find that the dust released during  the sublimation of ice from these objects might produce a detectable infrared excess around red giants of luminosities in the range 200 L$_{\\odot}$ to 300 L$_{\\odot}$ if the mass of the KBO's is at least 0.1 M$_{\\oplus}$.   To date, there is no strong evidence for such sublimating KBO's, but with improved data, it may be possible to detect  systems which are analogs to our Solar System's KBO's around first ascent red giants.   This work has been partly supported by NASA."
    },
    "0311/astro-ph0311303_arXiv.txt": {
        "abstract": "Six low-luminosity active galactic nuclei have been imaged at multiple frequencies from 1.7--43~GHz (2.3--15~GHz for three of the galaxies) using the Very Long Baseline Array.  In spite of dynamic ranges of about 100 in several frequency bands, all six galaxies remain unresolved, with size limits at 8.4~GHz of $10^3$--$10^4$ times the Schwarzschild radii of the black holes inferred at their galactic centers.  The galaxy spectra are roughly flat from 1.7 to 43~GHz, rather than steepening to classical optically thin synchrotron spectra at high frequencies.  Although the spectral slopes somewhat resemble predictions for advection-dominated accretion flows, the luminosities are too high for the black hole masses of the galaxies and the slight spectral steepening at high frequencies cannot be explained by standard simple models of such accretion flows.  In contrast, compact jets can accommodate the average spectral index, the relatively high radio luminosity, and the unresolved appearance, but only if the jets in all six galaxies are fairly close to our line of sight.  This constraint is in agreement with inclination angle predictions for five of the six AGNs based on the dusty torus unification model. ",
        "introduction": "}    The physical processes that control the accretion and outflow in low-luminosity active galactic nuclei (LLAGNs) are still not well understood.  In their study of the Palomar Seyfert sample, \\citet*{Ho_FS1997.0} found that almost 40\\% of bright ($B_T < 12.5$~mag) nearby galaxies have optical spectra indicating the presence of an AGN.  Although their line emission resembles scaled down versions of more luminous Seyfert nuclei, which are thought to be driven by accretion in thin disks, LLAGN properties at other wavelengths suggest that different processes are at work.  Whereas ``classical'' Seyfert galaxies have radio spectral indices $\\alpha \\approx -0.7$, where $S_\\nu \\propto \\nu^{+\\alpha}$, almost half of the low-luminosity Seyfert galaxies in the Palomar Seyfert sample have flat or positive spectral indices \\citetext{\\citealp{Ulvestad_W1989}; \\citealp{Ho_U2001}, hereafter \\citetalias{Ho_U2001}; \\citealp{Ulvestad_H2001.0}, hereafter \\citetalias{Ulvestad_H2001.0}}.  Similar results have been found for low-ionization nuclear emission regions (LINERs) and other AGNs in the Palomar sample (\\citealp{Nagar_ea2000,Nagar_ea2002}; \\citealp*{Filho_BH2000,Filho_BH2002}).  These flat-spectrum objects typically appear compact on arcsecond scales, and high resolution studies with the Very Long Baseline Array\\footnote{The VLBA is operated by the National Radio Astronomy Observatory, a facility of the National Science Foundation, operated under cooperative agreement by Associated Universities, Inc.} \\citep[VLBA;][]{Napier_ea1993} show that they remain compact even on milliarcsecond scales \\citep[hereafter \\citetalias{Ulvestad_H2001.1}]{Falcke_ea2000,Ulvestad_H2001.1}.  Following the work of \\citet{Ferrarese_M2000}, \\citet{Gebhardt_ea2000}, and others, it has become clear that both normal and active galaxies with bulges have supermassive black holes at their nuclear centers.  Accretion onto these black holes is thought to fuel the AGN process for both ``normal'' and low-luminosity AGNs. Since their black hole masses cover roughly the same range \\citep[see, for example, ][]{Ho_2002}, the mass of the central object cannot account for the difference in emission levels.  The explanation for why LLAGNs emit so little radiation may be either that they have far less mass accreting onto them, or that the dominant mode of the accretion mechanism itself must be different, or both.  The center of our own Milky Way Galaxy has been used as the testbed for many models that attempt to explain the physics of LLAGNs. Acceleration measurements of stellar orbits near the source Sagittarius~A$^\\ast$ (Sgr~A$^\\ast$) leave little doubt as to the existence of a supermassive black hole of $\\sim 3.3\\times 10^6~\\mathrm{M_\\Sun}$ at the Galactic Center \\citep{Ghez_ea2003,Schoedel_ea2003}.  Invisible in the optical and near-infrared, Sgr~A$^\\ast$ has a mildly inverted spectrum from 1 to 100~GHz \\citep{Falcke_ea1998}.  The luminosity of Sgr~A$^\\ast$ is extremely low, even compared with other LLAGNs.  Advection dominated accretion flows (ADAFs) have been used to explain the spectral energy distributions (SEDs) of Sgr~A$^\\ast$ and other LLAGNs during the past decade (\\citealp{Narayan_Y1994,Narayan_Y1995.1}; \\citealp*{Narayan_YM1995,Narayan_MQ1998}).  In the ADAF model, a hot plasma develops in which the protons decouple from the electrons, allowing the protons to carry most of the energy released by the accretion process into the black hole, limiting the amount that is radiated away.  Hot thermal electrons generate synchrotron radiation that is self-absorbed, resulting in a radio spectral index of $\\sim 0.4$ \\citep{Mahadevan_1997}.  Bremsstrahlung and inverse-Compton processes generate additional emission, producing a relatively well-defined spectrum from radio through X-ray frequencies.  An alternative hypothesis for the emission from LLAGNs is the compact-jet model \\citep{Falcke_B1995,Falcke1996,Falcke_B1999,Falcke_M2000}.  Plasma is ejected in a collimated beam away from the black hole, and expands sideways to form a cone-shaped emission region.  Synchrotron emission is initially self-absorbed, becoming progressively more transparent to lower frequencies as the material travels away from the black hole, producing an overall radio spectrum that can be quite flat.       With both ADAF (and other related low-radiation accretion flows) and jet models failing to account for the overall SEDs in LLAGNs (see, for example, \\citealp{Narayan_MQ1998}; \\citealp*{Yuan_MF2002}), several attempts have been made recently to combine the two models. \\citet{Yuan_MF2002} have reasonable success in explaining the spectrum of Sgr~A$^\\ast$ with this combined model.  Although this work has been extended to \\NGC{4258} \\citep{Yuan_ea2002}, rigorous testing of a large sample of LLAGNs in the nearby universe has yet to be performed. These LLAGNs are many orders of magnitude \\emph{more} luminous than Sgr~A$^\\ast$, but are also several orders of magnitude \\emph{less} luminous than ``classical'' AGNs.  They therefore occupy a luminosity range that has thus far had few detailed studies.  We have formed an LLAGN sample from the \\citetalias{Ho_U2001} study, selecting objects which have 5~GHz Very Large Array (VLA) peak flux densities between 5 and 30~mJy and spectral indices at milliarcsecond resolution of $\\alpha_{1.4}^{5} > -0.35$.  In this paper we discuss simultaneous, multifrequency radio imaging of six LLAGNs.  The purpose of these observations, and the point of the present paper, is to measure the sizes and radio spectra of LLAGNs within $10^4$~Schwarzschild radii of their black holes, and to determine what constraints these parameters can place on possible ADAF and jet models. ",
        "conclusions": "}  We have performed multifrequency, milliarcsecond-scale radio imaging of six LLAGNs with flat/inverted spectra.  All six galaxies have slightly rising spectra throughout the gigahertz range, and are unresolved on scales of $10^3$--$10^4$ Schwarzschild radii.  The radio powers are higher than predicted by simple ADAF models, but the radio powers and spectra are consistent with dominance by compact radio jets.  However, the unresolved nature of the radio emission implies that modified ADAF models including nonthermal electrons are still viable models.  Otherwise, jet models require the jets to be seen closer to end-on than to side-on in all six galaxies.  This result is consistent with the fact that at least five of the size galaxies are Type~1 galaxies, where unified schemes predict that our viewing angle is within $\\sim 45\\degr$ of the jet axis."
    },
    "0311/astro-ph0311135_arXiv.txt": {
        "abstract": "We report the detection of the first 2 microlensing candidates from the Wendelstein Calar Alto Pixellensing Project (WeCAPP). Both are detected with a high signal-to-noise-ratio and were filtered out from 4.5 mill. pixel light curves using a variety of selection criteria. Here we only consider well-sampled events with timescales of $1\\E{d} < \\tfwhm < 20\\E{d}$, high amplitude, and low $\\chi^2$ of the microlensing fit. The two-color photometry $(R,I)$ shows that the events are achromatic and that giant stars with colors of $(R-I)\\approx 1.1\\E{mag}$ in the bulge of M31 have been lensed. The magnification factors are 64 and 10 which are obtained for typical giant luminosities of $M_I=-2.5\\E{mag}$. Both lensing events lasted for   only a few days ($t_\\M{fwhm}^\\M{GL1}=1.7\\E{d}$ and $t_\\M{fwhm}^\\M{GL2}= 5.4\\E{d}$). The event GL1 is likely identical with PA-00-S3 reported by the POINT-AGAPE project. Our calculations favor in both cases the possibility that MACHOs in the halo of M31 caused the lensing events. The most probable masses, $0.08\\,M_\\odot$ for GL1 and $0.02\\,M_\\odot$ for GL2, are in the range of the brown dwarf limit of hydrogen burning. Solar mass objects are a factor of two less likely. ",
        "introduction": "Microlensing experiments are an ideal method to search for dark objects within and between galaxies.  A large number of microlensing events have been detected towards the Galactic bulge constraining the number density of faint stars in this direction \\citep{alard99, afonso99,alcock00a,udalski00}. Towards the LMC only 13-17 microlensing events have been reported so far \\citep{alcock00b}. If all this events are attributed to $0.5\\,M_\\odot$ MACHOs, the associated population of dark objects would contribute up to the 20\\% level to the dark matter content of the Milky Way \\citep{alcock00b}. However, both the relatively large size of the LMC relative to its distance and the nature of the lenses has cast doubt on this interpretation.  It is indeed likely that a large fraction of the microlensing events towards the LMC are due to self-lensing of stars within the LMC (\\citet{lasserre00,evans00} and references therein).  Studying microlensing events towards M31 allows to separate self-lensing and halo-lensing in a statistical way, since the optical depth for halo lensing is larger on the far side of M31. In M31 individual stars can not be resolved and one therefore has to use the pixellensing technique \\citep{crotts92,baillon} to follow the variability of sources blended with thousands of other sources within the same pixel. First detections of possible microlensing events were reported by several pixellensing experiments \\citep{crotts96,ansari99,auriere01, paulinhenriksson02,paulinhenriksson03,calchinovati03}. But since the candidate nature of only 5 of these events is convincing, no conclusions concerning the near-far asymmetry or the most likely dark matter lensing masses could be drawn yet.  The Wendelstein Calar Alto Pixellensing Project (WeCAPP, \\citealt{arno01}) started in 1997 with test observations. Since 1999 the bulge of M31 was monitored continuously during the time of visibility of M31. The analysis of our 4 year data will allow not only the identification of very short duration events (eg., in the 4$^\\M{th}$ year data of the combined field have been taken on 83 \\% of possible nights) but also the separation of long duration ($<20\\E{d}$) microlensing events from long periodic variables as Mira stars. For this letter we analyzed the short duration events ($\\tfwhm<20\\E{d}$) within one season of Calar Alto data and restricted the detection to high-signal-to noise, high-magnification events. We report our first 2 microlensing candidates of that type. ",
        "conclusions": "We presented the first two high ($S/N$), short timescale microlensing events from WeCAPP. GL1 is likely identical to PA-00-S3 found by POINT-AGAPE. Combining the data from AGAPE with ours shows that the error bar of the derived Einstein time scale becomes smaller by a factor of 5 compared to the individual error bar. This demonstrates the importance of a good time sampling of the events. We derived the colors of the lensed stars, the amplification factors and likely lens masses for both bulge/disk self-lensing and MACHO lensing. We showed that red giants are the likely source objects, while main sequence stars are highly implausible.  Self-lensing in the bulge can only be separated from MACHO lensing statistically.  Halo-lensing events show a spatial asymmetry because the optical depth for lensing events is higher for stars on the far side of M31 than on its near side \\citep{crotts92,theory}. In contrast, bulge self-lensing is symmetric.  The bulge self-lensing hypothesis yields lensing stars at or below the the main sequence turn-off of the M31 bulge. On the other hand, if the lensing events are caused by MACHOs, their masses are typically very low, most probable below $0.1\\,M_\\odot$. Masses in the range of $0.5\\,M_\\odot$ to $1\\,M_\\odot$ are more unlikely.  So far, we have  analyzed one observing season and restricted the lensing search to short-time, high-amplification events in order to avoid confusion with variable stars. The whole WeCAPP dataset will allow us to identify all variables and thus will enable a search for lower amplitude and longer duration microlensing events.  Decreasing the amplitude threshold will increase the detected rate of events in two ways. As the event rate is proportional to the inverse of the minimum required magnification $A_{0,\\M{min}}$ in the pixellensing regime we expect to detect more lensed giants. On the other hand lowering the amplification threshold could make it possible to detect also highly amplified main sequence stars \\citep{hangould} which exceed the evolved stars in the bulge of M31 by a factor of more than a hundred. How many more lensing events will be detected depends on the mass function of the lenses but we can expect at least a factor of a few \\citep{theory}.  Finally, the effects of time sampling and noise properties of our sample on the detectability of lensing events have to be taken into account. Results of the modeling of these effects for events of different durations and amplitudes using Monte-Carlo simulations will be presented in a future publication. With the full dataset we expect therefore to increase the number of lensing events to detect the predicted asymmetry of MACHO lensing or to rule out a significant MACHO population in the halo of M31."
    },
    "0311/astro-ph0311245_arXiv.txt": {
        "abstract": "We develop a formalism wherein the solution of the equations of equilibrium for a self-gravitating, magnetized interstellar cloud may be accomplished analytically, under the only hypotheses that ({\\it a}\\/) the cloud is scale-free, and ({\\it b}\\/) is electrically neutral. All variables are represented as series of scalar and vector spherical harmonics, and the equilibrium equations are reduced to a set of coupled algebraic e\\-qua\\-ti\\-ons for the expansion coefficients of the density and the magnetic vector potential. Previously known axisymmetric solutions are recovered and new non-axisymmetric solutions are found and discussed as models for pre-protostellar molecular cloud cores. ",
        "introduction": "{\\it ``Do you think that non-symmetric, magnetized, self-gravitating equilibria can exist?''} This question is a typical example of the classic, fundamental kind of problems that Frank likes to address. This contribution presents a summary of the results obtained by the writer in his effort to answer Frank's question. ",
        "conclusions": "Whether or not real cloud cores are triaxial ellipsoids that pulsate because of their intrinsic topological non-equilibrium, as suggested by the theoretical arguments discussed in the previous sections, the material presented in this contribution illustrates well the richness of possible implications concealed in a typical question asked by Frank."
    },
    "0311/astro-ph0311559_arXiv.txt": {
        "abstract": "{It is well known that the surface brightness of elliptical galaxies and of bulges of spiral galaxies is best fitted by the Sersic $r^{1/n}$ profile. It is thus interesting to explore the lensing properties of the Sersic model because of its wide range of applicability. To this aim, we evaluate the lensing potential, the deflection angle, the time delay between the images and the amplification for a circularly symmetric lens whose surface density is described by the Sersic profile. We estimate the same quantities also adding a shear term to the deflection potential in order to study (at the lowest order) the impact of deviations from radial symmetry or the contribution from other nearby lenses. Moreover, we investigate the systematic errors due to the use of a de\\,Vaucouleurs profile instead of the correct Sersic model also taking into account the presence of a dark halo. ",
        "introduction": "Elliptical galaxies present a striking regularity in their global luminosity distributions in the sense that, within a wide range of sizes, their surface brightness profiles can be described by simple functions. The de Vaucouleurs or $r^{1/4}$\\,law (\\cite{deV48}) was first proposed as a quite general function to fit the light profile of the elliptical galaxies and of the bulges of lenticular and spiral galaxies. However, it was soon realized that the Sersic $r^{1/n}$ profile, which generalizes the de Vaucouleurs law, is best suited to describe the surface brightness distribution of these systems (\\cite{CCD93,GC97,PS97}). The regular properties of ellipticals have been the subject of different approaches, concerning both photometric and spectroscopic parameters, resulting in interesting scaling relations. Some well known examples are the fundamental plane relating the effective radius $r_e$, the luminosity intensity $I_e$ at $r_e$ and the central velocity dispersion $\\sigma_0$ (\\cite{DD87,7S87}); the photometric plane linking the three parameters of the Sersic profile (\\cite{KWKM00,G02}) and the entropic plane (\\cite{Metal01}).  The regularity of elliptical galaxies is an attractive feature for cosmological applications of gravitational lensing since it makes it possible to reduce the degeneracy among lens parameters and hence the systematic uncertainty in the estimate of cosmological quantities. The time delay among the images in multiply imaged quasars systems may be used as a tool to determine the Hubble constant $H_0$ (\\cite{Ref64}) once a model for the lens has been fitted to the observed images configuration, flux and time delay ratios. The use of scaling relations makes it possible to narrow the parameter space expressing some of them as function of the remaining ones. On the other hand, it has been estimated that at least 80$\\%$ of gravitational lenses are elliptical galaxies (\\cite{TOG84,FT91}). Hence, it is important that statistical lensing studies, such as the number counts of lens systems, the distribution of image angular separations and of time delays, take into account the most correct model of elliptical galaxies in the analysis of the available data. Up to now, this has not been done since these studies have been concentrated on the dark halos rather than on the luminous components of the galaxies. Therefore, this latter component is usually neglected or simply modeled with the spherical Hernquist profile (\\cite{H90}) which approximately reproduces the de\\,Vaucouleurs law when projected to the lens plane. The only remarkable exception is the work by Maoz \\& Rix (1993) where the lens is modeled using the $r^{1/4}$\\,law directly, together with an isothermal dark halo. It should be interesting to repeat their analysis using the more realistic Sersic profile and the scaling relations we have quoted above. As a first step, one has to study the lensing properties of this model. This is the aim of the current work.  The paper is organized as follows. In Sect.\\,2, we determine the lensing potential and the deflection angle for a circularly symmetric lens described by the Sersic model. We are thus able to write the lens equations whose solutions are found in Sect.\\,3. Section\\,4 is devoted to the evaluation of the amplification of the images, while the effect of a shear term, due to deviations from circular symmetry or to nearby lenses, is investigated in Sect.\\,5. A study of the systematic errors due to the use of a de\\,Vaucouleurs profile mimicking the best\\,-\\,fit Sersic model to the lens surface brightness is presented in Sect.\\,6. We summarize and conclude in Sect.\\,7. ",
        "conclusions": "Simple analytical calculations show that elliptical galaxies form 80$\\%$ of the gravitational lenses. Hence, it is very interesting to investigate their lensing properties taking into account models which are as realistic as possible. The Sersic $r^{1/n}$ profile is well known to fit the surface brightness of both elliptical galaxies and bulges of lenticular and spiral galaxies. In this paper we have studied its lensing properties assuming circular symmetry of the surface mass distribution.  Under this hypothesis, we have derived analytical expressions for the lensing potential, the deflection angle and the image amplification. Solving the lens equations, we have found that a circularly symmetric Sersic lens produces two images of the source that are merged in the Einstein ring if $r_s = 0$, i.e. the source and the lens centre are perfectly aligned. One of the two images may be deamplified and hence, for some choice of $n$, a Sersic lens does not produce two observable images, an effect that must be taken into account in statistical studies such as lens number counts. No radial critical curve is formed by a circularly symmetric Sersic lens, while the tangential critical curve is the Einstein ring.  Adding a shear term to the lensing potential allows us a first order investigation of the impact of deviations from circular symmetry and to take into account the contribution of nearby lenses. We have solved the lens equations in this case assuming, without loss of generality, that the shear is on\\,-\\,axis. Due to the breaking of radial symmetry, the number and positions of the images formed depend now both on $r_s$ and $\\theta_s$ and not only on $r_s$ as in the symmetric case. We have found that an on axis source gives rise to two images, while two or four images may appear as the source moves away from the lens axis. The presence of the shear also changes the amplification of the images and the shape of the critical curves. The tangential critical curve is deformed into an ellipse, while an inner radial critical curve appears. The size of both curves depends on the Sersic model parameters and decreases with the exponent $n$.  \\begin{figure} \\centering \\resizebox{8.5cm}{!}{\\includegraphics{masserrsers+plum.eps}} \\hfill \\caption{Same as Fig.\\,\\ref{fig: masserriso}, but with $\\beta = -2$ for the halo model.} \\label{fig: masserrplum} \\end{figure}  Our results suggest that a Sersic lens can produce up to four images of a distant source if a shear term is added to the potential. This is in agreement with the recent result by Kochanek (2003) who was able to fit constant $M/L$ models to two quadruple QSOs, namely RXJ0911+0551 (\\cite{Ba97,Bu98}) and PG1115+080 (\\cite{W80,Ietal96}), using the de\\,Vaucouleurs profile, that is to say a Sersic lens with $n = 4$. However, a direct comparison is not possible. To fit RXJ0911+0551, one has to use two lensing galaxies modeled with the de Vaucouleurs profile, taking also into account the presence of the cluster which the lenses belong to. A similar situation also holds for PG1115+080 where a single de Vaucouleurs lens is unable to fit the images configuration and the contribution of a nearby cluster is strongly required to reproduce the observable quantities.  The Sersic profile has been shown to be a better fit to the surface brightness distribution of elliptical lensing galaxies than the de\\,Vaucouleurs law. Nonetheless, the de\\,Vaucouleurs model is usually used in lensing studies. We have thus investigated the systematic errors induced by this procedure. This analysis has shown that the mass\\,-\\,to\\,-\\,light ratio $\\Upsilon$ may be seriously under\\,- or overestimated (depending on $n$ being lower or larger than 4) when considering constant $M/L$ models. On the other hand, if a dark halo is added to the visible component of the lens galaxy, the use of a de\\,Vaucouleurs model instead of the correct Sersic one leads to a wrong estimate of the dark halo projected mass of the dark halo inside the Einstein ring. These considerations could suggest that the use of de\\,Vaucouleurs profile to model galaxies that are best fitted by the Sersic law must be avoided. Actually, these results should be treated with great caution since they have been obtained using simulated galaxies whose model parameters have been fixed with a procedure that is by no means unique. In particular, as described in detail in the Appendix, we have used the photometric plane (\\cite{KWKM00,G02}) to relate the Sersic parameters $(n, r_e, I_e)$. Actually, this relation has been determined from galaxies belonging to the Virgo and Fornax clusters, while we have extended it to higher redshift galaxies the validity of which is still to be proven. Henceforth, a different prescription to fix the model parameters could be implemented and a detailed study of how this affects the results is needed. However, the procedure we have adopted leads to Sersic parameters describing galaxies which look quite reasonable so that we are confident that our main results are trustworthy. Actually, the best strategy is to resort to a careful study of real lens systems in order to check whether the effects we have reported are indeed true, or a consequence of how we have chosen the model parameters.  The analytical expressions of the lensing potential and of the deflection angle here obtained for a circularly symmetric Sersic lens may be easily generalized to the flattened case following the procedure described, e.g., in Keeton (2001). The resulting potential, deflection angle, and amplification may then be used to try to fit the observed lens systems (both double and quadruple) in order to investigate if constant $M/L$ models are viable solutions or whether a dark matter component has to be invoked to reproduce the observable lensing quantities (images positions and flux ratios). Since the Sersic model parameters $(n, r_e, I_e)$ may be directly measured from detailed photometry, the only unknown lens quantity is the mass\\,-\\,to\\,-\\,light ratio $\\Upsilon$ and hence we can study eventual correlations among $\\Upsilon$ and the photometric parameters. On the other hand, we can assume that lensing galaxies obey the same scaling relations (such as the photometric plane and the fundamental plane) as the low redshift systems. If we model the lens as the sum of a luminous Sersic component and a dark halo, we can reduce the degeneracy among model parameters excluding those solutions which lead to Sersic parameters not allowed by the quoted scaling relations. Moreover, if the time delays among the images were measured, this approach should reduce the systematic error in the determination of the Hubble constant through the time delay method (\\cite{Ref64}). On the other hand, statistics of gravitational lensing, such as lenses number counts and distribution of images angular separation, are known to be powerful tools to discriminate among different cosmological models. The modeling of the lenses plays a key role in this kind of analysis and it is thus very important to use realistic lens models. The Sersic lens we have studied may thus be a first step towards a reanalysis of the available data in order to better investigate if it is possible to reconcile them with constant $M/L$ models and if the constraints on the cosmological parameters are changed by the use of more realistic models of elliptical galaxies. These questions will be addressed in a series of forthcoming papers."
    },
    "0311/astro-ph0311073_arXiv.txt": {
        "abstract": "We searched for spectral signatures from the exosphere of HD209458 in the UV and the near-IR with UVES and ISAAC at the VLT. We looked in  particular for the helium absorption feature at 10 830 \\AA\\, predicted to be among the strongest ones. The upper limit of the He I line derived is 0.5\\% at 3 $\\sigma$ for a 3 \\AA\\ bandwith. Planet-induced chromospheric activity search on HD209458 was also performed. ",
        "introduction": "The He I feature at 10 830 \\AA\\ is predicted to be one of the deepest absorption lines of the diffuse and extended exosphere in the optical and near-IR transmission spectrum The goal of the present study is to search for exospheric signatures in the spectrum of HD209458, during transits of the planet. Shkolnik et al. (2003) suggested that a chomospheric activity can be induced by the planet orbit. We searched for such a signature in stellar chromospheric lines. ",
        "conclusions": ""
    },
    "0311/astro-ph0311590_arXiv.txt": {
        "abstract": "Optical spectroscopy of the Type Ia supernova SN 2002ic obtained on 2003 June 27.6~UT, i.e., $\\sim$ 222 rest-frame days after explosion, is presented. Strong H emission indicates an interaction between the expanding SN ejecta and an H-rich circumstellar medium (CSM). The spectrum of SN~2002ic resembles those of SNe 1997cy and 1999E. The three SNe also have similar luminosities, suggesting that they are the same phenomenon and that the CSM is also similar. We propose a new classification, Type IIa SNe, for these events. The observed line profiles and line ratios are measured and discussed within the ejecta-CSM interaction scenario. The emission in H Balmer, [\\ion{O}{3}], and \\ion{He}{1} lines, and in permitted \\ion{Fe}{2} blends, resembles the spectra of the Type IIn SN 1987F and of Seyfert 1 galaxies. A high-density, clumpy CSM is inferred. Strong, very broad [\\ion{Ca}{2}]/\\ion{Ca}{2} and [\\ion{O}{1}]/\\ion{O}{1} emissions imply that not all the outer SN ejecta were decelerated in the interaction, suggesting that the CSM is aspherical. ",
        "introduction": "\\citet{ham03} reported strong \\ion{Fe}{3}, \\ion{Si}{2}, and \\ion{S}{2} features in the early-time spectra of SN~2002ic and classified it as a Type Ia supernova (SN Ia). However, strong H$\\alpha$ emission was also observed. The detection of H$\\alpha$ is unprecedented in an SN Ia. (For reviews on SN spectra, see \\citealt{fil97}). The emission was broad (FWHM $> 1000$ km~s$^{-1}$), suggesting that it was intrinsic to the SN. \\citet{ham03} suggested that it arose from the interaction between the SN ejecta and a dense, H-rich circumstellar medium (CSM), as in SNe IIn (e.g., \\citealt{chu91,che94}). If this interpretation is correct, SN~2002ic may be the first SN Ia to show direct evidence of the circumstellar (CS) gas ejected by the progenitor system, presenting a unique opportunity to explore the CSM around an SN Ia and the nature of the progenitor system.  In this Letter, we present optical spectroscopy of SN~2002ic obtained more than 200 days after explosion, and discuss it within the context of the ejecta-CSM interaction. ",
        "conclusions": "\\label{sec-disc}  The currently preferred model for SNe Ia is the thermonuclear explosion of a C+O white dwarf (WD) in a binary system, reaching the Chandrasekhar limit via either accretion from a normal companion (the SD scenario, which is generally favored) or merging with another WD (the DD scenario). (For recent reviews, see \\citealt{nom00,liv00}.)  The SD scenario predicts the presence of an H/He-rich CSM. The discovery of strong CSM interaction in SN~2002ic may prove that this scenario does exist in nature. SNe~1997cy and 1999E may also be CS-interacting SNe Ia. However, such events are rare. Solid observational evidence of CSM has not yet been found in other SNe Ia \\citep{lun03}.  Our spectral analysis suggests a high-density \\ion{H}{1}-emitting CSM ($n \\sim 10^8-10^9$~cm$^{-3}$); so this is probably clumpy. Assuming the CSM was formed in a progenitor wind, we relate the mass-loss rate $\\dot{M}$ and the wind velocity $u$ to the H$\\alpha$ luminosity, through $L({\\rm H_\\alpha})\\sim 1-5\\times 10^{39}$ $(\\dot{M}/10^{-2}{\\rm~M_\\sun~yr^{-1}})^2$ $(u/100 {\\rm~km~s^{-1}})^{-2}$ $r_{16}^{-1} f^{-1}$ ergs~s$^{-1}$, where $r_{16}$ is the radius of the CSM shock in units of $10^{16}$ cm and $f$ is the CSM filling factor. For $u\\sim 100$ km s$^{-1}$ and $f\\sim 0.01$, $\\dot{M}$ can be as high as $\\sim 10^{-2}$ M$_\\sun$~yr$^{-1}$. The case B emissivity used here may underestimate $L({\\rm H_\\alpha})$ by $\\sim 50\\%$, considering that most H$\\beta$ photons may have cascaded to H$\\alpha$.  We have identified high-velocity lines (FWHM $\\gtrsim 10^4$ km$^{-1}$) emitted in the ejecta of SN~2002ic. Based on the hydrodynamical model of T. Suzuki et~al. (2004, in preparation), we suggest an equatorially concentrated structure for the CSM to explain the coexistence of these lines with the strong CS interaction. Such a geometry is not unexpected for the mass loss from a binary system or for stars approaching the end of the asymptotic giant branch (AGB). A preexisting clumpy disk was also suggested by \\citet{wan04}, based on spectropolarimetry.  Is $\\dot{M}\\sim 10^{-2}$ M$_\\sun$~yr$^{-1}$ too high for the SD scenario? The highest $\\dot{M}$ observed in AGB stars is between $10^{-3}$ and $10^{-4}$ M$_\\sun$~yr$^{-1}$ \\citep{ibe95}; $\\dot{M}$ in our estimate is scaled as $u$. The bulk of the CSM may be near the equator and at low velocity ($\\sim 10$ km~s$^{-1}$), which cannot be resolved. The observed H$\\alpha$ P Cygni, dominated by the emission component, could be produced by some high-velocity CSM ($\\gtrsim 100$ km~s$^{-1}$) along the pole, which should be close to our line of sight. Similar wind patterns have been observed in some symbiotic stars (e.g., \\citealt{sol85}), which have been suggested as candidates for SN Ia progenitors.  Further observations and modeling are required to understand the nature of the CSM giving rise to this type of event, which may be classified as ``Type IIa.'' Suggestions include a common envelope \\citep{liv03}, WD accretion wind \\citep{hac99}, and the superwind from an AGB star exploding as a ``type 1.5'' event \\citep{ham03}."
    },
    "0311/astro-ph0311523_arXiv.txt": {
        "abstract": "We apply the tidal truncation model proposed by \\citet{neg01} to arbitrary Be/compact star binaries to study the truncation efficiency dependance on the binary parameters. We find that the viscous decretion disks around the Be stars could be truncated very effectively in narrow systems. Combining this with the population synthesis results of Podsiadlowski, Rappaport and Han (2003) that binary black holes are most likely to be born in systems with orbital periods less than about 30 days, we suggest that most of the Be/black-hole binaries may be transient systems with very long quiescent states. This could explain the lack of observed Be/black-hole X-ray binaries. We also discuss the evolution of the Be/black-hole binaries and their possible observational features. ",
        "introduction": "X-ray binaries consist of a neutron star (NS) or a black hole (BH) accreting from strong stellar winds or via Roche lobe overflow of the companion star. Depending on the masses of the optical companions, they are conventionally divided into low- and high-mass X-ray binaries. By measuring the radial velocity curve of the non-degenerate donor, the value of the mass function can be determined, which provides a lower limit on the mass of the accreting object. When it is combined with the information of the spectra of the donor and the orbital light curves, actual measurements of the mass can be obtained. In this way, currently the masses of 18 compact objects in X-ray binaries have been found to exceed the the maximum mass possible for a neutron star. These binaries are thought to be black-hole X-ray binaries (BHXBs) (see \\citet{mcc03} and references therein).  Among the 18 BHXBs, 3 are persistently bright X-ray systems and 15 are X-ray novae in which 6 show recurrent outbursts. All of the X-ray novae are low-mass X-ray binaries (LMXBs). The three persistent sources, i.e. Cyg X$-1$, LMC X$-1$ and LMC X$-3$ are high-mass X-ray binaries (HMXBs) with massive O/B companion stars. In the most recent catalogue of high mass X-ray binaries (HMXBs) edited by \\citet{liu00}, only 20 out of 130 HMXBs are O/B supergiant systems including the black hole binaries. Though roughly two thirds of HMXBs are Be/X-ray binaries, and X-ray pulsations have been found in about sixty Be/X-ray binary systems \\citep{zio02}, there are no acknowledged Be/BH binaries. The existence of Be/BH binaries becomes an open problem. The solutions may be related to two questions: Do they exist? if yes, how can they be observed?  The first question is connected with the birth-rate estimation of BH/massive X-ray binaries (BMXBs). Using Monte-Carlo simulation, \\citet{rag99} calculated the number and distribution of binary BHs with Be stars. They obtained the expected number of Be/BH binaries to be of order of unity per $20-30$ Be/NS systems. The Be/BH binaries were found to be highly eccentric with orbital periods lying in a range of 10 days to several years. The actual number of Be/BH binaries could be even higher since in their work, the lower limit of $50 \\, M_{\\sun}$ for the progenitor of the BH was adopted, based on the mass of the supergiant Wray 977 \\citep{kap95}. But current stellar evolution models suggest that stars with mass higher than about $20-25\\, M_{\\sun}$ will result in BHs\\citep{fry99, fry01}. More recently \\citet{pod03} did the binary population synthesis calculation with the mass of the BH progenitors in the range of $25\\, M_{\\sun}<M_{\\rm p}<45\\, M_{\\sun}$, and presented the calculated distribution of the orbital period and donor mass for BH binaries in their Fig.2. It is shown that BHs have large possibility to be formed in relatively narrow systems with orbital periods of less than tens of days. The formation rate of BMXBs estimated from their result is about $10^{-5}\\ \\rm {yr}^{-1}$. Although the massive donor stars are not constrained to be Be stars in their work, the conclusion of the orbital distribution may also be applied to Be/BH binaries, since the fraction of the Be stars to the massive stars is not small (see \\citet{coe00} and references therein). Nevertheless, theoretical estimation of this fraction is difficult. It is not yet clear how Be stars eject the circumstellar envelope. While rapid rotation plays a part in disc creation, it is not the sole cause \\citep{por99}. Whether the physical properties of the donor star (metallicity, instability in the photosphere induced by binary rotation, and etc.) have effects on producing the dense disk winds is still unknown \\citep{zam01}.  The observation of supergiant B[e]/BH system, e.g. XTE J$0421+560$/CI Cam \\citep{rob02}, may provide a further proof of the existence of Be/BH binaries, since theoretically sgB[e]/X-ray binaries would have less probability to exist than Be X-ray binaries due to the much larger mass of sgB[e] stars. Thus, the absence of Be/black-hole binaries might be due to some observational selection effects. In this work, we try to find the clue of the second question, that is, why we can hardly observe them. We would start from the interaction between a Be star and its compact companion.  Be/X-ray binaries are characterized by their transient nature. Recently, theoretical work on disc truncation in Be/X-ray binaries \\citep{oka01,oka02} was invoked to explain the X-ray outbursts in Be/X-ray binaries, and has got considerable success in the application to several systems \\citep{oka01}, especially 4U 0115+63/V635 Cas \\citep{neg01} and A0535+26 \\citep{hai03}. The key idea is as follows: The neutron star exerts a net negative tidal torque on the viscous decretion disk of the Be star, diminishing the action of the viscous torque outside some critical radius, and thus resulting in the truncation of the disk. The disc matter would then accumulate in the outer rings of the disk till overcoming of the truncation by the effects of global one-armed oscillations or disk warping, etc. The following sudden infall of high density disk matter onto the neutron star causes type II X-ray outbursts. During the episodes between the outbursts, the neutron star could hardly be observed because of low mass accretion rate or the propeller effect. On the other hand, if the tidal truncation is not very efficient, and the disc can extend beyond the Roche lobe of the Be star at periastron, the matter could be accreted onto the neutron star during the periastron passage, resulting in (quasi-)periodic type I bursts.  It is interesting to see what would happen when the neutron star in the above picture is replaced by a black hole. We extend the application of this theory to arbitrary systems to find the influence of mass ratio, orbital period, eccentricity and viscosity coefficient on the efficiency of disc truncation. The model and the necessary parameters are briefly described in section 2. The results are shown in section 3. Our discussion and conclusions are made in section 4. ",
        "conclusions": "We have extended the application of Be disk truncation model proposed by \\citet{neg01} and \\citet{oka01} to general Be binary systems, to study the dependance of truncation efficiency on some basic binary parameters, such as the stellar mass, viscosity coefficient, orbital period and eccentricity. The results not only confirm the trend that low viscosity and small eccentricity would lead to effective Be disk truncation \\citep{oka01}, but also show that most effective truncation would occur in relatively narrow systems.  The above results can help us to understand the observations of Be binaries. In the context of Be/NS binaries, the observed transient features of the wide ($\\Porb\\sim 17-263\\ \\rm {d}$) eccentric ($e\\sim 0.1-0.9$) systems can be well understood by the disk truncation effect. Furthermore, the results also presents a possible explanation for the absence of Be/BH binaries if they are formed in relatively narrow orbits without high eccentricity, as suggested by \\citet{pod03}. The relatively narrow orbit of $\\Porb<30 \\, {\\rm d}$ is supported by the result of binary population synthesis simulation work of \\citet{pod03} with the assumption of circular orbit. The knowledge of the eccentricity could be traced back to the formation of the BHs in the massive binaries. The BHs were born from the collapse of massive stars in two different ways (see \\citet{fry99} and reference therein ). (i) The massive stars directly collapse into black holes. (ii) After supernova explosion, a significant part of the stellar matter falls back onto the proto-neutron star forming a black hole. If there were no kick in the latter case, the eccentricity induced by the supernova can be estimated by $e=\\Delta M/\\Mf$ \\citep{bha91}, where $\\Delta M$ is the total mass loss from the binary in the supernova explosion, and $\\Mf$ is the final mass of the binary. Assuming the mass of the secondary star change little during the process, one can expect less $\\Delta M$ and larger $\\Mf$ in black hole binaries, resulting in smaller eccentricity compared to neutron star binaries. The spatial distribution and the kinematics of most of the black-hole binaries appear to be consistent with the assumption that small asymmetric kicks are imparted to the black holes \\citep{whi96}, with at least one exception GRO J1655-40, the first black hole system which has evidence of runaway space velocity of $112\\pm 18 \\, \\rm {km\\, s^{-1}}$ in an orbit of $e=0.34\\pm 0.05$ \\citep{mir02}.  Because of very effective tidal truncation on Be disks in the relatively narrow and low eccentric systems, the flow of the matter towards the BH will be effectively blocked during almost the whole orbital cycle. In such episode, the accretion could be mainly from the polar wind of the donor star. As suggested by \\citet{wat89}, the polar wind of the Be star probably resembles that in OB stars. Such low-density high-velocity wind can hardly form accretion disks around the BHs, and the accretion would follow the classical Bondi-Hoyle-Littleton (BHL) approximation. For a typical main-sequence $15 \\, M_{\\sun}$ main-sequence star in the system of $\\Porb=10 \\, {\\rm d}$, if the mass loss rate $\\Mwdot=10^{-9} \\, M_{\\sun}\\ {\\rm yr}^{-1}$ and the terminal wind velocity $v_{\\infty}=2000 \\, {\\rm km \\, s}^{-1}$ exist, the accreting compact star would have luminosity of about $\\Lx\\sim 10^{33} \\, \\rm{erg \\, s^{-1}}$. For the Be/BH binaries, the luminosity would be still less, since the polar wind just takes a small fraction of the mass loss from the donor star. The wind plasma accumulates in the outer rings of the decretion disc till one-armed oscillaiton instability - probably responsible for the V/R variability seen in the Be stars - occurs, which causes the almost total disruption of the Be disk. The resulted large mass-infall onto the black hole would lead to very luminous outbursts. However, no aspect of the truncation model implies that the large mass transfer must occur \\citep{oka02}. If the collapsed disk due to the dynamical instability falls back on to the more massive Be star, such narrow and low-eccentricity transients would be almost out of detection even for the most sensitive instruments, so it should be careful to deal with the Be stars unrelated to the X-ray sources.  For systems in which effective truncation occurs, e.g. XTE J$1543-568$ and 2S $1553-542$, the burst activities are very rare (see \\citet{oka02} and reference therein). During the long quiescent stage, Be stars could be identified from the Balmer (and sometimes other) line emission and the associated strong infrared excess. In the HMXB catalogue of \\citet{liu00}, there are probably 24 Be/X-ray systems where the nature of the compact star and the orbital period are undetermined. They are signed as Be/X-ray binaries because of their highly variable X-ray characteristics analogous to those of the well studied Be/NS transients, or the optical identification of the donor stars. Most of them have been observed once during the burst states, and then disappeared from X-ray detection because of very low luminosity. Among these undetermined Be/X-ray systems, XTE J$1739-302$ and AX J$0052.9-7158$ (SMC 32) can be excluded now since the donor star of the former has been identified as an O supergiant \\citep{smi03} and 167.8 seconds pulsations have been found in the latter \\citep{yok01}. To look for Be/BH binaries in the left systems, optical observations in the quiescent stage become important, from which we can confirm the nature of the donor star and get information of the orbital period and the velocity curve of the donor. The latter two are useful in estimating the dynamical mass of the compact star - the common way to determine whether the compact star is a black hole, though the measurements would be very difficult in Be/X-ray binaries, because of not only the relatively low velocity of the donor but also the variability of the circumstellar envelope of Be stars \\citep{rag99}. The accurate locations of some of the undetermined Be/X-ray binaries could be derived with Chandra, which will enable following up observations in the infrared/optical wavelength. Till now, we have found very limited data of these faint Be/X-ray binaries, and there is at least one system GRS $1736-297$ went undetected even in the Chandra observations \\citep{wil03}.  Even if the X-ray outbursts due to the large mass inflow from the erupted disks are observed, it is still difficult to distinguish between black holes and neutron stars, unless pulsations have been found in the latter. Although there is growing evidence that black hole X-ray binaries display radio emission when they are in the low/hard X-ray state \\citep{reig03}, this is not a decisive criterion. LS I$+61{\\degr}303$ is an example, whether the compact star is a neutron star \\citep{zam95} or a magnetized black hole \\citep{pun99} is still in debate.  Be stars can be distinguished from normal O/B stars by the special dense disk wind structures when they are on the main sequence. Once they evolve to fill the Roche lobe, they can hardly be distinguished. Due to the high mass transfer rate from the Roche lobe overflow, black holes in HMXBs would be extremely luminous, even well above the Eddington limit $\\LEdd$ for a stellar-mass black hole, resulted from either the beaming effect \\citep{kin01} and/or being genuinely super-Eddington \\citep{beg02}, and thus become ones of the ultra luminous X-ray sources (ULXs). This can also be seen in fig. 6. of \\citet{pod03}.   Finally, we conclude that by investigating the disk truncation efficiency in Be/compact star binary systems, we propose a possible explanation for the absence of the observation of Be/BH binaries. We show that if most of the Be/BH binaries are born in relatively short low eccentric orbit, the interaction of the binary stars will lead to very effective truncation on the decretion disk around the Be star, so that most of the Be/BH systems would appear as very low luminous X-ray sources."
    },
    "0311/astro-ph0311009_arXiv.txt": {
        "abstract": "In the favored progenitor scenario, Type Ia supernovae (SNe~Ia) arise from a white dwarf accreting material from a non-degenerate companion star.  Soon after the white dwarf explodes, the ejected supernova material engulfs the companion star; two-dimensional hydrodynamical simulations by \\cite{Marietta} show that, in the interaction, the companion star carves out a conical hole of opening angle 30$^\\circ$-40$^\\circ$ in the supernova ejecta.  In this paper we use multi-dimensional Monte Carlo radiative transfer calculations to explore the observable consequences of an ejecta-hole asymmetry.  We calculate the variation of the spectrum, luminosity, and polarization with viewing angle for the aspherical supernova near maximum light. We find that the supernova looks normal from almost all viewing angles \\emph{except} when one looks almost directly down the hole.  In the latter case, one sees into the deeper, hotter layers of ejecta.  The supernova is relatively brighter and has a peculiar spectrum characterized by more highly ionized species, weaker absorption features, and lower absorption velocities.  The spectrum viewed down the hole is comparable to the class of SN~1991T-like supernovae.  We consider how the ejecta-hole asymmetry may explain the current spectropolarimetric observations of SNe~Ia, and suggest a few observational signatures of the geometry.  Finally, we discuss the variety currently seen in observed SNe~Ia and how an ejecta-hole asymmetry may fit in as one of several possible sources of diversity. ",
        "introduction": "\\subsection{Asymmetry of Type Ia Supernovae} Some Type~Ia supernovae (SNe~Ia) are known to be aspherical; direct evidence for this comes from optical spectropolarimetric observations. Because a spherically symmetric system has no preferred direction, the polarization integrated over the projected supernova surface cancels -- detection of a non-zero intrinsic polarization demands some degree of asymmetry.  The measured intrinsic polarization of SNe~Ia is relatively small, but certainly detected in a few cases.  Pre-maximum observations of the normal Type~Ia SN~2001el using the ESO Very Large Telescope found an intrinsic polarization level of $\\sim0.3\\%$, which decreased at later epochs \\citep{Wang_pol_01el, Kasen_01el}. Intrinsic polarization of $\\sim 0.7\\%$ was also measured for the underluminous and spectroscopically peculiar SN~1999by \\citep{Howell_99by}.  The geometry of SNe~Ia must be closely tied to the supernova explosion physics and progenitor system, both of which are still under debate. But little is known about the shape of the ejecta.  For both SN~1999by and SN~2001el we do know that the bulk of the ejecta obeyed a nearly axial symmetry.  This is because in both cases, after subtraction of the interstellar polarization, the polarization angle was fairly constant over the entire spectral range (with the exception of an unusual high-velocity calcium feature in SN~2001el).  Most theoretical attempts at modeling the spectropolarimetry have so far assumed the ejecta was ellipsoidal \\citep{Howell_99by,Wang_96x,Jeffery_87a,Hoeflich_91}.  A shape like this might arise, for example, in the explosion of a rapidly rotating progenitor star.  Another potential cause of asymmetry in SNe~Ia is the binary nature of the progenitor system.  In the favored progenitor scenario (the \\emph{single-degenerate scenario}; see \\cite{Branch_progenitor} and references therein), SNe~Ia arise from a white dwarf accreting material from a non-degenerate companion star.  The companion may be either a main sequence star, a red-giant, or a subgiant; as it is close enough to be in Roche-lobe overflow, it subtends a substantial solid angle from the perspective of the white dwarf.  The supernova explosion occurs when the white dwarf has accreted enough matter that the densities and temperatures at the center are sufficient to ignite carbon, just below the Chandrasekhar limit.  The ejected supernova material moves at a few percent of the speed of light and soon after the explosion (from minutes to hours) engulfs the companion star.  In the impact it would not be surprising if a substantial asymmetry was imprinted on the supernova ejecta.  The ejecta-companion interaction has been studied with two-dimensional hydrodynamical models by \\cite{Fryxell_Arnett}, \\cite{Livne_binary}, and most recently and extensively by \\cite{Marietta}.  These studies were primarily concerned with the fate of the companion star, in particular how much hydrogen gets stripped from its outer envelope. Stripped hydrogen may appear as narrow Balmer emission lines in the supernova spectrum, which if observed might provide direct evidence of a binary progenitor system.  With the advance of spectropolarimetric observations, however, the nature of SN~Ia asphericity becomes another relevant test of the single-degenerate progenitor scenario.  In their hydrodynamical models, \\cite{Marietta} find that the impact with the companion star carves out a conical hole in the supernova ejecta.  The opening angle of the hole is 30$^\\circ$-40$^\\circ$, and because the ejecta is moving supersonically, the hole does not close with time. The final configuration is axially symmetric, as was seen in the polarization observations of SN~2001el.  In this paper we use multi-dimensional radiative transfer calculations to address the possibility of SNe~Ia having an ejecta hole asymmetry. We calculate the variation of the spectrum, luminosity, and polarization with viewing angle for the aspherical supernova near maximum light.  In contrast to the ellipsoidal models, the angular variations in an ejecta-hole geometry can be rather extreme, especially when one looks near the hole itself.  These variations would necessarily introduce some diversity into the observed properties of SNe~Ia.  The question is: exactly what sort of diversity arises in the ejecta-hole geometry, and does this fit in with the diversity already known to exist in SNe~Ia?  While SNe~Ia are considered to be a rather homogeneous class of objects, they do show some variety in their spectral and photometric properties.  The observed peak magnitudes of SNe~Ia vary by $\\sim 0.3$ mag, and the brightness is found to correlate with the width of the light curve \\citep{Phillips93}.  The spectra of SNe~Ia can be classified as either normal or peculiar \\citep{Branch_normal}.  The peculiar spectra have feature strengths at maximum light that differ from ``normal'' cases (such as SN~1981B), and are usually subdivided into two classes: SN~1991bg-like supernovae have a broad \\ion{Ti}{2} absorption trough not seen in the normals \\citep{Filippenko_91bg}; SN~1991T-like supernovae have weak or absent features from singly ionized species but noticeable \\ion{Fe}{3} lines \\citep{Filippenko_91T,Phillips_91T,Jeffery_92}.  Not all supernovae fit cleanly into the classification scheme.  In its pre-maximum spectra, SN~1999aa resembled SN~1991T, but by maximum light it had begun to look much more normal, with \\ion{Si}{2} and \\ion{Ca}{2} lines that were stronger than SN~1991T but weaker than normal \\citep{Li_peculiar}.  As such SN~1999aa is considered by some to be an intermediate link between the normal and the SN~1991T-like supernovae. Other observations have uncovered singular objects like SN~2000cx \\citep{Li_00cx} and SN~2002cx \\citep{Li_02cx}, that while resembling SN~1991T in some ways (weak \\ion{Si}{2}, strong \\ion{Fe}{3} lines) showed other peculiarities that were unique.  Additional spectral diversities include the abnormally high photospheric velocities of SN~1984A \\citep{Branch_84A} and the detached, high velocity features seen in several supernovae \\citep{Hatano_94d,Wang_pol_01el,Thomas_00cx}.  The diversity of SNe~Ia is thus multi-faceted, a point we return to in the conclusion. ",
        "conclusions": "\\subsection{Asymmetry and Diversity in SNe~Ia}  Despite the seemingly extreme nature of an ejecta-hole asymmetry, we find that the geometry is actually consistent with what is currently known about SNe~Ia, at least for the observables we have calculated. The variation of the peak magnitude with viewing angle is $\\sim 0.2$ mag in $B$, comparable to the intrinsic dispersion of SNe~Ia, and the level of polarization in the range observed ($0-0.8\\%$).  The spectrum of the supernova looks peculiar when viewed near the hole, but this peculiarity may fit in with the spectral diversity already known to exist in SNe~Ia. In addition, the polarization spectrum from some lines of sight is a qualitative match to that of SN~2001el.  An ejecta-hole asymmetry could therefore be one source of diversity in SNe~Ia, but of course not the only one.  The primary source of diversity in SNe~Ia is thought to be due to variations in the amount of \\Ni produced in the explosion.  Because SNe~Ia are powered by the radioactive decay of \\Ni and its daughter $^{56}$Co, different \\Ni production can explain the variety in SN~Ia peak magnitudes.  Greater \\Ni masses may lead to higher atmospheric temperatures and higher effective opacities, which may explain why brighter SNe~Ia have broader light curves \\citep{Hoeflich_WLR,Pinto_Eastman_I,Nugent_hydro}.  It has often been thought that the spectroscopic diversity of SNe~Ia fits into the same one-parameter \\Ni sequence \\citep{Nugent_sequence}. In this picture, SN~1991T-like supernovae occupy the overluminous end of the sequence, where the larger \\Ni mass leads to higher envelope temperatures and a higher ionization fraction.  This may explain the peculiar spectral appearance \\citep{Mazzali_91T,Jeffery_92}.  As the models in this paper show, there could be a second, physically very different route to the same sort of spectral peculiarities -- one could be peering down an ejecta hole.  In the ejecta-hole model, the spectrum shows some level of peculiarity for $\\theta \\la \\theta_H$ or about 12\\% of the time, although the peculiarities will only be very intense for views more directly down the hole ($\\theta \\la \\theta_H/2$ or $\\sim 3\\%$ of the time).  The observed rate of SN~1991T-like supernovae is $\\sim 3-5$\\% in the samples of both \\cite{Branch_peculiar} and \\cite{Li_peculiar}; the rate of SN~1991T/SN~1999aa-like supernovae is $20\\% \\pm 7\\%$ in the \\cite{Li_peculiar} sample.  Therefore it is possible that a substantial percentage of these peculiar supernovae could be the result of an ejecta-hole asymmetry.  In this paper we have chosen to compare the spectra emanating from the hole with SN~1991T only because it is the well-known prototype of a certain kind of spectral peculiarity.  Whether SN~1991T itself was an example of merely looking down an ejecta hole is debatable.  Initial estimates suggested that SN~1991T was as much as 0.7-0.8 magnitudes brighter in $B$ than normal, which is too much to be explained by the asymmetry alone \\citep{Fisher_91T}.  More recent Cepheid measurements of the distance to the host galaxy, however, show that SN~1991T was not really much brighter than a normal SN~Ia.  \\citep{Saha_91T} find a moderate overluminosity of 0.3 mag, although a value as high 0.6 mag cannot be ruled out due to large uncertainty in the dust extinction. This lower value for the brightness of SN~1991T calls into question whether the peculiar spectral appearance can still be explained alone by high envelope temperatures due to a larger \\Ni mass.  SN~1991T also had a rather broad light curve (\\dmf = 0.95 $\\pm$ 0.05; \\cite{Phillips99}), which is often taken as an indication of a large \\Ni mass.  Because we have not yet computed time-dependent models, we do not know exactly what effect an ejecta hole asymmetry will have on the light curve.  Because the hole acts as an energy leak, it probably lowers the net diffusion time, and we expect that the \\emph{integrated} light curve (i.e that summed over all viewing angles) will be narrower in a ejecta-hole model than a spherical model.  But the real question relevant to SN~1991T is not how the integrated light curve compares to a spherical model, but whether the light curve viewed down the hole is broader or narrower than that from other viewing angles.  In other words we need to know how Figure~\\ref{mags} -- the distribution of the total luminosity among viewing angles -- varies with time.  This is more difficult to intuit, because as the ejecta thins out and the asymmetry and opacities evolve with time, it is hard to say off-hand whether it will become more or less easy for photons to preferentially escape out the hole.  We leave the question for future work.  In any case, although the prototype SN~1991T did have a broad light curve, it is not clear whether a general correlation between light curve width and SN~1991T-like spectral peculiarities even exists \\citep{Howell_01ay}.  Several SNe~Ia have similar or broader light curves, and yet the spectrum is apparently normal -- at least eight such supernovae with \\dmf $< 1.0$ are listed in \\cite{Phillips99}, for example SN~1992bc (\\dmf = 0.87 $\\pm$ 0.05) and SN~1994ae (\\dmf = 0.86 $\\pm$ 0.05).  SN~2001ay also had a normal spectrum but an exceptionally broad light curve (\\dmf = 0.6-0.7; \\cite{Phillips_01ay}).  Among the supernovae with SN~1991T-like spectral peculiarities, there also appears to be diversity.  SN~1997br had a moderately broad light curve (\\dmf $= 1.00 \\pm 0.15$; \\cite{Li_97br}), but the light curve of SN~2002cx was on the narrow side (\\dmf $= 1.30 \\pm 0.09$; \\cite{Li_02cx}).  In another SN~1991T-like supernova the $B$-band light curve was lopsided -- SN~2000cx brightened much faster than SN~1991T (resembling the rise of the normal SN~1994D) but the decline was slow (\\dmf = 0.93 $\\pm$ 0.04; \\cite{Li_00cx}).  The examples make it clear that the connection between light-curve width and SN~1991T-like spectral peculiarities remains vague, and that more than one parameter of diversity needs to be identified.  The nebular spectra of SN~1991T may also suggest a large \\Ni production.  In the late time spectra, the iron emission lines of SN~1991T have larger velocity widths than in most SNe~Ia \\citep{Mazzali_nebular}.  Assuming the late time ionization/excitation conditions are similar in all SNe~Ia, this implies that the nickel/iron core in SN~1991T is larger than normal. Confusing this conclusion, however, is the fact pointed out by \\cite{Hatano_97br} that the \\ion{Si}{2} velocities in the post maximum spectra are among the \\emph{lowest} of all SNe~Ia.  If SN~1991T really did have a large inner \\Ni zone, one naively expects the zone of silicon and other intermediate mass elements to occur at especially high velocities (as for instance in the delayed detonation models of e.g. \\cite{Hoeflich_99by}).  To account for the low \\ion{Si}{2} velocities, some have invoked a late-detonation model for SN~1991T, which produces a layer of intermediate mass elements sandwiched between two nickel zones \\citep{Yamaoka_91T,Pilar_91T}.  Of course a lower \\ion{Si}{2} velocity is also naturally expected if one is looking down an ejecta hole.  It is possible that SN~1991T did have a relatively large \\Ni mass, rather than (or perhaps in addition to) being viewed down the ejecta hole.  However among other supernovae with SN~1991T-like peculiarities there is a good deal of diversity, and the large \\Ni mass explanation will not apply in all cases.  The most obvious case in point is SN~2002cx \\citep{Li_02cx}.  The spectrum of SN~2002cx resembled SN~1991T in that \\ion{Si}{2}, \\ion{S}{2}, and \\ion{Ca}{2} lines were weak, while \\ion{Fe}{3} was prominent, but the supernova was \\emph{underluminous} by $\\sim2$ mag.  The velocities of the absorption features were also unusually low ($v \\approx 7.000$ \\kms) \\citep{Branch_02cx}.  The singularity of the supernova led \\cite{Li_02cx} to consider alternative progenitor systems, although they conclude that no existing theoretical model could explain all the peculiarities.  On the other hand, there is also the possibility that we are seeing multiple channels of diversity operating at once -- one scenario to entertain now is that we are looking down the ejecta hole of a ``weak'' supernova that produced a small mass of $^{56}$Ni.  Such underluminous objects (e.g. SN~1991bg) typically have relatively low absorption velocities \\citep{Turatto_91bg,Mazzali_91bg,Hatano_97br} which would be further reduced by looking down the hole.  Despite the low luminosity, the spectrum might still appear hot and iron dominated if one is peering into the deeper layers, and (eventually) into the iron core.  Of course, the chance of seeing two distinct sources of diversity operating at once would be, like SN~2002cx, a very rare occurrence.  Whatever the final explanation for SN~2002cx, its singularity highlights the fact that the diversity of SNe~Ia is more complicated than a one-parameter sequence based upon $^{56}$Ni.  One can imagine many sources of variation contributing, including: different amounts and distributions of high-velocity material, different sizes and inclinations of an ejecta hole (or other) asymmetry, various explosion processes (e.g. pure deflagration, deflagration-to-detonation, late-detonation), metallicity variations and even the possibility of completely different progenitor channels (e.g. a double-degenerate scenario).  With so much going on at once, interpreting the observations becomes a daunting task, but in some sense the identification of multiple channels of diversity is welcome, considering that observations continue to turn up unusual SNe~Ia.  \\subsection{Observational Consequence of an Ejecta Hole}  The results of this paper suggest a few observational signatures of the ejecta-hole geometry.  First, the continuum polarization should be low for views directly down the hole, where the spectrum looks peculiar.  However because of the partial obscuration effect, the polarization spectrum should show large line peaks for views just away the hole ($10^\\circ < \\theta < 60^\\circ$), where the spectrum looks marginally peculiar or normal.  For views from the side ($\\theta \\approx 90^\\circ$), a relatively high continuum polarization should be correlated with a slightly dimmer supernova with normal spectral features and inverted P-Cygni line polarization features.  Another possible signature of the ejecta hole is ``lopsided'' P-Cygni flux profiles -- the view down the hole weakens only the absorption, not the emission feature, so one could look for a weak (or absent) absorption associated with noticeable emission.  The easiest place to look would be in the \\ion{Si}{2} 6150 and the \\ion{Ca}{2} IR triplet features of SN~1991T-like supernova.  Unfortunately the relative strength of absorption to emission depends also on the line source function, which is determined by the detailed excitation conditions in the atmosphere.  In general, because we recognize that an ejecta-hole asymmetry is only one of several possible sources of diversity in SNe~Ia, it will be difficult to isolate the geometrical effects from the other variations that may be operating.  The only hope is to collect a large sample of supernovae with well observed light curves, spectra and polarization, so that one might try to pull out the different trends.  In our calculations we have used a parameterized hole (half opening angle $40^\\circ$) in order to explore the essential observable consequences of the geometry.  The next step is to address the same questions using specific hydrodynamical models representing a wide variety of progenitor configurations.  The details of the progenitor system could potentially affect the size and shape of the hole. \\cite{Marietta} compute interactions using main sequence, subgiant and red giant companions and note that the variation in the hole asymmetry is not large.  This is because in all cases the companion star is near enough to have undergone Roche lobe overflow and always occupies a similar solid angle (the red giant is farther away but physically larger than a main sequence companion which is smaller but much closer).  However if the ratio of companion radius to separation distance is decreased for some reason, the size of the hole also decreases.  A larger sample of spectropolarimetric observations will help determine if SNe~Ia really do have an ejecta-hole geometry, and could constrain the hole opening angle if one exists.  While a hole smaller than $\\theta < 20^\\circ$ has only minor effects on the spectrum, luminosity and continuum polarization, it will still create substantial line peaks in the polarization spectrum when seen from some viewing angles.  If such signatures of the hole are not seen in future spectropolarimetric observations, this would have interesting consequences for the progenitors of SNe~Ia, or the hydrodynamics of the ejecta/companion interaction.  Finally, we mention that an asymmetry like an ejecta hole could have a number of subtle consequences on the use of SNe~Ia as standard candles for cosmology.  The asymmetry causes a $\\sim 20\\%$ dispersion in observed SNe~Ia peak magnitude.  If the asymmetry is identical in all supernova, this dispersion behaves like a statistical error (although a non-gaussian one) and can be averaged out by observing enough objects. The averaging out is not achieved, however, if one does not sufficiently sample \\emph{every} possible viewing angle, either because not enough supernovae are observed, or because those viewed down the hole are withheld from the sample due to concern over their spectral peculiarities.  In addition, if the nature, degree, or frequency of the asymmetry evolves with redshift (say because of evolving progenitor populations) the peak magnitude of SNe~Ia becomes a function of redshift.  One might also be concerned that the significant angular variation of the colors and spectrum may complicate extinction and K-corrections. The errors incurred from all these and other related systematic effects would be relatively small, but may need to be considered in the next generation of precision supernova cosmology experiments."
    },
    "0311/astro-ph0311379_arXiv.txt": {
        "abstract": "We analyse the observed correlation between galaxy environment and H$\\alpha$ emission line strength, using volume-limited samples and group catalogues of 24968 galaxies at $0.05<z<0.095$, drawn from the 2dF Galaxy Redshift Survey ($M_{b_J}<-19.5$) and the Sloan Digital Sky Survey ($M_r<-20.6$).  We characterise the environment by 1) $\\Sigma_5$, the surface number density of galaxies determined by the projected distance to the 5$^{\\rm th}$ nearest neighbour; and 2) $\\rho_{1.1}$ and $\\rho_{5.5}$, three-dimensional density estimates obtained by convolving the galaxy distribution with Gaussian kernels of dispersion 1.1 Mpc and 5.5 Mpc, respectively.  We find that star-forming and quiescent galaxies form two distinct populations, as characterised by their H$\\alpha$ equivalent width, \\ewha.  The relative numbers of star-forming and quiescent galaxies varies strongly and continuously with local density.  However, the distribution of \\ewha\\ amongst the star-forming population is independent of environment. The fraction of star-forming galaxies shows strong sensitivity to the density on large scales, $\\rho_{5.5}$, which is likely independent of the trend with local density, $\\rho_{1.1}$. We use two differently-selected group catalogues to demonstrate that the correlation with galaxy density is approximately independent of group velocity dispersion, for $\\sigma=200$--$1000\\kms$. Even in the lowest density environments, no more than $\\sim 70$ per cent of galaxies show significant H$\\alpha$ emission. Based on these results, we conclude that the present-day correlation between star formation rate and environment is a result of short-timescale mechanisms that take place preferentially at high redshift,  such as starbursts induced by galaxy-galaxy   interactions. ",
        "introduction": "The galaxy population today can be effectively described as a combination of two distinct types.  The first are red, morphologically early-type galaxies with little or no current star formation; the remainder are blue, late-type galaxies with active star formation.  This segregation has been known for a long time; however the superb data from the Sloan Digital Sky Survey has revealed how surprisingly distinct these two populations are, at least in terms of their colours \\citep{Strateva01_short,Baldry03}.  This segregation is known to be strongly mass-dependent, with the most massive galaxies being predominantly red, early-types \\citep{Kauffmann-SDSS1_short}. It is currently unclear whether galaxy evolution at the present day consists of galaxies changing from one type to another in a short time, or of evolution in properties within a given class.  Direct evidence of galaxy evolution comes from observations of galaxies at different redshifts;  this shows that, in the past,  the average star formation rate (SFR) was much higher \\citep{L96,madau,Wilson+02}, and the typical star-forming galaxy was more luminous \\citep{Cowie+99}.  Given the near-constancy of the SFR in the Milky Way \\citep[e.g.][]{R-P}, it seems likely that this downsizing effect is the main characteristic of the decline in global star formation: the most massive galaxies have recently stopped forming stars altogether, while less massive galaxies continue unhindered.  A more indirect form of evolution is observed as the change in galaxy populations as a function of their environment at a given epoch.  In particular, galaxies in dense environments (i.e. clusters) tend to have early-type morphologies \\citep[e.g.][]{Dressler,DML,Treu03} and low SFRs \\citep{B+97,B+98,PSG,P+99}.  One interpretation of this trend has been that the cluster environment causes galaxies to transform their properties as they move, pulled by gravity, from low density regions into the cluster centre.  However, recent work \\citep[e.g.][]{kap2,OHely,Kodama_cl0939}, especially that based on  the 2dF Galaxy Redshift Survey \\citep[2dFGRS, ][hereafter Paper~I]{MS03,2dF-sfr_short} and Sloan Digital Sky Survey \\citep[SDSS, ][hereafter Paper~II]{Sloan_sfr_short} has shown that the correlation between galaxy type and local density extends to very low densities, well beyond the region where the cluster is expected to have much influence. These works showed that there is a smooth dependence of SFR on local galaxy density, and identified a ``critical'' surface density of 1 Mpc$^{-2}$, where SFR correlations with environment first occur. This critical density is quite low, and corresponds to regions well outside the virialised cluster region; this suggests that galaxies may be pre-processed in groups, before they end up in clusters.  This had been anticipated \\citep[e.g.][]{ZM98,ZM00}, since galaxy-galaxy interactions are known to induce star formation \\citep[e.g.][]{Barton,Lambas} and such interactions should be common in groups.  There is evidence that these interactions lead to the build-up of elliptical galaxies with hot X-ray halos which will later be incorporated into clusters \\citep[e.g.][]{PBnat,MZ99,SNDS}. Groups are also expected to be important because they are the first level of the super-galactic hierarchy.  Models of galaxy formation assume the rate at which gas falls onto a galaxy depends on its environment \\citep[e.g.][]{Cole2000,HS03}.  In these models, satellite galaxies do not have their own supply of hot gas to replenish gas in the disk used to form stars.  Thus, star formation will gradually decline in any galaxy which is part of a larger halo \\citep{infall,Diaferio}; groups will be the first environment to demonstrate this effect.  \\subsection{Previous Work} Our understanding of galaxies in groups and lower density environments has been hindered by the difficulties in obtaining large, unbiased samples. Most of the work has been restricted to compact groups \\citep[e.g.][]{Hickson,RHF,IPV,Coziol,N+00,VM01,dlR01,KF03}, which are most likely to be physically bound systems.  However, these may be a special class of group, not representative of a typical stage of the hierarchy through which most galaxies evolve.  \\citet{PG84} analysed the morphology-density relation in groups selected from the CfA redshift survey \\citep{Cfa3} and demonstrated that the original relation found in clusters by \\citet{Dressler} extends continuously to lower densities.  Below a density of $\\sim 1$ Mpc$^{-3}$ ($M_{B(0)}<-17.5$) no further correlation is seen.  Similar trends have since been confirmed in numerous other studies \\citep[e.g.][]{ZM98,HO99,Tran,Treu03} .  Interestingly, \\citet{PG84} found the trend with local density is the same in both rich clusters and poor groups; the only difference is that poor groups typically sample lower density regions and, thus, have populations more dominated by late-type galaxies.  It has been argued based on this evidence that galaxy-galaxy interactions within the group environment are the mechanism responsible for the high fractions of early-type galaxies in clusters \\citep[e.g.][]{ZM98,ZM00,HO00}.  Less work has been done on the stellar populations of group galaxies, though there is good evidence that they are intermediate between those of the field and rich clusters \\citep{A+93,H+98,G+02,Tran,CNOC_groups2_short}. More recently, analysis of SFRs in galaxy groups selected from  the partially-completed 2dFGRS survey has been done by \\citet{MZ2dF}.  Based on this group catalogue, \\citet{Martinez} found that even the lowest-mass groups, with masses $\\sim10^{13}M_\\odot$, show reduced total SFRs relative to the field.   \\citet{D+02} claim that the spectral-type dependence on local density is only observed in groups more massive than $3\\times10^{13}M_\\odot$; however, this claim appears to arise mostly from the fact that their low-mass groups do not sample densities as high as found in the higher-mass groups.  \\subsection{The purpose of this paper} We analyse the local correlation between star-formation activity and environment using data obtained from the two largest galaxy surveys ever conducted: the 2dFGRS \\citep{2dF_colless,2dF_final} and the SDSS \\citep{SDSS_tech_short}.  In particular, we have selected a sample of galaxy groups from each survey, but in significantly different ways.  Since the definition of a group is partly subjective, and a variety of algorithms have been developed to find groups in redshift surveys, comparing results from the two catalogues allows us to investigate the sensitivity of our results to the way the group catalogue is constructed.  The purpose of this paper is to establish how star formation in the galaxy population, as characterised by the distribution of H$\\alpha$ emission, depends on environment.  Specifically, we will investigate whether the most important variable is 1) velocity dispersion of the embedding group or cluster; 2) local galaxy density, on scales $\\lesssim 1$ Mpc; or 3) large-scale structure, as parameterised by the density on $\\sim 5$ Mpc scales.  We will show that the number of actively star-forming galaxies depends on both the local and large-scale density, but that the properties of the star-forming galaxies themselves do not.  A summary of the paper follows.  In Section~\\ref{sec-data} we describe the details of the two datasets analysed herein, focusing on the construction of group catalogues, measurement of emission line strengths, and sample selection. Our results are presented in Section~\\ref{sec-results}.  We discuss the physical implications of our results in Section~\\ref{sec-discuss}, and draw some conclusions in Section~\\ref{sec-conc}. Throughout the paper we use a cosmology of $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=100h\\kmsmpc$, with $h=0.7$.  All distances are proper lengths in units of Mpc. ",
        "conclusions": "\\label{sec-conc} We have made a joint analysis of 24968 galaxies selected from the SDSS and 2dFGRS.  The distribution of \\ewha\\ among these galaxies is bimodal, consisting of an active population of galaxies with \\ewha$>4$ \\AA, and a quiescent population with no significant star formation at the present day.  We have investigated how this distribution depends on environment, as characterised by: \\begin{itemize} \\item The projected surface number density of galaxies $\\Sigma_5$, determined from the distance to the fifth-nearest neighbour. \\item The three-dimensional number density of galaxies measured on 1.1 Mpc and 5.5 Mpc scales, $\\rho_{1.1}$ and $\\rho_{5.5}$. \\item The velocity dispersion of the embedding structure, as determined from the group catalogues of \\citet{Eke-groups} and Nichol, Miller et al. (in preparation). \\end{itemize} We use these different measures of environment to establish the scales and structures on which the present-day galaxy population depends.  Our findings are summarised as follows:  \\begin{itemize} \\item[1.] The distribution of H$\\alpha$ line strength for the star-forming population, selected on \\ewha\\ or $(g-r)$ colour, does not itself depend strongly on environment.  Thus, it is unlikely that SFRs are gradually decreasing in a substantial number of star-forming galaxies in or near dense environments today. \\item[2.] The fraction of galaxies with \\ewha$>4$\\AA\\ decreases steadily with increasing local density.  There is evidence that it decreases more strongly at densities exceeding $\\Sigma_5\\sim1$ Mpc$^{-2}$, or $\\rho_{1.1}\\gtrsim 0.05$ Mpc$^{-3}$.  The persistence of this correlation at low densities means that ram--pressure stripping, at any redshift, cannot be the only  physical  mechanism at work. \\item[3.] The fraction of galaxies brighter than $M^\\ast+1$ with \\ewha$>4$\\AA\\ is never more than $\\sim 70$ per cent, even in the least dense environments explored here.  We have shown that this means the recent decline in globally-averaged star formation rate cannot be wholly due to the growth of large scale structure. \\item[4.] The emission-line fraction of a galaxy population appears to depend both on the local environment (on $\\sim 1$ Mpc scales) {\\it and} on the large-scale structure as parameterised by the density $\\sim 5.5$ Mpc scales. There is little further dependence on the velocity dispersion of the group or cluster in which the galaxy is embedded.  This result suggests that the composition of the galaxy population today is likely related {\\it indirectly} to its present environment \\citep[see also][]{ZM98,ZM00}. \\end{itemize}  The most likely physical explanations for the correlation between \\ewha\\ distribution and environment at $z=0$ are those which are effective over a large range of environment, affect the SFR on short ($<1$ Gyr) timescales, and were much more effective in the past.  As suggested by previous authors \\citep[e.g.][]{Z+96,MZ98,H+98,HO00}, one good candidate is starbursts induced by galaxy interactions, since close pairs of galaxies are the only environment known to directly provoke a physical transformation.  Such interactions will likely be more common at high redshift, and will have had more time to influence galaxies that end up in high density environments."
    },
    "0311/astro-ph0311465_arXiv.txt": {
        "abstract": "A number of very small isolated \\HII\\ regions have been discovered at projected distances up to 30 kpc from their nearest galaxy. These \\HII\\ regions appear as tiny emission line objects in narrow band images obtained by the NOAO Survey for Ionization in Neutral Gas Galaxies (SINGG). We present spectroscopic confirmation of four isolated \\HII\\ regions in two systems, both systems have tidal \\HI\\ features. The results are consistent with stars forming in interactive debris due to cloud-cloud collisions. The \\Ha\\ luminosities of the isolated \\HII\\ regions are equivalent to the ionizing flux of only a few O stars each. They are most likely ionized by stars formed in situ, and represent atypical star formation in the low density environment of the outer parts of galaxies. A small but finite intergalactic star formation rate will enrich and ionize the surrounding medium. In one system, NGC 1533, we calculate a star formation rate of $1.5\\times10^{-3}$ \\msun yr$^{-1}$, resulting in a metal enrichment of $\\sim1\\times10^{-3}$ solar for the continuous formation of stars. Such systems may have been more common in the past and a similar enrichment level is measured for the `metallicity floor' in damped Lyman-$\\alpha$ absorption systems. ",
        "introduction": "\\label{sect:intro}  \\HII\\ regions signifying the presence of highly ionizing OB stars are usually found in the luminous inner regions of galaxies \\citep[e.g.][]{Martin01}. \\HII\\ regions are also located in the faint outer arms of spirals \\citep[e.g.][]{Ferguson98b}, and as single or multiple star forming knots in narrow emission line dwarfs (\\HII\\ galaxies). In each case, new stars are formed in the vicinity of an existing stellar population. However, recent observations by \\cite{Gerhard02} have spectroscopically confirmed an isolated compact \\HII\\ region on the extreme outskirts of a galaxy (NGC 4388) in the Virgo cluster. Several luminous \\Ha-emitting knots have also been discovered in a compact group in the A1367 cluster \\citep{Sakai02}. In these cases it appears that the \\HII\\ regions are due to newly formed stars where no stars existed previously, albeit in a galaxy cluster environment.  As they evolve, OB stars increase the metal abundance in their local environment. Absorption line studies show that the intergalactic medium (IGM) and galaxy halos, including our own, are enriched \\citep[e.g.][]{Chen01,Tripp02,Collins03}. Isolated \\HII\\ regions provide a potential source for this enrichment. In situ star formation in the IGM offers an alternative to galactic wind models to explain metal enrichment hundreds of kilo-parsecs from the nearest galaxy.  Here we present a number of very small isolated \\HII\\ regions that have been discovered by their \\Ha\\ emission in the narrow band images obtained by the NOAO Survey for Ionization in Neutral Gas Galaxies (SINGG). SINGG is an \\Ha\\ survey of an \\HI-selected sample of nearby galaxies. The survey is composed of nearly 500 galaxies from the \\HI\\ Parkes All-Sky Survey (HIPASS, Barnes et al.  \\citeyear{Barnes01}; Meyer et al.  \\citeyear{Meyer03}), of these about 300 have been observed in \\Ha.  Since a gaseous reservoir is a prerequisite for star formation, SINGG measures a broad census of star formation in the local Universe.  The \\HII\\ regions appear as tiny emission line objects at projected distances up to 30 kpc from the apparent host galaxy. \\HII\\ regions are defined as isolated when they are projected at least twice the $\\mu_R =$ 25 ${\\rm mag\\, arcsec^{-2}}$ isophotal radius from the apparent host galaxy. This is typically much further than outer disk \\HII\\ regions in spiral galaxies \\citep{Ferguson98b}. In fact, for the systems discussed in detail here it is not totally clear whether the isolated \\HII\\ regions are even bound to the apparent host, hence we refer to them as ``intergalactic''.  Their high equivalent widths suggest they are due to newly formed stars where no stars existed previously.  In Section~2 spectra is presented for five isolated \\HII\\ regions candidates in three systems. These five candiates are referred to as the spectroscopically {\\it{detected}} emission line objects or \\HII\\ region candidates. All but one source has \\Ha\\ detected at a comparable recessional velocity to the nearest galaxy, two sources are also detected in \\oiii\\ as further confirmation.  These four objects are referred to as the spectroscopically {\\it{confirmed}} isolated \\HII\\ regions. Optical spectra and \\HI\\ distributions for all three systems are described in Sections~3$-$5. In Section~6 models of the underlying stellar population, scenarios for star formation, enrichment of the IGM, implications of the intergalactic star formation rate, and the possibility that isolated \\HII\\ regions are progenitors of tidal dwarf galaxies are discussed. H$_0$=75 \\kms Mpc$^{-1}$ is used throughout. ",
        "conclusions": "\\label{sect:conclusion}  The discovery of intergalactic \\HII\\ regions presented here and in other recent publications provides a small but finite source of enrichment and ionization of the IGM. In two cases the fact that these emission line objects are detected in both \\Ha\\ and \\oiii\\ rules out the possibility that they are background emitters. The \\Ha\\ luminosities imply that each isolated \\HII\\ region is ionized by 4$-$7 O stars. If these stars have formed in situ they represent atypical star formation in a low density environment. The low level of continuum emission from three of four confirmed isolated \\HII\\ regions suggests the stellar populations are very young and have formed where no stars existed previously. If part of a normal IMF, the corresponding total cluster mass would be $\\sim10^{3}$ \\msun. In two out of three systems, isolated \\HII\\ regions are associated with tidal \\HI\\ features, providing a reservoir of neutral gas. In one particular system, NGC~1533, the mass, distribution and velocity dispersion of the \\HI\\ suggests the rate of star formation ($1.5\\times10^{-3}$ \\msun yr$^{-1}$) could be sustained by the collision of clouds. This would result an increase in the metal abundance by $\\sim1\\times10^{-3}$ solar. This is the same abundance level as seen in the DLA `metallicity floor' \\citep{Prochaska03}. The amount of intergalactic high column density \\HI\\ and rate of collision-triggered intergalactic star formation may have been higher in the past. On-going investigations into the metallicities and underlying stellar population of these and other isolated \\HII\\ regions in the SINGG images will shed more light on their nature and origin."
    },
    "0311/astro-ph0311186_arXiv.txt": {
        "abstract": "Subarcsecond infrared and radio observations yield important information about the formation of super star clusters from their surrounding gas. We discuss the general properties of ionized and molecular gas near young, forming SSCs, as illustrated by the prototypical young, forming super star cluster nebula in the dwarf galaxy NGC~5253. This super star cluster appears to have a gravitationally bound nebula, and the lack of molecular gas suggests a very high star formation efficiency, consistent with the formation of a large, bound cluster. ",
        "introduction": "While the first suggestions that young analogues to globular clusters exist in the local universe came from ground-based imaging (Arp \\& Sandage 1985), the age of the super star cluster really dawned with the subarcsecond imaging capabilities of HST, which allows super star clusters (SSCs) to be resolved in nearby galaxies (e.g., O'Connell et al. 1995, Whitmore \\& Schweizer 1995) The Antennae galaxies alone have thousands of young (Myr) SSCs (Whitmore et al. 1999.) Some of these clusters are revealed only by infrared imaging (Vigroux et al. 1996); presumably these are the youngest regions.  Clues to the mystery of how large and potentially bound clusters form lie in the environments of the youngest, embedded regions, which are visible only in the radio and infrared.  That large extragalactic compact HII regions exist and can be detected was first inferred for the starburst in M83 by Turner, Ho, \\& Beck (1987) on the basis of a high Brackett line/radio continuum ratio, although their 8\\arcsec\\ (160 pc at M83) infrared beam did not allow them to define the nature of the unexpected Brackett excess (cf. Thompson 1987). With arcsecond mid-infrared imaging, Telesco and Gezari (1992) isolated a compact 12 $\\mu$m source that they suggested was a forming globular cluster in M82. Pina et al. (1992) and Keto et al. (1993) found a similar strong and compact mid-infrared source in the starburst galaxy, NGC~253.  High spatial resolution IR imaging and spectroscopy with HST and large ground-based telescopes and the expanded high frequency imaging capability at the VLA now allow the study of nebulae in the infrared and radio at subarcsecond resolution, corresponding to scales of $\\sim$1 pc in local galaxies. With this resolution one can now study the star formation processes in individual regions. Due to time and space limitations we will focus on the impact of these observations on issues in the formation of large bound clusters as exemplified by the well-studied  young embedded SSC in the dwarf galaxy NGC 5253. ",
        "conclusions": "High resolution radio and infrared observations in nearby galaxies have found luminous and compact nebulae excited by obscured super star clusters.  The IR emission of these nebulae can be a significant, or even dominant, part of the total galactic flux.  In NGC 5253,  there are so many young stars in such a small volume that the gravitational attraction appears to be slowing or stopping the rapid expansion of the nebula. The star formation efficiency also appears to be quite high in this galaxy, $\\sim 60-75$\\%, which is consistent with the formation of bound clusters.  Star formation in SSC-starburst galaxies does not much resemble that closer to home.  \\begin{figure} \\plotone{turner_j_fig3.eps} \\caption{CO in NGC 5253, from Meier et al. 2002. The white cross marks the supernebula. There are 5 clouds of gas and dust in this figure, marked A through E. Only cloud D is resident within the galaxy, while the other clouds are falling in. Cloud D has a mass of $\\sim 4 \\times 10^6~ \\rm M_\\odot.$} \\end{figure}"
    },
    "0311/astro-ph0311192_arXiv.txt": {
        "abstract": "We present high resolution ($\\leq 20$~mas) observations of the radio continuum emission at 1.4 GHz from three high-redshift quasars: J1053$-$0016 ($z=4.29$), 1235$-$0003 ($z=4.69$), and J0913$+$5919 ($z=5.11$), thereby doubling the number of $z>4$ radio-loud quasars that have been imaged at mas resolution. The observations were carried out with the Very Long Baseline Array (VLBA) of the NRAO.  All three sources are unresolved in these observations, with source size limits of a few mas. In all cases the flux densities measured by the VLBA are within 10$\\%$ of those measured with the VLA, implying that the sources are not highly variable on yearly timescales.  We find no indication for multiple images that might be produced by strong gravitational lensing on scales from 20 mas (VLBA) to a few arcseconds (VLA), to dynamic range limits of $\\sim 100$. ",
        "introduction": "Surveys such as the Sloan Digital Sky Survey (SDSS) \\citep{YOR00} and the Digitized Palomar Sky Survey \\citep{DJO99} have revealed large samples of quasi-stellar objects (QSOs) out to $z \\sim 6$. Studies by \\citet{FAN02} have shown that at such a high redshift we are approaching the epoch of reionization, the edge of the ``dark ages'', when the first stars and massive black holes were formed. Eddington-limit arguments suggest that the supermassive black holes at the center of these QSOs are on the order of $10^9~M{_\\odot}$. If the correlation between bulge and black-hole masses \\citep{GEB00,FEME00} also holds at these high redshifts, then these sources have associated spheroids with masses on the order of $\\sim 10^{12}~M_\\odot$. It is challenging to explain the formation of such massive structures on relatively short timescales {($\\sim$~1~Gyr)}. However, \\citet{WL02} estimate that almost one third of known quasars at $z\\sim6$ ought to be lensed by galaxies along the line of sight. If these quasars are indeed gravitationally lensed, the estimated masses of their associated spheroids could be smaller by up to an order of magnitude; this would allow a less efficient assembly process.  High resolution radio observations of high-redshift radio-loud quasars can be used to test for strong gravitational lensing by looking for multiple imaging on scales from tens of milliarcseconds (mas) to arcseconds. Also, Very Long Baseline Interferometry (VLBI) observations of core-jet radio sources in quasars over a large range in redshift have been used to constrain the cosmic geometry, under the assumption that such sources are (roughly) 'standard rulers' \\citep{GKF99}. In general, the high resolution of the VLBI observations permits a more detailed look at the physical structures in the most distant cosmic sources.  To date, only three radio loud quasars at $z > 4$ have been imaged at mas resolution \\citep{FRE97,FRE03}. In this paper we present Very Long Baseline Array (VLBA) observations of three more quasars at $z> 4$: SDSSp J105320.42$-$001649.7 at $z=4.29$, SDSSp J123503.04$-$000331.7 at $z=4.69$, and SDSSp J091316.56$+$591921.5 at $z=5.11$ (hereafter J1053$-$0016, J1235$-$0003, and J0913$+$5919). J1053$-$0016 was first identified by \\citet{IMH91} and \\citet{STD94} as BRI 1050$-$0000, while the other two sources were discovered in the five-color imaging data from the SDSS \\citep{FAN00,AND01}. All three quasars are known to be radio-loud \\citep{CAL01,PET03}  Throughout this paper, we assume a flat cosmological model with $\\Omega_{m}=0.3$, $\\Omega_\\Lambda=0.7$, and ${H_{0}=65}$~km~s$^{-1}$~Mpc$^{-1}$. In this model, 1~mas corresponds to about 7~pc at $z=4.5$. ",
        "conclusions": "The VLBA imaging of three $z> 4$ radio loud quasars (J1053$-$0016, J1235$-$0003, and J0913$+$5919) shows that these sources are smaller than a few mas in size, corresponding to physical scales of $\\le 20$ pc. On the other hand, all three sources show falling spectra at high frequency, and none of the sources are variable on yearly timescales.  These results suggests that the sources are likely core-jet radio sources, and hence may be used to extend to very high redshifts the studies of cosmic geometry using core-jet radio sources (Gurvits et al. 1999). Higher frequency VLBA observations at higher spatial resolution are planned to test this possibility.  We find no indication of multiple radio components in the fields of these sources on scales of 20 mas to a few arcseconds, to dynamic range limits of $\\sim 100$.  A similar conclusion was reached for three other $z>4$ radio loud quasars \\citep{FRE97,FRE03}. These results imply that at least these six high $z$ quasars are not strongly gravitationally lensed.  The compact nature of these quasars make them excellent candidates for future H~{\\footnotesize I} 21~cm absorption experiments to detect the neutral IGM in their host galaxies \\citep{FL02}, in particular the steep spectrum source J0913$+$5919. Such a search is currently under-way using the Giant Meter Wave Radio Telescope. Knowledge of source structure, as presented herein, is critical for both identifying potential candidates for H~{\\footnotesize I} 21~cm absorption searches, and for subsequent interpretation of the results."
    },
    "0311/astro-ph0311471_arXiv.txt": {
        "abstract": "It is demonstrated that not only gravity, but also neutrostriction forces due to optical potential created by coherent elastic neutron-neutron scattering can hold a neutron star together. The effect of these forces on mass, radius and structure of the neutron star is estimated. ",
        "introduction": "Interaction of neutrons with matter at low energies is characterized by optical potential \\begin{equation}\\label{1} V_o(\\rr)=\\frac{\\hbar^2}{2m}4\\pi n(\\rr)b, \\end{equation} where $m$ is neutron mass, $n(\\rr)$ is atomic density at a point $\\rr$, and $b$ is coherent s-wave neutron-nucleus scattering amplitude (see, for example~\\cite{ig0,ign}). It is important that though the amplitude $b$ is the result of short range strong interactions, the potential $V_o$ is the long range one.  Neutron-nucleus scattering amplitude $b$ is of the order of several fm, the density of matter in the earth conditions is of the order $10^{23}$ cm$^{-3}$, so the optical potential of matter is of the order $10^{-7}$ eV.  Interaction of neutrons with neutron matter is also described by eq. (\\ref{1}), and since the amplitude $b$ is negative, the potential (\\ref{1}) of the neutron matter is attractive. The attractive force $F_o=-\\na V_o(\\rr)$ is called neutrostriction.  Neutron-neutron s-wave scattering can take place only in singlet state, and the singlet amplitude $b_s$ at low energies is $b_s\\approx -18$ fm~\\cite{huhn}. The coherent amplitude obtained by averaging over all possible spin states of two neutrons is 4 times lower, therefore $b=-4.5$ fm. For such values of $b$ at star densities larger than $10^{36}$ neutrons/cm$^{3}$ the potential $V_o$ is larger than 1 MeV.  Let's compare the total gravitational and optical energies for a star of radius $R$, mass $M$ and uniform density \\begin{equation}\\label{1a} n=\\frac{M}{(4\\pi/3)R^3m}. \\end{equation} The total gravitational energy is \\begin{equation}\\label{1b} U_g=\\frac35G\\frac{M^2}{R}, \\end{equation} where $G$ is gravitational constant, while the optical energy is \\begin{equation}\\label{1c} U_{opt}=\\frac{\\hbar^2}{2m}4\\pi nbN=3\\frac{\\hbar^2}{2m^3}bM^2/R^3, \\end{equation} where $N$ is the total number of neutrons in the star. For amplitude $|b|=4.5$ fm we have $U_{opt}>U_g$, when~\\cite{conf} \\begin{equation}\\label{1d} R<R_0=\\sqrt{\\frac{5\\hbar^2|b|}{2m^3G}}=\\frac{\\hbar}{mc} \\sqrt{\\frac{5mc^2|b|}{2Gm^2}}\\approx 20{\\rm\\ km}, \\end{equation} which is independent of the star mass $M$. Using the parameter $R_0$ one can rewrite (\\ref{1c}) as \\begin{equation}\\label{1cc} U_{opt}=\\frac35GR_0^2\\frac{M^2}{R^3}. \\end{equation}  However, the s-wave amplitude $b_s$ is a constant only at low energies. In neutron stars, where energies are rather high, energy dependence of $b$ has to be taken into account. This dependence is well described by the theory of effective radius \\begin{equation}\\label{1d2a} \\frac{1}{b_s(E)}=\\frac{1}{b_s(0)}\\left(1-\\frac{1}{2}k^2ab_s(0)\\right), \\end{equation} where $a$ is the effective radius of neutron nucleus interaction, $a=1.2$ fm, $b_s(0)=-18$ fm, and $k^2=2mE/\\hbar^2$. One can now determine $b$ as the function of energy. This relationship is shown in (\\ref{1d2}) \\begin{equation}\\label{1d2} b(E)\\equiv\\frac{b_s(E)}{4}=\\frac{b(0)}{1+Qx^2}. \\end{equation} In (\\ref{1d2}) $x=k/k_c$, $k_c=mc/\\hbar=4.8\\cdot10^{13}$ cm$^{-1}$, and $Q=a|b_s(0)|k_c^2/2\\approx250$. In the degenerate neutron gas the most important is the energy at the Fermi level: $E=E_F=\\hbar^2k^2_F/2m$, where $k_F$ is the neutron wave number at the Fermi level. It is related to neutron density by equation $n=k_F^3/3\\pi^2$. Thus $b(E)=b(E_F)$ can be represented as a function of $n$: \\begin{equation}\\label{1d2x} b(E)\\equiv b(n)=\\frac{b(0)}{1+Qx^2}=\\frac{b(0)}{1+Q(n/n_c)^{2/3}}, \\end{equation} where a unit of density $n_c=k_c^3/3\\pi^2\\approx3.7\\cdot10^{39}$ cm$^{-3}$ is introduced.  In the next section the contribution of the optical potential to neutron star parameters $R$, $M$ and density distribution $n(r)$ is estimated. The neutron star is considered as a nonrotating spherical object composed of a degenerate neutron gas at zero temperature. Calculations were performed with the help of the Tolman-Oppenheimer-Volkov (TOV) equation~\\cite{re}, generalized by inclusion of neutrostriction forces. We do not take into account the short range nuclear forces because they come into play only at nuclear densities $n_N\\approx10^{38}$ cm$^{-3}$, while optical potential is the most important at $n\\ll n_N$.  In the third section the optical potential is compared to commonly used short range nuclear interactions, and in conclusion some effects are discussed, which can take place in neutron stars, if energy dependence $b(E)$ contains a resonance. ",
        "conclusions": "The neutrostriction forces are to be taken into account in calculation of neutron star. They play important role at low densities and small masses ${\\mathcal M}<M_\\odot$. More over they present many interesting problems worth of research for their own.  \\begin{enumerate} \\item It seems that neutron-neutron interaction in the degenerated neutron gas is eliminated because of the Pauli exclusion principle. However the Pauli exclusion principle, as correctly is pointed out in~\\cite{bom}, eliminates scattering process and imaginary part of the scattering amplitude, but it does not affect its real part. It means that because of the Pauli exclusion principle the optical potential in neutron stars becomes lossless. \\item The decrease of $b(n)$ with density is a source of pulsations, and it is interesting to investigate how well possible pulsation match parameters of the observed pulsars. \\item The pulsations are especially well understandable, if scattering contains a resonance, as is shown in fig. \\ref{ff3}. At some energy the scattering amplitude changes the sign, so for smaller density the optical potential is attractive, and for larger one it becomes repulsive. At the point $E$, where $b(E)=0$, pulsations arise naturally. \\begin{figure}[ht] {\\par\\centering\\resizebox*{10cm}{!}{\\includegraphics{resonanc.eps}}\\par} \\caption{\\label{ff3}Schematic energy dependence of scattering length in presence of resonance.} \\end{figure} \\item In the case of resonance we have a mechanism for star explosion. Indeed, if at contraction the density (and therefore Fermi energy) overcomes the resonant point, the strong repulsive optical energy changes to strong attractive one, and a huge energy is released. \\item Protons were not considered in this evaluation, however inclusion of protons will not change the arguments about optical potential, because the coherent neutron-proton scattering amplitude is nearly the same as the neutron-neutron one. At the same time, with protons one must also take into account neutron-proton resonances, which do certainly exist. Their presence can also provide a source of pulsations and explosions. \\item We considered only s-scattering amplitude. However, when the energy (or density) increases, one must also include p-, d- and higher harmonics. The question arises: how will they affect the results. \\item This paper discussed only neutron stars, but the notion of the optical potential is considerably more widely applicable. It can be used in other stars, in Bose-Einstein condensates, superfluidity and superconductivity, because everywhere we must take into account the coherent atom-atom and atom-electron scattering amplitude. \\end{enumerate}"
    },
    "0311/astro-ph0311537_arXiv.txt": {
        "abstract": "We show that a general, exact cosmological solution, where dynamics of scalar field is assigned by an exponential potential, fulfils all the issues of dark energy approach, both from a theoretical point of view and in comparison with available observational data. Moreover, tracking conditions are discussed, with a new treatment of the well known condition $\\Gamma>1$. We prove that the currently used expression for $\\Gamma$ is wrong. ",
        "introduction": "\\subsection{Motivations}  The recent discovery of the acceleration of the Universe \\cite{1,2} has produced an overwhelming number of papers on its possible explanations, alternative to a plain cosmological constant (see Refs. \\cite{3,4,4bis} for reviews). Historically, the most favourite one is the introduction of a scalar field, usually minimally coupled with gravity (but see Refs. \\cite{4bis1,4bis2,5,6,7,8,9,9bis} for other possible models). This poses the problem of the choice of a suitable potential. Moreover, as the equations are not exactly solvable in general, there is also a problem of qualitative analysis of the solutions and of the choice of initial conditions. Usually, very crude approximations are made, first of all the socalled \\textquotedblleft slow roll\\textquotedblright\\ condition, i.e., $\\dot{\\varphi}^{2}\\ll V(\\varphi )$.  The aim of this paper is to illustrate some subtleties involved by this situation, by means of a complete treatment of the general exact solutions of the Einstein equations in presence of both nonrelativistic matter and scalar field, with a suitable physical potential, and with a new treatment of the well known \\textquotedblleft tracking condition\\textquotedblright, $% \\Gamma>1$ .  One result is that the slow roll condition is indeed not necessary and that insisting on it can bring to lose very interesting models; closely connected to this, is a discussion on the meaning of the statement that the state equation $w\\equiv p_{\\varphi }/\\rho _{\\varphi }$ for the scalar field should be ``almost constant''. By means of an analytic example, we shall see that this expression is highly misleading, indeed.  Our choice for the potential is a particular class of exponential functions \\begin{equation} V(\\varphi)=B^{2}e^{-\\sigma\\varphi} \\,,  \\label{eq1} \\end{equation} with $B$ arbitrary and $\\sigma\\equiv\\sqrt{3/2}$ (in units $8\\pi G=1$). The reason for this special choice will be clear below.  First of all, we have to say that, although exponential potential is natural in higher dimension theories, supergravity and superstring models \\cite{9ter,9quater,9quinque}, it is not generally considered as feasible, even though it is often used at least as a first tool of analysis \\cite{10,12,11}. The usual objection to this potential is in fact that it cannot really be the right one, because \\textquotedblleft it scales approximately\\textquotedblright\\ like the matter component, so that, if the dark energy is dominant now, it should have been dominating also in the past, which is in contrast with BBN observations \\cite{13}. This potential was nonetheless proposed by some of us in Ref. \\cite{14}, where also some explanations of the reason why this is not true in our case are given. Anyway, in Sec. III the problem is fully explored, and we show why the above objection is wrong.  A second common objection is that it describes a very peculiar situation, due to a fine tuning of initial conditions. We insist on the point that the solution found in Ref. \\cite{14} is general and exact. Here, we start indeed from a particular choice, but then we go further and release it, showing that the same \\textquotedblleft final situation\\textquotedblright\\ is reached, starting from a wide set of \\textquotedblleft initial conditions\\textquotedblright. Again, we use quotation marks because the expressions are ambiguous and misleading, so that a preliminary discussion is needed to clarify what we mean.  A third possible objection is that this model could be in contrast with the observations. As a matter of fact, some papers have already appeared, successfully comparing this model with present day observations \\cite{15,16,17}. In Sec. II.B a first comparison with observational data is made, with very encouraging results. However, this paper is mainly concentrated on theoretical discussion. Our purpose is to show that the model we consider can emulate very well a standard model with dust plus a cosmological constant, well beyond any reasonable improvement of the observational techniques. In order to do this, a preliminary discussion on what is really measurable is again needed.  A second important result of our considerations consists in showing that many of the current statements about tracking solutions are based on an incorrect result and should therefore be revisited.  The paper is organized as follows. First, we examine some issues usually investigated when building a cosmological model with dark energy; key features like cosmic components, choice of the scalar field potential, measurability, and tracking behavior are therefore shortly sketched. Sec. II is then devoted to the exponential potential case and the general, exact solution that can be given to the cosmological equations. A comparison with a constant $\\Lambda $ model and observational data is made; also, there can be found an analysis of the scalar field equation of state and a more detailed one of the tracking behavior. In Sec. III, the features usually assigned to the tracking notion are critically examined, giving new insights into such a concept. Finally, concluding remarks are drawn in Sec. IV.  \\subsection{Preliminary discussion}  First of all, before going through the analysis of our potential, we think it is useful to discuss some general questions which should be of help in understanding better what follows.  \\subsubsection{How can we set the choice for the potential?}  At the moment there is no idea of which form the potential of the scalar field should take. The widely discussed inverse power law \\cite{4bis,20}, for instance, is certainly useful for a qualitative treatment of the problem, but only under the assumption of the slow rolling condition. It seems that the answer to this problem is simple in principle: the right form is the one that gives the best fit to the observed luminosity distance curve $d_{L}(z)$, and the fact that we cannot establish it firmly is only due to our limited observational precision. Indeed, it is possible to show that, whichever is the expression of $d_{L}(z)$, we can find, in principle, a suitable form for $V(\\varphi)$, which gives it back. Unfortunately, the value of $\\Omega_{0m}$ turns out to be \\textit{arbitrary}! The reason of this is due to the fact that the introduction of an arbitrary function in the model is equivalent to taking an infinite number of free parameters into account. Worst of all, this result is possible only if one has the \\textit{exact} functional form of $d_{L}(z)$. This is clearly impossible without already having a model in mind, so that one is led again at the starting blocks \\cite{20bis1,20bis2,20bis3,21,22}.  In our opinion, the only possible way to escape from this \\textit{impasse} is to ask for potentials which have some possible explanation in terms of fundamental physics or which exhibit some nice mathematical feature, which at least leaves the hope that it is due to some still unknown fundamental physical property. As for the first requirement, to our knowledge, the best candidates are the exponential potential, like that in Eq. (\\ref{eq1}) (but with $\\sigma$ undetermined), and polynomials of the form \\begin{equation} V(\\varphi)=\\frac{1}{2}m^{2}\\varphi^{2}+\\lambda\\varphi^{4} +\\text{\\textit{% hig. ord.}}  \\label{eq2} \\end{equation}  It is important to observe, however, that this is a generic expansion of an even function $f(\\varphi)$, with a suitable minimum, so that any function of this type will be good as well.  As for the second requirement, we shall see below that the potential proposed in Eq. (\\ref{eq1}), due to the special value of $\\sigma$, exhibits a Noether symmetry in the equations, which is just the reason why exact integration is possible. Also, the potential of Eq. (\\ref{eq2}) can be seen as the expansion of a generalized hyperbolic sine, with the same type of symmetry \\cite{23} (see below).  On the other hand, it is also possible to show that a sum of exponentials of the class in Eq. (\\ref{eq1}) plus a cosmological constant is essentially the most general type which can lead to exact integration in elementary way \\cite{24}.  \\subsubsection{What is really measurable?}  In the context chosen by us, we have a two-dimensional configuration space, with variables $\\left\\{ a,\\varphi \\right\\} $ plus the \\textit{velocities} $% \\left\\{ \\dot{a},\\dot{\\varphi}\\right\\} $, so that the total number of variables is four. Consequently, this is the maximum number of independent quantities which can be measured. As a matter of fact, none of these quantities is in fact practically measurable. In flat cosmologies, $a_{0}$ is not measurable \\textit{in principle}. This is connected with the choice $% k=0$, i.e., in terms of measurable quantities, $\\Omega _{tot.}=1$ exactly, which is clearly artificial. Instead of these rather abstract objects, we can also use derived quantities. One of them is, of course, $\\Omega _{0m}$. This parameter is very important, not only for its physical meaning, but also because its value can, in principle, be measured independently from the SNIa observations \\cite{25}. Much has also been said about the value of the acceleration parameter $q_{0}$; but, again, it is not directly measurable in practice. However, $q_{0}$ is also present in the second term of the series \\begin{equation} d_{L}(z)=H_{0}^{-1}\\left( z-\\frac{(1-q_{0})}{2}z^{2}+...\\right) \\text{ }\\,, \\label{eq3} \\end{equation}% so that it seems possible to measure it, after all. Unfortunately, truncation of Eq. (\\ref{eq3}) after the second term is valid only at rather low redshifts (say $z\\leq 0.1$). At these values, it is practically impossible to obtain sufficient precision, while at higher values the other terms cannot be neglected.  What is thus left? It seems reasonable to stick to what is for the moment firmly established, i.e., $h\\approx0.7$ and $\\Omega_{0m}\\approx0.3$, with $3\\sigma$ level $\\sim\\!10\\%$ and $\\sim50\\%$ errors, respectively\\footnote{These values are based on estimates not strictly dependent on the cosmological model.}. If we are very optimistic, we can expect an improvement of precision in the future of a factor 5, say. We shall see that even this could be not enough to discriminate among models.  Also, a third useful variable could be the age of the universe $t_{0}$, for which we have only a rather poor estimate $t_{0}>12$ Gyr.  \\subsubsection{Which is the precise meaning of the expression \\textquotedblleft tracking behavior\\textquotedblright ?}  One of the greatest advantages of the introduction of a scalar field consists in the possibility to avoid the fine tuning of initial conditions, obtaining the same final behavior \\textquotedblleft for a wide range of initial possibilities\\textquotedblright. The last sentence is intentionally rather ambiguous. In fact, the situation is often depicted as \\textquotedblleft similar but not exactly equal\\textquotedblright\\ to that of an attractor in the theory of dynamical systems.  Let us thus try to be more precise. The equations which we have to consider are the Einstein equations plus the Klein-Gordon equation for the scalar field (for sake of simplicity we write them directly in the flat case, but the following argument holds in the general case as well) \\begin{equation} 3H^{2}=\\rho _{\\varphi }+\\rho _{m}\\text{ } ,  \\label{eq4} \\end{equation}  \\begin{equation} 2\\frac{\\ddot{a}}{a}+\\left( \\frac{\\dot{a}}{a}\\right) ^{2}+\\frac{1}{2}\\dot{ \\varphi}^{2}-V(\\varphi )=0\\text{ } \\,,  \\label{eq5} \\end{equation}  \\begin{equation} \\ddot{\\varphi}+3H\\dot{\\varphi}+V^{\\prime }(\\varphi )=0\\text{ } \\,. \\label{eq6} \\end{equation}  Eq. (\\ref{eq4}) is a first integral of the other two. The integration constant is hidden into $\\rho _{m}$. For instance, in the case of dust, we have $\\rho _{m}=Da^{-3}$; again, the value of $D$ is not directly measurable, but we shall not in fact make use of this equation. Instead, we shall use Eqs. (\\ref{eq5}) and (\\ref{eq6}). The reason is that it can be shown that, when $\\rho _{m}$ is related to a pure dust component, they are derived by the following Lagrangian \\cite{26,27} \\begin{equation} L=3a\\dot{a}^{2}-\\frac{1}{2}a^{3} \\dot{\\varphi}^{2} +a^{3}V(\\varphi )\\text{} \\,.  \\label{eq7} \\end{equation} Eq. (\\ref{eq4}) is then obtained as conservation of the energy function $% E_{L}\\equiv \\left( \\partial L/\\partial \\dot{a}\\right) \\dot{a}+\\left( \\partial L/\\partial \\dot{\\varphi}\\right) \\dot{\\varphi}-L$.  A first consequence of this fact is that Liouville's theorem holds for this system. Thus, the phase space volume of possible initial conditions is conserved and has no attractors.  Nevertheless, the tracking behavior can be recovered by the following argument. Consider a set of independent variables for this system (not necessarily $\\{a,\\varphi ,\\dot{a},\\dot{\\varphi}\\}$) and a volume of possible initial conditions. During the evolution, the volume is deformed, and it is possible that its projection on a two(three)-dimensional subspace converges to a point or a line. It is clear now that the result is strongly dependent on the choices of variables and projection. It is also clear that the gain in information on the variables of the subspace is compensated by a loss of information on the variables of the complement.  In our specific case we can take advantage of the fact that, as said above, only a part of the possible variables is practically measurable. Therefore we can make the usual choice for the subspace, i.e., consider $\\{\\log a,\\log \\rho _{\\varphi }\\}$, and show that, for a large set of initial conditions of these two variables\\footnote{% Apparently, $a$ is constrained to be zero at the initial point, but in fact the models cannot be pushed up to the initial singularity, so that the initial time cannot be really zero. This gives some arbitrariness in the choice of initial $a$. (See next Section for details.)}, the orbits converge to a line. Another possibility is to consider $\\{H,\\Omega _{m}\\}$, and check that the orbits converge to a point.  It is also clear that only \\textit{observable} quantities should be taken into account. For any other parameter the concept of fine tuning is, in our opinion, simply meaningless.  The situation is in fact rather involved and will be discussed in Sec. II.D, by means of an example.  We think that the concept illustrated above is the only one which is relevant for the solution of the coincidence problem. However, in the literature the term \\textquotedblleft tracking\\textquotedblright\\ is often used in a different meaning. The idea is that a large set of initial conditions should converge towards a \\textit{nearly constant} $w$, well \\textit{before} the present epoch, and maintain this condition forever. We shall show in Sec. II.C that this situation, although of course possible, is not at all necessary. ",
        "conclusions": "We have presented above a class of physically meaningful exponential potentials, which allow general exact solution of the Friedman and Klein-Gordon equations. We have also shown that this solution, in the limits of validity of the model, meets all the requests which are generally made for a quintessence model, except one: the possibility of switching off the acceleration in the future, which is needed for the asymptotic freedom of the model \\cite{28}. This can indeed be obtained by a modification of the potential.  One possibility, studied in Refs. \\cite{31,32,33}, is to bring the value of $ \\sigma$ to the range $\\sqrt{2}<\\sigma<\\sqrt{3}$, so that the limit value for $w$ is raised to a value greater than $-1/3$. Another possibility is to pass to a combination of exponentials, which we have treated in another paper \\cite{23}.  This last example is particularly interesting, and it is worthwhile to summarize some results. Using the potential \\begin{equation} V(\\varphi )=(Ae^{\\sigma \\varphi /2}-Be^{-\\sigma \\varphi /2})^{2}\\,, \\label{doubpot} \\end{equation}% it is possible to obtain essentially the same situation as shown here (with a simple exponential potential) for the past, until present epoch. Towards the future, instead, we get that $w$ oscillates forever between $-1$ and $1$, but the oscillations of $\\rho _{\\varphi }$ turn out to be damped, so that it effectively scales like matter, and dark energy never reaches complete domination of dynamics. (See Ref. \\cite{23} for details, and Ref. \\cite{33bis} for a closely related situation.)  Coming back to the potential in Eq. (\\ref{eq1}), it is possible to guess a slightly positive spatial curvature. As said in the introduction, this is not excluded by the present data. It is then possible to add a curvature term into Eqs. (\\ref{eq4}), (\\ref{eq5}) and (\\ref{eq6}), and try to integrate them numerically. An appropriate choice for $\\Omega_{k}$ gives the plot of Fig. 12, where we see that the acceleration starts from negative values, reaches a regime of positive values, and then goes back to negative in the future.  We have now to shortly discuss which are the shortcomings of the model we have considered so far in this paper.  First of all, although exponential potentials are widely considered as possible candidates to some kind of exotic dark energy (for physical reasons which are \\textit{a priori} with respect to the framework here presented), we have found no explanation which justifies the particular class of exponential potentials we have chosen, except the nice feature of the Noether symmetry of the Lagrangian (\\ref{eq7}) (but see Refs. \\cite{26,27,29,30} for some discussion on this point).  Secondly, as we have repeatedly stressed, this model loses its validity when radiation is not negligible, and we have seen that a lot of interesting features occur just during that epoch. An extension is thus needed, but this must probably be done with numerical methods. Due to the very small values of some parameters, one must in fact be very careful to avoid numerical errors. In order to generate properly the plots of Sec. II, for instance, we had to work with $40$ digits floating numbers, and there was not cumulated error typical of numerical integration.  A study along these lines is presented in Ref. \\cite{33}, where they find a behavior very similar to the one depicted by us for a reasonable range of values of $\\sigma$ and initial conditions. They include radiation in numerical computation, but the procedure is somewhat different by ours.  Most of all, our discussion shows not only that this model can perfectly emulate a cosmological constant, but also gives no answers about the real possibility to discriminate, at least between these two models. It seems that we have shown that this discrimination is however impossible on the pure basis of luminosity distance observations. We have obtained this already known result \\cite{34} by showing that there is an important degeneracy in the $\\Omega_{m}$ parameter.  In principle, this parameter can be measured independently, but it appears very difficult to reduce its present error of a factor, say, $10$ in the near future. We have also shown elsewhere that the two models are equally compatible with the observations on the peculiar velocities of galaxies \\cite{35}. This kind of measure is also affected by large errors, is time consuming and can say something only at rather low redshift ($z\\approx 0.15)$. More promising seems to be the analysis of CMBR, as said above, and gravitational lensing, particularly of very distant QSO (see Ref. \\cite{36}, for instance, and references therein). Here the problem appears to be the presence of a lot of other parameters, leading to huge degeneracy.  In our opinion, however, the main result which we obtain in this paper is probably pedagogical: we have learned that some statements, taken for granted by the majority of people, can be misleading. Let us summarize some of them here again.  We have seen that the most interesting epoch in our model is when $w$ is mostly variable, but only with respect to $\\log a$. This shows that to say `` $w$ is necessarily almost constant'' is badly stated, and risks to be merely a prejudice.  The study of the tracker behavior is subtle and strongly depends on the representation chosen. This is so even in the case when a general exact solution is given. This does not mean that qualitative analysis cannot be done in less favourable situations, but that it should be very carefully and unambiguously presented.  The requirement $\\Gamma >1$ seems to be not necessary and, in any case, many widely accepted statements on exponential potentials and tracking solutions are indeed based on a result which has revealed to be incorrect.  In conclusion, we think that the final answer to the problem of the acceptable form for the potential is far to be reached, and that more insight on the physical motivations for the dark energy potential should be developed."
    },
    "0311/astro-ph0311584_arXiv.txt": {
        "abstract": "We have developed a three dimensional Monte Carlo photoionization code tailored for the study of Galactic H~{\\sc ii} regions and the percolation of ionizing photons in diffuse ionized gas. We describe the code, our calculation of photoionization, heating \\& cooling, and the approximations we have employed for the low density H~{\\sc ii} regions we wish to study. Our code gives results in agreement with the Lexington H~{\\sc ii} region benchmarks. We show an example of a 2D shadowed region and point out the very significant effect that diffuse radiation produced by recombinations of helium has on the temperature within the shadow. ",
        "introduction": "A major characteristic of the interstellar medium (ISM) is that it is clumpy, with filamentary structure over scales ranging from hundreds of parsecs through sub-AU cloudlets. The spatial density of electrons is approximately a power law over an astonishing 12 orders of magnitude \\citep{ars}, from $\\sim$300 pc to $<$0.1 AU in the ISM. The magnetic field of the Galaxy varies in a similar fashion on scales from 0.01 to 100 pc. Molecular clouds show a fractal (scale-free) spatial variation of the intensity of CO (e.g., \\citealt{fal}) and neutral H \\citep{sl}. Sub-parsec variations are shown by the variation of ionic column densities along different ISM sightlines towards individual stars in close binaries \\citep{lmb}.  Large area velocity resolved surveys of the molecular (e.g., \\citealt{sjc}), neutral (e.g., \\citealt{hb}), and ionized \\citep{r02} gas now allow us to probe the geometry, kinematics, ionization, and temperature structure of the ISM.  Our development of Monte Carlo radiation transfer techniques to study photon scattering, ionization, and radiative equilibrium dust temperatures in complex geometries is motivated by these and the many other observations that clearly require three dimensional modeling analyses. In this paper we describe improvements and extensions to our Monte Carlo photoionization code \\citep{wl} that will enable us to model the 3D temperature, ionization structure, emission line images, and spectra of low density H~{\\sc ii} regions and the Diffuse Ionized Gas (DIG), often called ``the warm ionized medium\" or ``Reynolds layer''. This diffuse ionized gas has a complicated structure and surely requires three-dimensional modeling, both as regards the sources of its ionization and the transfer of the ionizing radiation. How can Lyman continuum photons from O and B stars percolate from their midplane origins to ionize the high latitude gas? A non-uniform ISM is required that provides low density paths for ionizing photons to reach the halo.  Using a ``Str\\\"omgren volume'' ionization technique and a statistical approach to clumping, \\citet{mc} showed that O and B stars can ionize the DIG (see also \\citealt{ds94}; \\citealt{dsf}). Using Monte Carlo simulations for hydrogen ionization, \\citet{c02} have investigated ionization structures in a fractal ISM as proposed by \\citet{elm}. In Elmegreen's picture, the observed H$\\alpha$ emission may arise from the ionized surfaces of fractal clouds. This suggestion clearly needs to be investigated with 3D radiation transfer techniques.  Other features of the DIG that may require 3D analyses are the temperature structure and relative ionization structure of He and H. The temperature of the DIG, as determined from line intensity ratios (e.g., \\citealt{rht}), appears to require additional heating over and above that provided by photoionization heating \\citep{m00}. Among the possible heating mechanisms, dissipation of energy in a turbulent medium has been proposed \\citep{ms}, definitely 3D! Another question is, can clumping/shadowing explain the apparent lack of He$^0$ in H$^+$ regions \\citep{rt}? In this scenario, direct stellar photons are prevented from ionizing He in shadow regions behind dense clumps, but diffuse ionizing photons from H and He recombinations can percolate into the shadow regions and ionize H, but not He.  Our work on the analysis of scattered light in fractal reflection nebulae \\citep{mww} highlighted the crucial role of geometry when trying to determine dust properties from scattered light observations. Much of what we know about abundances of elements comes from analysis of spectra utilizing one-dimensional photoionization techniques (see e.g. \\citealt{s02}).  Over the last decade several three dimensional photoionization codes have been developed.  The code presented by \\cite{bdg} uses the ``on the spot'' approximation for the diffuse radiation field, while that of \\cite{gruen} employs an adaptive mesh grid and an iterative procedure to determine the contribution of the diffuse radiation field.  Monte Carlo techniques naturally include diffuse radiation within three dimensional systems and Monte Carlo photoionization codes have been described by \\citet{erc} for the three-dimensional case and \\citet{olr} for the one-dimensional case.  We have extended our basic hydrogen-only Monte Carlo photoionization code \\citep{wl} to treat the photoionization, heating, \\& cooling for the complex geometry expected in low density H~{\\sc ii} regions and the diffuse ISM. Our code is similar to those described by \\citet{olr} and \\citet{erc}, but calculates the rates of photoionization and heating ``on the fly'', rather than calculating the mean intensity of the radiation field (see also Lucy 1999). It also treats reprocessing of radiation in a different manner to that described by \\citet{olr}, using ``photon packets'' rather than ``energy packets.''  In \\S2 we describe our assumptions and approximations to model the DIG. \\S3 describes the Monte Carlo technique, our method of discretizing the radiation field into ``photon packets'' rather than ``energy packets,'' and how we determine the diffuse radiation field using ratios of recombination rates. \\S4, \\S5, \\& \\S6 deal with our photoionization and heating/cooling algorithms, \\S7 presents the atomic data used in our code, and \\S8 compares our code with standard H~II region benchmarks and gives an example of a 2D calculation. ",
        "conclusions": ""
    },
    "0311/astro-ph0311067_arXiv.txt": {
        "abstract": "The Wide-field X-ray Monitor (WXM) is one of the scientific instruments carried on the High Energy Transient Explorer 2 (HETE-2) satellite launched on 2000 October 9. HETE-2 is an international mission consisting of a small satellite dedicated to provide broad-band observations and accurate localizations of gamma-ray bursts (GRBs). A unique feature of this mission is its capability to determine and transmit GRB coordinates in almost real-time through the burst alert network. The WXM consists of three elements: four identical Xe-filled one-dimensional position-sensitive proportional counters, two sets of one-dimensional coded apertures, and the main electronics. The WXM counters are sensitive to X-rays between 2 keV and 25 keV within a field-of-view of about 1.5 sr, with a total detector area of about 350 cm$^2$. The in-flight triggering and localization capability can produce a real-time GRB location of several to 30 arcmin accuracy, with a limiting sensitivity of $10^{-7}$ erg cm$^{-2}$. In this report, the details of the mechanical structure, electronics, on-board software, ground and in-flight calibration, and in-flight performance of the WXM are discussed. ",
        "introduction": "\\label{sect:intro}  The origin and nature of Gamma-ray bursts (GRBs) has been one of the outstanding mysteries of astrophysics since the phenomenon was discovered in the 1970s. The discovery of GRB afterglows by the BeppoSAX X-ray satellite~\\citep{1997Natur.387..783C} produced a breakthrough in the study of GRB. Since the discovery, follow-up observations of many GRB afterglows have been successfully performed.  The observations have furnished strong evidence that at least the long class of GRB originates from explosions of super-massive stars at cosmological distances. The HETE-2~\\citep{wh.2003.grr} is the first satellite dedicated to GRB observations and is intended to provide the astronomical community with prompt, high-accuracy GRB positions --- a few tens of arcminutes to a few tens of arcseconds, with delay times of 10 s to a few hours. This capability not only increases the number of detected GRB counterparts, but also provides an opportunity for multi-wavelength observations of GRBs in their very early phase.  Such observations were nearly impossible prior to the HETE-2 era.  The HETE-2 instrument complement consists of three scientific instruments: the French Gamma Telescope (FREGATE), the Wide-field X-ray Monitor (WXM), and the Soft X-ray Camera (SXC). FREGATE~\\citep{wh.2003.atteia} consists of four identical scintillation detectors.  It is sensitive to photons with energies from 6 to 400 keV. It is responsible for triggering and spectroscopy in the hard X-ray and Gamma-ray energy ranges. Its field of view is 70$^{\\circ}$ (half width at zero maximum) and its total detection area is 160 cm$^2$. The SXC~\\citep{wh.2003.joel} is a 1-D coded mask system using MIT-LL CCID-20's as the detecting elements, and is sensitive to photons with energies from 500 eV to 14 keV, with 2\\% resolution at 6 keV and with a position resolution of 10$''$ for bright bursts (1 Crab; 10 s). SXC's total detection area is 75 cm$^2$ and its field of view is $65^{\\circ}$ (full width at zero maximum). The WXM, on the other hand, has a detection area of 352 cm$^2$ and a field of view of $80^{\\circ}$ (FWZM). Thus, the WXM is more sensitive to weak bursts than the SXC, and the two X-ray instruments play a complementary role to each other, balancing detection sensitivity against location resolution. ",
        "conclusions": "We have calibrated the alignment between the WXM and the HETE-2 spacecraft aspect system. We estimate the systematic uncertainty of WXM locations of GRB to be 2$'$ for each X and Y direction at the 68\\% level.  The  corresponding two-dimensional systematic error is 4$'$ (90\\% confidence level). We calibrated the gain scale of the proportional counters using observations of the Crab Nebula and using the signal from the calibration radio-isotopes. At present, the systematic uncertainties in fitting a power-law spectrum model to the Crab spectrum are: $\\pm 0.1$, $\\pm 0.2\\times 10^{22}$ cm$^{-2}$ and 10\\% for power-law index, absorption column density, and absolute flux, respectively. Improvements in the spectral calibration are in progress. The in-orbit flight localization performance of the WXM was examined based on XRB observations and Monte Carlo simulations. The sensitivity of the flight localization is estimated to be $> 10^{-7}$ erg~cm$^{-2}$, and the estimated false alert rate of real-time localizations is 20\\% for $1.5\\times 10^{-7}$--$2 \\times 10^{-6}$ erg~cm$^{-2}$. Based on the recent GRB localizations, the rate of GRB localization is 28 per year, and the real-time localization rate is approximately one per month."
    },
    "0311/astro-ph0311251_arXiv.txt": {
        "abstract": "{ In this paper we present Very Large Array (VLA) 1.4 GHz (21 cm) observations of the region between the centres of A3558 and A3562, in the major cluster merger complex of the Shapley Concentration. Our final catalogue includes a total of 174 radio sources above the flux density limit of 0.25 mJy b$^{-1}$. By cross-correlation with optical and spectroscopic catalogues we found 33 optical counterparts belonging to the Shapley Concentration. \\\\ We investigated the effects of cluster merger on the radio emission properties of the galaxy population by means of the radio source counts and the radio luminosity functions (RLF). We found that the radio source counts are consistent with the field source counts. The RLF of elliptical and S0 galaxies in the region surveyed here, is consistent with the ``universal'' RLF  for early--type galaxies. This result suggests that the deficit in radio galaxies found in our previous works over the whole A3558 chain, is entirely due to the cluster A3558. A population of faint radio galaxies (logP$_{1.4~GHz}$(W Hz$^{-1}$) \\ltsim  22) is also found. Half of these objects are also blue, suggesting that starburst is the main mechanism driving the radio emission. Finally, we detected 14 spiral radio galaxies, whose ratio between radio and optical emission is similar to those found in galaxies located in rich and dynamically evolved clusters. \\\\ Our results are briefly discussed in the light of the age and stage of the merger in the A3558 cluster complex. ",
        "introduction": "Evidence is accumulating that the interaction processes between clusters of galaxies, known as cluster mergers, may significantly affect the radio emission characteristics of the cluster galaxy population. In particular, the evolution of galaxies and the properties of their nuclear and/or star forming activity are undoubtely influenced by the interaction with the environment. In this frame, cluster merging and group accrection seem to play an important role, but it is not yet completely understood how the merging environment affects the nuclear and starburst emission in galaxies. Owen et al. (\\cite{owen99}) suggested that merging may trigger the radio emission, both in the form of nuclear activity and starburst phenomena. Burns et al. (\\cite{burns94}) interpreted the presence of post-starburst galaxies in the X--ray and radio bridge connecting the Coma cluster and the NGC 4839 group as consequence of a recent merging process between these two structures. Using numerical simulations, Bekki (\\cite{bekki99}) concluded that the tidal gravitational field of a cluster merger may drive a considerable transfer of gas to the central regions of galaxies, leading to enhanced star formation activity or feeding the central engine of active galactic nuclei. Vollmer et al. (\\cite{vollmer01b}) suggested that a local burst of star formation could be due to re-accretion of gas after a ram pressure stripping event, which is thought to be responsible for the HI deficit of spiral galaxies in the central regions of nearby clusters with respect to the field galaxies of the same morphological type and optical size (Gunn \\& Gott, \\cite{gunn72}; Bothun \\& Dressler, \\cite{bothun86}). On the other hand, Fujita et al. (\\cite{fujita99}) and Balogh et al. (\\cite{balogh98}) claimed that gas stripping resulting from ram pressure is important in preventing gas supply to the central regions of the galaxies and may suppress star formation. Finally Venturi et al. (\\cite{venturi00}, hereinafter V2000) found evidence that merging may anticorrelate with the radio emission, possibly switching off previously existing radio galaxies, or temporarily inhibiting the nuclear radio activity. \\\\ A major problem in our current knowledge of the role of cluster mergers on the radio emission and in the interpretation of the observations is the lack of statistics. Only few clusters of galaxies were deeply imaged over their whole extent (i.e. A2125, Dwarakanath \\& Owen, 1999; A2255, Miller \\& Owen, 2003 (hereinafter MO03); A2256, Miller et al., 2003; A2645, Owen et al., 1999; the Shapley supercluster chains  A3528, A3558 and A3571, Venturi et al. 2001, 2000 and 2002 respectively), and the complexity of the phenomenon clearly needs a much larger observational support. Cluster mergers evolve over a timescale of $\\sim$ Gyr, while typical ages of radio sources are of the order of few times 10$^7$ -- 10$^8$ years, therefore it is of crucial importance to identify the effects of the various stages of a merger event. For this reason it is essential to carry out deep radio observations of a large number of galaxy clusters, at different merger stages. \\\\ In order to study the effects of a major cluster merger event after the first core--core encounter on the nuclear and starburst radio emission, in this paper we present 1.4 GHz Very Large Array (VLA) observations of the cluster A3562 and the poor groups SC 1329$-$313 and SC 1327$-$312, at the eastern side of the A3558 cluster complex. In the following we will refer to the area surveyed as A3558--C. This work is part of a larger project which aims to study the influence of the ongoing merger in the core of the Shapley Concentration on the radio/optical properties of the cluster galaxies, both from a statistical point of view and through a detailed analysis of the physical properties of the extended radio galaxies. In section \\ref{sec:intro} we briefly overview the optical and X--ray properties of A3558--C, with emphasis on the cluster A3562 and on the poor groups SC 1329$-$313 and SC 1327$-$312; the observations and data reduction are presented in Section \\ref{sec:obs}; in Section \\ref{sec:sample} we present the 1.4 GHz radio sample and the source counts; in Section \\ref{sec:ident} we deal with the optical identifications and in Section \\ref{sec:prop} we present the general properties of the radio galaxies in this region; the results of the analysis of the radio luminosity function (RLF) for the early-type and late-type galaxy populations are given in Section \\ref{sec:RLF}; discussion and conclusions are given in Section \\ref{sec:disc}.  We assume $H_0 = 100$ km s$^{-1}$ Mpc$^{-1}$ and $q_o$=0.5. If we define h=H$_o$/100, at the average redshift of the Shapley Concentration ($z=0.05$) this leads to a linear scale of 1 arcsec = 0.67 h$^{-1}$ kpc. We will assume $S \\propto \\nu^{-\\alpha}$.   \\vskip 1 truecm \\noindent ",
        "conclusions": "\\label{sec:disc}  In this paper we presented deep 1.4 GHz VLA observations of the major cluster merger in the A3558 complex. The area under study, A3558--C, defined as the region between the centres of the two Abell clusters A3558 and A3562, including also the two groups SC 1329--213 and SC 1327--312, is thought to have recently experienced a major merger between two massive clusters. What we see now is expected to be the result of the first core--core encounter. The observational properties in this region, together with numerical simulations, suggest that A3558 is the main cluster, while the whole chain, beyond A3558 itself, is the remains of the colliding cluster (Bardelli et al. 2002). \\\\ A picture is emerging, in which radio AGN and starburst activity, radio relics and halos, or the lack thereof, are signature of cluster mergers at different stages. In particular, Venturi et al. (2002) proposed an evolutionary merger sequence to account for the diversity in the radio properties of the three main merging cluster systems in the core of the Shapley Concentration; Kempner \\& Sarazin (2001) postulated that radio halos and relics may form at different times during mergers; MO03 explained the different fraction of radio emitting galaxies (starburst and AGNs) in a cluster sample as due to different merger stages. With this study we have a unique opportunity to connect the observed properties in the radio band with a well defined cluster merger stage. This is of crucial importance for a better understanding of the complex effects of cluster mergers, with particular emphasis to the stage of the merger. \\\\ The most relevant results of our analysis can be summarised as follows:  \\begin{itemize}  \\item[(i)] a faint radio halo is found at the centre of A3562, whose properties are consistent with the idea that it is a young source at the beginning of the reacceleration phase (V2003) induced by a recent merger event;  \\item[(ii)] the origin of the extended emission in the radio galaxy J1332--3146a, associated with the dominant galaxy in the group SC 1329--313, is unclear. No radio jets are present in the nuclear region of J1332--3146a, and the nuclear radio component has steep spectrum. It is possible that this extended emission has actually cluster origins (see Section 6.2), being either (a) a ``bright'' area of a very low brightness bridge of radio emission, connecting A3562 and SC 1329--313, or (b) a ``revived'' radio emission region, where pre--existing old electrons were reaccelerated;  \\item[(iii)] a large number of radio sources associated with A3558--C galaxies was found, i.e. 33 objects, most of them with radio emission at low power levels. Our analysis (see Section 6.3) suggests that 26/33 radio galaxies are candidate starbursts. Among them, 11 also show a blue excess;  \\item[(iv)] the total number of radio AGNs detected in this region is consistent with the expectations from the RLF of LO96, suggesting that the cluster merger has not affected the probability of an early--type galaxy to develop a nuclear radio source;  \\item[(v)] the distribution of the radio/optical ratios for the spirals in A3558--C is similar to what is found in rich and evolved environments.  \\end{itemize}  The main question is if and how this wealth of observables is connected to the cluster merger in this region. Beyond the noticeable finding of the radio halo in the centre of A3562 (V2003), the radio properties of the Shapley galaxies in this region may contain important pieces of information.  V2000 showed that the RLF for early type galaxies in the whole A3558--C shows a deficit of radio galaxies compared to the ``universal'' RLF presented in LO96 (V2000), over the whole power range. On the other hand, the RLF for AGNs presented here matches the expectations of LO96; furthermore the outskirts of the chain (A3556 and A3562) contain the largest fraction of Shapley radio galaxies (see also Section 2). This suggests that A3558 itself, the most massive cluster in the chain, is the main responsible for the lack of radio sources in the RLF found in V2000. A possible explanation is the key role of A3558 in the merger, i.e. it is experiencing the most dramatic effects of the merger, being the result of the interaction of core regions of the two colliding clusters. If our interpretation is correct, then the role of cluster merger on the radio emission from AGNs may be many--fold, depending on the age and strength of the merger. \\\\ Optical photometric information on the faint population of radio galaxies in A3558--C confirms that at least 50\\% of these objects are most likely starburst candidates. However, a radio/optical analysis carried out following MO03 provides only weak evidence that the fraction of radio emitting galaxies in A3558--C is higher than in non--merging environments. \\\\ The statistical results on the late--type galaxies show that the radio emission in the A3558--C spirals is similar to those in rich and dynamically evolved clusters. This suggests that the radio emission in spiral galaxies may be one of the first ``parameters'' to react to a cluster merger event.  In conclusion, V2003 showed that the radio halo at the centre of A3562 is consistent with a reacceleration phase which started $\\sim$ a few $10^8$ years ago. This ongoing merger is therefore {\\it advanced}, in the sense that the core--core encounter has already taken place, but it is still young if compared to the total duration expected for a cluster merger, i.e. 10$^9$ yr. Our study suggests that on this timescale, the effect of cluster merger on the radio emission from cluster galaxies is many--fold. In particular, we found only marginal evidence of enhanced radio emission of starburst origin, but we found significant enhanced radio emission from spirals. We argue that the role of cluster merger on the nuclear activity in early--type galaxies is a complex phenomenon, since observational evidence in the whole A3558 cluster chain clearly shows that the deficit in radio galaxies found in V2000 is entirely due to the cluster A3558, which has experienced the most violent consequences of the merger."
    },
    "0311/astro-ph0311366.txt": {
        "abstract": "\\noindent We use a principal components analysis of radio-selected (3CRR, 6CE and 7CRS) AGN datasets to define two parameters related to low-frequency ($151$ MHz) radio luminosity $L_{151}$ and [OIII] luminosity $L_{{\\rm [OIII]}}$: a parameter $\\alpha$ encoding the $L_{151} - L_{{\\rm [OIII]}}$ correlation and a parameter $\\beta$ encoding scatter about this correlation. We describe methods for constructing generalized luminosity functions (GLFs) based on $\\alpha, \\beta$, redshift and schemes for unifying quasars and radio galaxies. These GLFs can be used to generate radio luminosity functions (RLFs) which improve on those of \\citet{wrlf}, mostly because they incorporate scatter and are therefore much smoother.  Luminosity-dependent unified schemes (e.g. a receding-torus scheme) have been invoked to explain the low quasar to radio galaxy fraction at low $\\alpha$ and the differences in emission-line luminosities of radio quasars and radio galaxies. With the constraints of the 3CRR, 6CE and 7CRS datasets and radio source counts, our GLF approach was used to determine whether a receding-torus-like scheme is required if there are two populations of radio sources: one at low $\\alpha$, consisting of `starved AGN'; the other at high $\\alpha$ consisting of `Eddington-tuned AGN'. Because of the overlap between these two populations and the effects of the $\\beta$ parameter, schemes with or without a receding torus can produce a low quasar fraction at low $\\alpha$ and differences in [OIII] luminosity between radio galaxies and quasars. The receding torus may be a physical process important in one or more populations of radio sources, but this is not yet proved either by the quasar fraction or the emission-line properties of radio-selected samples. ",
        "introduction": "\\label{sec:intro}  It has been recognised for some time that there is a strong positive correlation between the extended-radio luminosities and narrow-emission-line luminosities of 3C radio sources \\citep{bh, retal, mcarthy}, and by combining the 3C sample with the 7C Redshift Survey (7CRS), \\citet{wem} found that these correlations were not primarily due to redshift. This suggests that the sources of narrow-line emission and radio jets are linked, with explanations for this link ranging from the effects of environment \\citep{dp93}, to jet-disk symbiosis \\citep{rs91, fb}.  Unified schemes for radio galaxies and radio quasars propose that they are the same objects viewed at different angles between their radio axis and their line of sight. An obscuring torus is invoked to hide the nucleus at large angles to the jet axis i.e. in radio galaxies. Emission from narrow-line regions is believed to be broadly independent of the jet-axis orientation, as it is emitted from a region larger than the extent of the torus. The opening angle of the torus $\\Theta_{\\rm trans}$ marks the transition from the object being viewed as a quasar to a radio galaxy. We therefore expect for simple, arguably naive, unified schemes, where $\\Theta_{\\rm trans}$ is constant, that the distribution of emission-line luminosities should be similar for radio quasars and radio galaxies. \\citet{jb} found that [OIII] luminosities of radio quasars are $\\sim 5-10$ times more luminous than radio galaxies. In similar comparisons of radio quasars and radio galaxies, \\citet{hbf} found no difference using [OII], and \\citet{jr} found no difference in [OIII] at high redshift. In summary, some studies comparing the narrow emission line strengths of radio quasars and radio galaxies are seemingly in agreement with the predictions of the simplest unified schemes, whilst others are not.  These seemingly contradictory, results can perhaps be understood in the context of a `receding torus' model for AGN. This model \\citep{law, hgd} proposes that the inner radius $r$ of the obscuring torus is determined by the radius at which dust sublimates, scaling as $r \\propto L_{\\rm phot}^{0.5}$. Assuming that the half-height $h$ of the torus is independent of $L_{\\rm phot}$, then the half-opening-angle of the torus, $\\theta = \\tan^{-1} r / h$, will be larger in the higher-$L_{\\rm phot}$ objects. This means that more luminous objects, with higher $L_{\\rm phot}$, are more likely to be viewed as quasars so that quasar fraction rises systematically with $L_{\\rm phot}$. It also means that orientation-independent quantities which scale with $L_{\\rm phot}$ will be higher on average for quasars than for radio galaxies. \\citet{law} found that the fraction of narrow-line objects in the low-frequency-selected 3CR sample decreased with increasing radio luminosity, and that narrow-line objects have weaker [OIII] at a fixed radio luminosity. He argued that this was inconsistent with the simplest unified schemes but that it could be explained by a cone-angle dependence on luminosity, e.g a receding-torus model. \\citet{simp} argued that the [OIII] emission line is a much better indicator of $L_{\\rm phot}$ than [OII], explaining why differences between quasars and radio galaxies are more obvious in [OIII] than [OII]. \\citet{sconf} revisited the arguments of \\citet{simp} and corrected a small error (compare Fig.~2 of \\citealp{sconf} with Fig.~4 of \\citealp{simp}). He concluded that, allowing for a spread in $h$, the receding torus model predicts that radio quasars should be a factor of a few brighter in $L_{\\rm phot}$ than radio galaxies in samples exhibiting a wide range of quasar fractions.  \\citet{wsubmm} found that radio quasars have higher submillimetre luminosities by a factor of $\\sim 4$ than radio galaxies of the same radio luminosity and redshift, a factor which cannot be reduced below $\\sim 2$ by accounting for possible synchrotron contamination. This result supports the idea that the simplest unified schemes, where $\\Theta_{{\\rm trans}}$ is constant, are not an adequate description of the relationship between radio galaxies and quasars: submillimetre emission comes from cool dust grains in optically thin regions, and therefore radiates isotropically; moreover, such emission could not be obscured by a torus even if it were emitted from regions close to the nucleus. Submillimetre luminosity is therefore an orientation-independent quantity which might scale closely with $L_{\\rm phot}$ either because the cool dust is heated directly by the quasar or because it is heated by a starburst whose luminosity scales with $L_{\\rm phot}$. We conclude that the submillimetre study of \\citet{wsubmm} is in quantitative agreement with the receding torus model of \\citet{sconf}.   Other arguments for a receding-torus-like model include the evidence that the fraction of lightly reddened $z \\sim 1$ 3C quasars decreases with increasing radio luminosity in agreement with the higher fraction of lines of sight expected to graze the torus at lower $L_{\\rm phot}$ (at fixed $h$) in the receding torus model (\\citealp*{hgd}; \\citealp{srl}). Also the near-infrared spectral indices of quasars from the 3CR sample are correlated with luminosity, whereas the optical spectral indices are uncorrelated with the quasar luminosity or orientation, so that the strength of the `big red bump', relative to the ionizing continuum, appears to be less in the more luminous objects \\citep{sr,sconf}.  However, there is significant scatter in the relationship between radio and emission-line luminosities and a receding torus is not the only way of explaining the differences in the emission-line properties of radio quasars and radio galaxies or the luminosity-dependence of the quasar fraction (the fraction of objects that show broad emission lines). Two-population radio luminosity functions (RLFs) have been used to provide a best-fit to the 3CRR, 6CE and 7CRS radio source redshift surveys and radio source counts \\citep{wrlf}. It is possible that a two-population model with a simple unified scheme in one population, combined with the effects of scatter, could mimic the effects of a receding torus in producing both emission-line differences between radio galaxies and radio quasars and the gradual rise in quasar fraction with emission-line luminosity. \\citet{wqf} found a drop in the quasar fraction of AGN at low luminosity, postulating that their results are consistent with either the emergence of a second population of low-luminosity radio sources, which lack a well-fed nucleus, or a gradual decrease in $\\Theta_{{\\rm trans}}$ with decreasing radio luminosity.  Traditionally, radio galaxies have been divided into two populations, FRI and FRII, \\citep{fr}, based on radio structure but correlated with radio luminosity such that above $\\log_{10}(L_{151} / {\\rm W\\,Hz^{-1}\\,sr^{-1}}) \\sim 25$, objects are FRII, whereas below that critical $L_{151}$, they are FRI. \\citet{ccc02} found that the FRII population is inhomogeneous and that not all of them can be unified with quasars. They find a low-radiative-efficiency accretion, weak or absent broad-line emission and a lack of a significant nuclear absorbing structure for weak-jet, low-ionisation narrow-line galaxies. For broad-line objects and obscured high-ionization narrow-line galaxies they see or infer intense ionizing emission, powerful jets and a torus-like absorber. In low-radio-frequency-selected samples, the fraction of objects with observed broad lines changes rapidly from $\\sim 0.4$ above $\\log_{10}(L_{151} /{\\rm W\\,Hz^{-1}\\,sr^{-1}}) \\sim 26.5$ to $\\sim 0.1$ below this. The less luminous population from \\citet{wrlf} is composed of FRIs and FRIIs with weak/absent emission lines and their more luminous population of strong-emission-line FRII radio galaxies and quasars. A two-population model like this is motivated by the fact that the presence or absence of a quasar nucleus (as shown by emission-line strength) seems likely to be connected to the properties of a central engine of the radio source. \\citet{ka97}, and others before them, have argued that radio structure is influenced by the environment. It seems likely that a high jet power is necessary for a highly collimated non-dissipative jet, thus all high luminosity sources are FRII, but as lower jet powers are reached, a jet is more likely to disrupt resulting in a FRI source. The exact environmental density and homogeneity would determine the radio luminosity threshold of FRI/FRII within the lower luminosity population. \\citet{lo96} found that the FRI/FRII division is proportional to the square of the optical luminosity of the host galaxy, which is plausibly related to such environmental effects, although other causes have been proposed. This simple picture is shown to be incomplete by the discovery by \\citet{br} of an optically powerful quasar with FRI radio structure. It also seems that there is evidence that some FRI sources in 3C have hidden quasar nuclei \\citep{cao}, and one clear example of a 3CR FRI with broad lines (3C386, \\citealp{s96}) is known. In this paper we follow \\citet{hl}, \\citet{laing} and \\citet{jw} in assuming that the FRI/FRII division is strongly influenced by environmental effects and is much less fundamental than a division based on accretion properties of the central engines.  It is well known (e.g. \\citealt{dp90}) that the evolution in comoving space density of radio galaxies depends on both $L_{151}$ and $z$. In two-population models this can be expressed as differential density evolution of these populations \\citep{wrlf}. \\citet{sb} have criticized this approach on the basis of the artefacts it produces when the populations join.  The primary motivation of this work was to investigate in a quantitative manner the differences between the emission-line properties of radio-loud quasars and radio galaxies alongside the luminosity dependence of the quasar fraction. In this paper we will quantify the three-dimensional distribution in [OIII] emission-line luminosity $L_{{\\rm [OIII]}}$, $151$ MHz luminosity $L_{151}$ and redshift $z$ using a new method applied to complete samples. A principal components analysis (PCA) is used to find the axis which causes the most differentiation between the objects, and it is possible that it can identify parameters which are more physically meaningful than the attributes of each data point (radio and emission-line luminosity) used in the PCA. We carry out a PCA and define a generalized luminosity function (GLF) in this new parameter space. The GLF is then used to make proper comparisons between receding-torus and two-population models, allowing for cosmic evolution of the populations.  The paper is organized as follows. In Sec.~\\ref{sec:data} the data are described. In Sec.~\\ref{sec:pca}, the principal components analysis is described. In Sec.~\\ref{sec:model}, simple one-population GLFs are described to enable a comparison between unified schemes with and without a receding torus. The effects of using  two-population GLFs are presented in Sec.~\\ref{sec:2pop}, and in Sec.~\\ref{sec:discussion} we compare our results with previous studies and discuss possible meanings for the new parameters found by the PCA. We assume throughout that ${\\rm H_0} = 70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}, \\Omega_{\\rm M} = 0.3$ and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "A new approach to investigating unified schemes has been presented, based on a principal components analysis of the 3CRR, 6CE and 7CRS complete samples. Generalized luminosity functions have been derived based on these principal components: $\\alpha$ which encodes the $L_{151} - L_{\\rm [OIII]}$ correlation and $\\beta$ which encodes the scatter about this correlation. The main advantage of this new approach has been that the unified scheme parameters have been found by taking into account the intrinsic scatter in the $L_{151} - L_{\\rm [OIII]}$ correlation. The main conclusions to be drawn from this analysis are as follows.  \\begin{enumerate} \\item A receding-torus model is strongly favoured  over a simple non-luminosity-dependent unified scheme for GLFs with one population of radio sources.  \\item With the extra constraint of 6C and 7C radio source counts and the normalization of the local RLF, two-population GLFs were derived. The GLFs give rise to RLFs which are very similar but smoother than those of \\citet{wrlf}. There is very little difference in likelihoods between a GLF with a receding torus in both populations, a GLF with a torus opening angle that does not vary with ionizing luminosity in both populations, and a GLF with a receding-torus in the high-$\\alpha$ population and a constant-opening-angle torus in the low-$\\alpha$ population.  \\item Two-population models reproduce the radio survey data well, and can provide a natural explanation for the rise in quasar fraction with emission-line luminosity and the emission-line differences between radio quasars and radio galaxies. The receding-torus may be a feature in both populations but this is not yet proved. \\end{enumerate}"
    },
    "0311/astro-ph0311414.txt": {
        "abstract": "This article gives an analysis of the behavior of polarizing grids and reflecting polarizers by solving Maxwell's equations, for arbitrary angles of incidence and grid rotation, for cases where the excitation is provided by an incident plane wave or a beam of radiation. The scattering and impedance matrix representations are derived and used to solve more complicated configurations of grid assemblies. The results are also compared with data obtained in the calibration of reflecting polarizers at the Owens Valley Radio Observatory (OVRO). From these analysis, we propose a method for choosing the optimum grid parameters (wire radius and spacing). We also provide a study of the effects of two types of errors (in wire separation and radius size) that can be introduced in the fabrication of a grid. ",
        "introduction": "The literature on wire grids is abundant and they have been studied with different techniques and for numerous applications. Most of the analysis were however restricted to special cases of incident field and grid orientations. The more general and arbitrary situation seems to have been first studied by Wait (see \\citet{Wait 1955a} and \\citet{Larsen}). This problem is addressed again in this paper and follows a line of analysis fairly similar to the one used by Wait. Our treatment is, however, more general in that we do not assume that the wires of the grid are induced with only a longitudinal current; we will indeed show that an azimuthal component is also present. We also solve for the induced current by considering the tangential components of both the electric and magnetic fields at the surface of the wires.  This analysis is carried out in the next two sections and will serve as our basis for the treatment of the reflecting polarizer (section \\ref{sec:RP}) and the introduction of the scattering and impedance matrix representations for a grid (section \\ref{sec:imp}) which will in turn enable us to briefly discuss more complicated systems. These matrices will be particularly useful in allowing us to define what will be called the principal axes of a grid. These are two orthogonal and independent directions of polarization in the plane of the incident radiation along which an arbitrary electric field can be decomposed and shown to scatter without cross-polarization. With this representation at hand, it will then be possible to derive a set of optimal parameters (wire radius and spacing) to be used in the selection of a grid. We will also present an analysis of the effects of random errors that can be introduced in the fabrication of grids, the results obtained will then be compared to experimental results previously published by \\citet{Shapiro}.  The last section will be dedicated to the study of the more subtle impacts that the nature of the incoming radiation can have on the response of a grid assembly such as a reflecting polarizer (section \\ref{sec:RP}). Although limited to this particular case, our discussion could possibly apply to other types of instruments. We have also included at the end (Appendix B) a list of the symbols used in the different equations. ",
        "conclusions": "In this paper, a general solution for the analysis of polarizing grids was presented; it is valid for arbitrary angles of incidence and of grid rotation. With it and the scattering matrix representation that derives from it, basically any configuration or system of grids can be analyzed as long as some assumptions concerning the wire radius and spacing are respected ($\\lambda>40\\, a$ and $d>4\\, a$). This is not a severe restriction as most grid currently available satisfy those conditions, we refer the reader to \\citet{Chambers 1986, Chambers 1988} for cases where a larger size of wire is needed. Our analysis also allowed us to define a set of optimum values for both the wire radius and spacing as specified by the following equations:  \\begin{eqnarray*} a & \\simeq & \\left[\\frac{\\lambda^{5}}{\\left(1-\\alpha^{2}\\right)^{4}\\pi^{7}\\sigma Z_{o}}\\right]^{\\frac{1}{6}}\\\\ d & \\simeq & 2\\pi a\\,.\\end{eqnarray*}   We provided an analysis of the effects that two types of random errors can have on the performance of a grid. It was shown that errors in the wire spacing were the most important and could have some impact on the amount of unwanted polarization transmitted through a grid. In that respect, our model showed to be in good agreement with the experimental results of \\citet{Shapiro}.  Comparisons with experimental data obtained in the calibration of a reflecting polarizer used at the OVRO were also presented and predictions from our model are in good agreement with it. The only discrepancies appeared in the nature of a resonance, more precisely its width. But we have shown that it could be accounted for by including in the analysis a proper treatment of the effects of the nature of the incident radiation on the response of the polarizer.  We are grateful to J. B. Shapiro and E. E. Bloemhof for their permission to use their previously published experimental results. We wish to thank the staff of the Owens Valley Radio Observatory, and O. P. Lay for numerous discussions and suggestions. The Owens Valley Radio Observatory is funded by the National Science Foundation under Contract No. AST 96-13717 and the polarimetry project at OVRO through NASA grant NAG5-4462. M. H. work was supported in part by a grants from FCAR and the d\u00e9partement de physique de l'Universit\u00e9 de Montr\u00e9al.  \\pagebreak  \\appendix"
    },
    "0311/astro-ph0311317_arXiv.txt": {
        "abstract": "We have used high resolution {\\it WFPC2-HST} and wide field ground-based observations to construct a catalog of blue straggler stars (BSS) in the globular cluster 47 Tuc spanning the entire radial extent of the cluster.  The BSS distribution is highly peaked in the cluster center, rapidly decreases at intermediate radii, and finally rises again at larger radii. The observed distribution closely resembles that discovered in M3 by Ferraro et al (1993,1997).  To date, complete BSS surveys covering the full radial extent (\\hst\\ in the center and wide field CCD ground based observations of the exterior) have been performed for only these two clusters.  Both show a bimodal radial distribution, despite their different dynamical properties.  BSS surveys covering the full spatial extent of more globular clusters are clearly required to determine how common bimodality is and what its consequence is for theories of BSS formation and cluster dynamics. ",
        "introduction": "\\label{sec:intro} Blue straggler stars (BSS) are found primarily in star clusters, where they appear as a sparsely populated extension of the main sequence above the turnoff point. Superficially they appear to be main sequence stars with masses larger than expected for the cluster at its turnoff age. There are two proposed mechanisms thought to produce BSS: the first is mass exchange in a binary system and the second is the merger of two stars induced by stellar interactions (either single or in binaries) in a dense stellar environment. Globular cluster cores are obvious targets for BSS not only because of their high stellar density, but because primordial binaries in the clusters may well have sunk to the cores.  With the advent of the Hubble Space Telescope (\\hst) it became possible for the first time to search dense cluster cores for BSS. Searches in the ultraviolet (UV) can be particularly effective, not because BSS are especially bright in the UV, but because the red giants are faint in the UV. The photometric blends which mimic BSS in visible CMDs are far less problematic in the best UV CMD planes, and it is possible to obtain complete BSS samples in even the densest cores. With \\hst, the densest cluster cores are {\\it the obvious place} to search for BSS, and there have been few systematic searches for BSS in the outer parts of clusters in recent years. Indeed only one cluster, M3, has been adequately surveyed for BSS over its entire radial extent \\citep{f93,f97}.  Since it is a massive cluster with high central density, 47~Tuc is an obvious target for a BSS search, and its core has been the subject of several earlier investigations \\citep{gua92,demarchi47bss,ferraro47bss,knigge47}. These studies have covered only a fraction of the core either because they used the Faint Object Camera or just the PC of WFPC2. Albrow et al. (2001) found BSS in their WFPC2 search for eclipsing biniaries in the core of 47~Tuc. Blue stragglers have also been studied outside the core of 47 Tucanae by Sills et al. (2000) who also modeled the BSS formation rate. Kaluzny et al. (1998) note several BSS candidates in their CMD of 47~Tuc, and there are clear BSS candidates in the Str\\\"omgren photometry of Grundahl et al. (2002). Only a small part of the cluster exterior was covered by these studies.  With new wide field imagers on ground based telescopes at sites with excellent seeing it is now possible to perform BSS surveys which yield complete samples over the entire extent of the cluster. Here we present results covering both the full cluster core using \\hst\\ WFPC2, and the exterior of the cluster using the wide field imager at the ESO 2.2\\,m telescope. ",
        "conclusions": "F97 argued that the bimodal distribution of BSS in M3 was evidence that two formation scenarios were at action in the same cluster: the Exterior BSS (EBSS) arising from mass transfer in primordial binaries and the Central BSS (CBSS) arising from stellar interactions which lead to mergers. As had been earlier noted by Bailyn \\& Pinsonneault (1995), the EBSS and CBSS luminosity functions differed in the sense theoretically expected for two mechanisms.  One difficulty with primordial binaries as a source for the EBSS is that one ordinarily would expect that they would have settled to the cluster centers since relaxation times are typically less than 1~Gyr \\citep{djor-physparm}. Simulations to address the question of how many primordial binaries can form in the external regions of the cluster and how long primordial binaries with mass approaching twice the turnoff mass could remain in the outer parts of a cluster are now in progress (M. Mapelli et al 2003, in preparation).   Sigurdsson \\etal\\ (1994) offered another explanation for the bimodal BSS distribution in M3. They suggested that the EBSS were formed in the core and then ejected into the outer regions by the recoil from the interactions. Those binaries which get kicked out to a few $r_c$ rapidly drift back to the center of the cluster due to mass segregation, leading to a concentration of BSS near the center and a paucity of BSS in the outer parts of this region.  More energetic kicks will take the BSS to larger distances; these stars require much more time to drift back toward the core and may account for the overabundance of BSS at large distance.   How does the discovery of a bimodal radial distribution of BSS in 47~Tuc affect these arguments? Dynamically 47~Tuc is quite  different from  M3: its central density is 40 times larger \\citep{pm93}, and its core is rich in X-ray binaries, millisecond pulsars, and other interaction products \\citep{ferraro47bss,grindlay01, camilo00,e03}. We have already shown \\citep{bss6} that CBSS populations cannot be characterized by a simple parameter like collision rate, so the fact that CBSS in 47~Tuc and M3 have comparable specific frequencies should not come as a surprise.  It should be noted that an BSS specific frequency upturn similar to that observed in 47 Tuc and M3 has been detected also in M55 (Zaggia et al. 1997).  This is of particular significance because M55 has a low central density---0.1 that of M3. Note that although the data set presented by Zaggia et al. for M55 are ground based observations and covered only a quadrant of the cluster, there is at least preliminary evidence that a upturn is present also in a low density cluster like M55. We note that the effect found by Zaggia et al in M55 could be even stronger since ground-based observations tend to hide BSS in the central region of the cluster. For example ground based observations of the center of M3 led Bolte et al. (1993) to claim depletion of BSS in the center whereas \\hst\\ observations eventually found a peak. M55 is surely less concentrated than M3, but still intermediate crowding conditions and the presence of highly saturated giants in the center can produce some level of incompleteness at the BSS level.  The implications of the bimodality and relatively large populations of EBSS are not yet clear beyond the fact that EBSS may be fairly common and form in diverse clusters. If the BSS bimodality in M55 is confirmed, we will be faced with the result that clusters with central densities ranging over a factor of 400 can produce a bimodal BSS distributions. Any dynamical model like that of Sigurdsson et al. would seem more applicable to 47~Tuc than M3. It would be a real stretch to reach the dynamical conditions in M55.  Given these factors it seems even more likely that primordial binaries play a fundamental role in the BSS formation mechanism. Better information on the abundance of primordial binaries in clusters and their survival rate at a given distance from the cluster center is essential.  In 47 Tuc a recent extensive search for binary stars in the core (Albrow et al 2001) has led to an overall binary frequency of $14\\%\\pm 4\\%$. Their Fig.~21 shows that the binaries appear to be significantly more centrally concentrated with respect to the normal cluster stars but less concentrated than the BSS. Hence, if the cluster is in dynamical equilibrium, the CBSS are among the heaviest visible stellar components in the core of 47 Tuc. Albrow et al. also note that the W~UMa stars form a sequence in the CMD which like that which would be expected for equal mass binaries. The brightest lie in the BSS region of the CMD. Since W~UMa systems are contact binaries they may well merge to become BSS.  Unfortunately, binary populations have been measured in only a few additional clusters (most notably NGC~6752, NGC~288), and they probably have diverse origins: {\\it (i)} in the core of NGC~6752 \\citep{ruba1} the binaries are probably not primordial, since NGC~6752 is such a dynamically evolved cluster \\citep[and references there-in]{ferraro6752bh}; {\\it (ii)} in the low density cluster NGC~288 \\citep{bellazzini288bin}, the binaries are probably primordial. Despite the much longer relaxation time of NGC~288, the binaries have all settled to within the inner half light radius; The technique used by Rubenstein \\& Bailyn (1997) and Bellazzini et al. (2002) requires only modest \\HST\\ observation time and should be applicable to any cluster. Even though such studies of the binary populations are painful, we need more.   Clearly, it is desirable to have full BSS surveys for more clusters. It may well be a fluke that the first two clusters to be fully surveyed have bimodal distributions. We have pointed out \\citep{bss6} that all properly studied clusters have CBSS. The existence of EBSS in M3 \\citep{sandage53} marked the discovery of BSS. Similarly there is clear evidence for EBSS in ground based CMDs of 47~Tuc \\citep{sills00,grundahl47,kaluzny47}. Quite possibly all clusters with significant EBSS populations are bimodal. On the other hand, high quality CMDs of other clusters like M5 \\citep{sandquistm5} and M80 \\citep{brocatom80} show evidence for few if any EBSS. If more careful studies of such clusters do not reveal a yet undetected EBSS population, we will have to search for a mechanism which produces bimodal radial BSS distributions in some clusters but not in others.   The population of BSS discovered here opens a new window for the investigation of the origin and the formation mechanism of BSS in GGCs. In fact, this discovery  suggests that the {\\it peculiar} radial distribution firstly found in M3 is much more {\\it common} that was thought. Possibly in a few years we will refer to Figure 5 as the typical radial BSS distribution in stellar aggregates. In the meanwhile let us still   refer to it as {\\it peculiar} until its generality can be finally addressed.  We are now exploiting the wealth of information in the large dataset presented here by performing an extensive spectroscopic survey of the BSS population in 47 Tuc using the new multi-fiber Spectrograph FLAMES at the ESO-Very Large Telescope. These data should be highly informative. For instance, radial velocity determinations for a significant number of EBSS would help to clarify their origin: if the EBSS were ejected from the cluster center their velocity dispersion should be low because they are expected to be observed near the apocenter of an highly elliptical orbit; if EBSS arise from primordial binaries the velocity dispersion might be close to the Keplerian velocity (circular orbits which avoid the cluster center might allow the binaries to remain in the exterior) and show small variations due to low mass companions. The same spectra could lead to many mass estimates: so far the only BSS with a measured mass ($1.7 \\pm 0.4\\,\\msun$) is in 47~Tuc and was obtained with \\hst\\ \\citep{shara-bssmass}."
    },
    "0311/astro-ph0311207_arXiv.txt": {
        "abstract": "The Major Atmospheric Gamma ray Imaging Cherenkov Telescope (MAGIC) is in commissioning phase and will start to become fully operative by the end of 2003. Located at {\\it El Roque de los Muchachos} in La Palma (Canary Islands, Spain), it has the largest reflector area (17 m diameter) of all the existing Cherenkov telescopes. New technologies have been used to reduce the energy threshold for gamma-ray detection to about 30 GeV. Due to its characteristics, the catalog of very high energy sources will considerably increase with the MAGIC observations, anticipating exciting results for the near future. An overview of the telescope, its current status and first results, together with a highlight of the scientific research is presented. ",
        "introduction": "Gamma ray astronomy has provided in the las few years spectacular results, led by the success of the CGRO ({\\it Compton Gamma Ray Observatory}) satellite mission which revealed that the high energy Universe was more exciting than expected. During its lifetime, the EGRET telescope onboard the CGRO provided an all-sky survey above 100 MeV consisting of 271 sources (Hartman et al. 1999). Apart of the remarkable detection of 66 AGN Blazars, the vast majority of these sources are still unidentified. Almost during the same period, the ground-based gamma ray astronomy has also been developed. The technique consists of the detection of the atmospheric Cherenkov light emitted by the particle showers initiated by gamma radiation on entering into the atmosphere. This Cherenkov flash lasts for a few {\\it ns}. The Whipple Cherenkov telescope opened this ground-based era by the detection of the Crab Nebulae in 1989. Then, extragalactic Blazar sources not seen by satellites were also discovered by using this technique. Other instruments based on the same principle have confirmed the Whipple results (HEGRA, CANGAROO, CAT) and nowadays about a dozen gamma ray sources have been detected to emit in the TeV energy range. The observations done by satellites measured well below around 10 GeV, while the existing Cherenkov telescopes have detected sources above 300 GeV. This energy gap, which is still virtually unexplored, is really important to understand which is the origin of the cut-off on the spectra which will explain the lack of sources measured by the Cherenkov telescopes in comparison to EGRET.  The MAGIC (Major Atmospheric Gamma Imaging Cherenkov) telescope was designed in 1998 (Barrio et al. 1998) with the main goal of being the Imaging Atmospheric Cherenkov Telescope (IACT) with the lowest gamma energy threshold possible with the technological improvements affordable and based on the experience acquired with the first generation of Cherenkov telescopes. By using this detection technique, which provides much large effective areas (and much superior flux sensitivity) than satellite detectors, good angular resolution, acceptable energy resolution and a well tested capability to separate gammas from backgrounds, eventually a plethora of new sources will be discovered since for most of known sources the energy spectrum is of power-law nature and therefore they should exhibit a much higher flux in that energy region than at higher energies.   \\begin{figure}[h] \\plotfiddle{DSCN0011.eps}{5,5cm}{0}{27}{27}{-115.}{0.} \\caption{View of the MAGIC Telescope in August 2003. Almost all the mirrors are already on place.} \\end{figure}  The 17 m diameter f/D=1 MAGIC telescope is the largest of the new generation of IACTs. MAGIC is located in the Canarian island of La Palma (28.8 N, 17.9 W) at the {\\it Roque de los Muchachos} observatory (ORM), 2200 m above sea level. Its 241 m$^2$ parabolic dish is composed of 964 $49.5\\times 49.5$ cm$^2$ all-aluminum spherical mirror tiles mounted on a lightweight ($< 10$ ton) carbon fiber frame. The parabolic shape was chosen to minimize the time spread of the Cherenkov light flashes on the camera plane, which allows to reduce the rate of fake events induced by night-sky background light. Mirrors are grouped in panels of three or four, which can be oriented during the telescope operation through a novel active mirror control system to correct for the possible deformations of the telescope structure.  The camera is made of 577 good quantum efficiency, fast photomultipliers with hemispherical photocatode that allows for light double-crossing. Each photomultiplier is coupled to a small light collecting cone to maximize the active surface of the camera. An special wavelength-shifting coating provides red extended sensitivity and allows for light-trapping, which increases the photomultiplier effective quantum efficiency. The total field of view of the camera is of about 4$\\deg$. The photomultiplier signals are transmitted to a distant Control House ($\\sim$150 m) by using analog optical fiber signals. Signals are processed by a multilevel trigger system and 300 MHz FADC are used for pulse digitalization.  While all other new generation Cherenkov telescopes aim at the improvement of sensitivity and energy resolution in the 100 GeV regime by using stereoscopic systems of relatively small (10 m) telescopes, MAGIC, through the choice of a single, larger reflector, will achieve the lowest energy threshold among IACTs, of about 30 GeV. Its altazimuth mount can point to anywhere in the sky in about 20 seconds, a unique feature which is essential for the study of transient events like GRBs. ",
        "conclusions": "The MAGIC telescope is in its final commissioning phase and it is expected to start regular observations by the end of this year. If the telescope behaves as expected, it will soon be able to provide exciting results on a wide variety of astrophysical phenomena."
    },
    "0311/astro-ph0311177_arXiv.txt": {
        "abstract": "Using {\\it Far Ultraviolet Spectroscopic Explorer} (FUSE) and {\\it Hubble Space Telescope} (HST) observations of the QSO PG\\,1259+593, we detect \\ion{D}{1} Lyman-series absorption in high velocity cloud Complex~C, a low-metallicity gas cloud falling onto the Milky Way.  This is the first detection of atomic deuterium in the local universe in a location other than the nearby regions of the Galactic disk. We construct a velocity model for the sight line based on the numerous \\ion{O}{1} absorption lines detected in the ultraviolet spectra. We identify 8 absorption-line components, two of which are associated with the high velocity gas in Complex~C at $\\approx -128$ and $\\approx -112$ \\kms. A new Westerbork Synthesis Radio Telescope (WSRT) interferometer map of the \\ion{H}{1} 21\\,cm emission toward PG\\,1259+593 indicates that the sight line passes through a compact concentration of neutral gas in Complex~C.  We use the WSRT data together with single-dish data from the Effelsberg 100-meter radio telescope to estimate the \\ion{H}{1} column density of the high velocity gas and to constrain the velocity extents of the \\ion{H}{1} Lyman-series absorption components observed by FUSE.  We find N(\\ion{H}{1}) = $(9.0\\pm1.0)\\times10^{19}$ cm$^{-2}$, N(\\ion{D}{1}) = $(2.0\\pm0.6)\\times10^{15}$ cm$^{-2}$, and N(\\ion{O}{1}) = $(7.2\\pm2.1)\\times10^{15}$ cm$^{-2}$ for the Complex~C gas (68\\% confidence intervals). The corresponding light-element abundance ratios are D/H = $(2.2\\pm0.7)\\times10^{-5}$, O/H = $(8.0\\pm2.5)\\times10^{-5}$, and D/O = $0.28\\pm0.12$. The metallicity of Complex~C gas toward PG\\,1259+593 is approximately 1/6 solar, as inferred from the oxygen abundance [O/H] = $-0.79\\pm^{0.12}_{0.16}$. While we cannot rule out a value of D/H similar to that found for the local ISM (i.e., D/H $\\sim 1.5\\times10^{-5}$), we can confidently exclude values as low as those determined recently for extended sight lines in the Galactic disk (D/H~$< 1\\times10^{-5}$).  Combined with the sub-solar metallicity estimate and the low nitrogen abundance, this conclusion lends support to the hypothesis that Complex~C  is located outside the Milky Way, rather than inside in material recirculated between the Galactic disk and halo. The value of D/H for Complex~C is consistent with the primordial abundance of deuterium inferred from recent {\\it Wilkinson Microwave Anisotropy Probe} (WMAP) observations of the cosmic microwave background and simple chemical evolution models that predict the amount of deuterium astration as a function of metallicity. ",
        "introduction": "Observations of the abundance of deuterium relative to hydrogen (D/H) in different environments provide insight into the evolution of the light elements in the universe.  With the excellent concordance in the estimates of the cosmic baryon density from measurements of D/H in low-metallicity quasar absorption-line systems and measurements of the cosmic microwave background, the cosmic baryon density and the primordial value of D/H are now tightly constrained (Burles, Nolett, \\& Turner 2001; O'Meara et al.\\ 2001; Spergel et al.\\ 2003).  Thus, it should be possible to test chemical evolution models by examining the progression of D/H with time (see Lemoine et al.\\ 1999 and Olive, Steigman, \\& Walker 2000 for recent discussions).  The abundance of deuterium is expected to decrease with time since there are no known sources of deuterium capable of increasing the cosmic abundance significantly (Epstein, Lattimer, \\& Schramm 1976; see also Prodanovi\\'c \\& Fields 2003). Many chemical evolution models predict moderate levels of deuterium destruction by stellar nucleosynthesis, typically less than a factor of 3--5 (Clayton 1985; Edmunds 1994; Steigman \\& Tosi 1995; Tosi et al.\\ 1998). More recent models suggest that slightly lower levels of astration are also possible (e.g., Chiappini, Renda, \\& Matteucci 2002).   Unfortunately, measurements of D/H are particularly difficult, and there have been relatively few opportunities to directly measure the detailed changes in the abundance of deuterium as a function of metallicity or time.  A key piece of information missing in discussions of the evolution of the light element abundances with time is the behavior of D/H in environments with metallicities between those of the high-redshift systems (typically $Z\\lesssim0.01Z_\\odot$) and those of gas in the local neighborhood of the Sun (typically $Z \\sim Z_\\odot$).  There are several reasons why it it is important to determine the D/H ratio in a wide variety of galactic and extragalactic environments. First, there have been few high-precision estimates of D/H at moderate to high redshifts ($z \\sim 2-4$), where the amount of stellar processing of deuterium is presumably low, as evidenced by low metallicity (e.g., O'Meara et al.\\ 2001; Pettini \\& Bowen 2001 and references therein). The measurements that have been made appear to yield conflicting values for the primordial abundance of deuterium, with the observed values depending on the type of system observed (e.g., Lyman-limit or damped Ly$\\alpha$ systems -- see Pettini \\& Bowen 2001).  Second, estimates of D/H in many locations yield a more global perspective of the chemical history of gas at different epochs than is possible from a few isolated measurements.  Chemical evolution models seeking to describe the general evolution of the light-element abundances need a large sample of measurements to avoid systematic problems encountered by relying upon data for only a few types of environments. Third, although measurements of deuterium in nearby gas clouds  imply a relatively constant value of D/H within the local interstellar medium (ISM; Moos et al.\\ 2002 and references therein), substantial variations may exist in D/H and D/O over distances of only a few hundred parsecs (Jenkins et al.\\ 1999; Sonneborn et al.\\ 2000; Hoopes et al.\\ 2003; H\\'ebrard \\& Moos 2003). If a sufficiently large number of high-precision D/H and D/O measurements can be made in a diverse set of nearby environments, it may be possible to understand the exact causes of this variability and the degree to which galactic chemical evolution and accretion of intragroup gas clouds  influence the scatter in the observed ratios, both locally and at high redshift. This goal is a major science driver for the {\\it Far Ultraviolet Spectroscopic Explorer} (FUSE) mission (Moos et al.\\ 2000).  To bridge the gap in D/H between low and high metallicity environments, we have obtained an extensive set of FUSE, {\\it Hubble Space Telescope} (HST), and interferometric \\ion{H}{1} 21\\,cm observations of the quasar PG\\,1259+593 behind high velocity cloud (HVC) Complex~C.  The HVC is located at least 3.5 kpc from the Galactic plane (Wakker 2001), well beyond all Milky Way clouds with current D/H determinations. Unlike previous investigations of D/H and D/O in either high-redshift clouds or the local ISM, we know which gaseous system is responsible for the high-velocity \\ion{D}{1} Lyman-series absorption observed toward PG\\,1259+593.  A global description of the gas in Complex~C is available from both emission and absorption-line measurements.  The neutral gas in Complex~C has been mapped extensively in \\ion{H}{1} 21\\,cm emission (see Wakker 2001 and references therein) and low-ionization absorption (e.g., Wakker et al.\\ 1999; Richter et al.\\ 2001b; Collins, Shull, \\& Giroux 2003). The ionized gas in Complex~C has been investigated in absorption by Sembach et al.\\ (2003) and Fox et al.\\ (2003), and in emission by Tufte, Reynolds, \\& Haffner (1998) and Wakker et al.\\ (1999). Complex~C is an excellent site to determine D/H for  comparisons with the high-redshift values because it is chemically young (Richter et al.\\ 2001b; Collins et al. 2003; Tripp et al.\\ 2003) and has a metallicity (10--25\\% solar) lower than that of the general ISM of the Milky Way and higher than that of intergalactic clouds at high redshifts.  In this paper we describe these new measurements and the resulting D/H and D/O ratios in Complex~C.   In \\S2 we describe the FUSE and HST Space Telescope Imaging Spectrograph (HST/STIS) absorption-line observations and the \\ion{H}{1} 21\\,cm interferometer observations.  Section~3 contains a short summary of the properties of the PG\\,1259+593 sight line.  In \\S4 we outline the methods and general assumptions used to determine the column densities of \\ion{H}{1}, \\ion{D}{1}, and \\ion{O}{1} in Complex~C. We determine the \\ion{H}{1} column density from interferometric \\ion{H}{1} 21\\,cm emission data and use the FUSE and HST/STIS ultraviolet absorption-line data to determine the \\ion{O}{1} and \\ion{D}{1} column densities. Sections 5, 6, and 7 contain descriptions of these determinations for oxygen, hydrogen, and deuterium, respectively, and provide estimates of the various errors. The column densities and error ranges are summarized and discussed in \\S8, and comments on future progress appear in \\S9. We summarize the results of the study in \\S10. ",
        "conclusions": "We summarize the Complex~C column densities of \\ion{H}{1}, \\ion{D}{1}, and \\ion{O}{1} in Table~9, together with the $1\\sigma$ (68\\% confidence interval) and $2\\sigma$ (95\\% confidence interval) error estimates based on the descriptions of errors given above.  Systematic errors dominate the uncertainties for all three species.  We also list the  D/H, D/O, and O/H column density ratios, along with the propagated errors from the column density determinations. A summary of the light element abundance ratios for Complex~C and other environments is given in Table~10, along with references for the ratios. The environments considered in this table include the local ISM, the ISM of the Galactic disk at distances of 200--1000 pc from the Sun, and high-redshift absorption line systems (both Lyman-limit systems and damped Ly$\\alpha$ systems).  We list both D/H and D/O when estimates of [O/H] are available. These comparisons are not meant to be exhaustive summaries of all the available data in the literature; rather they are synopses of several studies given to provide insight into how the Complex~C values compare to typical values found elsewhere. The data in Table~10 are shown graphically in Figure~17, where we plot D/H as a function of metallicity determined through measurements of [O/H], or [Si/H] if [O/H] is not available.  For the Lyman-limit systems where [Si/H] is used, the metallicity may be somewhat uncertain because large ionization corrections are necessary to convert estimates of N(\\ion{Si}{2})/N(\\ion{H}{1}) into estimates of N(Si)/N(H). Various other metallicity indicators are available, but [O/H] tends to be the best for the reasons described in \\S5 (see also the detailed discussion presented by Timmes et al.\\ 1997). Our estimate of D/H in Complex~C rules out values of D/H greater than $3.3\\times10^{-5}$ or less than $1.1\\times10^{-5}$ with a reasonable degree of confidence ($\\sim95$\\%).  Higher values over-produce the amount of \\ion{D}{1} absorption in the higher order Lyman series lines, and lower values under-produce the amount of absorption expected in the Ly$\\delta$ line. Both limits have interesting implications.  Most models of Galactic chemical evolution predict that astration of deuterium should result in a factor of $\\sim1.5-3$ decline in the cosmic abundance of deuterium from the Big Bang to the present time (Edmunds 1994; Tosi et al.\\ 1998; Chiappini et al. 2002), though higher levels of astration may be possible if special conditions, such as a Galactic wind, are invoked to counter over-production of elements heavier than helium (see, e.g., Scully et al.\\ 1997).  For a simple closed-box model with the assumption of instantaneous recycling, a deuterium astration efficiency of 60\\%, and standard oxygen yields, Pagel (1997) finds that the there should be only a few percent reduction of deuterium relative to its primordial abundance if [O/H] $\\lesssim -1$, as found for Complex~C.  In Figure 17, we show the results for a simple chemical evolution model from Fields et al.\\ (2001) having (D/H)$_p$ and little destruction of deuterium at low metallicities from Population~III stars.  More complicated chemical evolution models requiring bimodal episodes of star formation to explain the scatter in D/H values measured at high redshift can be found in that work. Our $2\\sigma$ upper limit of D/H $< 3.3\\times10^{-5}$ suggests  that the value of D/H produced by Big Bang nucleosynthesis cannot be too much larger than this upper limit unless the assumptions about the destruction of deuterium and production of oxygen are invalid. Similar arguments have been made to infer a primordial value of (D/H)$_p \\lesssim 3\\times10^{-5}$ from determinations of D/H in high-redshift, low-metallicity systems (Tytler et al.\\ 2000; O'Meara et al.\\ 2001; Kirkman et al.\\ 2003). Values of (D/H)$_p \\lesssim 3\\times10^{-5}$ imply values of the baryon density, $\\Omega_bh^2$, and the baryon-to-photon ratio, $\\eta$, in concordance with estimates derived from cosmic microwave background (CMB) measurements.  Using the standard Big Bang nucleosynthesis predictions for the abundance of deuterium (Burles et al. 2001), a value of (D/H)$_p \\lesssim 3\\times10^{-5}$ implies $\\Omega_bh^2 \\gtrsim 0.02$ and $\\eta \\gtrsim 5.6\\times10^{-10}$.  The CMB estimates cluster near values of $\\Omega_bh^2 \\approx 0.022-0.024$ (Netterfield et al.\\ 2002; Pryke et al.\\ 2002; Spergel et al.\\ 2003), and imply (D/H)$_p = (2.62\\pm^{0.19}_{0.20})\\times10^{-5}$ (Spergel et al. 2003; see also Steigman 2003).  The region between the dashed horizontal lines in Figure~17 is the $\\pm2\\sigma$ range of (D/H)$_p$ allowed by the recent {\\it Wilkinson Microwave Anisotropy Probe} (WMAP) observations.  With the PG\\,1259+593 data, we can confidently rule out very high D/H values for Complex~C, like those claimed for some intergalactic systems [e.g., D/H $= (2.0\\pm0.5)\\times10^{-4}$ at $z\\approx0.7$ toward PG\\,1718+4807; Webb et al.\\ 1997].  This is important because the Complex~C measurement is the only other measurement of D/H outside the Milky Way at $z < 1.0$. Such high values are also inconsistent with the CMB measurements, and recently the value for the $z=0.7$ absorber toward PG\\,1718+4807 has been challenged on other grounds (e.g., Kirkman et al.\\ 2001 -- but see Crighton et al.\\ 2003).   On the low end of the D/H range, we cannot exclude the possibility that D/H in Complex~C is similar to D/H in the local ISM (D/H $\\sim1.5\\times10^{-5}$, Moos et al. 2002;  Linsky 2003).  However, we  can  rule out D/H values as low as those found for some Galactic disk sight lines extending beyond the local ISM (i.e., at $d \\gtrsim 100$ pc). Values of D/H less than $1.0\\times10^{-5}$ have been determined for two extended sight lines in the Galactic disk using FUSE data.  Hoopes et al.\\ (2003) find D/H = $(0.85\\pm^{0.34}_{0.34})\\times10^{-5}$ ($2\\sigma$) toward HD\\,195965 ($d \\sim 800$ pc), and D/H = $(0.78\\pm^{0.52}_{0.25})\\times10^{-5}$ ($2\\sigma$) toward HD\\,191877 ($d \\sim 2200$ pc).  Using IMAPS data, Jenkins et al.\\ (1999) find D/H = $(0.74\\pm^{0.19}_{0.13})\\times10^{-5}$ (90\\% confidence) toward $\\delta$\\,Ori\\,A ($d \\sim 500$~pc); a similar result was found by Laurent, Vidal-Madjar, \\& York (1979) using Copernicus data. The low values found for these sight lines indicate that more astration of deuterium may have occurred in these regions than in the local ISM or in Complex~C. The local star formation histories of the regions explored toward $\\delta$\\,Ori\\,A and HD\\,195965 are roughly consistent with this idea (see Hoopes et al.\\ 2003 and Jenkins et al.\\ 1999), but the number of sight lines is still too small to draw general conclusions.  Eventually, with a large enough number of sight lines, it may be possible to determine whether refinements to  chemical evolution models are needed.   The value of D/O = $0.28\\pm0.12$ in Complex~C is very different from the D/O ratio in the local ISM or the disk of the Milky Way. H\\'ebrard \\& Moos (2003) find D/O = $0.038\\pm0.002$ for white dwarf and subdwarf sight lines within $\\sim150$ pc of the Sun.  Most of this difference can be accounted for by the different metallicities of the local ISM and Complex~C.  Note that even without the \\ion{H}{1} measurement, the D/O ratio strongly suggests that the metallicity in Complex~C is lower than solar.   The Complex~C result is more similar to the value of D/O = $0.37\\pm0.03$ found for the $z=3.025$ absorber toward Q\\,0347-3819 (Levshakov et al. 2002) than it is to the value for the Galactic disk.  However, values of D/O for  two other high-redshift systems are an order of magnitude or more larger than the Complex~C value.  We note that determinations of N(\\ion{O}{1}) in the high-redshift systems are challenging since the weaker \\ion{O}{1} lines occur within the \\ion{H}{1} Ly$\\alpha$ forest.   As estimates of D/H in the high-redshift systems have fallen in recent years, a few local measurements with high values of D/H have received more attention.  For example, the value of D/H = $(2.18\\pm^{0.36}_{0.31}) \\times10^{-5}$  for the $\\gamma^2$ Vel sight line (90\\% confidence; Sonneborn et al.\\ 2000) is indistinguishable from the values for the high-redshift ($z>2$) systems toward HS\\,0105+1619 and Q\\,1243+3407, and perhaps higher than the value for Q\\,2206--199 (see Table~10 and references therein). Preliminary FUSE results for the IX~Vel sight line, which lies $\\sim1\\degr$ from $\\gamma^2$ Vel, indicate that the value of D/H for the first $\\sim100$ pc of the $\\gamma^2$ Vel sight line is similar to that in the local ISM (Blair et al.\\ 2003). If this result is confirmed by upcoming FUSE observations, then the D/H ratio in the Vela region must be even higher than the sight line average.  Note, however, that the $\\zeta$~Pup sight line, which is in the same region of the sky and extends further than the $\\gamma^2$~Vel sight line has a D/H ratio similar to that of the local ISM (see Table~17). One possible explanation for a high value in the ISM toward $\\gamma^2$ Vel is infall of ``primordial'' (metal-poor, D-rich) gas.  Clouds as large as Complex~C could contribute significant amounts of deuterium to regions of the Galactic disk; a mass of M(\\ion{H}{1}) $\\gtrsim(1.2-3.0)\\times10^6$ M$_\\odot$ and D/H = $2.2\\times10^{-5}$ imply that Complex~C contains $\\gtrsim(26-66)$ M$_\\odot$ of deuterium. The effects of infalling low-metallicity gas have been considered by various authors in the context of deuterium astration and potential solutions to the well-known ``G-dwarf'' problem (Tosi 1988; Matteucci \\& Francois 1989; Tosi et al.\\ 1998; Wakker et al.\\ 1999).   However, a serious problem with this interpretation for the $\\gamma^2$~Vel region is that the gas along the sight line does not show a corresponding reduction in the ratios of metals to hydrogen as would be expected if infall were the explanation for the high D/H ratio. Furthermore, the constancy of the oxygen and krypton abundances within several hundred parsecs of the Sun (Meyer, Jura, \\& Cardelli 1998; Cartledge et al.\\ 2001)  supports the idea that the nearby gas is reasonably well mixed and has similar heavy element abundances. Any explanation for the high D/H ratio in the Vela region that incorporates infall must also include processing of the gas to explain the relatively ``normal'' metal abundances in the region.  Other possibilities also exist for explaining the variations of D/H in the Milky Way.  For example, D could be depleted onto dust grains (Jura 1982).  From simple thermodynamic considerations, Draine (2003) has shown that differences in the binding energies of D and H on the peripheries of polycyclic aromatic hydrocarbons (PAHs) could produce extreme enrichment of deuterium in bound form, similar to what has been found for some simple interstellar molecules.  He argues that it is plausible that roughly 20\\% of the H atoms bound to PAHs are replaced by D.  If this is correct, he estimates that there are sufficient PAH binding sites in a volume of normal interstellar material to account for a decrease of free D atoms by one part in $10^5$ of H (i.e., roughly the size of the reduction observed along some lines of sight).  In some places, the atomic medium can revert to the intrinsic D/H when the PAHs are destroyed by the recent passage of a shock, which could explain why D/H seems to vary.  Reinforcing the concept that variable D abundances might be explained by Draine's proposal is the recent detection of the CD bond stretch features at 4.4 and 4.67 $\\mu$m toward the Orion Bar reported by Peeters (2002).  She found ${\\rm D/H} = 0.17\\pm 0.07$ in bound form.  This detection, although not very far above the noise, indicates that Draine's estimate of a 20\\% replacement of H by D atoms may be approximately correct.  For any gas with an overall metallicity as low as that of Complex C, the concentration of PAHs would be so low that there would be no measurable reduction in the concentration of free deuterium atoms, which is consistent with our finding for the sight line toward PG\\,1259+593.  Finally, Mullan \\& Linsky (1998) have proposed that deuterium production in stellar flares may cause some variations in D/H locally, but Prodanovi\\'c \\& Fields (2003) have recently reconsidered this issue and find that flare production of deuterium is not a significant source of \\ion{D}{1}  on galactic scales.  They concur with Mullan \\& Linsky that local variations by stellar flare production cannot be ruled out. Dust depletion and stellar flare production of deuterium are unlikely to alter the present value of D/H in Complex~C substantially. Complex~C contains no known stars, and analyses of its metal content indicate that it contains little dust.   The metallicity of $Z/Z_\\odot\\approx0.17$ we derive for Complex~C through the ratio of \\ion{O}{1} and \\ion{H}{1} is very similar to the recent determination of the metallicity in Complex~C toward 3C\\,351 by Tripp et al.\\ (2003). They find [O/H] = $-0.76\\pm^{0.23}_{0.21}$.  Both metallicity estimates are slightly higher than the value of $Z/Z_\\odot\\sim0.1$ inferred for the Mrk~290 line of sight through Complex~C based on measurements of \\ion{S}{2} and \\ion{H}{1} (Wakker et al.\\ 1999).  The metallicity estimates may depend on environment and may change slightly as a function of position within Complex~C; Collins et al.\\ 2003 find a range from 0.1 to 0.25 solar, but some of these estimates are complicated by uncertain ionization corrections for \\ion{S}{2}. One possible explanation for the slightly varying abundances in Complex~C might be that the gas is interacting with the Galactic corona or high Galactic halo, as indicated by measurements of highly ionized gas associated with the complex (Sembach et al.\\ 2003; Fox et al.\\ 2003).  Mixing of the diffuse regions of Complex~C with the more metal abundant gas of the Galactic thick disk and halo would tend to increase the observed abundances in these regions, while the abundances in denser clumps of gas, as seen toward PG\\,1259+593, should reflect more closely the original (pre-interaction) abundances in the cloud.   It is interesting that our best-fit value of D/H found for Complex~C is similar to that found for high-redshift damped Ly$\\alpha$ systems (see Pettini \\& Bowen 2001 and Kirkman et al.\\ 2003).  The \\ion{H}{1} column density of N(\\ion{H}{1}) $=9.0\\times10^{19}$ cm$^{-2}$ is sufficient to produce damping wings on the \\ion{H}{1} Ly$\\alpha$ and Ly$\\beta$ lines.  If observed in isolation from the Milky Way by an outside observer, Complex~C  would produce  \\ion{H}{1} and \\ion{D}{1} absorption similar to that seen in these higher redshift systems.  Complex~C also has a low nitrogen abundance inferred from the ratio of \\ion{N}{1} to \\ion{H}{1}, [N/H] $\\approx -1.9$ toward PG\\,1259+593 (Richter et al.\\ 2001b; Collins et al. 2003), about a factor of 10 lower than [O/H].    (Initial work on the amount of \\ion{N}{2} present indicates that photoionization of \\ion{N}{1} in the \\ion{H}{1} regions is not able to account for the deficit in the derived value of [N/H].) Commensurate values of [N/H] have been found for other Complex~C sight lines (Collins et al. 2003; Tripp et al. 2003). Some damped Ly$\\alpha$ systems show similar deficits of nitrogen relative to oxygen, which several authors have explained as a nucleosynthetic effect resulting from delayed release of nitrogen between massive and intermediate mass stars or a top-heavy or truncated initial mass function (Prochaska et al.\\ 2002; Molaro 2003).  At the metallicity of Complex~C, the low N/O ratio observed is probably more dependent upon the secondary nature of nitrogen production than these other processes (see Figure~6 in Pettini et al. 2002). The metallicity, abundance pattern, D/H ratio, and lower distance limit all indicate that Complex~C is an external system (or the remains of an external system) falling into with the Milky Way rather than gas ejected from the Galactic disk.      We report the detection of \\ion{D}{1} Lyman-series absorption in high velocity cloud Complex~C, a low-metallicity gas cloud falling onto the Milky Way.  This is the first detection of atomic deuterium in the local universe in a location other than the nearby regions of the Galactic disk.  The primary results of our study are as follows:   \\begin{enumerate} \\item{We combine data from the {\\it Far Ultraviolet Spectroscopic Explorer}, {\\it Hubble Space Telescope}, and Westerbork Synthesis Radio Telescope to derive the column densities of \\ion{H}{1}, \\ion{D}{1}, and \\ion{O}{1} in the direction of PG\\,1259+593 ($l=120\\fdg56, b = 58\\fdg05$).}   \\item{We use the FUSE and HST/STIS absorption-line data to construct a velocity model for the sight line based on the numerous \\ion{O}{1} absorption-lines detected.  We identify 8 absorption-line components common to all three species (\\ion{H}{1}, \\ion{D}{1}, and \\ion{O}{1}) along the sight line. Two of these components, at $v_{\\rm LSR}\\approx -128$ and --112 \\kms, are identified with Complex~C. The remaining lower-velocity components are identified with the Intermediate Velocity Arch, the Milky Way ISM, and a positive intermediate-velocity cloud in this direction.}  \\item{Our WSRT interferometer map of the \\ion{H}{1} 21\\,cm emission toward PG\\,1259+593 indicates that the sight line passes through a compact concentration of neutral gas in Complex~C. We use the WSRT data having a resolution of $\\sim1\\arcmin$ together with single-dish data from the Effelsberg 100-meter radio telescope to estimate the \\ion{H}{1} column density in Complex~C and to constrain the velocity extents of the Complex~C \\ion{H}{1} Lyman-series absorption components observed by FUSE.}  \\item{We use a combination of analysis techniques to estimate the column densities of \\ion{H}{1}, \\ion{D}{1}, and \\ion{O}{1} in the absorbing regions along the sight line.  We combine line-profile fitting analyses and curve-of-growth techniques with $\\chi^2$ minimization codes to estimate the most likely value and uncertainties on the column density of each species. The use of multiple analysis methods provides important consistency checks on the results of the individual methods.  The good agreement found from independent analyses by members of our team confirm that the formal errors are reliable representations of the uncertainties involved in the analyses.}  \\item{For each species, we describe the uncertainties involved in the estimation of the column density.  Systematic errors dominate the uncertainties, especially for \\ion{D}{1}.  In estimating the uncertainties on D/H, we have included the possibility that some of the absorption in the vicinity of the \\ion{D}{1} lines might be caused by a weak, high-velocity \\ion{H}{1} feature  near --190 \\kms.}  \\item{For the Complex~C absorption, we find N(\\ion{H}{1}) = $(9.0\\pm1.0)\\times10^{19}$ cm$^{-2}$, N(\\ion{D}{1}) = $(2.0\\pm0.6)\\times10^{15}$ cm$^{-2}$, and N(\\ion{O}{1}) = $(7.2\\pm2.1)\\times10^{15}$ cm$^{-2}$.  The uncertainties quoted are $1\\sigma$ errors (68\\% confidence intervals);  two sigma error estimates can be found in Table~9. The corresponding light element abundance ratios are D/H = $(2.2\\pm0.7)\\times10^{-5}$, O/H = $(8.0\\pm2.5)\\times10^{-5}$, and D/O = $0.28\\pm0.12$.}   \\item{ The metallicity of the Complex~C gas in this direction is approximately 1/6 solar, as inferred from the oxygen abundance of [O/H] = $-0.79\\pm^{0.12}_{0.16}$~($1\\sigma$).  This metallicity is in excellent agreement with a recent high-quality determination of [O/H] in Complex~C toward 3C\\,351, which is located $\\sim30\\degr$ from PG\\,1259+593.  Two other Complex~C sight lines analyzed by Collins et al. (2003) yield similar metallicities (Mrk~279: [O/H] = $-0.71\\pm^{0.36}_{0.25}$;  Mrk~817: [O/H = $-0.59\\pm^{0.25}_{0.17}$). }  \\item{While we cannot rule out a value of D/H for Complex~C similar to that found for the local ISM (i.e., D/H~$\\sim1.5\\times10^{-5}$), we can reasonably exclude values as low as those determined recently for extended sight lines in the Galactic disk (D/H~$< 1\\times10^{-5}$), provided that the absorption observed is really \\ion{D}{1} and not very high velocity \\ion{H}{1}.  This conclusion lends support to the hypothesis that Complex~C is a manifestation of gas originating outside the Milky Way, rather than material recirculated between the Galactic disk and halo.}  \\item{A value of D/H $\\approx 2.2\\times10^{-5}$ for Complex~C is consistent with the primordial abundance of deuterium inferred from recent WMAP observations of the cosmic microwave background, combined with simple chemical evolution models that predict the amount of deuterium astration with time.  Reducing the errors on the D/H estimate in Complex~C toward PG\\,1259+593 will be difficult, but may eventually be possible with a large space telescope equipped with a high-resolution ultraviolet spectrograph.}  \\item{The number of sight lines for which \\ion{D}{1} can be observed in Lyman-series absorption beyond a few hundred parsecs of the Sun appears to be small because of \\ion{H}{1} velocity crowding. We considered the data for several other sight lines through Complex~C and found that they are unsuitable for a reliable \\ion{D}{1} column density determination.  We continue to examine FUSE data for \\ion{D}{1} absorption in other HVCs.}  \\end{enumerate}    \\noindent {\\bf Acknowledgments.} We thank Van Dixon for answering numerous questions about the {\\tt CALFUSE} pipeline software, and we thank Max Pettini and Jeff Linsky for useful comments on the manuscript. This work is based on data obtained for the Guaranteed Time Team by the NASA-CNES-CSA FUSE mission operated by the Johns Hopkins University. We acknowledge the many contributions of the members of the FUSE D/H working group and the tireless efforts of the FUSE Operations Team. This work is also based in part upon observations made with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. The Westerbork Synthesis Radio Telescope is operated by the ASTRON (Netherlands Foundation for Research in Astronomy) with support from the Netherlands Foundation for Scientific Research (NWO).  The Effelsberg Telescope belongs to the Max Planck Institute for Radio Astronomy in Bonn, Germany. We are indebted to Robert Braun for assistance with the Westerbork observations. Partial financial support has been provided by NASA contract NAS5-32985 and Long Term Space Astrophysics grants NAG5-3485 (K.R.S.) and NAG5-7262 (J.M.S.). B.P.W.\\ was supported by NASA grants NAG5-9179, NAG5-9024, and NAG5-8967.  T.M.T.\\ appreciates support for this work from NASA Long Term Space Astrophysics grant NAG5-11136 as well as NASA grant GO-08695.01-A from the Space Telescope Science Institute.  P.R.\\ was supported by the Deutsche Forschungsgemeinschaft. French co-authors were supported by CNES.  G.H.\\ and R.F.\\ acknowledge use of the {\\tt Owens.f} software developed by Martin Lemoine and the French FUSE Team. \\clearpage \\newpage"
    },
    "0311/astro-ph0311388_arXiv.txt": {
        "abstract": "{We present a model of a pulsar wind nebula evolving inside its associated supernova remnant. The model uses a hydrodynamics code to simulate the evolution of this system when the pulsar has a high velocity. The simulation distinguishes four different stages of pulsar wind nebula evolution: the supersonic expansion stage, the reverse shock interaction stage, the subsonic expansion stage and ultimately the bow shock stage. The simulation bears out that, due to the high velocity of the pulsar, the position of the pulsar is off-centered with respect to its pulsar wind nebula, after the passage of the reverse shock. Subsequently the pulsar wind nebula expands subsonically untill the event of the bow shock formation, when the motion of the pulsar becomes supersonic. The bow shock formation event occurs at roughly half the crossing time, when the pulsar is positioned at 0.677 times the radius of the supernova remnant blastwave, in complete agreement with analytical predictions. The crossing time is defined by the age of the supernova remnant, when the pulsar overtakes the blastwave bounding the supernova remnant.  The results of the model are applied to three supernova remnants: N157B, G327.1-1.1 and W44. We argue that the head of the pulsar wind nebula, containing the active pulsar, inside the first two systems are not bounded by a bow shock. However, in the case of W44 we argue for a scenario in which the pulsar wind nebula is bounded by a  bow shock, due to the supersonic motion of the pulsar. ",
        "introduction": "A supernova remnant (SNR) is the relic of a supernova explosion, which injects an energy of $\\sim 10^{51}$ erg into the surrounding medium. The dynamics of a young SNR expanding into the interstellar medium (ISM) is determined by two shocks: a forward shock which propagates into the ISM and is being decelerated by sweeping up material from the ISM, and a reverse shock which results from the high pressure behind the forward shock and propagates back into the freely expanding ejecta of the SNR (McKee 1974; McKee \\& Truelove 1995).  A young SNR becomes dynamically more interesting in those cases where the collapse of the progenitor star, preceding the supernova explosion, yields a pulsar: a rapidly rotating neutron star. In those cases the dynamics of the central region of the SNR is dominated by the continuous injection of energetic particles by a relativistic pulsar wind, driven by the spin-down energy of the pulsar. The pulsar wind is terminated by a strong MHD shock (Rees \\& Gunn 1974), and drives a pulsar wind nebula (PWN) in the interior of the young SNR. The dynamics of the PWN is coupled to the evolution of the SNR, because the total energy release over the pulsar's lifetime is small ($\\sim 10^{49}-10^{50}$ erg) compared with the total mechanical energy of the SNR ($\\sim 10^{51}$ erg).  Several authors (Reynolds \\& Chevalier 1984; van der Swaluw et al. 2001; Blondin et al. 2001; Bucciantini et al. 2003) have considered the evolution of a centered PWN inside an evolving SNR. In these systems, the initial stage of the PWN is the supersonic expansion stage: the pulsar wind bubble is bounded by a strong PWN shock propagating through the freely expanding ejecta of the SNR. A transition to the reverse shock interaction stage takes place when the reverse shock collides with the PWN shock, and subsequently crushes the pulsar wind bubble (van der Swaluw et al. 2001, Blondin et al. 2001). This reverse shock interaction stage is characterised by an unsteady expansion of the pulsar wind bubble (van der Swaluw et al. 2001), due to the reverberations from the violent collision between the reverse shock and the PWN shock. The expansion of the PWN proceeds subsonically when these reverberations have vanished. The expansion is subsonic because the surroundings of the PWN have been reheated by the passage of the reverse shock: the PWN shock, bounding the swept-up shocked ejecta around the hot pulsar wind bubble, has disappeared.   In this paper we discuss the evolution of PWNe inside SNRs, for pulsars with a constant kick velocity: initially the PWN starts its expansion at the center of the SNR, however the pulsar motion will move the pulsar wind cavity along as it moves through the SNR interior. We present results from hydrodynamical simulations for such a system, which distinguishes all three evolutionary stages mentioned above, i.e. the supersonic expansion stage, the reverse shock interaction stage and the subsonic expansion stage. However at the end of the simulation an additional stage can be distinguished, when the head of the PWN, containing the active pulsar, deforms into a bow shock, due to the supersonic motion of the pulsar.  The supersonic expansion stage in the simulation shows a PWN which is off-centered with respect to the twofold shock structure of the SNR, due to the kick velocity of the pulsar. Therefore the timescale on which the reverse shock collides with the complete shock surface, bounding the PWN, can be a significant fraction of the total lifetime of the PWN when the reverse shock interaction stage starts. We use a semi-analytical approach to estimate the timescale of the collision process, which is shown to scale roughy with the pulsar velocity. The results from the simulation are in almost complete agreement with these semi-analytical calculations. Due to the high velocity of the pulsar, its position inside the PWN is strongly off-centered after the passage of the reverse shock. Ultimately, at the end of the simulation, when the pulsar is approaching the shell of its SNR, the head of the PWN, containing the active pulsar, is deformed into a bow shock at a time and position, in complete agreement with analytical predictions made by van der Swaluw et al. (1998).   We determine the evolutionary stage of the PWNe inside the SNRs N157B, G327.1-1.1 and W44, using our model. For the first two systems, we argue that the position of the pulsar at the head of its PWN, is a result of the passage of the reverse shock and the high velocity of the pulsar: the head of these PWNe are not bounded by a bow shock. Therefore these PWNe are either in the {\\it reverse shock interaction stage} or the {\\it subsonic expansion stage}. The PWN inside the SNR W44 however, is shown to be a good candidate for having a bow shock nebula around its pulsar. ",
        "conclusions": "We have considered the case of a PWN interacting with a SNR, for which the associated pulsar is moving at a high velocity through the interior of its SNR. The model we discussed made use of a hydrodynamics code. One could distinguish four different stages in the simulation: the supersonic expansion stage, the reverse shock interaction stage, the subsonic expansion stage and the bow shock stage. Below we summarise the most important results from our model: \\begin{itemize} \\item The reverse shock interaction stage starts when the reverse shock collides with the front of the PWN shock. The timescale on which the reverse shock collides with the {\\it complete} PWN shock, bounding the pulsar wind bubble, scales with the pulsar velocity. This timescale is a significant fraction of the lifetime of the PWN, when the reverse shock interaction stage starts, for PWNe containing a high velocity pulsar. \\item The high velocity of the pulsar results in an off-centered position of the pulsar with respect to its pulsar wind bubble, after the passage of the reverse shock. \\item The morphology of a PWN, after the passage of the reverse shock, consists of a twofold structure which ultimately expands subsonically inside the relaxed Sedov-Taylor SNR: \\par 1) a roughly spherically symmetric relic PWN \\par 2) a head, containing the pulsar, directed towards the SNR shell \\item The formation of a bow shock around the head of the PWN occurs at half the crossing time, when the position of the pulsar $R_{\\rm psr}$ with respect to the SNR shock $R_{\\rm snr}$ equals $R_{\\rm psr}\\simeq 0.677 R_{\\rm snr}$. We derived a lower limit for the pulsar velocity for the bow shock formation to occur while the SNR is in its Sedov-taylor stage. \\end{itemize} From our model it follows that a SNR containing a pulsar at the head of its PWN does not imply the presence of a bow shock, bounding the head feature of the PWN. We discussed three SNRs and interpreted their morphology, using the results from our model. For the SNRs N157B and G327.1-1.1 we argued that they contain PWNe which do not have a bow shock. The PWN inside the SNR W44 seems to be a good candidate for having a bow shock around the head of its PWN."
    },
    "0311/astro-ph0311561_arXiv.txt": {
        "abstract": "The outflow in SDSS J0300+0048 has the highest column density yet reported for a broad absorption line quasar. The absorption from different ions is also segregated as a function of velocity in a way that can only be explained by a disk wind outflow.  Furthermore, the presence of the such large column densities of gas at the high observed outflow velocities may be incompatible with purely radiative acceleration. MHD contributions to the acceleration should be considered seriously. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311082_arXiv.txt": {
        "abstract": "{  In the framework of the ESO Large Programme ``First Stars'', very high-quality spectra of some 70 very metal-poor dwarfs and giants were obtained with the ESO VLT and UVES spectrograph. These stars are likely to have descended from the first generation(s) of stars formed after the Big Bang, and their detailed composition provides constraints on issues such as the nature of the first supernovae, the efficiency of mixing processes in the early Galaxy, the formation and evolution of the halo of the Galaxy, and the possible sources of reionization of the Universe.  This paper presents the abundance analysis of an homogeneous sample of 35 giants selected from the HK survey of Beers et al. (1992; 1999), emphasizing stars of extremely low metallicity: 30 of our 35 stars are in the range $-2.7 <{\\rm [Fe/H]}< -4.1$, and 22 stars have ${\\rm [Fe/H]} < -3.0$. Our new VLT/UVES spectra, at a resolving power of $R\\sim45,000$ and with signal-to-noise ratios of 100-200 per pixel over the wavelength range 330 -- 1000 nm, are greatly superior to those of the classic studies of McWilliam et al. (1995) and Ryan, Norris, \\& Beers (1996).  The immediate objective of the work is to determine precise, comprehensive, and homogeneous element abundances for this large sample of the most metal-poor giants presently known. In the analysis we combine the spectral line modeling code ``Turbospectrum'' with OSMARCS model atmospheres, which treat continuum scattering correctly and thus allow proper interpretation of the blue regions of the spectra, where scattering becomes important relative to continuous absorption ($\\lambda < 400$~nm). We obtain detailed information on the trends of elemental abundance ratios and the star-to-star scatter around those trends, enabling us to separate the relative contributions of cosmic scatter and observational/analysis errors.  Abundances of 17 elements from C to Zn have been measured in all stars, including K and Zn, which have not previously been detected in stars with [Fe/H] $< -$3.0. Among the key results, we discuss the oxygen abundance (from the forbidden [OI] line), the different and sometimes complex trends of the abundance ratios with metallicity, the very tight relationship between the abundances of certain elements (e.g., Fe and Cr), and the high [Zn/Fe] ratio in the most metal-poor stars.  Within the error bars, the trends of the abundance ratios with metallicity are consistent with those found in earlier literature, but in many cases the scatter around the average trends is {\\it much} smaller than found in earlier studies, which were limited to lower-quality spectra. We find that the cosmic scatter in several element ratios may be as low as 0.05~dex.  The evolution of the abundance trends and scatter with declining metallicity provides strong constraints on the yields of the first supernovae and their mixing into the early ISM. The abundance ratios found in our sample {\\it do not} match the predicted yields from pair-instability hypernovae, but are consistent with element production by supernovae with progenitor masses up to 100 M$_{\\sun}$. Moreover, the composition of the ejecta that have enriched the matter now contained in our very metal-poor stars appears surprisingly uniform over the range $-4.0 \\le {\\rm [Fe/H]} < -3.0$. This would indicate either that we are observing the products of very similar primordial bursts of high-mass stars, or that the mixing of matter from different bursts of early star formation was extremely rapid. In any case, it is unlikely that we observed the ejecta from individual (single) supernovae (as has often been concluded in previous work), as we do not see  scatter due to different progenitor masses. The abundance ratios at the lowest metallicities ($-4.0 \\le {\\rm [Fe/H]} \\le -3.0$) are compatible with those found by McWilliam et al. (1995) and later studies. However,  when elemental ratios are plotted with respect to Mg, we find no clear slopes below [Mg/H] = --3, but rather, a plateau-like behaviour defining a set of initial yields.  \\keywords {Galaxy: abundances -- Galaxy: halo -- Galaxy: evolution -- Stars: abundances -- Stars: Population II -- Stars: Supernovae -- reionization} } ",
        "introduction": " ",
        "conclusions": "\\begin {figure*} \\begin {center} \\resizebox  {5.8cm}{2.90cm} {\\includegraphics {abgmg-c.ps} } \\resizebox  {5.8cm}{2.90cm} {\\includegraphics {abgmg-o.ps} } \\resizebox  {5.8cm}{2.90cm} {\\includegraphics {abgmg-na.ps} } \\resizebox  {5.8cm}{2.90cm} {\\includegraphics {abgmg-al.ps} } \\resizebox  {5.8cm}{2.90cm} {\\includegraphics {abgmg-si.ps} } \\resizebox  {5.8cm}{2.90cm} {\\includegraphics {abgmg-ca.ps} } \\resizebox  {5.8cm}{2.90cm} {\\includegraphics {abgmg-ti.ps} } \\resizebox  {5.8cm}{2.90cm} {\\includegraphics {abgmg-sc.ps} } \\resizebox  {5.8cm}{2.90cm} {\\includegraphics {abgmg-k.ps} } \\resizebox  {5.8cm}{2.90cm} {\\includegraphics {abgmg-cr.ps} } \\resizebox  {5.8cm}{2.90cm} {\\includegraphics {abgmg-mn.ps} } \\resizebox  {5.8cm}{2.90cm} {\\includegraphics {abgmg-fe.ps} } \\resizebox  {5.8cm}{2.90cm} {\\includegraphics {abgmg-co.ps} } \\resizebox  {5.8cm}{2.90cm} {\\includegraphics {abgmg-ni.ps} } \\resizebox  {5.8cm}{2.90cm} {\\includegraphics {abgmg-zn.ps} } \\caption {Abundance ratios [X/Mg] plotted vs. [Mg/H]. The scale is the same for all the plots: the amplitude in the coordinate [X/Mg] is 2.2 dex. For Mn the symbols are the same as in Fig. \\ref{abCr-Mn}} \\label {abMg-SiZn} \\end {center} \\end {figure*}   \\subsection{Comparison with previous studies of very metal-poor stars}  The trends of the relations [X/Fe] vs. [Fe/H] reported in the present study are generally in agreement with previous results in the literature, such as those of McWilliam et al. (\\cite{MPS95}), Ryan et al. (\\cite{RNB96}), and Norris, Ryan, \\& Beers (\\cite{NRB01}). However, the much smaller scatter of the ratios [X/Fe] is a notable result of the greatly improved spectra obtained for this study.  Carretta et al. (\\cite{CGC02}) observed three very metal-poor MS or TO stars and two extremely metal-poor giants ([Fe/H] $< -3.0$) with Keck spectra having a figure of merit F (Norris et al. \\cite{NRB01}) larger than 600. The elements from Mg to Fe were analyzed. Within a (rather large) scatter the two analyses are compatible, except for a few elements, in particular Cr (their slope of [Cr/Fe] vs. [Fe/H] seems to be steeper than ours).  \\subsection{Mg as an alternative ``reference element''}\\label{sec-refelem}  Iron is a convenient ``reference element'' in high-resolution spectral analyses, because it has by far the largest number of usable lines and is represented by two ionization states. Iron may not be the best choice as a tracer of Galactic chemical evolution, however, since its nucleosynthesis channels are not very well understood, and are not necessarily even unique (e.g., explosive nucleosynthesis, Si burning in massive SNe II, or SNe Ia).  From the point of view of the chemical evolution of the Galaxy oxygen would be a better choice (see Wheeler, Sneden \\& Truran \\cite{WST89}), as oxygen is the most abundant element after H and He, it comes from a single source, and its abundance should not be significantly affected in the explosive phase. However, its well-known observational difficulties would considerably degrade the accuracy of the derived trends: Oxygen could be measured in only 21 of our programme stars and the uncertainties on its abundance are large (see in particular Sect. \\ref{sec-errors} and the error bars in Fig. \\ref{ab-o}).  Mg or Ca might be good alternatives. The abundances of these elements are accurately determined, and they are also formed mainly in massive SNe. We choose Mg rather than Ca here (Fig. \\ref{abMg-SiZn}), because Mg is more robust in the sense that its production is dominated by hydrostatic burning processes and it is also less affected by explosive burning and by ``fallback'' (Woosley \\& Weaver \\cite{WW95}). Note that Shigeyama \\& Tsujimoto (\\cite{ST98}) also recommended Mg rather than Fe as a useful reference element, following much the same logic. We note, however, that the ``plateau''-like behavior of [Mg/Fe] with increasing iron abundance in the range $-$4 to $-$2 implies that Mg and Fe have parallel early nucleosynthesis histories. Therefore, the trends of elemental ratios with [Fe/H] found in section \\ref{CZn-Fe} should survive if [Fe/H] is replaced by [Mg/H] as a metallicity indicator.  Yet, the new diagrams (see Fig. \\ref{abMg-SiZn}) present a few notable differences. One is expected -- the scatter is never as low as in some of the earlier diagrams because [Mg/H] is determined less accurately than [Fe/H], being based on 8 lines instead of 150. The other -- new -- result is very interesting: rather than a roughly linear variation over the full range $-4.0<$ [Mg/H] $<-2.5$, there is a hint that all abundance ratios are flat in the interval $-4.0<$ [Mg/H] $<-3.0$, with something qualitatively different occurring at higher metallicity.  This pattern is in fact what is expected if the first SNe are primordial and have specific yields: the plateau between $-4.0$ and $-3.0$ would then reflect a pure zero-metallicity type of SNe, whereas at higher metallicity we may observe a mix of primordial and non-primordial SNe. At still higher metallicities ($\\ge$ --2.0), the scatter is expected to decrease again because so many SN precursors are involved that any differences average out.  We now consider a few elements of particular interest.   $\\bullet$ Carbon\\\\ Karlsson \\& Gustafsson (\\cite{KG01}) have statistically simulated the chemical enrichment of a metal-poor system, assuming that the stellar yields are one-dimensional functions of the progenitor mass of the supernovae and the masses of the supernovae are distributed according to a Salpeter IMF. In particular, they computed the distribution of the abundance ratio [C/Mg] vs. [Mg/H] in a hypothesized sample of 500 XMP stars (their Fig. 3), adopting the yields of either Woosley \\& Weaver (\\cite{WW95}) or Nomoto et al. (\\cite{NHT97}). Our Fig. \\ref{abMg-SiZn} for C is compatible with their Fig. 3a (yields of Woosley \\& Weaver), but not with their Fig. 3b (yields of Nomoto et al.).  On the other hand, both in their simulation and as observed in our present sample, the [C/Mg] ratio seems to decrease with increasing [Mg/H]. Following Karlsson and Gustafsson (\\cite{KG01}), this effect could be the result of different supernova masses operating at different metallicities.  SNe producing a high [C/Mg] ratio produce only small amounts of Mg; on the contrary, SNe producing a low [C/Mg] ratio also produce substantial Mg.  Karlsson \\& Gustafsson show that the patterns they predict become barely visible (or invisible) in an observational sample smaller than N $\\approx 500$ and affected by uncertainties of the order of 0.1 dex. We have examined how their [C/Mg] vs. [Mg/Fe] diagram would appear for our sample (Fig. \\ref{abmg-C-Mg}). Not only is no fine structure visible, but our diagram is considerably more extended in the vertical direction, strongly suggesting that the scatter in [C/Mg] is not explained by the theoretical yields.   \\begin {figure} \\begin {center} \\resizebox  {6.cm}{6.cm} {\\includegraphics {abmg-C-Mg.ps}} \\caption {[C/Mg] vs. [Mg/Fe] for all stars of our sample with temperatures higher than 4800 K (and thus presumably not mixed). Filled circles represent the most metal-poor stars ([Fe/H] $\\leq -3.0$), open circles the stars with [Fe/H]$> -3.0$.  The variation of [C/Mg] is larger than expected from the computations of Karlsson \\& Gustafsson (\\cite{KG01}), suggesting that the scatter in [C/Mg] is not explained by the theoretical yields. } \\label {abmg-C-Mg} \\end {center} \\end {figure}    $\\bullet$ Sodium and Aluminium\\\\ [0cm] [Na/Fe] and [Al/Fe] exhibit very different behaviours as functions of [Fe/H] (Fig. \\ref{abFe-NaAl}). However, when Mg is used as a reference element (Fig. \\ref{abMg-SiZn}), the behavior of these elements appears rather similar, and a plateau at the lowest metallicities appears for Na as well as for Al: below $\\mathrm{[Mg/H] = -3.0}$, [Na/Mg] remains constant at about [Na/Mg] = 0.9 (in fact this plateau appears also in Fig. \\ref{abFe-NaAl}:  [Na/Fe] $\\approx 0.7$ for [Fe/H]$<-3.4$). However, the rise in [Na/Mg] at higher metallicity is not seen for Al.  The discrepant position of CS~22952--015 in the diagrams of [Na/Mg] and [Al/Mg] vs. [Mg/H] will be discussed in Sect. \\ref{pec-stars}.  $\\bullet$ Iron peak elements\\\\ There is no clear slope of [Cr/Mg] vs. [Mg/H], as was seen for [Cr/Fe] vs. [Fe/H]. This raises the suspicion that the slope in the latter diagram may be an artifact due to different NLTE corrections for the two elements as a function of metallicity (Th\\'evenin \\& Idiart \\cite{TI99}). These corrections have not been applied, as they are not known for Cr and have not been published line by line for Fe. We return to this question below.  No significant slope is found for [Mn/Mg], (in agreement with absence of slope for [Mn/Fe]). Similarly, for [Mg/H]$< -3.0$ no clearly significant slope is found for the other elements (Fe, Co, Ni, Zn) .   \\subsection{Cosmic vs. observational scatter in the abundance ratios} \\label{sec-scatter} A key motivation for the present programme was to explore to what extent the scatter in the observed abundance ratios is due to observational error, and to what extent it reflects physical conditions in the early Galaxy when these stars were formed. McWilliam et al. (\\cite{MPS95}) and McWilliam \\& Searle (\\cite{MS99}) already noted that the scatter in some of their diagrams of [X/Fe] vs. [Fe/H] could be entirely accounted for by observational errors. The issue was summarized by Ryan, Norris, \\& Beers (\\cite{RNB96}) as follows. ``The abundance patterns, especially those of Cr, Mn, or Co, raise the following question: why should all halo supernova ejecta around this epoch that possess a particular [Cr/Fe] ratio (or [Mn/Fe] ratio or [Co/Fe] ratio) {\\it subsequently} form into stars of the same [Fe/H]? Put differently, how do the ejecta know how much interstellar hydrogen to combine with?''.  Our observations were designed to achieve twice the spectral resolution and 3--4 times the S/N ratio of the earlier data in order to test this very point. With our much lower observational errors, we can conclude that while the scatter for C, Na, Mg, Al, and Si is probably real, the very small scatter of Ca, Cr, and Ni {\\it still} do not leave room for the existence of an intrinsic scatter!  Consider the case of chromium, which has the lowest observed scatter (r.m.s. 0.05 dex). The problem mentioned earlier is very acute and derives from the {\\it simultaneous} absence of scatter and presence of of a slope of [Cr/Fe] versus metallicity.  Although one can argue that the amount of  hydrogen swept up by the ejecta is mainly determined by the energy of the explosion of the SN (Cioffi et al. \\cite{CMB88}), thus relating the abundance ratios produced by the SN to the final [Fe/H] of the enriched gas, it is still difficult to believe that there is so little room for noise in the mixing process. However, if the slope is in fact spurious (e.g., due to neglected differential NLTE corrections between Cr and Fe), as suggested by the diagram of [Cr/Mg] vs. [Mg/H] (Fig. \\ref{abMg-SiZn}), the problem vanishes. One would simply conclude that Cr and Fe are produced together, independent of the metallicity of the SN progenitor, and  the amount of mixing cannot be localized anymore along the metallicity axis. Until detailed NLTE computations for Cr and Fe become available we cannot decide if this interpretation is correct.  For all the elements discussed here, our results show that the scatter of their production ratios is very small, far below the values derived earlier. This implies that we are observing either the ejecta of fairly large bursts of massive stars, so the sampling of the IMF is reasonably good, or the result of several events promptly mixed by strong turbulence.   \\subsection{The nature of the first supernovae}  Theoretical work  (Bromm et al. \\cite{BFC01}  and references therein) predict that the first stellar generation is made of very massive stars,  with masses above 100 M$_{\\odot}$, because zero-metal matter lacks adequate cooling mechanisms for fragmenting down to classical supernova-progenitor masses. Results of WMAP (Kogut et al. \\cite{KSB03}) on an early reionization of the Universe have triggered further claims (Cen \\cite{Cen03}) of a very massive stellar generation.This has  very important implications for early 'stellar' nucleosynthesis. According to current models, such stars either end up as pair-instability supernovae, or as collapsed black holes, in the latter case with no contribution to the metal enrichment of the ISM (Heger \\& Woosley \\cite{HW02}; Umeda \\& Nomoto \\cite{UN02}). A comparison of the yields of Heger \\& Woosley with our results show a clear disagreement, in particular the predicted strong odd-even effect, not seen in our observations, and  a strong decline of Zn with metallicity also not observed. Nakamura et al. (\\cite{NUI01}) and Nomoto et al. (\\cite{NHT97}) have computed yields of SNe with a progenitor mass of 25 M$_{\\odot}$ of very high energy (also called hypernovae). Their predicted yields have some positive features at very low metallicity, such as the high [Zn/Fe] and [Co/Fe], and a low [Mn/Fe], as observed. However they also predict a lack of [O/Fe] enhancement, in clear disagreement with our observations.  Our conclusion is that classical SNe are still the best candidates to explain our observational results.   \\subsection{Peculiar objects} \\label{pec-stars} \\subsubsection{CS~22949--037}  The highly peculiar abundances of CS~22949--037 were studied in detail by Depagne et al. (\\cite{DHS02}). They may be explained by a single progenitor or by an enrichment event dominated by massive SNe II with substantial fallback; this applies also to the similar star CS~29498--043 analysed by Aoki et al. (\\cite{ANR02a}, \\cite{ANR02b}).  Tsujimoto \\& Shigeyama (\\cite{TS03}) propose another interpretation of CS~22949--037 and CS~29498--043. The high [Mg/Fe] ratio could be due to a low-energy explosion, strong enough to eject the layers containing the light elements, but ejecting little iron and other iron-peak elements. In this interpretation, the lack of Fe relative to the lighter elements O and Mg is also associated with a large fallback on the remnant, but due to a low explosion energy rather than a large mass of the collapsed core as in the model adopted by Depagne et al. (\\cite{DHS02}). However, the normal Cr/Mn/Fe/C/Ni ratios observed in this star would rather suggest a normal explosion energy, the larger fallback being due to a larger mass of the single or multiple progenitors.  Further exploration of the precise abundance patterns of these two putative low-energy supernova descendents should be quite interesting. This interpretation would link the exceptional cases of the most extreme metal-poor stars with rather low-mass progenitors. It would, however, conflict with the usual interpretation which associate the more massive progenitors with both the earliest explosions and the largest volume of hydrogen swept up, resulting in low metallicity in the ISM. The problem clearly requires further investigation.   \\subsubsection{CS~22952--015}  CS~22952--015 is known to be Mg deficient (McWilliam et al., \\cite{MPS95}). In the abundance ratio plots vs. [Mg/H] (Fig. \\ref{abMg-SiZn}), it does indeed appear Fe-rich, but its most notable characteristic is the large values of [Na/Mg] and [Al/Mg]. An interesting point is that this effect is not seen as clearly in the diagrams of [Na/Fe] and [Al/Fe] vs. [Fe/H] (Fig. \\ref{abFe-NaAl}), because all three elements Na, Al, and Fe are enhanced relative to Mg in CS~22952--015.  \\subsubsection{CS~22169--035}  CS~22169--035 appears to be particularly deficient in Ti (Fig. \\ref{ab-mgti}), but in fact all the ratios [Mg/Fe], [Si/Fe], [Ca/Fe], [Co/Fe], [Ni/Fe], and [Zn/Fe] are low as well. When Mg is used as the reference element, this star has a normal position in the diagrams, and the abundance anomalies are most simply characterised as a deficiency of Fe.   \\subsection{Yields of the first supernovae}  With Mg chosen as the reference element (Fig. \\ref{abMg-SiZn}), the most metal-poor stars in our sample ([X/H]$<-2.9$) define a plateau at abundance ratios [X/Mg] corresponding to the yields of the first supernovae, thus providing constraints on these yields. The mean value of [X/Mg] of each plateau is given in Table \\ref{yieldsuper} and represent our best estimate of the yields from the first metal producers in the Galaxy.   \\begin {table}[t] \\caption {Mean [X/Mg] ratios for stars with [Mg/H]$<-2.9$, corresponding to the yields of the first supernovae. The r.m.s. is the scatter around the mean.} \\label{yieldsuper} \\begin {center} \\begin {tabular}{lrrrr} \\hline &      &$r.m.s.$&  n    \\\\ $[$Na/Mg] & $-$0.84 & 0.22 & 11   \\\\ $[$Al/Mg] & $-$0.33 & 0.15 & 13   \\\\ $[$Si/Mg] & $+$0.21 & 0.14 & 14   \\\\ $[$K/Mg]  & $-$0.14 & 0.14 & 13   \\\\ $[$Ca/Mg] & $+$0.06 & 0.09 & 14   \\\\ $[$Sc/Mg] & $-$0.17 & 0.14 & 14   \\\\ $[$Ti/Mg] & $-$0.01 & 0.09 & 14   \\\\ $[$Cr/Mg] & $-$0.63 & 0.09 & 14   \\\\ $[$Mn/Mg] & $-$0.65 & 0.20 & 14   \\\\ $[$Fe/Mg] & $-$0.21 & 0.10 & 14   \\\\ $[$Co/Mg] & $+$0.13 & 0.17 & 14   \\\\ $[$Ni/Mg] & $-$0.23 & 0.13 & 14   \\\\ $[$Zn/Mg] & $+$0.15 & 0.19 & 14   \\\\ \\hline \\end {tabular} \\end {center} \\end {table}   In Fig. \\ref{compayields}a, these mean values are compared to the values predicted by Woosley \\& Weaver (\\cite{WW95}) for their zero-metal supernova models 15A, 25B and 35C (progenitor masses 15, 25 and 35 $M_{\\odot}$). The Woosley \\& Weaver values of [X/Mg] for the elements from C to Ca have been enhanced in order to bring their predicted value of [Fe/Mg] into agreement with our observations.  In Fig. \\ref{compayields}b we compare with the zero-metal models of Chieffi \\& Limongi (\\cite{CL03}) for 15, 20, 35, 50 $M_{\\odot}$ stars. The predicted [X/Mg] values were already adjusted by Chieffi \\& Limongi to obtain [Mg/Fe]$\\approx +0.45$). These comparisons suggest that some adjustment of the models is indeed required.  The connection between the abundances observed in these very old stars and those observed in intergalactic clouds -- in particular the damped Ly$_\\alpha$ systems (DLA) -- will be further discussed in subsequent papers.  \\begin {figure} \\begin {center} \\resizebox  {7.cm}{3.5cm} {\\includegraphics {woosley-nous.ps} } \\resizebox  {7.cm}{3.5cm} {\\includegraphics {limongi-nous.ps} } \\caption {Mean [X/Mg] values for [Mg/H]$<-3$ (Fig. \\ref{ab-mgti} and Table \\ref{yieldsuper}) compared to the model yields by: {\\bf a)} Woosley and Weaver for 15 (dotted line), 25 (dashed), and 35 $M_{\\odot}$ SNe (full), and {\\bf b)} Limongi and Chieffi for 15 (dotted line), 20 (dotted-dashed), 35 (dashed), and 50 $M_{\\odot}$ (full). } \\label{compayields} \\end {center} \\end {figure}"
    },
    "0311/nucl-th0311055_arXiv.txt": {
        "abstract": "An effective field theory technique that combines the standard nuclear physics approach and chiral perturbation theory is applied to the $hen$ process, ${}^{3}\\mbox{He}+n\\rightarrow {}^{4}\\mbox{He}+ \\gamma$. For the initial and final nuclear states, high-precision wave functions are generated via the variational Monte Carlo method using the Argonne $v_{14}$ potential and Urbana VIII  trinucleon interactions, while the relevant transition operators are calculated up to ${\\cal O}(Q^4)$ in HB$\\chi$PT. The imposition of the renormalization condition that the magnetic moments of ${}^{3}\\mbox{He}$ and ${}^{3}\\mbox{H}$ be reproduced allows us to carry out a parameter-free calculation of the $hen$ cross section. The result, $\\sigma=(60\\pm 3\\pm 1)~\\mu b$, is in reasonable agreement with the experimental values,  $(54\\pm 6)~\\mu b$ and $(55\\pm 3)~\\mu b$. This agreement demonstrates the validity of the calculational method previously used for estimating the reaction rate of the solar $hep$ process. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311269_arXiv.txt": {
        "abstract": "We examine the ability of the Sunyaev-Zel'dovich effect (SZE) statistics proposed by Lee (2002) using numerical simulations. The statistics describe the distribution of the peak heights in the noise-contaminated SZE sky maps, and provide an analysis technique to sort out the noise-contamination effect and estimate the number density of real clusters. The method is devised to be suitable for the interferometric SZE observations in drift-scanning mode like AMiBA experiment. We apply the proposed method to a set of realistic SZE sky maps constructed from large-scale cosmological simulations, and show that the method indeed allows us to estimate the number density of clusters efficiently. The efficiency of the method is demonstrated in two aspects: (i) it can reconstruct the number density of clusters even though the clusters contribute only a small fraction ($\\sim 10 \\%$) of the total number of peaks in each SZE map; (ii) it can count even those clusters whose amplitudes are low enough to be compatible with the noise level while conventional method of cluster identification using the threshold cutoff cannot count them properly within the same observation time. Thus, the proposed method may be useful for the study of cluster abundance exploiting future interferometric SZE surveys such as AMiBA. ",
        "introduction": "The thermal Sunyaev-Zel'dovich effect (SZE) by galaxy clusters represents the spectral distortion of the cosmic microwave background radiation (CMBR) through the inverse Compton scattering by the hot intra-cluster electrons \\citep{sun-zel72}.  Since the SZE is independent of the redshift, it provides a unique tool to detect high-redshift clusters. Thus, the blank-field SZE survey can measure the formation epoch and the abundance evolution of galaxy clusters, which can be used as a probe to put constraints on the cosmological parameters \\citep{bar-etal96,hen97,bah-fan98,via-lid99,bor-etal99,fan-chi01, gre-etal01,mol-etal02}.  The SZE flux diminution is directly proportional to the cluster Comptonization parameter $y$, defined as \\beq y(\\bth) = \\frac{\\sigma_{T}k_{B}}{m_{e}c^{2}} \\int n_{e}(l\\hbth)T_{e}(l\\hbth)dl, \\label{eqn:compton} \\eeq where $\\bth$ represents the location vector along the line-of-sight ($\\hbth \\equiv \\bth/\\vert \\bth\\vert$) and the integral is over the CMBR path in the line-of-sight direction. Hence the SZE survey at a fixed frequency basically provides a map of the $y$-parameter. The number density of galaxy clusters, however, cannot be directly measured from the observed SZE maps because of the presence of noise.  Conventionally, signals from the noise-contaminated maps are identified as the local peaks whose flux amplitudes (for the SZE observation, proportional to the peak height of the Comptonization parameter, $y$) significantly exceed the noise level. Therefore, to detect lower-amplitude signals using the conventional approach, one has to decrease the noise level by increasing the observational time. Unfortunately, the necessary observational time increases inversely proportional to the squares of the imposed noise level. It is clearly inefficient to use the conventional signal-selection method for the estimate of the cluster number counts, because of the high costs of the interferometric observations.  Recently \\citet{lee02} proposed a statistical strategy to estimate efficiently the number density of galaxy clusters from the noise-contaminated SZE maps constructed from the interferometric observation in drift-scanning mode, with the AMiBA experiment as a target model. By {\\it the efficient estimate},  we mean that the method is capable of estimating the number density of clusters which contribute only a small fraction of the total number density of peaks in the SZE maps, counting properly even those clusters whose peak heights are low enough to be compatible with the noise level without increasing the observational time. The key feature of the method is to exploit the property that the noise-contamination is described by a Gaussian statistic. Then the cluster number density can be estimated by simply deconvolving a one-dimensional Gaussian function and the zero-th order approximation to the cluster number density.  Although the method has been shown to render reasonable results for idealized Monte-Carlo simulations \\citep{lee02}, it is necessary to test its ability by applying it to more realistic simulations before using it in real SZE surveys.  In this {\\it Letter}, we test the ability of the SZE statistical strategy using the outputs of a large cosmological simulation. In $\\S 2$, we provide a concise overview of the model experiment AMiBA and the SZE statistics. In $\\S 3$, we describe the numerical simulation, explain how to apply the statistical strategy to the simulated SZE maps, and compare the reconstructed cluster number densities with the original ones. In $\\S 4$, we summarize the results, and give concluding remarks. ",
        "conclusions": "We have tested the validity of the SZE statistics proposed by \\citet{lee02} using numerical simulations. The statistical method is devised for the interferometric SZE observations in drift scanning mode, targeting at AMiBA as a model experiment, and applied to estimate the number density of cluster signals that occupy only a small fraction ($\\sim 10 \\%$) of the noise-contaminated SZE maps more efficiently than the conventional method based on the threshold cut-off. The method is developed to model the dominant noise contamination effect on the SZE peak distribution by a Gaussian scatter and to deconvolve the noise scatter and the SZE peak distribution. The stability of the deconvolution process has been achieved by convolution with a Wiener optimal filter estimated from the measurable power spectrum of the SZE peaks.  We have applied the method to a large number of realistic SZE sky maps, and found that it indeed works for all realizations, giving a fairly good estimate of the cluster number density, including even those clusters whose peak heights are compatible with noise-levels. In the conventional approach, those low-amplitude peaks are usually all disregarded as noise, so that a large fraction of low-amplitude cluster signals are missing. Our method allows us to count those missing low-amplitude clusters as quickly and accurately as possible. Hence, we conclude that it will provide an efficient method to determine the number density of galaxy clusters in the future interferometric SZE surveys."
    },
    "0311/astro-ph0311096_arXiv.txt": {
        "abstract": "The COBE-DMR 4-year maps displayed a strong non-Gaussian signal in the ``inter-scale'' components of the bispectrum: their observed values did not display the scatter expected from Gaussian maps. We re-examine this and other suggested non-Gaussian features in the light of WMAP. We find that they all disappear. Given that it was proved that COBE-DMR high noise levels and documented systematics could at most {\\it dilute} the observed non-Gaussian features, we conclude that this dataset must have contained non-negligible undocumented systematic errors. It turns out that the culprit is a combination of QuadCube pixelization and data collected during the ``eclipse season''. ",
        "introduction": "The possibility of non-Gaussianity in the COBE-DMR 4 year maps led to a rather protracted  story. Using ``single-$\\ell$'' bispectrum analysis, \\citet{fmg} found strong evidence for non-Gaussianity in the anisotropy of the Cosmic Microwave Background (CMB) temperature.   This detection was followed by similar claims by \\citet{nfs98} and \\citet{pvl98}, and caused considerable consternation among theorists (see \\cite{nvs} for a discussion). Further work, however, showed that claims other than those based on the bispectrum could not be reproduced \\citep{muk}. The bispectrum claims were confirmed by~\\cite{bt99}.  Later~\\cite{banday} cast serious doubts upon the cosmological origin of the observed signal. A systematic was identified which removed the observed ``single-$\\ell$'' bispectrum  signal -- the so-called eclipse season data, which should never have been used. However, when an extension to ``inter-$\\ell$'' bispectrum components was sought, a new non-Gaussian signal was found for inter-$\\ell$ separations of $\\Delta l=1$~\\citep{joao}. Specifically, their observed values were found to concentrate uncannily close to zero instead of displaying the scatter expected from Gaussian maps.   This signal could not be blamed on any systematic effects studied by~\\cite{banday} or otherwise \\citep{joao}. It was also proved that the instrument high noise levels could at most dilute the observed signal, in spite of the noise correlations and anisotropy \\citep{joao}. It was also found that the observed signal did not extend to higher inter-$\\ell$ separations $\\Delta l>1$ \\citep{haav}.  What shall we make of these claims in the light of the recent observations \\citep{wmap} by the Wilkinson Microwave Anisotropy Probe (WMAP)? In this paper we show that they don't survive the new data, which displays consistency with Gaussianity on the angular scales probed by COBE. This might seem at first suprising, given that the WMAP data is considerably less noisy than the COBE-DMR dataset, and that noise can at most hide a non-Gaussian signal. However one should never forget the issue of systematics. Even though documented COBE-DMR systematics were shown not to correlate with the observed non-Gaussian signal, the new data allows us to cast a new look at the problem. We find that a highly non-subtle combination of QuadCube pixelization systematics and the ``eclipse data''  are  to be blamed for the observed effect.  This, we believe, closes the story. The moral is clear: care must be exercised regarding similar claims currently being made with WMAP. ",
        "conclusions": "The conclusion is that a combination of pixelization effects and the ``eclipse'' systematic is behind the inter-$\\ell$ non-Gaussian effect previously reported for COBE maps. Curiously if we take a difference map (COBE minus WMAP) the anomalous alignment of the $J^{3}_{l}$ is not present (see top panel of Fig.~\\ref{fig3}). On the scales we are considering such maps picture COBE noise, plus COBE systematics, minus WMAP systematics (assumed to be small). Thus one would expect the COBE $J^{3}_{l}$ signal to be enhanced in the difference map. The fact that it is not results from the non-linearity of the statistic being used, plus the subtle interplay of signal and systematic via the QuadCube pixelization scheme.  In this paper we have concentrated on $J^3_\\ell$, because the single scale bispectrum anomalies have been explained long ago. However we have checked that no new anomalies on large angular scales emerge in the WMAP data.  We reserve to a future publication a complete study of the bispectrum of WMAP of smaller angular scales. To our mind the work of \\cite{ngwmap} is just the beginning."
    },
    "0311/astro-ph0311575_arXiv.txt": {
        "abstract": "Simulations of large-scale structure in the universe have played a vital role in observational cosmology since 1980's in particular. Their important role will definitely continue to be true in the 21st century. Rather the requirements for simulations in the precision cosmology era will become more progressively demanding; they are supposed to fill the missing link in an accurate and reliable manner between the ``initial'' condition at z=1000 revealed by WMAP and the galaxy/quasar distribution at z=0 - 6 surveyed by 2dF and SDSS. In this review, I will summarize what we have learned so far from the previous cosmological simulations, and discuss several remaining problems for the new millennium. ",
        "introduction": "Cosmological $N$-body simulations started in late 1970's, and since then have played an important part in describing and understanding the nonlinear gravitational clustering in the universe.  As far as I know, the cosmological N-body simulation in a comoving periodic cube, which is quite conventional now, was performed for the first time by Miyoshi \\& Kihara (1975) using $N=400$ particles.  Figure 1 plots the evolution of the number of particles employed in cosmological N-body simulations. Here I consider only the ``high-resolution'' simulations including Particle-Particle, Particle-Particle--Particle-Mesh, and tree algorithms which are published in refereed journals (excluding, e.g., conference proceedings). I found that the evolution is well fitted by \\begin{eqnarray} \\label{eq:Nyear} N= 400 \\times 10^{0.215({\\rm Year}-1975)} , \\end{eqnarray} where the amplitude is normalized to the work of Miyoshi \\& Kihara (1975).  Just for comparison, the total number of CDM particles of mass $m_{\\rm CDM}$ in a box of the universe of one side $L$ is \\begin{eqnarray} N= \\frac{\\Omega_\\CDM \\rho_{\\rm cr} L^3}{m_\\CDM} \\approx 10^{83} \\left(\\frac{\\Omega_\\CDM}{0.23} \\right) \\left(\\frac{L}{1 h^{-1}{\\rm Gpc}}\\right)^3 \\left(\\frac{1 {\\rm keV}}{m_\\CDM}\\right) \\left(\\frac{0.71}{h}\\right) . \\end{eqnarray} If I simply extrapolate equation (1) and adopt the WMAP parameters (Spergel et al. 2003), then the number of particles that one can simulate in a $(1 h^{-1}{\\rm Gpc})^3$ box will reach the real number of CDM particles in December 2348 and February 2386 for $m_\\CDM= 1{\\rm keV}$ and $10^{-5}{\\rm eV}$, respectively. I have not yet checked the above arithmetic, but the exact number should not change the basic conclusion; simulations in the new millennium will be {\\it unbelievably} realistic.  \\begin{figure}[tbh] \\begin{center} \\leavevmode\\epsfxsize=7.50cm \\epsfbox{Nyear.eps} \\caption{Evolution of the number of particles in ``high-resolution'' cosmological N-body simulations. \\label{fig:Nyear}} \\end{center} \\end{figure} ",
        "conclusions": "It turned out that I was able to review the cosmological simulations only for a time-scale of decades. Still the material may be heavily biased, which I have to apologize for the organizers and possible readers of these proceedings.  A millennium is definitely too long for any scientist to make any reliable prediction for its eventual outcome. Thus the number of simulation particles that I predicted in Introduction might sound ridiculous, but in reality the progress in the new millennium may be even more drastic that whatever one can imagine.  For instance, it is unlikely that one still continues to use currently popular particle- or mesh-based simulation techniques over next hundreds of years. In that case, the number of particles may turn out to be a totally {\\it useless} measure of the progress or reliability of simulations. Nevertheless I believe that a historical lesson that I learned in preparing this talk will be still true even at the end of this millennium; good science favors the prepared mind, not the largest simulation at the time.   \\bigskip"
    },
    "0311/astro-ph0311605_arXiv.txt": {
        "abstract": "We explore possible correlations between light profile shapes, as parameterized by the S\\'ersic index $n$ or the concentration index $C_{re}(1/3)$, and relevant stellar population parameters in early-type galaxies.  Mean luminosity weighted ages, metallicities and abundance ratios were obtained from spectra of very high signal--to--noise and stellar population models that synthesize galaxy spectra at the resolution given by their velocity dispersions $\\sigma$, in combination with an age indicator (H$\\gamma_{\\sigma}$) that is virtually free of the effects of metallicity. We do not find any significant correlation between $n$ (or $C_{re}(1/3)$) and mean age or metallicity, but a strong positive correlation of the shape parameters with Mg/Fe abundance ratio. This dependence is as strong as the Mg/Fe -- $\\sigma$ and $C_{re}(1/3)$ -- $\\sigma$ relations. We speculate that early-type galaxies settle up their structure on time-scales in agreement with those imposed by their Mg/Fe ratios. This suggest that the global structure of larger galaxies, with larger Mg/Fe ratios and shorter time-scales, was already at place at high $z$, without experiencing a significant time evolution. ",
        "introduction": "\\label{Sec:intro}  A successful galaxy formation model should be able to explain the main scaling relations followed by early-type (E/SO) galaxies. These galaxies show up a tight relation between the size (by means of the effective radius $r_e$), surface brightness, and central velocity dispersion, $\\sigma$, which is known as the Fundamental Plane (Djorgovski \\& Davis 1987; Dressler et al. 1987). Early-type galaxies also follow the Color-Magnitude Relation, CMR, i.e. brighter galaxies look redder (Visvanathan \\& Sandage 1977; Bower et al. 1992). The CMR and the Mg$_2$--$\\sigma$ relation (e.g. Bender, Burstein \\& Faber 1992; Colless et al. 1999) links the mass of a galaxy, through its luminosity, to its constituent stellar populations. However, little is known about the relation between the morphological shape and the stellar populations of early-type galaxies (Conselice 2003), most likely because these two fields rely on methods that have been traditionally disconnected.  From a stellar population point of view, measured colors and absorption line-strengths are translated to mean ages and metallicities via stellar population synthesis models. Using these models the CMR is often interpreted as a metallicity sequence (e.g., Arimoto \\& Yoshii 1987), i.e. elliptical galaxies are old and passively evolving systems (e.g., Ellis et al. 1997; Stanford, Eisenhardt \\& Dickinson 1998). However, this view is questioned by some studies, which show evidence for a significant intermediate age population in some ellipticals (e.g., Gonz\\'alez 1993; J\\o rgensen 1999).  Element ratios provide further constraints on the star formation histories (e.g. Worthey 1998). These studies have shown compelling evidence for a systematic departure from solar element ratios of some elements as a function of $\\sigma$ (Trager et al. 2000). Recent modeling efforts have predicted spectra of stellar populations at moderately high resolution (Vazdekis 1999; Schiavon et al. 2002; Bruzual \\& Charlot 2003), leading to more accurate stellar population parameter estimates (Vazdekis \\& Arimoto 1999; Vazdekis et al. 2001a; Kuntschner et al. 2002).  From the point of view of galaxy structure, the shape of the overall surface brightness as parametrized by S\\'ersic (1968) index, $n$, correlates with the global properties of early-type galaxies: total luminosity, central surface brightness and effective radius, $r_e$ (Graham, Trujillo \\& Caon 2001 and references therein). Nowadays it is clear that the above luminosity--dependent departures from de Vaucouleurs profile are real, as clearly proven by the existence of strong correlations between $n$ (or equivalently galaxy central concentration of light $C_{re}(1/3)$ (Trujillo, Graham \\& Caon 2001)) and the photometric--independent properties,  $\\sigma$ (Graham, Trujillo \\& Caon 2001) and galaxy's supermassive black hole mass (Graham et al. 2001b).  Aiming at exploring possible correlations between galaxy light profiles and their stellar populations we combine in this letter the results obtained from applying these modeling developments to high quality spectroscopic data. We find a tight correlation between [Mg/Fe] ratio and the shape parameters.  \\begin{deluxetable*}{rrrrrrrrrr} \\tablecaption{Galaxy measurements \\label{tbl-1}} \\tablewidth{0pt} \\tablehead{\\colhead{NGC}&\\colhead{n}&\\colhead{Ref}&\\colhead{$C_{re}(1/3)$}&\\colhead{$\\sigma$(${\\rm km~s^{-1}}$)}& \\colhead{Age (Gyr)}& \\colhead{[M/H]}& \\colhead{[Fe/H]}& \\colhead{[Mg/H]}& \\colhead{[Mg/Fe]} } \\startdata \\multicolumn{10}{c}{Virgo galaxies}\\\\[4pt] 4239 & 0.67 & 2 & 0.17 &  82 &  5.5$^{+ 1.9}_{-1.7}$&-0.21$^{+0.11}_{-0.10}$& -0.21$^{+0.09}_{-0.09}$& -0.16$^{+0.15}_{-0.14}$& 0.05$^{+0.10}_{-0.09}$ \\\\[2pt] 4339 & 3.50 & 1 & 0.44 & 142 & 12.6$^{+ 5.3}_{-2.8}$&-0.06$^{+0.10}_{-0.06}$& -0.17$^{+0.07}_{-0.07}$& +0.13$^{+0.15}_{-0.15}$& 0.30$^{+0.11}_{-0.11}$ \\\\[2pt] 4365 & 6.08 & 1 & 0.55 & 245 & 20.4$^{+ 3.1}_{-6.4}$&-0.02$^{+0.10}_{-0.08}$& -0.28$^{+0.16}_{-0.08}$& +0.21$^{+0.21}_{-0.16}$& 0.49$^{+0.08}_{-0.11}$ \\\\[2pt] 4387 & 1.92 & 1 & 0.34 & 105 & 15.1$^{+ 5.4}_{-4.8}$&-0.20$^{+0.11}_{-0.11}$& -0.22$^{+0.10}_{-0.08}$& -0.09$^{+0.19}_{-0.10}$& 0.13$^{+0.12}_{-0.05}$ \\\\[2pt] 4458 & 2.55 & 1 & 0.40 & 105 & 16.6$^{+ 3.9}_{-5.6}$&-0.31$^{+0.13}_{-0.05}$& -0.40$^{+0.05}_{-0.10}$& -0.05$^{+0.16}_{-0.07}$& 0.35$^{+0.14}_{-0.03}$ \\\\[2pt] 4464 & 2.61 & 1 & 0.38 & 135 & 21.0$^{+ 2.3}_{-7.1}$&-0.32$^{+0.12}_{-0.06}$& -0.40$^{+0.09}_{-0.07}$& -0.03$^{+0.12}_{-0.13}$& 0.37$^{+0.06}_{-0.09}$ \\\\[2pt] 4467 & 1.54 & 2 & 0.28 &  75 & 16.0$^{+ 5.2}_{-5.3}$&-0.23$^{+0.13}_{-0.12}$& -0.35$^{+0.13}_{-0.08}$& -0.08$^{+0.18}_{-0.13}$& 0.27$^{+0.08}_{-0.08}$ \\\\[2pt] 4472 & 6.27 & 1 & 0.55 & 306 & 10.4$^{+ 9.5}_{-4.1}$& 0.07$^{+0.17}_{-0.20}$& -0.12$^{+0.15}_{-0.23}$& +0.46$^{+0.16}_{-0.37}$& 0.58$^{+0.04}_{-0.17}$ \\\\[2pt] 4473 & 5.29 & 1 & 0.52 & 180 & 10.3$^{+ 3.1}_{-2.1}$& 0.15$^{+0.07}_{-0.11}$& -0.06$^{+0.06}_{-0.08}$& +0.48$^{+0.08}_{-0.16}$& 0.54$^{+0.04}_{-0.10}$ \\\\[2pt] 4478 & 1.98 & 1 & 0.33 & 135 & 10.0$^{+ 2.7}_{-2.8}$& 0.04$^{+0.15}_{-0.12}$& -0.08$^{+0.12}_{-0.06}$& +0.23$^{+0.14}_{-0.17}$& 0.31$^{+0.05}_{-0.14}$ \\\\[2pt] 4489 & 2.50 & 4 & 0.37 &  73 &  3.1$^{+ 0.6}_{-0.8}$& 0.20$^{+0.18}_{-0.16}$& +0.20$^{+0.15}_{-0.15}$& +0.23$^{+0.17}_{-0.21}$& 0.03$^{+0.06}_{-0.10}$ \\\\[2pt] 4551 & 1.89 & 1 & 0.32 & 105 &  7.2$^{+ 1.5}_{-0.9}$& 0.19$^{+0.07}_{-0.09}$& +0.07$^{+0.05}_{-0.10}$& +0.41$^{+0.07}_{-0.12}$& 0.34$^{+0.05}_{-0.05}$ \\\\[2pt] 4621 & 6.14 & 1 & 0.55 & 230 & 10.3$^{+ 2.4}_{-1.8}$& 0.26$^{+0.08}_{-0.08}$& -0.04$^{+0.04}_{-0.04}$& +0.64$^{+0.05}_{-0.14}$& 0.60$^{+0.02}_{-0.11}$ \\\\[2pt] 4697 & 3.70 & 3 & 0.45 & 168 & 11.3$^{+ 4.5}_{-2.7}$& 0.06$^{+0.10}_{-0.11}$& -0.10$^{+0.06}_{-0.10}$& +0.30$^{+0.08}_{-0.17}$& 0.40$^{+0.04}_{-0.09}$ \\\\[2pt] \\multicolumn{10}{c}{Field galaxies}\\\\[4pt] 584 & 4.70 & 2 & 0.50 & 217 &  8.6$^{+ 1.9}_{-2.3}$&  -                    & -0.03$^{+0.04}_{-0.08}$& +0.31$^{+0.06}_{-0.10}$&+0.34$^{+0.07}_{-0.05}$ \\\\[2pt] 720 & 4.96 & 3 & 0.51 & 247 & 10.9$^{+ 7.0}_{-2.7}$&  -                    & -0.16$^{+0.08}_{-0.08}$& +0.58$^{+0.12}_{-0.19}$&+0.74$^{+0.07}_{-0.14}$ \\\\[2pt] 821 & 4.00 & 2 & 0.47 & 199 & 12.3$^{+12.5}_{-5.6}$&  -                    & -0.15$^{+0.17}_{-0.17}$& +0.26$^{+0.18}_{-0.23}$&+0.41$^{+0.04}_{-0.08}$ \\\\[2pt] 1700 & 5.99 & 2 & 0.54 & 233 &  6.4$^{+ 3.1}_{-1.7}$&  -                    & +0.04$^{+0.08}_{-0.09}$& +0.37$^{+0.09}_{-0.15}$&+0.34$^{+0.09}_{-0.09}$ \\\\[2pt] 3379 & 4.64 & 3 & 0.49 & 201 & 17.8$^{+ 7.4}_{-6.9}$&  -                    & -0.29$^{+0.08}_{-0.10}$& +0.19$^{+0.19}_{-0.10}$&+0.48$^{+0.12}_{-0.03}$ \\\\[2pt] 5831 & 4.72 & 3 & 0.50 & 166 &  5.9$^{+ 3.5}_{-1.9}$& 0.32$^{+0.15}_{-0.13}$& +0.14$^{+0.15}_{-0.16}$& +0.56$^{+0.06}_{-0.18}$&+0.42$^{+0.04}_{-0.10}$ \\\\[2pt] 7454 & 2.55 & 4 & 0.38 & 112 &  6.5$^{+ 3.3}_{-1.8}$&  -                    & -0.20$^{+0.09}_{-0.10}$& -0.07$^{+0.13}_{-0.14}$&+0.13$^{+0.06}_{-0.06}$ \\enddata  \\tablecomments{The numbers in column 3 indicate the $n$ parameter sources: 1 Caon, Capaccioli \\& D'Onofrio (1993). 2 HST F555W archive images. 3  Fits to the profiles of Peletier et al. (1990). 4 V--band images from J. Blakeslee (private communication).}  \\end{deluxetable*} ",
        "conclusions": "\\label{Sec:discussion}  To test how strong is the new correlation between the shape parameters and [Mg/Fe], we show in Fig.~{\\ref{Fig:c1s}} the fits obtained for the already known $C_{re}(1/3)$ -- $\\sigma$ and [Mg/Fe] -- $\\sigma$ relations. We use the same galaxies plotted in Fig.~{\\ref{Fig:c1a}} and Fig.~{\\ref{Fig:c1m}}. The [Mg/Fe] -- $n$ relation shows a correlation value as strong as the ones plotted in Fig.~{\\ref{Fig:c1s}}.  \\begin{figure}[t] \\centerline{\\psfig{file=f3.eps,width=3.6in}} \\caption{Galaxy velocity dispersion versus $n$ and $C_{re}(1/3)$ and [Mg/Fe]. Symbols have same meaning as in Fig.~{\\ref{Fig:c1a}} \\label{Fig:c1s}} \\end{figure}  It is well known that early-type galaxies show non-solar abundance ratios (e.g., Trager et al. 2000). The [Mg/Fe] ratio, which spans a super-solar range (see Fig.~{\\ref{Fig:c1m}}), provides us with important clues on the star-formation time-scales. The $\\alpha$ elements, such as Mg, are ejected to the ISM by SN~II, which explode on short time-scales ($\\le10^6-10^7$ yr). On the other hand, SN~Ia are the main producers of Fe peak elements, which require time-scales around $\\sim$1~Gyr. It has been proposed that if the duration of the star formation is shorter than SN~Ia time-scales the [Mg/Fe] should be larger than 0 (e.g., Worthey et al. 1992). Other possibilities have also been proposed, such as a top-heavy initial mass function (IMF), which could also be combined with a star formation that stopped before the products of SN~Ia could be incorporated into the stars (e.g., Worthey et al. 1992; Vazdekis et al. 1996).  The fact that  the central ($r_{e}/10$) $\\sigma$ and [Mg/Fe] measurements are strongly correlated (see lower panel of Fig.~{\\ref{Fig:c1s}}) suggests that the innermost regions of the more massive galaxies, with larger potential wells, stopped their star formation on shorter time-scales. Furthermore since the light concentration index strongly correlates with [Mg/Fe] suggests that early-type galaxies settled up their global (and not only the innermost regions) structure on time-scales according to their [Mg/Fe] ratios. This scenario would be reinforced if no [Mg/Fe] radial gradients were present. Almost nothing is known about [Mg/Fe] gradients outward galaxy effective radii, there are indications that this ratio is roughly constant at least out to 1~$r_e$ (Halliday 1999; Mehlert et al. 2003). Mehlert et al. suggest that the constraint of short formation time-scales is a global feature and not only applies to galaxy centers.  We do not see a clear correlation between $n$ and galaxy mean luminosity weighted ages but [Mg/Fe] suggesting that the formation time-scale, rather than burst peppering as a function of time, is the main process linked to galaxy global structure. Since our galaxies are old and the time scales of the larger galaxies very short, the new correlation suggests that the global structure of these galaxies was already at place at high $z$, without experiencing a significant time evolution. This result is supported by Chiosi \\& Carraro (2002) models, which suggest that the gravitational potential well (and therefore mass and shape of the galaxy) dictates the efficiency and extention of the star formation. The favoured picture is that early-type galaxies are assembled at high $z$ (e.g., Fassano et al. 1998; van Dokkum \\& Ellis 2003), defining a tight sequence in the CM diagram at least out to $z \\sim$1.2 (Blakeslee et al. 2003), and their stellar populations evolve passively (e.g., Im et al. 2002; Schade et al. 1999). Furthermore, Mobasher et al. (2003) find that early-type galaxies show higher rest-frame B-band concentration indices than late-type spirals out to $z \\sim 1$, which suggest that by that redshift value early-type systems have already developed large central light concentrations. Such high light concentrations are predicted by both the monolithic scenario by means of a rapid gas collapse which turns into stars, and within hierarchical galaxy formation framework via mergers of stellar and gaseous systems that take place at high z. Conselice (2003) shows that the light concentration is a key parameter to trace the past evolutionary history of galaxy formation.  This research will benefit from studies reaching deeper gradients of line-strengths and abundance ratios, and from theoretical efforts linking them to the light profiles (Angeletti \\& Giannone 2003). Including other element ratios in this analysis would provide further clues, since each element is expelled to the ISM at different epochs, and because line-strengths do not only depend on abundance but on other stellar population parameters. S\\'anchez-Bl\\'azquez et al. (2003) found that the CN is stronger in Virgo than in Coma galaxies. The near-IR Ca{\\sc II} triplet, which is sensitive to both the Ca abundance and the IMF (Vazdekis et al. 2003), decreases as a function of $\\sigma$ (Saglia et al. 2002; Cenarro et al. 2003)."
    },
    "0311/astro-ph0311119_arXiv.txt": {
        "abstract": "We report on the detection of Main Sequence stars belonging to the recently identified Canis Major galaxy in a field located at $\\simeq 4.2\\degr$ from the center of the stellar system. With Main Sequence fitting we obtain a distance modulus $(m-M)_0=14.5\\pm 0.3$ to the dwarf, corresponding to a distance of $D_{\\sun}\\simeq 8.0 \\pm 1.2$ kpc, in full agreement with previous estimates based on the photometric parallax of M-giants. From the comparison with theoretical isochrones we constrain the age of the main population of the Canis Major system in the range $\\sim 4-10$ Gyr. A blue plume of likely younger stars (age $< 1-2$ Gyr) is also identified. The available Colour Magnitude Diagrams of open clusters that may be projected onto the main body of Canis Major are also briefly analyzed. The position, distance and stellar population of the old open clusters AM-2 and Tombaugh~2 strongly suggest that they are physically associated with the Canis Major galaxy. Using our own photometry and data from 2MASS and the GSC2.2 surveys we demonstrate that the claim by Momany et al. that the CMa overdensity is entirely due to the Galactic warp is not supported by the existing observations, once all the available pieces of information are taken into account. It is shown that the CMa overdensity clearly emerges at a heliocentric distance of $\\sim 8$ kpc above any overdensity possibly produced by the Galactic warp. ",
        "introduction": "The process of hierarchical merging \\citep{wr78,wf91} is generally accepted as the driving mechanism of the formation of giant galaxies. The study of the local (e.g. Galactic or in the Local Group) relics of such processes provides an unprecedentedly detailed insight of the current cosmological model as well as a formidable testbed for theories of galaxy formation. The case of the Sagittarius dwarf galaxy \\cite[Sgr dSph]{s1,s2} is emblematic, since we are presently observing it during its disruption into a giant stream contributing to the continued assembly of the Galactic halo with stars \\cite[see][and references therein]{ibata01,ibata02,ivez,newberg,majewski} and globular clusters \\citep{bell03a,bell03b,bell03c}. A similar case has been reported in M31 \\citep{m31s1,m31s2} and systematic searches in galaxies outside the Local Group are beginning \\citep{poh}.  An even more interesting case is provided by the recently discovered ring structure \\cite[the Monoceros Ring, hereafter the Ring, for brevity][]{newberg,yanny,ibaring,majewski,rochap} near the Galactic plane. The more recent observational results \\citep{crane,martin} suggests that the Ring is the stream remnant of an in-plane accretion that may be contributing to the build-up of the thick disk \\cite[see also][]{helmi,abadi}.  In particular \\citet[][hereafter Pap-I]{martin}, while tracing the whole extent of the Ring using M-giants from the 2MASS database\\footnote{See Cutri et al. (2003), Explanatory Supplement to the 2MASS All Sky Data Release, http://www.ipac.caltech.edu/2mass/releases/allsky/doc/explsup.html} discovered a large overdensity of \"Ring-like\" M-giants in the Canis Major constellation. The elliptical shape, the overall structure of this overdensity and its coincidence with a noticeable grouping of globular clusters \\citep{bell03c} led \\citet{martin} to the conclusion that it is the relic of the dwarf galaxy whose disruption generated the Ring. The discovered relic (hereafter Canis Major galaxy, CMa) is approximately centered at galactic coordinates (l,b)$\\sim(240^o,-7^o)$, it has a FWHM extension in the sky of $\\sim 12 \\degr$ in the latitude direction and it is located at $\\sim 7$ kpc from the Sun (i.e., $(m-M)_0\\sim 14.25$), as obtained from the photometric parallax to M-giants introduced by \\citet{majewski}. The mass, luminosity and characteristic dimensions of Canis Major appear quite similar to those of the Sgr dSph. The similarity with Sgr extends to other relevant characteristics: the two objects seem to host a similar grouping of globular clusters, and the upper RGB of both galaxies is dominated by stars of similar colour (M-giants), which are present in similar numbers in both galaxies. The analogy between the two systems may provide a useful guideline in the study of the newly discovered galaxy.  \\citet{martin} obtained a near infrared (NIR) Colour Magnitude Diagram (CMD) of the discovered structure (e.g., their Fig.~8) showing a Red Giant Branch (RGB) extending to $(J-K_S)_0\\simeq 1.3$ and a pronounced Red Clump (RC) at $K_{S0}\\ge 12.5$ and  $(J-K_S)_0\\simeq 0.6$. A Blue Plume (BP) of stars (of more uncertain nature) is also detected around $(J-K_S)_0\\simeq 0.3$. The Main Sequence (MS) of the CMa population is beyond the reach of 2MASS photometry. The degeneracy among distance, age, metallicity and reddening (which is particularly high and highly variable at such low latitudes) prevents any firm conclusion about the nature of the stellar population in the CMa structure. Furthermore the distance scale based on photometric parallax of M-giants is intrinsically uncertain \\citep{bell03b,majewski} and it is likely affected by a systematic effect leading to the underestimate of true distances \\citep{martin,newberg2}. Hence, the detection of fainter stars belonging to the newly discovered galaxy, the access to alternative distance indicators than M-giants and the study of its stellar population in other wavelength ranges (than the NIR) are highly desirable to confirm the discovery presented in Pap-I and to obtain a clearer characterization of the system.  Here we present the first detection of the Main Sequence (MS) of the CMa system in a wide field located $\\sim 4.2\\degr$ apart from its center, from deep B and V photometry of very recent ESO-WFI images taken from the ESO archive. We also take advantage of the fact that the low latitude region of sky covered by the CMa galaxy hosts a number of galactic open clusters (OC), most of them in the foreground. Thus we searched for existing wide field CCD photometry of these open clusters to look for evidence of a background population that may be related to the CMa galaxy. Unfortunately, modern wide field photometries of open clusters are quite rare since, typically, a field of few arcmin is sufficient to characterize the population of these clusters and the background population is an undesired contaminant for these kind of studies. Nevertheless, some interesting clues of the presence of CMa stars in these fields have been obtained, providing useful indications on the spatial extent of the system and guidance for future observations.  The list of the fields considered in the current study is reported in Table~1, where we report the position of the fields in Galactic coordinates, the field of view covered by the available observations, the angular distance from the center of CMa, an estimate of the reddening and the source of the data. The position of the various fields with respect to the CMa system is sketched in Fig.~1. Most of the photometry (catalogues in electronic form) used in this paper have been retrieved from the WEBDA\\footnote{\\tt http://obswww.unige.ch/webda/webda.html} database \\citep{mermi}. The adopted reddening values are taken from the maps by \\citet[][SFD98]{schlegel}, if not otherwise stated. All over the paper we adopt the reddening laws by \\citet{dean} for the (Cousins') I band and by \\citet{rl85} for all the other passbands.  We anticipate that the data presented in this paper are not sufficient to definitely break the above quoted {\\em distance/age/metallicity/reddening} degeneracy and that reliable spectroscopic abundances and dedicated optical photometry are still needed to obtain a fully consistent picture of the main characteristics of the galaxy. On the other hand, the results presented here should be regarded as a significant progress in the understanding of the CMa system as well as a confirmation of its existence.  Criticizing the results presented in \\citet{martin} and in a previous preliminary version of the present paper circulated as a preprint, \\citet[][hereafter M04]{moma4} claimed that the CMa overdensity is completely due to the North-South asymmetry produced by the Galactic warp \\cite[see][and references therein]{yusi,kuij}. We discuss this point in Sect.~4, demonstrating that the M04 hypothesis is clearly confuted by several independent observational evidences.  \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Explored Fields.} \\begin{tabular}{@{}lrrcccc@{}} \\hline Field Name & l$\\degr$ & b$\\degr$ & f.o.v. & dist\\footnote{Angular distance from the centroid of CMa, in degrees.}&E(B-V)\\footnote{From the reddening maps by \\citet{schlegel}} & Ref. \\\\ \\hline F-XMM   & 244.2 & -8.2 & $34\\arcmin\\times 33\\arcmin$ &4.2$\\degr$ &0.14 & This work\\\\ NGC~2477 (B,V) & 253.6 & -5.8 & $34\\arcmin\\times 33\\arcmin$ & 13.6$\\degr$&0.77 & \\citet{moma} \\footnote{The (V,I) photometry by \\citet{kas} has been also considered (f.o.v=$14.7\\arcmin\\times 14.7\\arcmin$).}\\\\ Tombaugh~1& 232.3 & -7.3 & $7.0\\arcmin\\times 7.0\\arcmin$ & 7.7$\\degr$&0.43& \\citet{tom1}\\\\ Berkeley~33& 225.4 & -4.6 & $10.5\\arcmin\\times 10.5\\arcmin$ & 14.8$\\degr$&0.75& \\citet{be33}\\\\ NGC~2204   & 226.0 & -16.1& $14.7\\arcmin\\times 14.7\\arcmin$& 15.9$\\degr$&0.10&\\citet{kas} \\\\ Melotte 66 & 259.6 & -14.3& $14.7\\arcmin\\times 14.7\\arcmin$& 20.2$\\degr$&0.22&\\citet{kas} \\\\ NGC~2243   & 239.5 & -18.0& $15.0\\arcmin\\times 15.0\\arcmin$& 10.5$\\degr$&0.07 & \\citet{2243kal}\\\\ &{\\bf Open Clusters}\\footnote{Possibly associated with CMa, see Sect.~5.}&  & & & & \\\\ Arp-Madore 2& 248.1&-5.9 & $5.5\\arcmin\\times 5.5\\arcmin$& 8.3$\\degr$ &0.73&\\citet{am2} \\\\ Tombaugh~2  & 232.9& -6.9& $10.4\\arcmin\\times 10.4\\arcmin$& 6.9$\\degr$&0.39 & \\citet{tom2}\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}   \\begin{figure} \\includegraphics[width=84mm]{fig1.ps} \\caption{Relative positions on the sky (Galactic coordinates) of the CMa galaxy and of the fields considered in the present analysis. The center of CMa is represented as a cross, the two concentric circles have radius r$=6^o$ deg and r$=12^o$, corresponding to the HWHM and FWHM of the distribution of CMa M-giants as derived in Pap-I, respectively. The clusters that are possible members of the CMa system are represented as open stars. } \\end{figure}  \\begin{figure*} \\includegraphics[width=168mm]{fig2.ps} \\caption{CMD of the $\\simeq 0.5\\degr \\times 0.5\\degr$ XMM field, located at $\\simeq 4.2\\degr$ from the center of the CMa system. The error bars at $B-V\\simeq 2.2$ are the average photometric errors as a function of magnitude. The dotted, dashed and continuous lines indicate the locus of the CMD where the joint Completeness Factor (CF) is 80\\%, 50\\% and 5\\%, respectively. The joint CF as a function of V magnitude (computed at $B-V=1.0$) is shown in the inset. } \\end{figure*} ",
        "conclusions": "We have detected, for the first time, the Main Sequence of the main stellar population of the newly discovered galaxy in Canis Major (Pap-I), from photometry of a wide field located at $\\simeq 4.2\\degr$ from its center (the XMM field). From the analysis of the CMD of the XMM field  we conclude that:  \\begin{itemize}  \\item The stellar population {\\em observed} in the XMM field has no counterpart in the R03 Galactic model.  \\item The distance modulus of CMa is $(m-M)_0\\simeq 14.5\\pm 0.3$, where the reported uncertainty  mainly reflects the uncertainty in the actual average metallicity of the system, that we assumed to lie in the range $-0.3\\le [M/H]\\le -0.7$, according to \\citet{martin,crane}. The corresponding distance is $D_{\\sun}\\simeq 8.0 \\pm 1.2$ kpc.  \\item The age of the main population of CMa lies in the range $4\\sol$ age$\\sol 10$ Gyr, slightly depending on the assumed metallicity.  \\item There is a sparse population of younger stars (age$\\sol 1-2$ Gyr) that is also associated with the CMa system (Blue Plume).  \\end{itemize}   Tentative detections of stars related to the CMa/Ring system are presented for the fields surrounding the open clusters NGC~2477, Tombaugh~1 and Berkeley~33. The cases of two open clusters (AM~2 and Tombaugh~2) that may be physically associated with the CMa system have also been discussed. The results presented in this study provide independent confirmation, full support and complement the results of \\citet{martin}, that described for the first time the CMa stellar system.  The claim by \\citet{moma4} that the CMa overdensity could be entirely due to the South/North asymmetry produced by the Galactic warp has been demonstrated to lack the essential observational basis. Using the F-XMM data we have shown that the warp model encapsulated in the R03 model (used by M04 to substantiate their claim) fails to reproduce both the star counts {\\em and} CMD morphology observed in this field, located at $\\simeq 4.2\\degr$ from the center of the CMa system. Using Red Clump and M giant stars from 2MASS we have shown that (a) the CMa overdensity is much larger than that produced by the warp, and (b) it is found only at a well defined distance not compatible with the warp, fully confirming the results of Pap-I and demonstrating that if all the available information is taken into account the CMa overdensity clearly emerges {\\em above} the ``noise'' provided by the Galactic warp. Using GSC2.2 data we demonstrated that the Blue Plume population is a distinctive characteristic of the CMa system that has no comparable counterpart neither in Northern Control Fields or in fields sampling the real warp.  A better characterization of the stellar content and star formation history of the CMa system must wait for dedicated photometry and spectroscopy fully suited for the study of this difficult target, immersed in very contaminated and extincted regions of the sky. Here we tried to obtain the best possible characterization of the system from existing and publicly available data, as an immediate follow-up of the discovery.  Future searches for CMa stars in the background of other targets may prove extremely useful in characterizing the spatial distribution of the galaxy and, in particular, of the associated Ring. The relevant change of view is that what may be perceived as an undesired contaminant of a given observation may hide more astrophysical information than the target itself. This is true either if the background object is an unknown stellar system (as CMa or Sgr) either if it is a poorly characterized Galactic component, as the stellar warp of the Milky Way \\citep{binney,kuij}. The lessons provided by the Sgr dSph (and Stream) and by the CMa galaxy (and Ring) suggests to all of us a renewed attitude in the analysis of astronomical data, to avoid missing relevant pieces of information hidden in the background of our main targets."
    },
    "0311/astro-ph0311355_arXiv.txt": {
        "abstract": "The mechanisms which could lead to chemo-thermal instabilities and fragmentation during the formation of primordial protostars are investigated analytically. We introduce new analytic approximations for \\HH cooling rates bridging the optically thin and thick regimes. These allow us to discuss chemo--thermal instabilities up to densities when protostars become optically thick to continuum radiation ($n\\equiv\\rho/m_H\\lsim10^{16-17}\\;{\\rm cm^{-3}}$).  During the proto-stellar collapse instabilities are active in two different density regimes. In the well known ``low density'' regime ($n\\sim10^8-10^{10}\\;{\\rm cm^{-3}}$), instability is due to 3-body reactions quickly converting atomic hydrogen into H$_2$. In the ``high density'' regime ($n\\gsim 10^{14}\\;{\\rm cm^{-3}}$), another instability is triggered by the strong increase in the cooling rate due to \\HH Collisional Induced Emission (CIE). In agreement with the three dimensional simulations, we find that the ``low density'' instabilities cannot lead to fragmentation, both because fluctuations are too small to survive turbulent mixing, and because their growth times are too slow. The situation for the newly found ``high density'' instability is analytically similar. This continuum cooling instability is as weak as ``low density'' instability, with similar ratios of growth and dynamical time scales, as well as allowing for the necessary fragmentation condition $t_{cool}\\lsim t_{dyn}$. Because the instability growth timescale is always longer than the free fall timescale, it seems unlikely that fragmentation could occur in this high density regime. Consequently, one expects the first stars to be very massive, not to form binaries nor harbour planets.  Nevertheless, full three dimensional simulations are required to be certain.  Such 3D calculations could become possible using simplified approaches to approximate the effects of radiative transfer, which we show to work very well in 1D calculations, giving virtually indistinguishable results from calculations employing full line transfer. This indicates that the effects of radiative transfer during the initial stages of formation of primordial proto--stars are local corrections to cooling rather than influencing the energetics of distant regions of the flow. ",
        "introduction": "In current scenarios for the formation of the first stars, molecular hydrogen plays a prominent role, providing the only cooling mechanism for metal-free gas at $T\\lsim 10^4\\;{\\rm K}$. \\HH cooling is responsible for both the collapse of the gas inside the small ($M_{\\rm halo} \\sim 10^6\\;\\msun$) cosmological halos where primordial star formation is believed to occur (\\eg Couchman \\& Rees 1986; Tegmark \\etal 1997, Abel \\etal 1998), and for the fragmentation of the gas itself into ``clumps'' with masses of $\\sim 100-1000\\;\\msun$, as is seen in the simulations of Abel, Bryan \\& Norman (1998, 2000) (see also Bromm, Coppi \\& Larson 2002). The fate of these objects is unclear. Several previous studies, based on analytical arguments (\\ie stability analysis) or single-zone models, led to different conclusions about the fragmentation properties of primordial clouds. One important instability-triggering mechanism was first suggested by Palla, Salpeter \\& Stahler (1983; hereafter PSS83): the onset of 3-body \\HH formation at number densities $n\\gsim 10^8\\;{\\rm cm^{-3}}$ (where $n\\equiv\\rho/m_{\\rm H}$) causes a fast increase in the cooling rate, possibly leading to fragmentation on a mass scale of $\\sim 0.1\\msun$. A similar result was obtained by Silk (1983; hereafter S83) by applying a criterion for ``chemo-thermal instability'' (first discussed by Sabano \\& Yoshii, 1977) to an object where 3-body \\HH formation is important.  Numerical simulations have recently shed new light upon the issue, and in particular Abel, Bryan \\& Norman (2002; hereafter ABN02) were able to reach the 3-body reactions density regime finding that these reactions actually lead to the formation of a single fully molecular core (with a mass of $\\sim 1\\;\\msun$) at the centre of each clump. They point out that in their simulations they see several chemo-thermally unstable regions, but that the instabilities do not lead to fragmentation because turbulence efficiently mixes the gas, erasing each fluctuation before it can grow significantly. Recently, Omukai \\& Yoshii (2003; OY03) found a somewhat different explanation by improving the original S83 investigation and pointing out that instabilities growth was too slow.  The further evolution of such an object is still uncertain: due to the high molecular fraction, many of the numerous \\HH roto-vibrational lines (accounting for most of the cooling) become optically thick just after the formation of the molecular core. For this reason, it becomes impossible to estimate the \\HH cooling rate without a complete treatment of radiative transfer, which currently can not be included into full 3-D hydrodynamical simulations. Simpler 1-D studies, such as those in Omukai \\& Nishi (1998; hereafter ON98) and Ripamonti \\etal (2002; hereafter R02), can follow the further evolution, giving useful predictions about the physical processes driving the collapse, or about the properties of the small hydrostatic protostellar core which finally forms in the centre; on the other hand, their intrinsic spherical symmetry prevents them from giving direct predictions about fragmentation.  An interesting result of ON98 and R02 is that during the collapse, the protostellar object experiences a phase when cooling is dominated by the effects of \\HH Collision-Induced Emission (CIE), the opposite process of the more commonly known Collision-Induced Absorption, or CIA. Both of them find that, at the beginning of the CIE-cooling phase, the molecular core is optically thin to CIE-produced continuum photons, and there is a very fast increase in the cooling rate until the continuum optical depth exceeds unity. However, neither of these previous studies pointed out that during this phase the core conditions resemble the ones that, at the onset of 3-body \\HH formation, led to the chemo-thermal instability: the core is optically thin (which is believed to be a necessary condition for fragmentation; see Rees 1976) and the cooling rate is undergoing a dramatic increase.  In the next sections, we will investigate more thoroughly the issue of chemo-thermal instability, with special attention to the effects of CIE cooling. In section 2 we will describe the main cooling processes, giving a brief account of CIE emission, and some useful approximations for \\HH line cooling. In section 3 we examine the conditions for the insurgence of chemo-thermal instabilities, and argue whether they can cause fragmentation. Finally, we discuss the results in section 4. ",
        "conclusions": "We have investigated the chemical and thermal instabilities of primordial metal-free gas during its collapse towards the formation of a protostar. We have focused on the Collision Induced Emission (CIE) and on its effects on fragmentation. We point out that even when \\HH line cooling becomes optically thick at densities $n\\gsim10^{10}\\;{\\rm cm^{-3}}$ this does not necessarily set a minimum mass scale for fragmentation due to an opacity limit. At higher densities ($n\\gsim10^{14}\\;{\\rm cm^{-3}}$) the optically thin continuum (CIE) cooling dominates, and the collapsing protostellar cloud can once again fulfill the necessary (but not sufficient) fragmentation criterion $t_{\\rm cool}<t_{\\rm dyn}$.  We present accurate approximations to cooling and chemical rates that are particularily useful for analytical studies. A detailed analysis of the chemo--thermal stability of collapsing primordial gas leads us to conclude that the thermal instability arising from optically thin continuum cooling is of similar strength as the three body H$_2$ formation instability discussed previously (PSS83; S83; OY03). However, our analytical analysis of the low density instability confirms the full 3D simulations results of Abel et al.~(2002) in showing that formally the instability is very strong yet it cannot grow sufficiently fast to lead to independently collapsing fragments.  In hindsight, it may not be too surprising that this instability does not lead to fragmentation. It fulfills the necessary (but not sufficient) condition of leading to a cooling time shorter than the dynamical time. An unstable patch of gas will, however, grow initially only iso--barically as it is compressed from the surrounding higher pressure regions. So although the cooling rate can increase dramatically as one increases the H$_2$ fraction by a factor one thousand the minimum temperature the gas cools to cannot.  Consequently, correspondingly small density fluctuations are formed which do not become gravitationally unstable and become mixed back in into the slightly warmer surrounding regions.  Given the efficiency of collision induced emission cooling (Figure~\\ref{cooling_comparison}) and its relatively small effect on the temperature in the one dimensional results one is tempted to conclude that it will also not lead to fragmentation. This is most likely a fair assessment since again the growth times are long and the initial iso--baric growth leads to only small overdensities. However, in a fully three--dimensional calculation the CIE cooling might be able to cool a disk forming around the primordial proto--star sufficiently far in order to become gravitationally unstable and form stellar, or even planetary size companions (see e.g. Boss 1993 and Boss 2002). Given the dramatic accretion rates of these proto--stars, however, it is not clear whether a sufficiently massive disk may form in the first place preventing a conclusive answer.  Despite the simplicity of our derivation the analytical rates for the \\HH line cooling rate in the optically thick regime and optically thick continuum cooling they are in remarkably good agreement with the radiative transfer calculations of R02 and ON98 that followed the detailed transfer of hundreds of molecular hydrogen rotational and vibrational lines.  We have shown that implementing these simple approximations into the one dimensional Lagrangian hydrodynamics code developed by R02 leads to virtually indistinguishable results over all relevant density and temperature ranges. This clearly demonstrates that at least in the initial phases of primordial proto--stellar formation the effects of radiative transfer are in essence a local correction to cooling rather than a transport of energy to distant regions of the hydrodynamic flow. This is an important result which, given the analytical and purely local cooling correction factors derived here should one allow to follow the proto--stellar collapse in three dimensions all the way to stellar densities employing the recently developed versions of 128bit adaptive mesh refinement techniques of Bryan, Abel \\& Norman (2001) or similar high dynamic range hydrodynamic codes."
    },
    "0311/astro-ph0311449_arXiv.txt": {
        "abstract": "{We present a perspective of strange quark matter research in the 1991-2001 decade focused on astrophysics topics, with particular attention to open problems. We outline the basic concepts and developments in the field, paying attention to the established research and promising perspectives. An analysis of the whole literature (experimental searches, theory, astrophysical/cosmological) serves to point out some general trends, supported by a fairy complete statistical sample of published works. } ",
        "introduction": "\\bigskip  The history of sciences register many cases of (almost) simultaneous works giving rise to an entirely new area of research. It is also quite common to be able to track backwards the branching of an existing area, generally much after those initial events. A less common situation is to witness that process from the start, and additionally being able to collect virtually all the papers which in turn help to, identify the main trends. Strange quark matter (SQM) belongs to this interesting category. Because of the quick development of several ideas about SQM, a benchmark event (Workshop on SQM in Physics and Astrophysics, Aarhus University, Denmark 1991) \\cite{AA}  made possible to join a large fraction of the active researchers in the field. Even today, and taking into account a considerable growth in the number of the latter, it would be possible to repeat the achievement (a few conferences have covered a major portion of SQM field in the following years, although other current topics were a main part of the program). Actually most of the colleagues that contributed to SQM physics/astrophysics are alive and still active, while several others enter the field each year.  In addition, and as a case study for the history of sciences, SQM presents some peculiarities worth noting. Perhaps the most notorious is that after decades of existence it is still completely speculative, yet vigorously pursued in several directions. How did this situation affected the form in which the research is done? Which are the main avenues streets and sidewalks? The purpose of this review is not to solve these problems, but rather to present an assessment of assorted astrophysics, with the aim of formulating quite clearly a few persistent questions about SQM which often remain hidden or implicit. We have no intention of writing a full review (many excellent ones have appeared over the years, see the reference list), but rather to offer a personal perspective of the evolution we have witnessed and future trends. In the present work we will focus on the last decade (1991-2001) astrophysical developments mainly. Some facts will be substantiated by a brief (but almost complete) compilations of the bibliography of the period, available at the address www.iagusp.usp.br$\\setminus\\sim$mpallen in .txt form which may be considered as the electronic companion to the Aarhus compilation. ",
        "conclusions": ""
    },
    "0311/astro-ph0311480_arXiv.txt": {
        "abstract": "Systematic effects on HI and H$\\alpha$ long-slit observations make a measurement of the inner slope of the dark matter density distribution difficult to determine. Halos with constant density cores and ones with $r^{-1}$ profiles both appear consistent with the data, although constant density cores generally provide better fits. High-resolution, two-dimensional velocity fields remove most of the systematic effects, yet as a result of noncircular and random motions the inner slopes still cannot be accurately measured.  Halo concentration parameters provide a more useful test of cosmological models because they are more tightly constrained by observations.  The concentration parameters for LSB galaxies appear consistent with, but on the low end of the distribution predicted by CDM. ",
        "introduction": "Studies of dark matter in spiral galaxies have been plagued by uncertainties in the mass-to-light ratio (M/L) of the stellar disk, resulting in a large degeneracy in the mass modeling (e.g., van Albada et al. 1985). This degeneracy can largely be avoided by studying dwarf and low surface brightness (LSB) galaxies. Even though the central parts of the observed rotation curves can in principle be explained by scaling up the contribution of the stellar disks (Swaters 1999; Swaters, Madore, \\& Trewhella 2000), the inferred stellar M/Ls are much higher than expected from population synthesis modeling (e.g., Bell \\& de Jong 2001).  For reasonable M/Ls, these galaxies are dominated by dark matter at all radii, making dwarf and LSB galaxies ideal for studying the properties of dark matter.  The power law slopes of the central dark matter density distributions $\\rho(r)\\propto\\rho^{-\\alpha}$ are of particular interest. Cosmological simulations indicate $\\alpha$ depends on the nature of dark matter. A measurement of $\\alpha$ may thus provide constraints on the nature of dark matter and theories of galaxy formation. Unfortunately, it is difficult to measure the inner slope $\\alpha$ accurately, both from simulations (e.g., Power et al. 2003; White, this volume; Navarro, this volume), and from observations, as is discussed in detail below (see also de Blok, this volume). ",
        "conclusions": ""
    },
    "0311/astro-ph0311163_arXiv.txt": {
        "abstract": "Currently $\\sim$36 different absorption bands have been detected in the infrared spectra of cold, dense interstellar and circumstellar environments.  These are attributed to the vibrational transitions of $\\sim$17 different molecules frozen on dust grains. We review identification issues and summarize the techniques required to extract information on the physical and chemical evolution of these ices. Both laboratory simulations and line of sight studies are essential. Examples are given for ice bands observed toward high mass protostars, fields stars and recent work on ices in disks surrounding low mass protostars. A number of clear trends have emerged in recent years. One prominent ice component consists of an intimate mixture between H$_2$O, CH$_3$OH and CO$_2$ molecules. Apparently a stable balance exists between low temperature hydrogenation and oxidation reactions on grain surfaces.  In contrast, an equally prominent ice component, consisting almost entirely of CO, must have accreted directly from the gas phase. Thermal processing, i.e. evaporation and crystallization, proves to be readily traceable in both these ice components.  The spectroscopic signatures of energetic processing by cosmic rays and high energy photons from nearby protostars are weaker and not as well understood. A fundamental limitation in detecting complex, energetically produced (and also some simple) species is blending of weak features in the spectra of protostars.  Sophisticated techniques are required to extract information from blended features. We conclude with a summary of key goals for future research and prospects for observations of ices using future instrumentation, including SIRTF/IRS. ",
        "introduction": "The infrared spectra of continuum emitting sources (stars, protostars, evolved stars, active galactic nuclei) extincted by large columns of cold dust in molecular clouds, protostellar envelopes, disks, shells, or circumnuclear rings, show a wealth of absorption features. Most of these are attributed to the vibrational modes of volatile solid state molecules ($T_{\\rm subl}\\leq 100$ K).  Over half a century ago (Bates \\& Spitzer 1951) it was realized that the formation of 'ices' is catalyzed by the cold surfaces of silicate and carbon-based dust particles originating from the outflows of novae, supernovae, and stars on the asymptotic giant branch (see reviews by Patzer et al., Jura et al., and Clayton et al. in these proceedings). When such particles enter the dense interstellar medium, atoms and molecules strike and stick to the surfaces of these cold dust grains. At average molecular cloud densities ($\\sim 10^4~{\\rm cm^{-3}}$), atoms and molecules accrete from the gas at the rate of about one per day (e.g., Tielens \\& Allamandola 1987). The sticking coefficients for most of these species at 10 K are close to unity with the exception of He and H$_2$. Atoms (such as H, D, C, O) are mobile on the surface and are able to find a compatible reaction partner.  Hydrogen atoms migrate rapidly across the surface by quantum-mechanical tunneling (Manico et al. 2001). Heavy atoms (e.g., C and O) can diffuse by thermal hopping (e.g., Tielens \\& Allamandola 1987). Efficient adsorption and diffusion of species allows chemical reactions on the surface to occur and to form simple and complex molecular species by, for example, exothermic hydrogenation reactions (e.g. Tielens \\& Hagen 1982; Charnley et al. 1992).  \\begin{figure} \\plotfiddle{fig1.eps}{310pt}{0}{65}{65}{-140pt}{0pt} \\caption{A grain in space. Possible formation and evolution processes of the icy mantles are indicated.  Interstellar grains are schematically represented by spheres, although this is likely unrealistic (e.g. \\S 5).} \\end{figure}  At high atomic abundances, for example early in the cloud life time or at cloud edges, an H$_2$O--rich ice mantle is thus rapidly formed. Since hydrogenated species like H$_2$O, and also CH$_3$OH and NH$_3$ have large dipole moments, these types of ices, including trace species diluted in them, are referred to as `polar' ices in the literature (Fig.~1a).  At low atomic abundances and low temperatures, i.e. deeper in the cloud, the accreting gas is primarily molecular and an `apolar' ice, rich in CO and possibly N$_2$ or O$_2$ is formed. Depending on the ice formation history (e.g.  increasing densities and molecular abundances as a function of time, or migration of grains in a turbulent cloud) apolar and polar ices may be present on the same grain in an onion-like structure (Fig. 1b), or at different locations along the same line of sight.  Polar and apolar ices can be distinguished spectroscopically (different band profiles), and can help trace the thermal and chemical evolution of ices, clouds, and comets (apolar ices evaporate at $\\sim$20 K and polar ices at $\\sim$90 K).  The high sticking coefficients would suggest that accretion leads to extreme depletions in the molecular gas on time-scales of about a few 10$^5$ years (e.g., Brown \\& Charnley 1990). However, astronomical observations of molecular gas in the dense interstellar medium indicate also efficient desorption of material from the grain surface. Several desorption mechanisms have been proposed including evaporation following grain heating by cosmic rays, mantle explosions, and ejection upon molecule formation (see Willacy \\& Millar 1998).  As dense clouds are often regions of star formation, dynamical events also have been proposed for removing and reforming the molecular ices within the accretion time-scale; these include sputtering in low-velocity shock waves (Bergin et al. 1998; Charnley et al. 2001) and grain-grain collisions (Markwick et al.  2000; Dickens et al. 2001).  The interaction between the interstellar gas and grains and the prevailing conditions in the line of sight are dominant factors in grain evolution. After icy grain mantles are formed through gas condensation and surface chemistry, the material can be hit by energetic photons and cosmic rays, which triggers dissociation and recombination mechanisms, and thermal processing which may result in crystallization, clathrate formation, polymerization and acid base chemistry (Fig. 1c; Ehrenfreund \\& Schutte 2000).  Determining the importance of the above-mentioned ice formation and evolution scenarios is a key topic in astrophysics and relates directly to studies of the origin and evolution of solar system ices (e.g. Mumma et al. 2003). In this review we discuss the observational status of interstellar ices. We present an up-to-date list of detected absorption features and discuss their identification and uncertainties in \\S 2. We discuss how information on ice origin and evolution can be derived from observations by comparison with laboratory simulations and line of sight characteristics (\\S 3). Detailed band profile studies of ices in different environments and the significance of thermal processing are discussed in \\S 4.  In \\S 5 we touch upon the importance of observing isotopes in the ices and deriving the shape of interstellar grains. Subsequently we discuss the observational evidence for evolution of interstellar ices from a chemical perspective (\\S 6).  Finally, we summarize outstanding issues and what can be expected from observations of interstellar ices with future instrumentation (\\S 7).  \\begin{figure} \\plotfiddle{fig2.eps}{490pt}{0}{90}{90}{-170pt}{0pt} \\caption{Mid-infrared ISO/SWS spectra of protostars showing the richness of ice absorption features. (a) The low mass protostar Elias 29 in $\\rho$ Oph at a spectral resolution of $R$=400 (Boogert et al. 2000a).  The bar on top indicates the wavelength coverage and spectral resolution of the SIRTF/IRS instrument (\\S 7). (b) The rich mid-infrared spectra of the massive protostars NGC 7538 : IRS9 (top) and W 33A (bottom). Data taken from Whittet et al. (1996) and Gibb et al. (2000). Identifications labeled in square brackets are uncertain.} \\end{figure} ",
        "conclusions": "Increasingly sensitive observations of interstellar ices have revealed several prominent, and often unexpected, trends, providing clues to the evolutionary pathway of interstellar to solar system ices. One such trend is the remarkable simplicity of the interstellar ice spectrum.  Ices consist mostly of simple species, likely the result of atomic O and H addition reactions on dust surfaces.  A stable balance between some of these reactions is reflected in the remarkably constant solid CO$_2$/H$_2$O abundance ratio (17\\%) throughout the Galaxy and the commonly observed intimate mix of these species with CH$_3$OH. The solid CO band reflects an entirely different ice formation environment, but again shows little compositional variation throughout the Galaxy.  This apolar ice phase is remarkably CO--rich. The similarity of ice composition in sight-lines sampling a variety of environments proves that to first order the production of complex species by energetic processing of simple ices is minor. However, at a low level energetic processing can presently not be excluded. Perhaps the variations of CH$_3$OH and `XCN' abundances, and the `6.0 \\mum\\ excess' (Gibb \\& Whittet 2002) are linked to locally increased photon or cosmic ray fluxes. If organic residues indeed form from ices, the inefficiency of energetic processing in dense clouds requires that multiple cycling of grains between the dense and diffuse interstellar medium takes place (Greenberg et al. 2000). Finally, although energetic processing is weak, the effects of thermal processing of ices, i.e. evaporation, crystallization and segregation, are commonly observed. In fact, ice bands can serve as tracers of the thermal history of circumstellar environments.  Many key issues on the formation and evolution of interstellar ices are still poorly known, however. In order to finally understand interstellar ice evolution a concerted effort is required, involving a variety of observational and laboratory techniques and related theory projects. Astronomical spectra sample different grain populations in one line of sight. Line of sight conditions need to be determined from complementary information, including radio observations of gas phase species and imaging of the star-forming region. Particularly exciting is the prospect of studying the properties of ices in disks that are well characterized by millimeter-wave interferometry (CARMA, ALMA, SMA) and infrared ro-vibrational CO line studies (Boogert et al. 2002b, Najita et al. 2003). Also, study of ice features in lines of sight with known UV radiation fields will shed more light on the importance of energetic processing of ices (e.g. Shenoy et al. 2003).  The importance of high spectral resolution in the infrared to separate weak, blended ice bands and gas and telluric lines, cannot be overstated. Such instruments are now available on large ground based telescopes like Keck and VLT. In spectral regions severely hindered by the earth atmosphere, such as the important 5-8 \\mum\\ region, high spectral resolution will be provided by spectrometers mounted on the airborne SOFIA observatory.  Further in the future the planned JWST and proposed ABE missions will become guaranteed milestones in the area of interstellar ices.  For example, these instruments allow an improved assessment of the significance of the apolar species N$_2$ and O$_2$ in the ices. Also, they will be a major leap forward in the continued search for the weak signatures of energetic processing.  In the near future, sensitive space based observations will be possible with the SIRTF/IRS spectrometer.  Its spectral resolution is limited (Fig. 2), but, in particular for the CO$_2$ bending mode, sufficient for detailed laboratory analysis tracing thermal processing and molecular complexity toward large samples of low mass protostars (Fig. 7). Linking the CO$_2$ analysis with ground based observations of CH$_3$OH is needed to understand the chemical history of this molecule and the large abundance fluctuations in different sight-lines.  This is of particular interest because CH$_3$OH is a starting point of complex molecule formation.  Finally, continued laboratory work will provide much needed chemical reaction efficiencies in astrophysically relevant environments. Laboratory simulations are also necessary to determine the microscopic structure of interstellar ices, which regulates key processes such as outgassing (e.g. Jenniskens \\& Blake 1996) and coagulation behavior (e.g. Chokshi et al. 1993). Perhaps related to this is the question why the interstellar ice bands are best fit with a `CDE'-type shape distribution (\\S 5)? Other major issues that have to be settled in the laboratory include the presence of PAH molecules and ions in ices (Gudipati \\& Allamandola 2003), the identification of the 6.85 \\mum\\ band, and the physical interaction between the components of inhomogeneous ices (Collings et al. 2003)."
    },
    "0311/astro-ph0311213_arXiv.txt": {
        "abstract": "{ We investigate the scaling relations of bulge and disk structural parameters for a sample of 108 disk galaxies. Structural parameters of individual galaxies are obtained from two-dimensional bulge/disk decomposition of their H-band surface brightness distributions\\thanks{Table 1 is only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/}. Bulges are modelled with a generalized exponential (S\\'ersic) with variable integer shape index $n$. We find that bulge effective scalelengths $r_e^B$ and luminosity $M^B$ increase with increasing $n$, but disk properties are independent of bulge shape. As Hubble type $T$ increases, bulges become less luminous and their mean effective surface brightness $<\\mu_e^B>$ gets fainter; disk $<\\mu_e^D>$ shows a similar, but much weaker, trend. When bulge parameters ($<\\mu_e^B>$, $r_e^B$, $M^B$) are compared with disk ones ($<\\mu_e^D>$, $r_e^D$, $M^D$), {\\it they are tightly correlated for $n=1$ bulges}. The correlations gradually worsen with increasing $n$ such that $n=4$ bulges appear {\\it virtually independent} of their disks. The Kormendy relation, $<\\mu_e^B>$ vs. $r_e^B$, is shown to depend on bulge shape $n$; the two parameters are tightly correlated in $n=4$ bulges ($r=0.8$), and increasingly less so as $n$ decreases; disk $<\\mu_e^D>$ and $r_e^D$ are well correlated ($r=0.7$). Bulge-to-disk size ratios $r_e^B/r_e^D$ are independent of Hubble type, but smaller for exponential bulges; the mean $r_e^B/r_e^D$ for $n=1$ bulges is 4 times smaller than that for $n=4$, with a spread which is 9 times smaller. Strongly barred SB galaxies with exponential bulges are more luminous than their unbarred counterparts. Exponential bulges appear to be closely related to their underlying disks, while bulges with higher $n$ values are less so; $n=4$ bulges and their disks apparently have no relation. We interpret our results as being most consistent with a secular evolutionary scenario, in which dissipative processes in the disk are responsible for building up the bulges in most spirals. ",
        "introduction": "This is the fifth paper in a series based on near-infrared imaging of a sample of disk galaxies in the Pisces-Perseus supercluster. Already in Paper III of the series (Moriondo et al. \\cite{moriondo99b}), we explored the relations between structural parameters of bulges and disks, but in the context of distance measurements and the consequent definition of their Fundamental Planes. For this reason, that paper: {\\it i})~made only use of the subset of galaxies (40) with measured rotation curves, and {\\it ii})~postponed to future work the discussion of the impact of scaling relations of bulges and disks on their formation and evolution. In the present paper, we will address this issue.  The traditional picture is that bulges formed well before the disk of a galaxy, in a ``monolithic collapse'', either as the first rapid stage of the collapse of a galaxy-size density perturbation, as in the Milky Way collapse model of Eggen et al. (\\cite{eggen62}), or through the infall of enriched gas from the star-forming halo (Carney et al. \\cite{carney90}). ``Early'' bulge formation is also predicted within the framework of hierarchical cold-dark matter (CDM) scenarios of structure formation and evolution (Kauffmann et al. \\cite{kauffmann93}; Baugh et al. \\cite{baugh96}). In this case, the bulge originates from the merging of stellar disks within the merging halos, while the disk that is observed today is assembled through gas accretion at later epochs. A more recent variation of this scenario is the growth of galactic bulges through merging of dense satellites (Aguerri et al. \\cite{aguerri01}). Evidence in favor of ``early'' bulge formation rises from studies of the stellar populations in the Milky-Way bulge (Ortolani et al. \\cite{ortolani95}) and in bulges of nearby galaxies (Jablonka et al. \\cite{jablonka96}). In addition, bulge velocity fields (Kormendy \\& Illingworth \\cite{ki82}) and their high central surface brightness seem to be imprints of a formation process that involved dissipational collapse (e.g., Carlberg \\cite{carlberg}) at epochs earlier than the formation of the disks and, perhaps, of the stellar halo.  More recently, it has been suggested that the bulges in (some) disk galaxies can be formed after the disk, through secular evolutionary phenomena. Numerical simulations (Combes \\& Sanders \\cite{combes81}; Combes et al. \\cite{combes90}; Norman et al. \\cite{norman96}) have shown that a spheroidal stellar component can originate from dissipationless processes such as the thickening or destruction of a bar. Moreover, when gas is included in the simulations, dissipation accelerates the evolution, giving rise to galaxies which are more centrally concentrated (Junqueira \\& Combes \\cite{junqueira}). Indeed, some of the observed bulge characteristics such as triaxial or peanut shape and cylindrical rotation (e.g., Shaw \\cite{shaw93}; L\\\"utticke et al. \\cite{lutticke}) are predicted by such models. However, gas is not necessary for dissipation, as pointed out by Zhang in a series of papers (Zhang \\cite{zhang96}, \\cite{zhang98}, \\cite{zhang99}). Secular evolution may also operate through outward transport of angular momentum through a collective dissipative process linked to spiral density waves. Observational evidence in favor of the secular scenario include the similarity of bulge and disk color indices (e.g., Balcells \\& Peletier \\cite{balcells94}), the mass-to-luminosity ratios of the components (Widrow et al. \\cite{widrow}), the disk-like kinematics observed in  many bulges (Kormendy \\cite{kormendy93}), and the correlation of bulge and disk scalelengths (de Jong \\cite{dejong96}). Indeed, when a simple exponential is used to fit the bulge surface brightness profile, a strong correlation is found between bulge and disk scalelengths (Courteau et al. \\cite{courteau96}), seemingly supporting the secular formation scenario for bulges. The correlation was later shown to depend critically on the use of a simple exponential as a bulge fitting function, since when generalized exponentials were used instead, the mean scalelength ratio depends on the bulge shape index (Moriondo et al. \\cite{moriondo98a} -- hereafter MGH; Graham \\cite{graham01}). Thus it appears, albeit for small samples, that the ratio of bulge and disk scalelengths may not be constant over a range of Hubble types as previously claimed. This is of importance not only for the chronology of bulge and disk formation, but also for the meaning of the Hubble sequence itself in terms of photometric observables (scale-free Hubble sequence).  To further investigate bulge and disk structural parameters, and to better constrain their connection with formation scenarios, we have performed two-dimensional (2D) bulge/disk decompositions on a large sample of disk galaxies in the Pisces Perseus supercluster (see Moriondo et al. \\cite{moriondo99a}, hereafter Paper I). These decompositions are performed on image data in the H band because the near-infrared (NIR) wavelengths are the most direct tracer of the stellar mass (e.g., Rix \\& Rieke \\cite{rix}, Gavazzi \\cite{gavazzi93}). Moreover, NIR emission is much less affected by recent episodes of star formation which typically dominate disk morphology at optical wavelengths. Finally, the effects of dust extinction are mitigated at 1.65$\\,\\mu$m (H) since $A_H$ is roughly a factor of 8 lower than $A_V$. We describe the sample and the imaging data in Sect. \\ref{sample}, together with the bulge/disk decompositions. Section \\ref{results} presents our results, and we interpret them in Sect. \\ref{discussion} in the context of various evolutionary scenarios. ",
        "conclusions": "}  First, we summarize the main results obtained in the previous section. \\begin{itemize} \\item Bulge effective scalelengths $r_e^B$ increase with increasing $n$; $n=1$ bulges are smaller than those with $n>1$ (Fig. \\ref{fig:bd_n}). Neither bulge and disk $<\\mu_e>$ nor disk $r_e^D$ vary with $n$. Bulge luminosity $M^B$ increases with $n$, while disk luminosity $M^D$ is independent of $n$ (Fig. \\ref{fig:bd_n}). In particular, exponential bulges are smaller and less luminous than those with $n>1$. \\item Bulge $<\\mu_e^B>$ gets brighter as $T$ decreases (Fig. \\ref{fig:sbmean_t}). Disk $<\\mu_e^D>$ shows a similar, although weaker, trend. Bulge luminosity $M^B$ increases with decreasing $T$ (Fig. \\ref{fig:sbmean_t}), but only disks with T$\\geq$7 are significantly less luminous than their early-type counterparts. Neither bulge nor disk scale lengths $r_e^B$, $r_e^D$ vary significantly with Hubble type (Fig. \\ref{fig:sbmean_t}). \\item When bulge parameters ($<\\mu_e^B>$, $r_e^B$, $M^B$) are compared with disk ones ($<\\mu_e^D>$, $r_e^D$, $M^D$), {\\it in all cases they are tightly correlated for $n=1$ bulges} (Figs. \\ref{fig:sb_n}, \\ref{fig:logre_n}, \\ref{fig:bdabs_n}). The correlations gradually worsen with increasing $n$ such that $n=4$ bulges appear {\\it virtually independent} of their disks. \\item The scatter in the correlation between bulge $<\\mu_e^B>$ and $r_e^B$ (Fig. \\ref{fig:sbre_n}) is shown to depend on bulge shape $n$; the two parameters are tightly correlated in $n=4$ bulges ($r=0.8$), and increasingly less so as $n$ decreases. Disk $<\\mu_e^D>$ and $r_e^D$ are well correlated ($r=0.7$). \\item Luminosity ratios $B/D$ are smaller for $n=1$ bulges than for $n>1$, and increase with decreasing $T$, but with a surprisingly large spread (Fig. \\ref{fig:bdre_nt}). The value of $B/D$ is governed by the bulge luminosity, since disks have the same average luminosity independently of their shape index $n$. Bulge-to-disk size ratios $r_e^B/r_e^D$ are also smaller for $n=1$ bulges, but are independent of Hubble type. The value of $r_e^B/r_e^D$ for $n=1$ bulges is 4 times smaller than that for $n=4$, with a spread which is 9 times smaller. \\item As already known, early-type spirals tend to be more luminous than late types (Fig. \\ref{fig:absnt}); except for $T\\geq$7, this is due to the increasing contribution of the bulge. \\item Strongly barred galaxies with exponential bulges are more luminous than their unbarred counterparts. The difference tends to vanish with increasing $n$. There is no analogous trend with $T$. \\end{itemize}  \\subsection{Bulge formation and evolution \\label{formation}}   The better correlations of bulge and disk properties for exponential bulges suggest that $n=1$ bulges and their disks are closely related. The tightness of the relation worsens as $n$ increases, so that when $n=4$, bulges appear independent of their disks. In contrast, the (Kormendy) relation between bulge SFB and scalelength is very tight for $n=4$ bulges, and increasingly noisy for lower $n$ values. This implies that $n=4$ bulges are in some sense more structurally defined than those with lower $n$.  In the Introduction, we sketched the competing scenarios for bulge formation; here we use the results of our study in an attempt to distinguish among them. In exclusively ``early'' bulge formation scenarios, in which the bulge formed before the disk, it is difficult to understand the close relation we observe between exponential bulges and their disks. This relation and its {\\it gradual} degradation as $n$ increases is not easily explained by either the monolithic collapse scenario or CDM merging of primordial disks. If initial conditions were everything, such as in a monolithic collapse, bulge and disk properties should be statistically unrelated. The same would be true if stable disk growth depended on satellite accretion after bulge formation from merging disks, such as in the CDM merging scenarios. It may be that such processes operate for some (higher $n$) bulges, but they are not a good explanation for the global properties of low $n$ ones. Nor can they explain the gradual worsening of bulge/disk correlations with $n$.  A close relation between bulges and disks is, however, foreseen by the inside-out formation scheme of van den Bosch (\\cite{vandenbosch}). In this scenario, the collapsing halo forms unstable disks which, through bar instabilities, rapidly become ``bulges''. A disk is able to form only after the bulge becomes large enough to stabilize the ensuing bar. Nevertheless, in this scenario, it is difficult to explain the existence of luminous galaxies with very low $B/D$ ratios. If the presence of a disk depends on the prior buildup of a sufficiently luminous (massive) bulge, then we would not expect to find galaxies with luminous disks and relatively small bulges. Indeed, the existence of pure disk galaxies ($\\sim$2\\% of our sample, but see Gavazzi et al. \\cite{gavazzi00}) would appear to contradict this scenario.  A more natural explanation of our observations is the secular evolution scheme in which $n=1$ bulges are relatively ``young'' structures, having formed ``recently'' from their disks through internal gravitational instabilities. This process must be dissipative because of phase-space arguments (e.g., Wyse \\cite{wyse}), but may not exclusively involve bar-driven gas inflow (e.g., Pfenniger \\& Norman \\cite{pfenniger}; Combes et al. \\cite{combes90}; Friedli \\& Benz \\cite{friedli93}; Norman et al. \\cite{norman96}). The evolutionary scheme proposed by Zhang in a series of papers (Zhang \\cite{zhang96}, \\cite{zhang98}, \\cite{zhang99}) predicts that ``as a galaxy evolves through the spiral-induced collective dissipation process, its surface density distribution (i.e., disk plus bulge) will become more and more centrally concentrated, together with the buildup of an extended outer envelope''. Moreover, in a given galaxy, the observed spread in the angular momentum distribution reflects the degree to which secular evolutionary processes have been at work. Since stable stellar systems require a greater degree of velocity anisotropy as $n$ decreases (Ciotti \\& Lanzoni \\cite{ciotti97}), $n=1$ bulges must be kinematically less isotropic than higher $n$ ones. If this anisotropic component is related to tangential motion, and thus to angular momentum, such bulges must be --in the Zhang scenario-- {\\it younger}.  Let us therefore hypothesize that exponential bulges are indeed younger than those of higher $n$, independently of the Hubble type of the galaxy in which they reside. In other words, the first spheroids to be formed from their disks would be bulges with $n=1$. Such a hypothesis is supported by the blue V-H colors of $n=1$ bulges, relative to those with $n=4$ (Carollo et al. \\cite{carollo01}). It would also be consistent with the good correlation of exponential bulge properties with those of their disks as we have seen in Sect. \\ref{bd}. If indeed the bulge $n$ increases with time, as predicted by the Zhang models, then our observation that $r_e^B$ and $M^B$ increase with $n$ implies that they also increase with time. The good $n=4$ correlation of $<\\mu_e^B>$ with $r_e^B$ would then be a consequence of their older age and better defined ``relaxed'' (virialized?) structure, and their consequent better consistency with the FP. The degradation of the correlations between bulge and disk parameters with increasing $n$ would also imply that larger-$n$ bulges are increasingly more distant in time from their disks. In such a scheme, the bulge and disk branches in Fig. \\ref{fig:bdmag_m} could be interpreted as the ``evolutionary tracks'' of these components.  \\subsection{The Hubble sequence \\label{hubble}}  Appealing as this picture may be, the trends with Hubble type still need to be explained. Consistently with previous work, we find that higher $n$ values are more common in early-type spirals. Conversely, all (our sample) galaxies with $T\\geq$7 have $n=1$ bulges. Therefore, there is some indication that $n$ and Hubble type are (albeit weakly) correlated, and may both be associated with age or evolutionary state. Since the degree of evolution is expected to be partly determined by galaxy mass (Zhang \\cite{zhang99}), early type spirals which are also more luminous (massive) should in some sense be more evolved than late types. That {\\it surface brightness increases with decreasing Hubble type for bulges and to a lesser extent for disks} is consistent with this expectation. For a given (constant) rate of star formation from the formation epoch to the present day, and a sufficient gas reservoir, we would expect higher stellar surface densities in galaxies which have been forming stars for longer. If, on the other hand, star-formation rates were to depend on (gas$+$stellar) mass, then we would expect higher mass systems to have higher stellar densities. Either way, it is difficult to reconcile the higher stellar surface densities in early type spirals (Fig. \\ref{fig:sbmean_t}) with an earlier evolutionary state (or younger age); they are probably more evolved (i.e., older) systems. Early-type spirals also tend to have lower gas fractions (e.g., Roberts \\& Haynes \\cite{roberts}) than later types, presumably a result of having converted the available gas into stars.  If late-type spirals ($n=1$ bulge) are younger than early types (higher $n$), the question of bars and their effect on evolution must be addressed, since in local samples of galaxies, $>$75\\% of very late-type spirals ($T\\geq$7) are barred (Ho et al. \\cite{ho}; Hunt \\& Malkan \\cite{hunt99}). The excess of bars in very late-type systems may contribute to their morphological evolution as proposed by many groups (e.g., Combes et al. \\cite{combes90}; Pfenniger \\& Friedli \\cite{friedli91}; Friedli \\& Benz \\cite{friedli93}; Hasan et al. \\cite{hasan}; Norman et al. \\cite{norman96}). Numerical models show that bulge-like structures with an exponential $n=1$ profile can be produced from bar-like perturbations that induce inward radial inflow of disk matter (Pfenniger \\& Friedli \\cite{friedli91}; Zhang \\& Wyse \\cite{zhangwyse}). Our finding that strongly barred $n=1$ bulges tend to be more luminous than unbarred ones could reflect such evolution.  Indeed, this bar-driven evolutionary scheme has been advocated by Courteau et al. (\\cite{courteau96}) and MACH03 to explain the relative constancy of $B/D$ size ratios they find in (the Type I profiles of) their late-type spiral sample. However, they claim that the ratio $r_e^B/r_e^D$ is constant for all spiral types, with the implication that the Hubble sequence is scale free. One of the main results of our study is that several of the trends observed are clearly detected as a function of bulge shape $n$, rather than $T$; we interpret this as due to the complex interplay of the spiral classification criteria. With a sample containing few early-type ($T<2$) spirals, and thus few high-$n$ bulges, several of the trends we found would not be detectable. Hence, we find a wider range of $n$, in agreement with other groups (Andredakis et al. \\cite{andredakis95}; MH01; Graham \\cite{graham01}), and also a wider range of $r_e^B/r_e^D$. We also find that the bulge-to-disk size ratio changes significantly with $n$. While such variations with $n$ indicate that the Hubble sequence is not scale free, they do not necessarily disprove the secular evolution scenario. However, the scenario must be expanded to include the genesis of bulges with different (higher) $n$'s, together with their distinct properties.  Bar-driven inflow as the unique mechanism of secular evolution may not be sufficient to explain the full range of properties in the Hubble sequence of spirals. The analytical models of Zhang \\& Wyse (\\cite{zhangwyse}) based on viscous evolution of bars find no galaxies with $B/D \\sim 1$; their bulges are not sufficiently luminous compared to their disks. Moreover, the secular-evolution simulations of Bouwens et al. (\\cite{bouwens}) which rely on similar assumptions also produce bulges which are too small and too blue. The collective stellar dissipation scheme proposed by Zhang (\\cite{zhang99}) may also be insufficient. While qualitatively consistent with our findings, this scheme produces effects (e.g., buildup of central concentration) which are much smaller than those we observe, perhaps because the effects of gas are not included.  Evidently, some other mechanism is needed, such as mergers (e.g., Baugh et al. \\cite{baugh96}) or accretion episodes (e.g., Steinmetz \\& Navarro \\cite{steinmetz}). However, even mergers in hierarchical clustering scenarios are not able to generate the range of bulge/disk properties in our sample. Scannapieco \\& Tissera (\\cite{scannapieco03}) use hydrodynamical simulations to assess the effects of mergers on disk galaxy structure. While they are able to reproduce the bulge-to-disk size relation for the $n\\leq2$ bulges in our sample and in that of MACH03, the tight correlation of $r_e^B/r_e^D$ for $n=1$ bulges is not explained. Moreover, they fail to predict the strong increase of $r_e^B$ with $n$ (Fig. \\ref{fig:bd_n} and other work), since their models predict the opposite trend. Nevertheless, they successfully reproduce the $B/D$ ratio as a function of bulge shape, although their $n$ values tend to cluster around 1, unlike the early-type spirals in our sample which tend to be best fit by larger $n$ values (Fig. \\ref{fig:T_n}). On the other hand Aguerri et al. (\\cite{aguerri01}) are able to reproduce $n=4$ bulges with satellite accretion, but only with satellites as massive as the bulge itself.  \\subsection{Conclusions \\label{conclus}}  The hypothesis that emerges from our study is essentially the following: \\begin{itemize} \\item Disks are formed first, during an initial collapse, and have properties which depend on the mass and angular momentum of the dark-matter (DM) halo. \\item Bulges are then born first as $n=1$ systems, through secular transfer of disk material to the central regions, probably by bar instabilities. \\item With an evolutionary rate that depends on initial conditions (gas+DM mass, angular momentum), $n$ of the bulge increases, with consequent growth of $r_e^B$, $M^B$, and $B/D$, perhaps by means of the collective dissipation mechanism described by Zhang (\\cite{zhang99}), together with the effects of gas transfer. \\item As $n$ gets larger and the bulge builds up, Hubble type gets earlier, with consequent brightening (SFB) of bulges and (less so) disks, and an increase of $M^B$, $B/D$. \\item Since $n=1$ bulges are relatively unevolved systems, in a state of formation at the present epoch, they are structurally not as relaxed as larger $n$ bulges, and present a large scatter in the $<\\mu_e^B>-\\log\\,r_e^B$ (Kormendy) relation. Their stellar luminosity is enhanced by the presence of a strong bar driving the evolution. The $n=1$ bulges that we see today are those systems which have had long evolutionary times, probably because of their lower luminosity (mass). The evolutionary state of the disk would also be a function of $n$, since the entire system would be expected to have evolved faster when $n$ is larger. \\item Mergers of bulge/disk systems and accretion of gas probably affect the buildup of the bulge, but are probably important only for the creation of high-$n$ bulges. \\item The correlation between ``evolutionary state indicators'' $n$ and $T$ is not expected to be perfect because of the complex nature of spiral classification, and because of variations in galaxy mass, angular momentum, bar strength, gas fraction, dark halo merging history, and perhaps other parameters. \\item The gradual growth of a bulge as $n$ increases is implied by the continuity we observe in bulge properties as a function of $n$. All trends are tight for $n=1$, and steadily worsen with increasing $n$, except for the Kormendy relation which is best followed by structurally evolved systems ($n=4$). \\end{itemize}  One prediction of such a scenario is a certain spread in stellar age for high-$n$ bulges. Detailed simulations are necessary to assess whether this prediction is in conflict with the level of synchronicity in the star-formation history observed in spheroids (Bower et al. 1992). We also would expect that exponential bulges, prevalent in late-type spirals and closely connected with their disks, would be more dominated by rotation than their higher $n$ counterparts. Low-luminosity spheroids are known to have a higher degree of rotational support than more luminous ones (Davies et al. 1983), so it may be that this is true also for $n=1$ bulges relative to those of higher $n$.  Clearly more modelling is needed to understand the full range in properties along the Hubble sequence, and what they imply for bulge formation. However, secular evolution appears to be a strong driver of the Hubble sequence. Our observations are compatible with the hypothesis that $n$ and $T$ are both rough measures of bulge age, and that bulges are, at least initially ($n=1$), formed secularly from disks."
    },
    "0311/astro-ph0311025_arXiv.txt": {
        "abstract": "{We report on the discovery of a very low mass (VLM) star in the solar neighborhood, originally identified as an optical counterpart of a flaring X$-$ray source detected in the \\ros\\ All-Sky survey. Optical spectroscopy and infrared photometry consistently reveal a spectral type of $ M9 \\pm 0.5$ and a distance of $\\approx$~11$\\pm$~2~pc. The optical counterpart of \\obj\\ shows a large proper motion of 1.08$\\pm$0.06~$\\arcsec$/year. \\obj\\ is the fourth object among the VLM stars displaying a huge X$-$ray flare, reaching the unprecedent value of $L_X/L_{\\rm bol}~\\simeq$~0.1. ",
        "introduction": "Imaging sky surveys performed during the last decade at different wavelengths (optical, near--infrared) have uncovered a large population of very low--mass (VLM) stars and brown dwarf candidates in star forming regions, open clusters and in the general galactic field.  According to an estimate presented by Reid et al. (\\cite{reid2003}) the current 10~pc sample is only $\\simeq $~75\\% complete and the level of incompleteness is somewhat higher in the southern sky and with the latest type objects.  With spectral types of M7 and later, these objects, sometimes called ``ultracool M dwarfs'', are so faint optically that even nearby ones eluded searches based on the optical, near--infrared and high proper motion sky surveys, and, therefore a detection of a new VLM star in the solar vicinity is of vital importance in the determination of fundamental parameters such as the luminosity function, the mass function, the kinematics of those stars and other properties.  The strength of H${\\alpha}$-emission increases as cooler, lower-mass stars are considered (Hawley et al. \\cite{haw96}; Gizis et al. \\cite{gizis00}). The emission is up to 100\\% at spectral types M7-M9. The observed decline in H${\\alpha}$-emission for old ($>$1\\,Gyr) L-type dwarfs in the field suggests that the activity borderline is close to  objects of substellar masses (Gizis et al. \\cite{gizis00}; Basri \\cite{basri01}; Mohanty \\& Basri \\cite{mohanty}).  Analyzing X$-$ray properties of VLM stars based on \\ros\\ observations, Fleming et al.~(\\cite{Fleming93}) showed that there is no apparent decrease in X$-$ray flux (hence coronal heating efficency) down to spectral type M5, as was suggested by Mullan (\\cite{mullan}) analyzing \\einst\\ data.  Flares as another property  of chromospheric and coronal activity represent a phenomenon which is important in ultracool dwarfs, because at these low temperatures (and masses) these objects are fully convective, and variability is often attributed to rotational modulation of star spots produced by magnetic activity. In solar-type stars it is believed to be due to the so-called $\\alpha \\Omega$ dynamo. This mechanism no longer operates in VLM stars and brown dwarfs, but as these objects are fully convective, a turbulent dynamo could come into operation (see Chabrier \\& Baraffe \\cite{chabrier00a} and references therein). This means that the change in interior structure is expected to result in a change of the field sustaining dynamo, and, therefore X$-$ray emission properties of late-type stars.  Flares were reported for a number of VLM stars, observed in ultraviolet, X$-$ray, and radio frequencies, as well as in optical spectra or via  photometry. Bursts of X$-$ray emission interpreted as coronal flares were detected from the long-known M8V dwarf \\object{vB 10} (Fleming et al. \\cite{vb10}), the M9V dwarf \\object{LHS~2065} (Schmitt \\& Liefke \\cite{lhs2065}) by \\ros\\ and the M9 brown dwarf \\object{LP~944$-$20} (Rutledge et al. \\cite{lp944}) by \\chan\\ observations.  \\object{vB 10} showed a far-ultraviolet flare observed by Linsky et al. (\\cite{linsky}) using the GHRS on the \\hst . Radio emission has been observed from \\lp\\ and some other ultracool dwarfs (Berger et al. \\cite{berger01}, Berger \\cite{berger02}). From the intensity of the continuous emission and flaring at 8.5\\,GHz detected in these objects, Berger (\\cite{berger02}) concluded that the emission mechanism is most likely synchrotron radiation and not of thermal origin.  This indicates the presence of magnetic fields and electron column densities similar to those inferred for earlier-type flaring M dwarfs.  H${\\alpha}$-flares were observed on several occasions, e.g. in the M8.5 \\object{APMPM~J2354$-$3316CM9.5} (Scholz et al. \\cite{scholz2003a}) and M9.5 dwarfs \\object{2MASSW~J0149090$+$295613} (Liebert et al. \\cite{liebert99}), \\object{BRI~0021$-$214} (Reid et al. \\cite{reid99})  and even in the cooler L5-dwarf \\object{2MASS~J01443536$-$0716142} (Liebert et al. \\cite{liebert03}).  Thus, it has become apparent that the level of quiescent chromospheric and coronal activity as measured by the H${\\alpha}$ emission line and X$-$ray fluxes relative to the total bolometric flux generally declines with spectral type later than M7.  However, {\\emph no} clear change in flare activity diagnostics is found in VLM stars (Fleming et al. \\cite{Fleming95}; Mokler \\& Stelzer \\cite{mokler}; Mart{\\'\\i}n \\& Bouy \\cite{martin}).  We performed a systematic search for X$-$ray variability in the \\ros\\ All-Sky survey data (Hambaryan et al. \\cite{skymi}). Our technique is customized for the detection of flare-like variability. Here we report on the identification of the optical counterpart of the strong flaring X$-$ray source \\obj\\ as a new ultracool dwarf star in the solar neighborhood by optical spectroscopy and astrometry. ",
        "conclusions": "\\label{discussion}  The object \\obj\\ was detected in the RASS due to the presence of a huge X$-$ray flare. This was uncovered by our time variability analysis of all RASS-BSC sources. The quiescence X$-$ray flux is consistent with the background level.  Our follow-up optical spectroscopy and proper motion study allowed us to determine the spectral type, M$9\\pm0.5$, and the absolute magnitude $M_K= 10.1\\pm 0.4$ of the most likely optical counterpart of the flaring X$-$ray source. It turned out that this object is a nearby, $d \\approx 11\\pm2$\\,pc, high proper motion, $\\mu =$ 1.08$\\pm$0.06$\\arcsec$/yr, previously unknown member of the solar neighborhood.  In order to compare the X$-$ray properties of \\obj\\ with other ultracool stars we estimated $\\log{(L_{X}/L_{bol})}$. From a spectral fit (Sect.~\\ref{xray_fl}) we estimated time averaged and flare peak X$-$ray luminosities of $L_{X}^{mean} \\approx (0.7-1.9) \\times 10^{28} {\\rm \\,\\,erg\\,\\,s^{-1}}$ and $L_{X}^{peak} \\approx (1.0-1.6) \\times 10^{29} {\\rm \\,\\,erg\\,\\,s^{-1}}$, using unabsorbed fluxes in the passband $0.1-2.4$\\,\\,keV. The luminosity uncertainties were determined by a combination of uncertainties of the distance and of the parameters of the spectral fit.  For an estimate of an upper limit of the X$-$ray luminosity of \\obj\\ in its quiescent state we converted from upper limit count rates (0.011 and 0.015) to fluxes by means of an energy conversion factor $ECF \\approx 7.0\\times 10^{-12} {\\rm \\,\\,erg\\,\\,cts^{-1} cm^{-2}}$ which was computed by us in XSPEC assuming a Raymond-Smith plasma (model {\\emph raymond \\/}) with temperature ${\\rm kT = 0.25 \\,\\,keV}$ and absorption column density N$_H = {\\rm 2.4} \\times 10^{-18} {\\rm cm^{-2}}$ (see also H{\\\"u}nsch et al.~\\cite{huensch99}, Fleming et al.~\\cite{fleming03}). Using a distance of $11.1 \\pm 2.0$\\,pc we arrived at an upper limit value of $L_{X}^{qui.} \\approx (1.6-2.0) \\times 10^{27} {\\rm \\,\\,erg \\,\\,s^{-1}}$ for the quiescent X$-$ray emission of \\obj.  Dahn et al.~(\\cite{dahn}) list bolometric magnitudes for two M8.5, five M9 and three M9.5 dwarfs in the immediate solar neighborhood\\footnote{\\scriptsize $M_{bol}=13.39$ for \\object{CTI~0126+28}; $M_{bol}=13.77$ for \\object{T513-46546}; $M_{bol}=13.33$ for \\object{BRI~1222-12}; $M_{bol}=13.44$ for \\object{T868-110639}; $M_{bol}=13.47$ for \\object{LHS~2065}; $M_{bol}=13.65$ for \\object{LHS~2924}; $M_{bol}=14.22$ for \\object{LP~944-20}; $M_{bol}=13.33$ for \\object{BRI~0021-02}; $M_{bol}=13.40$ for \\object{2M~0149+29}; and $M_{bol}=13.81$ for \\object{PC~0025+04}.}. An average gives $M_{bol}=13.58\\pm0.28$ magnitudes.  This leads to the  $L_{bol} \\approx (0.9-1.5) \\times 10^{30} \\rm{erg\\,\\,s^{-1}}$. Therefore, we estimate the ratios $\\log{(L_{X}/L_{bol})} \\approx -1.1 $ at the flare peak and an upper limit of $\\log{(L_{X}/L_{bol})} \\approx -3.0 $ in the quiescent state which can be compared with measurements of other late-type dwarfs.  In Fig.~\\ref{fig1_im} we display the dependence of X$-$ray activity on spectral type for M dwarfs, using data from the RASS catalogue of nearby stars (H\\\"unsch et al. \\cite{huensch99}). To derive bolometric magnitudes we used the relation \\begin{displaymath} M_{bol} = 7.315 \\times (B-V) - 2.238, \\end{displaymath} (Mullan  \\cite{mullan}) which has a typical scatter of 0.36 mag.  We also include in this diagram the only four known examples of VLM stars with detected X$-$ray emission in the solar neighborhood. All these stars have quiescent X$-$ray fluxes which are compatible with zero. They all show pronounced X$-$ray flares, the newly discovered object \\obj\\ being the most luminous among them.  The \\ros\\  surveys of both field and cluster stars show a maximum value of about $L_{X}/L_{bol}\\simeq 10^{-3}$. Since this ratio may be regarded as an efficiency measure of the coronal heating processes and of the activity related phenomena, the maximum value indicates a still not well understood saturation phenomenon.  While the final word on the quiescent X$-$ray emission down to the bottom of the main sequence at spectral type M8/9 is not known, (see also the discussions in Fleming et al. \\cite{Fleming95,fleming03}; Mokler \\& Stelzer \\cite{mokler}; Mart\\'\\i n \\& Bouy \\cite{martin}) the flare observations of the very late M dwarfs suggest that their X$-$ray luminosity is independent of the spectral subclass.  \\begin{figure*} \\resizebox{12cm}{!}{\\includegraphics{0378.f5}} \\hfill \\parbox[b]{55mm}{ \\caption{ The X$-$ray emission level of late-type dwarfs in the solar neighborhood (see, for details text). Filled symbols indicate H$\\alpha$ emission line stars, open symbols indicate stars without reported line emission (see H\\\"unsch et al.~\\cite{huensch99}). For the ultracool dwarfs (\\object{vB10}, \\object{LHS~2065}, \\object{LP~944-20} and \\obj) the quiescent, flare and time averaged X$-$ray luminosity levels are indicated. } \\label{fig1_im} } \\end{figure*}  Any nearby ultracool dwarf, like this new one, is a promising target for direct imaging searches for sub-stellar companions, because the primary is nearby (i.e. high physical separation) and faint (i.e. no overshining companions)\\footnote{\\scriptsize See e.g. Neuh\\\"auser et al (\\cite{ralph_denis}) on \\object{DENIS-P~J104814.7$-$395606} and Scholz et al. (\\cite{scholz2003b}) on $\\epsilon$~Indi system.}."
    },
    "0311/astro-ph0311554_arXiv.txt": {
        "abstract": "{We present an algorithm developed particularly to detect gravitationally lensed arcs in clusters of galaxies. This algorithm is suited for automated surveys as well as individual arc detections. New methods are used for image smoothing and source detection. The smoothing is performed by so-called anisotropic diffusion, which maintains the shape of the arcs and does not disperse them. The algorithm is much more efficient in detecting arcs than other source finding algorithms and the detection by eye. ",
        "introduction": "Gravitational lensing has turned out to be a universal tool for very different astrophysical applications. In particular, lensing by massive clusters of galaxies is extremely useful for cosmology. The measurement of various properties of the magnified and distorted images of background galaxies (``arcs and arclets'') provides information on the cluster as well as on the population of faint and distant galaxies. Many of these distant background galaxies (up to redshifts of $z \\approx 5$ , \\cite{FRANX97}) could not be studied with the largest telescopes if they were not magnified by the gravitational lensing effect.  Some of these distant galaxies are particularly useful for the study of galaxy evolution \\cite[]{SEITZ98,PETTINI00}. As these background galaxies are free of selection effects, because they lie serendipitously behind massive clusters, they are ideal targets for a population study of distant galaxies \\cite[]{FORT97}. Gravitationally lensed arcs also provide a way to measure the total mass and the dark matter in clusters \\cite[]{FORT94, WAMBSGANSS98, Mellier99}. As galaxy clusters can be considered to be fair samples of the universal mass fractions, such determinations probe cosmological parameters like $\\Omega_{\\rm tot}$,  $\\Omega_{\\rm matter}$,  and $\\Omega_{\\rm baryon}$. A third, very important application of gravitational lensing in clusters is the determination of the frequency of arcs (=arc statistics). This is a strong criterion in order to distinguish between different cosmological models \\cite[]{BARTELMANN98, KAUFMANN00}. Therefore detections of gravitationally lensed arcs are of high importance for astrophysics and cosmology.  Ideal for cosmological studies are systematic searches. A successful arc search was performed with the X-ray luminous cluster sample of the EMSS \\cite[]{LUPPINO99}. More searches are under way, which not only cover larger areas than the previous survey, but they also go much deeper, i.e. fainter galaxies can be detected.   The first arcs were detected only in 1986 \\cite[]{LYNDS86, SOUCAIL87} because they are very faint and very thin structures. Under non-ideal observing conditions (e.g. bad seeing) they are easily dispersed and disappear into the background. Even with ideal conditions they are not easy to detect because they are often just above the background level. In order to remove the noise and make faint structures better visible usually smoothing is applied. Unfortunately in the case of such thin structures as arcs the smoothing procedure often leads to a dispersion of the few photons so that the arcs are difficult to detect at all. To prevent this dispersing we suggest an algorithm that automatically smooths only along the arcs and not perpendicular to them, so-called ``anisotropic diffusion''. The subsequently applied source finding procedure extracts all the information from the sources necessary to distinguish arcs from other sources (i.e. edge-on spirals or highly elliptical galaxies). This new algorithm is much more efficient in finding gravitationally lensed arcs than existing source detection algorithms,  because it is optimized just for this purpose.  In Sect.~2 the algorithm is explained with its four different steps. In Sect.~3 examples of detected arcs are presented. Sect.~4 outlines the differences to existing source finding software and the advantages for arc detection. In Sect.~5 we draw conclusions on the applicability and usefulness of the new algorithm. ",
        "conclusions": "We proposed a new algorithm for the detection of gravitationally lensed arcs on CCD images of galaxy clusters, performing an edge-based object detection on the filtered image together with an automatic selection of arcs.\\\\ The algorithm consists of several steps: \\begin{enumerate} \\item Histogram modification: We focus on the typical range of the arcs intensities. \\item Filtering: we use anisotropic diffusion, which enhances thin and elongated objects such as arcs. \\item Segmentation: the principal directions as determined in the diffusion process are exploited in the later segmentation process and the edge enhancing eases border allocation. \\item Selection: after detection thin and elongated objects are selected. \\end{enumerate}  Comparing our algorithm with the software package ``{\\tt SExtractor}'' we find that both algorithms rely on \\emph{multi-level} strategies for feature extraction in astronomical data. The emphasis in {\\tt SExtractor} lies on the extraction of stars and galaxies while our algorithm is designed to extract elongated arcs. Both multi-level strategies are implemented rather differently; the major differences are the filtering methods used and the determination of the segmentation threshold: {\\tt SExtractor} relies on Gaussian convolution filtering; the threshold for segmentation depends on the estimated background. Our algorithm relies on anisotropic diffusion filtering; the segmentation threshold is determined by border allocation using the same principal directions as in the diffusion process and avoiding the effect of merging close-by objects.  The new algorithm is particularly well suited for the detection of arcs in astronomical images. It can be both applied to automated surveys as well as to  individual clusters.  The algorithm (implemented in C) will be provided to the public. Feel free to contact Frank Lenzen (Frank.Lenzen@uibk.ac.at).\\\\"
    },
    "0311/astro-ph0311248_arXiv.txt": {
        "abstract": "Semi-empirical molecular dynamics is used to simulate hydrocarbon grain sputtering and collisions which are extremely difficult to study experimentally. This microscopic and dynamic approach is particularly suited to high velocity impacts, where target destruction occurs far from equilibrium. A wide variety of processes are encountered, depending on grain size and velocity: vaporization, fragmentation, atomic implantation, sticking, elastic recoil, atomic chemisorption, H abstraction and H$_{2}$ formation, etc. The impact-velocity threshold is about 10 km/s but complete grain destruction requires much higher velocities and nearly equal grain sizes. The main outcome of strong collisions is vaporization, i.e. formation of small molecules rather than solid fragments. Most of the impact energy is carried away by these molecules in the form of kinetic energy, with the consequence that the destruction efficiency of collisions is considerably reduced. ",
        "introduction": "Strong collisions between IS (InterStellar) dust grains are commonly expected to alter significantly, if not outright determine, the grain size distribution, mainly by destroying big grains thereby creating small ones, but perhaps also by creating large grains through coagulation of smaller ones (see Wickramasinghe \\cite{wic67}, Bierman and Harwitt  \\cite{bier80},  Spitzer  \\cite{spi78}). If grain destruction is predominant in the ISM (IS Medium), then, grains larger than 1000 \\AA{\\ }are expected to be almost depleted; since they are observed to be quite abundant, the far-reaching conclusion may also be drawn, that grains must also be able to grow even in such a tenuous environment (see Boulanger and Cox  \\cite{bou98}; Miville-Deschenes et al. \\cite{miv02}). Also, if collisions induce vaporization, new molecules can be created, thus contributing to the rich population of observed molecules in the ISM. Finally, even if shattering does not occur upon collision, intense heating may cause structural changes in the grains. These expectations are significant enough to have triggered a lasting interest in grain collisions for at least half a century (see Salpeter  \\cite{sal74}; Barlow and Silk \\cite{bar77}; Cowie \\cite{cow78}; Spitzer  \\cite{spi78} and Draine and Salpeter  \\cite{dra79}) and produced a large body of literature. Since grains of relevant size are very difficult to produce, handle and accelerate in the laboratory, it is not surprising that this literature is essentially theoretical in nature, although, of course, it does rely upon whatever experimental results are available.  Earlier treatments considered vaporization to be the dominant result of grain collisions, and to be complete beyond the relative velocity at which the projectile kinetic energy was equal to the total dissociation energy. Assuming the collision is totally inelastic and all the kinetic energy is available for dissociation, this velocity is independent of grain size and estimated at 3 km/s for icy dielectrics up to 8 km/s for iron particles (Spitzer  \\cite{spi78}). More recent papers on the subject have relied upon condensed-state shock theory in order to predict effects of grain collisions (see Borkowski and Dwek  \\cite{bor95}, Tielens et al. \\cite{tie94}, Seab \\cite{seab90}). This approach enhances shattering and fragmentation over vaporization, with a consequent lowering of grain destruction threshold down to relative grain velocities of about 1 km/s, and complete target disruption at 100 km/s. Admittedly, however, such treatments apply best to macroscopic bodies and are not easily extended to the grains of interest here. Moreover, little data is available on shocks in materials similar to those that are thought to be the constituents of IS grains.  The theoretical and experimental obstacles encountered in the field of small grain collisions, even outside the realm of astronomy, and the development of material and chemical computation software have prompted a flurry of numerical simulations which shed a striking new light on the subject (see Cleveland and Landeman  \\cite{cle92}; Schulte \\cite{schu95}; Aoki et al. \\cite{aok97}). Of course,because of limitations on available computer memory, this approach is best suited to relatively small systems, but these are precisely those we are interested in. Moreover, the composition and structure of the interacting systems may be tailored to suit the problem at hand. Also, various relevant microscopic phenomena can be followed quantitatively in time and space with  exquisite detail: atomic adsorption, abstraction and implantation, molecule formation, grain cratering, shattering or vaporization, etc (see Davenas \\cite{dav93}). Again, by contrast with shock theory, which subsumes thermal quasi-equilibrium,  the time scales of computations can be short enough (10$^{-16}$s) to accomodate the out-of-equilibrium phenomena that are expected to occur upon high-velocity impacts of small systems.  These potentialities are taken advantage of to revisit grain collisions in the present paper. The scope here is limited to carbonaceous material similar to the presently envisioned constituent of IS grains: partly aromatic (6- and 5-C rings), partly aliphatic (C-H chains) and more or less hydrogenated. Their sizes range from 1 to 60 atoms and their velocities from 1 to 300 km/s. The penetration range of \\emph{single} C and H atoms in this material is first determined as a function of impact velocity. This already gives some insight into the mechanism and extent of material destruction upon collision. More quantitatively, a cross-section for bond dissociation, and the corresponding threshold, are determined as a function of impact velocity.  Impacts of \\emph{clusters} of the same model material are then analyzed, which provides the average dissociation energy of the material, as well as rudimentary statistics of product sizes. These results are finally used to estimate the ranges of grain size and velocity in which strong size modifications are to be expected. Generally speaking, these estimates are less pessimistic than is usually the case for the survival of C-rich grains in the ISM. However, no discussion is offered here of the detailed effects of collisions on grain size distributions. ",
        "conclusions": ""
    },
    "0311/astro-ph0311281_arXiv.txt": {
        "abstract": "We have analyzed five X-ray outbursts of the neutron-star soft X-ray transient \\aql\\ and investigated the timing properties of the source in correlation with its spectral states as defined by different positions in the color-color and hardness-intensity diagrams. The hard color and the source count rate serve as the distinguishing parameters giving rise to three spectral states: a low-intensity hard state, an intermediate state and a high-intensity soft state. These states are respectively identified with the extreme island, island and banana states that characterize the atoll sources. The large amount of data analyzed allowed us to perform for the first time a detailed timing analysis of the extreme island state. Differences in the aperiodic variability between the rise and the decay of the X-ray outbursts are found in this state: at the same place in the color-color diagram, during the rise the source exhibits more power at low frequencies ($\\simless$ 1 Hz), whereas during the decay the source is more variable at high frequencies ($\\simmore$ 100 Hz). The very-low frequency noise that characterizes the banana-state power spectra below 1 Hz cannot be described in terms of a single power law but a two-component model is required. In two outbursts a new 6-10~Hz QPO has been discovered and tentatively identified with the normal/flaring branch-like oscillation observed only at the highest inferred mass accretion rates. We have compared the spectral and timing properties of \\aql\\ with those of other atoll and Z sources.  Our results argue against a unification scheme for these two types of neutron-star X-ray binaries. ",
        "introduction": "\\aql\\ belongs to the general group of systems known as low-mass X-ray binaries (LMXB). These systems consist of a neutron star or a black hole orbiting around a late-type star (later than A). The X rays are the result of the accretion of material from the companion onto the compact star. Mass transfer is thought to occur via Roche lobe overflow, and hence proceeds via an accretion disk. In \\aql, the compact star is a neutron star and the mass-donating companion is a V=21.6 K7V star, located at an estimated distance of 2.5 kpc (Chevalier et al. 1999).  An orbital period of 18.95 hours has been suggested (Chevalier \\& Ilovaisky 1991; Welsh, Robinson \\& Young 2000).  The high-energy radiation is characterized by 1.5--2 month long transient X-ray outbursts during which the X-ray luminosity can increase by more than three orders of magnitude. They are thought to be due to thermal instabilities in the accretion disk (e.g. van Paradijs 1996). Their recurrence time and duration vary but typical values are $\\sim$~200 days and 40--60 days, respectively ($\\check{\\rm S}$imon 2002). The source also displays Type I bursts (Zhang, Yu \\& Zhang 1998) that last a few tens of seconds and are interpreted as runaway thermonuclear burning of matter on the surface of the neutron star.   Various schemes have been proposed to categorize the LMXBs:  X-ray spectral behavior as a function of intensity and the requirement of a low-energy blackbody component in the X-ray spectra (Parsignault \\& Grindlay 1978; Naylor \\& Podsiadlowski 1993), cluster analysis of a large number of source characteristics (Ponman 1982), detailed X-ray spectral fits (White \\& Mason 1985), X-ray hardness-intensity and color-color diagrams (Schulz, Hasinger \\& Tr\\\"umper 1989), aperiodic variability at very low frequencies (Reig, Papadakis \\& Kylafis 2003).  Most relevant for the purpose of this paper is the classification scheme in terms of the rapid aperiodic variability and the patterns that these sources display in  color-color diagrams (Hasinger \\& van der Klis 1989). In this scheme LMXBs are divided into two different subclasses, Z and atoll sources.  The three spectral branches that make up the Z shape are called horizontal branch, normal branch and flaring branch and the two structures that occur in atoll sources are known as the island and the banana (van der Klis 1989). At the lowest count rates an extension of the island state is sometimes seen with a harder spectrum and stronger band-limited noise than the \"canonical\" island state (van Straaten, van der Klis \\& M\\'endez 2003 and references therein). The term {\\em extreme island state} has been used to designate such state (Prins \\& van der Klis 1997; Reig et al. 2000). The spectral and timing properties of atoll sources in this state are reminiscent of those seen in the low/hard state of black hole systems (van der Klis 1994a, 1994b; Berger \\& van der Klis 1998; Olive et al. 1998; Belloni, Psaltis \\& van der Klis 2002). Both Z and atoll sources move through these patterns continuously, that is, without jumping from one branch to another, although the source motion along the spectral tracks is much slower in atoll sources when they are in the island state than in the banana state and in Z sources in all states. The classification of Aquila X--1 in this scheme was investigated by Reig et al. (2000) --- see also Cui et al. (1998) ---, who studied the correlated X-ray timing and spectral variations of \\aql\\ and presented evidence for its classification as an atoll source, exhibiting all classic atoll source states.  This scheme has been revisited by Muno, Remillard \\& Chakrabarty (2002) and Gierli\\'nski \\& Done (2002). They reported that three transient atoll sources which display a wide dynamic range in intensity ($F_{max}/F_{min} \\simmore 100$; 4U 1608--52, 4U 1705--44 and \\aql\\ trace out three-branch patterns in the color-color diagram similar to those of Z sources, and suggested this may be a general feature. Their study was based on the spectral properties only. However, as pointed out by Hasinger \\& van der Klis (1989) it is difficult to make a clear distinction between Z and atoll sources on the basis of the motion through the color-color diagram alone. Analysis of the rapid aperiodic variability to study the noise components in different regions of the color-color diagram and the actual time scales of the motion of the source through the diagram are important to identify source type and state.  In this work we have studied the spectral and timing properties of \\aql\\ in order to investigate whether atoll and Z sources can be unified into one classification scheme and provide new insights into the poorly known extreme island states of atoll sources. We find that although the color-color diagram shows a branch whose topology is similar to the normal branch of Z sources, neither the timing properties nor the motion in the color-color diagram of \\aql\\  agree with those of Z sources. ",
        "conclusions": "We have investigated the timing and spectral properties of the soft X-ray transient \\aql\\ during five X-ray outbursts covering a period of more than five years. Three spectral states show up in the color-color diagram of \\aql: a low/hard state, identified with an extreme island state, a high/soft state corresponding to the classic banana state and an intermediate state associated with the island state of atoll sources. We have found that the hard color plays a crucial role in determining the timing properties of the extreme island state. Although the overall distribution of these states in the color-color diagram may resemble a Z, neither the motion in the CD, nor the typical time scales through the different branches of the CD, nor the range of X-ray luminosities, nor the timing properties at very low count rates (the existence of the extreme island state) are compatible with the classical Z-source behavior."
    },
    "0311/astro-ph0311138_arXiv.txt": {
        "abstract": "{We present the results of a 47-ks Chandra-ACIS observation of the old open cluster M\\,67. We detected 25 proper-motion cluster members (including ten new sources) and 12 sources (all new) that we suspect to be members from their locations close to the main sequence (1$<$$B$--$V$$<$1.7). Of the detected members, 23 are binaries. Among the new sources that are members and probable members are four spectroscopic binaries with $P_{\\rm orb}$$<$12 d, two contact binaries and two periodic photometric variables with $P_{\\rm ph}$$<$8.4 d. Their X-rays are likely the result of coronal activity enhanced by tidally locked rapid rotation.  The X-rays of the new source S\\,997, a blue straggler in a wide eccentric orbit, are puzzling.  Spectral fits show that the X-rays of the brightest sources S\\,1063 (a binary with a sub-subgiant), S\\,1082 (a triple blue straggler with a close binary) and S\\,1040 (a circular binary of a giant and a cool white dwarf), are consistent with coronal emission. We detected a new bright source that must have brightened at least about ten times since the time of the ROSAT observations. It is not clear whether its faint blue optical counterpart belongs to M\\,67. We discuss the possibility that this source is a low-mass X-ray binary in quiescence, which would be the first of its kind in an open cluster. In addition to cluster members, we detected about 100 background sources, many of which we identify with faint objects in the ESO Imaging Survey. ",
        "introduction": "\\label{intro}  The X-rays of late-type stars are believed to result from a magnetic dynamo driven by convection and rotation. As single stars spin down as they age, their X-ray emission is expected to decrease accordingly. Observations of young ($<$600 Myr) open clusters indeed show a decline of the average X-ray luminosity with age \\citep{rand97}. Consequently, X-ray observations of {\\em old} ($>$1 Gyr) open clusters are especially suited to look for special stars among the cluster members, and above all to look for interacting binaries.  The efficiency to detect binaries in old open clusters with X-ray observations is illustrated by the results of ROSAT observations of IC\\,4651, NGC\\,6940, NGC\\,752, M\\,67 and NGC\\,188 \\citep{belltagl,belltagl98,bellverb,bellea93,bellea}. In these clusters, more than half of the X-ray sources with a cluster member as optical counterpart, are identified with a binary of some sort. Mainly they are active binaries in which tidal interaction forces the stars into rotation that is synchronized to the orbit, thereby maintaining rapid rotation even at high ages. Typical X-ray luminosities of late-type active binaries in the ROSAT 0.1--2.4 keV band are 10$^{28}$--10$^{32}$ erg s$^{-1}$, where the lowest luminosities are in general found for active main-sequence stars (BY\\,Dra systems) and the highest for active giants and subgiants \\citep[RS\\,CVn binaries;][]{dempea,dempea3}. In this paper, we refer to both types as RS\\,CVn systems. The probability to detect the brightest systems in young clusters -- with higher main-sequence turnoff masses -- is smaller as more massive (sub)giants evolve more rapidly. Among the counterparts in old open clusters are also binary blue stragglers and a single, rapidly rotating FK\\,Com giant (believed to be the product of a merged binary).  Until now, ROSAT also detected one cataclysmic variable (discovered in the optical) and one hot white dwarf, but due to their optical faintness and the magnitude limits of optical studies, such sources are much more difficult to identify.  Thus, in principle, observing old open clusters in X-rays is a way to study the properties of their binary populations, characterized by the clusters' properties like age and richness. Reversely, the number and properties of the binaries in a cluster are important for the dynamical evolution of a cluster \\citep{hutea}.  X-ray observations can be useful to point out close-binary candidates among stars that are difficult to include in systematic radial-velocity searches for binaries, for example among stars on the faint part of the main sequence.  In practice, the number of old open clusters that is not too distant or too reddened to be observed in X-rays, and for which the necessary information exists to complement the X-ray data (like proper-motion, radial-velocity and photometric data), is rather limited. One of the few suitable clusters for an X-ray study is M\\,67, a 4 Gyr-old cluster that has been investigated in detail in the optical. M\\,67 is a rich cluster, is relatively nearby ($\\sim$820 pc) and nearly unreddened ($E(B-V)=0.04(2)$), with a metallicity close to solar \\citep{saraea}. Proper-motion studies to establish cluster membership have been done by e.g.~\\citet{san}, \\citet{giraea} and \\citet{zhaoea}; in this paper we use membership information from \\citet{giraea} that reaches down to $V\\approx16$.  ROSAT-PSPC observations of M\\,67 are presented by \\citet{bellea93,bellea}.  The faintest source detected in the combined ROSAT images has an X-ray flux of about 9~10$^{-15}$ erg cm$^{-2}$ s$^{-1}$ (0.1-2.4 keV), which corresponds to a luminosity of 7~10$^{29}$ erg s$^{-1}$ at a distance of 820 pc. Here, we present the results of a recent Chandra observation of M\\,67 with a limiting flux about 40 times lower.  With the increased sensitivity of this observation compared to those of the ROSAT observations, we can detect fainter sources, but single cluster stars remain out of reach. Moreover, Chandra's subarcsecond positional accuracy avoids confusion when looking for optical counterparts in the crowded field of the cluster. We describe the data reduction and source detection in Sect.~\\ref{data}, the cross\\-identification with optical and ROSAT sources in Sect.~\\ref{iden}, and the spectral and variability analyses in Sect.~\\ref{ana}. Sect.~\\ref{disc} discusses individual systems in more detail. ",
        "conclusions": "\\label{conc}  We detected 25 cluster members, including ten new sources. Most of these members are binaries, which is not unexpected as old single stars are faint in X-rays. Almost all binaries with $P_{\\rm orb}$$<$ 12 d in the ACIS field of view are detected (Fig.~\\ref{binaries}); their X-ray luminosities decrease with increasing orbital period. For detached binaries with main-sequence or subgiant primaries, this period corresponds to the period where the transition from circularized to eccentric orbits occurs in M\\,67 \\citep[$P_{\\rm cut}=12$ d, ][]{lathmathea}.  The detection of most binaries with $P_{\\rm orb} < P_{\\rm cut}$ is therefore consistent with the explanation of their X-rays: the activity is enhanced by rapid synchronous rotation. Two binaries with $P_{\\rm orb}<P_{\\rm cut}$ are not detected. For S\\,810 ($P_{\\rm orb}$=10.2 d) this could be due to the high upper limit for $L_{\\rm X}$. S\\,1284 ($P_{\\rm orb}$=4.2 d) is a binary of a late-A blue straggler while the lower limit on the secondary mass is compatible with a white dwarf and a late-type dwarf \\citep{milolath92}. The eccentric orbit ($e=0.21$) indicates that strong tidal interaction is not taking place.  Surprisingly, we detected one cluster member, S\\,1281, that shows no radial-velocity variations. Also the detection of the systems with $P_{\\rm orb} \\gg 12$ d, in which strong tidal interaction is not expected, remains a mystery. Our spectral fits indicate that the X-rays of the sub-subgiant S\\,1063 and the giant plus white dwarf S\\,1040, are coronal.  The twelve X-ray sources near the main sequence are candidates for being close binaries, but their membership and binary status should be established by spectroscopy. We also detected a group of sources with optically faint blue counterparts, including a new bright source that must have brightened at least ten times since the time of the ROSAT observations. While we expect that most of them are background sources, some could be interacting binaries in M\\,67."
    },
    "0311/astro-ph0311412_arXiv.txt": {
        "abstract": "A broad class of dark energy models, which have been proposed in attempts at solving the cosmological constant problems, predict a late time variation of the equation of state with redshift. The variation occurs as a scalar field picks up speed on its way to negative values of the potential. The negative potential energy eventually turns the expansion into contraction and the local universe undergoes a big crunch. In this paper we show that cross-correlations of the CMB anisotropy and matter distribution, in combination with other cosmological data, can be used to forecast the imminence of such cosmic doomsday. ",
        "introduction": "\\label{intro}   Observational cosmology has yielded several surprises, of which the most perplexing is the discovery of a smooth dark energy (DE) component which has come to dominate the universe at recent epochs, causing cosmic acceleration \\cite{Riess,Perlmutter}.  The nature of this component is still a matter of speculation, and a very important challenge for the coming years will be to determine its origin and physical properties. Recently, several groups \\cite{xray,nolta,fosalba,scranton,afshordi} have reported positive results for the cross-correlation between the CMB power spectrum and that of different LSS surveys, providing further evidence for the existence of DE. In this paper, we shall try and illustrate how such cross-correlation may help in unveiling some of the properties of DE, focusing on the observational signatures of a model with a time dependent DE equation of state.  The simplest interpretation of the dark energy is in terms of a cosmological constant, with equation of state $p_D=-\\rho_D$. The cosmological constant, however, raises two puzzles of its own.  First, there is a glaring discrepancy between the observed dark energy density $\\rho_D$ and the huge values of the cosmological constant suggested by particle physics models. Second, the observed $\\rho_D$ is comparable to the matter density $\\rho_m$.  This is the notorious time coincidence problem: why do we happen to live at the epoch when the dark energy starts dominating?  It has long been suggested that both puzzles may find a natural explanation through anthropic selection effects, in scenarios where $\\rho_D$ is a random variable, taking different values in different parts of the universe.  The proposed selection mechanism is very simple \\cite{Linde84,AV95,Efstathiou95,Weinberg87,MSW,GLV,Bludman}. The growth of density fluctuations leading to galaxy formation effectively stops when $\\rho_D$ comes to dominate over the matter density.  In regions where $\\rho_D$ is greater, it will dominate earlier, and thus there will be fewer galaxies (and therefore fewer observers).  A typical observer should then expect to find herself in a region where $\\rho_D$ dominates at about the epoch of galaxy formation (which is close to the present time).  Much larger values of $\\rho_D$ would yield no galaxies at all, while much smaller values are unlikely due to the smallness of the corresponding range of $\\rho_D$, assuming that all values of $\\rho_D$ are {\\it a priori} equally likely \\cite{Weinberg87,GV00}.  A simple implementation of this idea is obtained by assuming that the dark energy is due to a scalar field $\\phi$ with a very flat potential $V(\\phi)$ \\cite{Linde86,GV00}.  The values of $\\phi$ are randomized by quantum fluctuations during inflation, resulting in a variation of $\\phi$ with a characteristic scale much greater than the present Hubble radius.  Galaxy formation is possible only in regions where $V(\\phi)$ is in a narrow range near $V=0$.  One expects that the potential in this range is well approximated by a linear function \\cite{GV00,GV03,DT03}, \\begin{equation} V(\\phi)= -s\\phi, \\label{linear} \\end{equation} where $s\\equiv -V'(0)$ and we have set $\\phi=0$ at $V=0$.  The slope $s$ should be sufficiently small, so that the variation of $\\phi$ is not fast on the present Hubble scale.  Quantitatively, this can be expressed as the slow roll condition, \\begin{equation} s\\lesssim 3 H_0^2 M_p, \\label{slowroll} \\end{equation} where $M_p=1/\\sqrt{8 \\pi G}$ is the reduced Planck mass, $G$ is Newton's gravitational constant, and $H_0$ is the present Hubble expansion rate.  In models with a single DE field, and in the absence of ad-hoc adjustments, it has been argued in \\cite{GV03} that the slow roll condition (\\ref{slowroll}) is likely to be satisfied by excess, by many orders of magnitude, rather than marginally.  In this case, $\\phi$ remains nearly constant on the Hubble scale, and the effective equation of state for the dark energy is $w\\equiv p_D/\\rho_D\\approx -1$, with a very high accuracy. However, a different situation may be expected in multi-field models, where the slope of the potential $s$ is itself a random variable \\cite{mistake,GLV03} (the role of the field $\\phi$ in multi-field models is played by the variable in the direction of $\\nabla V(\\phi_a)$ in the field space.) The observed value of the slope may then be influenced by anthropic selection.  A very large slope would cause a big crunch much before any observers can develop. Hence, in cases where the prior distribution favors large $s$, the most probable values of the slope would be the ones for which the slow roll condition (\\ref{slowroll}) is only marginally satisfied \\cite{GLV03} (see also \\cite{DT03}).  We thus have some motivation to consider a model where the dark energy is due to a scalar field with a linear potential (\\ref{linear}) and a slope $s$ marginally satisfying the slow roll condition (\\ref{slowroll}). A marginal value of the slope implies that the big crunch is imminent in about 10 billion years from now.  The model can therefore be called a ``doomsday model''. A salient feature of this scenario is that the equation of state of dark energy changes significantly at low redshift, when the correlation between the large scale structure evolution and the CMB temperature anisotropies develop. Hence, we may expect that the analysis of such cross correlation may reveal a time varying equation of state $p_D = w(z) \\rho_D$, where $z$ is the redshift. This will be the subject of the present paper.  Prospective constraints on cosmic doomsday based on future determinations of the dimming of distant supernova were discussed in Ref.\\cite{KKLLS03}. The analysis shows that SNIa observations, in combination with CMB and weak lensing data, have an impressive potential for constraining the equation of state parameter. However, the constraints reported in \\cite{KKLLS03} still show a considerable degeneracy amongst models with the same ``average'' $\\langle w \\rangle$ ([see Eq. (\\ref{average}) below). Interestingly, as we shall see, the ISW-matter cross correlation breaks this degeneracy, offering the possibility of telling a true doomsday model from a model with a constant $w=\\langle w \\rangle$.  The methods presented here can obviously be used in a more general context, provided that there is significant evolution of $w$ at low redshifts. A rather common assumption in phenomenological studies of dark energy is to consider the simplest case of a constant $w$. This is partially motivated by degeneracies such as the one we just discussed above, which also occur in the angular spectrum of CMB anisotropies \\cite{caldwell} as well as in the linear matter power spectrum. The analysis of CMB/LSS cross correlations with a constant $w$ (including $w=-1$) was considered in Refs.~\\cite{turok96,peiris00,cooray,bean}. Here we shall drop this assumption, since the variation of $w$ with redshift may provide a very exciting clue to the nature of dark energy, as discussed above.   In our calculation we shall adopt a top-down approach, starting from the primordial spectrum of fluctuations. This differs from previous studies where the starting point is the present matter power spectrum (which is evolved backward in order to find its correlation with CMB). Our approach unifies the treatment of CMB and matter power spectra, and is more convenient for taking full account of fluctuations in the dark energy. Moreover, since all perturbations are evolved numerically with the CMBFAST code \\cite{cmbfast}, we do not resort to the frequently used approximate analytical expressions for the growth function, or the also commonly used small angle approximation \\cite{Limber}. The details of our calculation are reported in the Appendix.  The paper is organized as follows. In Section~\\ref{evolution} we describe the evolution of the universe according to the doomsday model. We also show that the corresponding CMB and present matter power spectra are virtually identical to those which are obtained in a model with a constant $w$, equal to the average $\\langle w \\rangle$ for the doomsday model. In Section~\\ref{temp-matter} we study the matter/CMB temperature cross-correlation, and show that it can be used to break this degeneracy. Our conclusions are summarized in Section~\\ref{summary}. ",
        "conclusions": "\\label{summary}  Current efforts to understand the dark energy component of our universe often employ dynamical scalar fields. The dynamics of such scalar fields lead to variation of the dark energy equation of state parameter, $w$. We have considered astrophysical signatures of a varying equation of state, $w(z)$, in the context of doomsday models. There are two main issues that we have considered. First, can we tell if the average value of $w$ is different from $-1$. Secondly, what astrophysical signatures are sensitive to the time variation of $w$?  To address both issues, we have considered autocorrelations of the CMB temperature anisotropies, the matter power spectrum, and cross-correlation of the anisotropy with matter fluctuations. The WMAP data is well reproduced by models with different average value of $w$ by suitably adjusting the Hubble parameter, and therefore the CMB data cannot by itself discriminate amongst different models. For a given CMB, the present linear matter power spectrum is sensitive to the average value of $w(z)$, but not to its time dependence (see Fig. \\ref{pvsk}). Finally, the temperature-matter (TM)cross-correlation is sensitive to both the average value of $w$ and to the time  variation (see Figs. \\ref{multiw}, \\ref{multicl} and \\ref{ampl}).  The TM cross-correlation requires accurate determination of both CMB anisotropies and the matter fluctuations. Ongoing CMB observations have been very successful at producing data with very small error bars and it is likely that the situation will improve even further in the coming years. Accurate surveys of the matter fluctuations are likely to be more challenging since there are unknowns such as the bias factor. However, since the CMB anisotropies and large-scale structure originate from the same density fluctuations, it is quite possible that a combination of the data can significantly reduce the uncertainties. We have shown that the optimal strategy for detecting a doomsday variation of $w(z)$ is to use a survey in the redshift range $[0.2,0.4]$.  Observations of supernovae over the next ten years are also likely to provide information on $w(z)$, as discussed in Ref. \\cite{KKLLS03}. The advantage of planned supernovae observations are that they provide direct information on the Hubble expansion rate which is closely related to $w(z)$. Furthermore, the data will be of good quality, with small error bars. However, the observations may not be too sensitive to rapid variation of $w(z)$ at recent redshifts, the hallmark of the doomsday scenario. The TM cross-correlations, on the other hand, are sensitive to this time variation and may perhaps be used together with the supernovae observations to further constrain $w(z)$."
    },
    "0311/gr-qc0311020.txt": {
        "abstract": "This paper discusses the motivation and general design elements of a new fundamental physics experiment that will test relativistic gravity at the accuracy better than the effects of the second order in the gravitational field strength, $\\propto G^2$. The Laser Astrometric Test Of Relativity (LATOR) mission uses laser interferometry between two micro-spacecraft whose lines of sight pass close by the Sun to accurately measure deflection of light in the solar gravity.  The key element of the experimental design is a redundant geometry optical truss provided by a long-baseline (100 m) multi-channel stellar optical interferometer placed on the International Space Station (ISS). The spatial interferometer is used for measuring the angles between the two spacecraft and for orbit determination purposes. In Euclidean geometry, determination of a triangle's three sides determines any angle therein; with gravity changing the optical lengths of sides passing close by the Sun and deflecting the light, the Euclidean relationships are overthrown. The geometric redundancy enables LATOR to measure the departure from Euclidean geometry caused by the solar gravity field to a very high accuracy.  LATOR will not only improve the value of the parameterized post-Newtonian (PPN) parameter $\\gamma$ to unprecedented levels of accuracy of 1 part in 10$^{8}$, it will also reach ability to measure effects of the next post-Newtonian order ($\\propto G^2$) of light deflection resulting from gravity's intrinsic non-linearity.  The solar quadrupole moment parameter, $J_2$, will be measured with high precision, as well as a variety of other relativistic effects including Lense-Thirring precession.  LATOR will lead to very robust advances in the tests of Fundamental physics: this mission could discover a violation or extension of general relativity, or reveal the presence of an additional long range interaction in the physical law.  There are no analogs to the LATOR experiment; it is unique and is a natural culmination of solar system gravity experiments. ",
        "introduction": "Einstein's general theory of relativity (GR) began with its empirical success in 1915 by explaining the anomalous perihelion precession of Mercury's orbit, using no adjustable theoretical parameters.  Shortly thereafter, Eddington's 1919 observations of star lines-of-sight during a solar eclipse confirmed the doubling of the deflection angles predicted by GR as compared to Newtonian-like and Equivalence Principle arguments.  From these beginnings, GR has been verified at ever higher accuracy. Thus, microwave ranging to the Viking Lander on Mars yielded accuracy  $\\sim$0.1\\% in the tests of GR \\cite{viking_shapiro1,viking_reasen,viking_shapiro2}. The astrometric observations of quasars on the solar background performed with Very-Long Baseline Interferometry (VLBI) improved the accuracy of the tests of gravity to $\\sim$0.03\\% \\cite{eubanks97,RoberstonCarter91,Lebach95}. Lunar Laser Ranging (LLR),  a continuing legacy of the Apollo program, provided $\\sim$0.01\\% verification of GR via precision measurements of the lunar orbit \\cite{Ken_LLR68,Ken_LLR91,Ken_LLR98,Ken_LLR30years99,Ken_LLRgravitomagnetic88,Ken_LLR_PPNprobe03,JimSkipJean96,Williams_etal_2001}. Finally, the recent experiments with the Cassini spacecraft improved the accuracy of the tests to $\\sim$0.0023\\% \\cite{iess_etal_1999,cassini_ber, cassini_and}. As a result, by now not only the ``non-relativistic,'' Newtonian regime of GR is well understood, but also the first ``post-Newtonian'' approximation is also well-studied, making GR the standard theory of gravity when astrometry and spacecraft navigation are concerned.  However, the continued inability to merge gravity with quantum mechanics, and recent cosmological observations indicate that the pure tensor gravity of GR needs modification.  Progress in scalar-tensor extensions of gravity which are consistent with present cosmological models \\cite{damour_nordtvedt1,damour_nordtvedt2,DPV02,DPV02b} motivate new searches for very small deviations of relativistic gravity in the solar system, at levels of 10$^{-5}$ to 10$^{-7}$ of the post-Newtonian effects or essentially to achieve accuracy that is compatible to the size of the effects of the second order in the gravitational field strength ($\\propto G^2$).  This will require a several order of magnitude improvement in experimental precision from present tests. The ability to measure the first order light deflection term at the accuracy comparable with the effects of the second order is of the utmost importance for the gravitational theory and is the challenge for the 21st century fundamental physics.  %************************************************** \\begin{table}[t!] \\caption{Comparable sizes of various light deflection effects in the solar gravity field. The value of deflection angle calculated on the limb of the Sun ($b=R_\\odot$); the corresponding delay is given for a 100 m interferometric baseline proposed for LATOR. \\label{tab:eff}} \\vskip 5pt \\begin{tabular}{|c|c|c|c|} \\hline & & Deflection &\\\\[-3pt] Effect  & Analytical Form      & angle, $\\mu$as  & Delay, pm \\\\ \\hline \\hline &&&\\\\[-12pt] First  Order  & $2(1+\\gamma )\\frac{M}{b}$ & $1.75\\times 10^6$   & $0.849 $~mm     \\\\ &&&\\\\[-12pt] \\hline &&&\\\\[-12pt] Second Order   & $\\{[2(1+  \\gamma)-\\beta+\\frac{3}{4} \\delta]\\pi- 2(1+\\gamma)^2\\}\\frac{M^2}{b^2} $ &  3.5 & 1697      \\\\ &&&\\\\[-12pt] \\hline &&&\\\\[-12pt] Frame-Dragging   & $\\pm 2(1+\\gamma)\\frac{J}{b^2}$ & $\\  \\pm0.7$   &     $\\pm339$ \\\\ &&&\\\\[-12pt] \\hline &&&\\\\[-12pt] Solar Quadrupole  & $2(1+\\gamma )J_2\\frac{M}{b^3}$ & 0.2   &   97  \\\\[4pt] \\hline \\end{tabular} \\end{table} %********************End Table******************  When the light deflection in solar gravity is concerned, the magnitude of the first order effect as predicted by GR for the light ray just grazing the limb of the Sun is $\\sim1.75$ arcsecond (asec) (for more details see Table \\ref{tab:eff}). (Note that 1 arcsec $\\simeq5~\\mu$rad; when convenient, below we will use the units of radians and arcseconds interchangeably.) The effect varies inversely with the impact parameter. The second order term is almost six orders of magnitude smaller resulting in  $\\sim 3.5$ microarcseconds ($\\mu$as) light deflection effect, and which falls off inversely as the square of the light ray's impact parameter \\cite{EpsteinShapiro80,FishbachFreeman80,RichterMatzner82_1,RichterMatzner82_2,Ken_2PPN_87}. The relativistic frame-dragging term\\footnote{Gravitomagnetic frame dragging is the effect in which both the orientation and trajectory of objects in orbit around a body are altered by the gravity of the body's rotation.  It was studied by Lense and Thirring in 1918.} is $\\pm 0.7 ~\\mu$as, and contribution of the solar quadrupole moment, $J_2$, is sized as 0.2 $\\mu$as (using theoretical value of the solar quadrupole moment $J_2\\simeq10^{-7}$ ). The small magnitudes of the effects emphasize the fact that, among the four forces of nature, gravitation is the weakest interaction; it acts at very long distances and controls the large-scale structure of the universe, thus, making the precision tests of gravity a very challenging task.  The LATOR  mission will directly address the challenges discussed above. The test will be performed in the solar gravity field using optical interferometry between two micro-spacecraft.  Precise measurements of the angular position of the spacecraft will be made using a fiber coupled multi-chanelled optical interferometer on the ISS with a 100 m baseline. The primary objective of LATOR will be to measure the gravitational deflection of light by the solar gravity to accuracy of 0.1 picoradians (prad) ($\\sim0.02 ~\\mu$as), which corresponds to $\\sim$10 picometers (pm) on a 100 m interferometric baseline.  In conjunction with laser ranging among the spacecraft and the ISS, LATOR will allow measurements of the gravitational deflection by a factor of more than 3,000 better than had recently been accomplished with the Cassini spacecraft. In particular, this mission will not only measure the key PPN parameter $\\gamma$ to unprecedented levels of accuracy of one part in 10$^8$, it will also reach ability to measure the next post-Newtonian order ($\\propto G^2$) of light deflection resulting from gravity's intrinsic non-linearity. As a result, this experiment will measure values of other PPN parameters such as parameter $\\delta$ to 2 part in $10^3$ (never measured before; see Eq.~(\\ref{eq:metric}) and discussion thereafter), the solar quadrupole moment parameter $J_2$ to 1 part in 20, and the frame dragging effects on light due to the solar angular momentum to precision of 1 parts in $10^2$.  The LATOR mission technologically is a very sound concept; all technologies that are needed for its success have been already demonstrated as a part of the JPL's Space Interferometry Mission (SIM) development. (In fact, accuracy of 5 pm was already demonstrated in our SIM-related studies.)  The LATOR concept arose from several developments at NASA and JPL that initially enabled optical astrometry and metrology, and also led to developing expertize needed for the precision gravity experiments. Technology that has become available in the last several years such as low cost microspacecraft, medium power highly efficient solid state and fiber lasers, and the development of long range interferometric techniques make possible an unprecedented factor of 3,000 improvement in this test of GR possible. This mission is unique and is the natural next step in solar system gravity experiments which fully exploits modern technologies.  LATOR will lead to very robust advances in the tests of fundamental physics:  this mission could discover a violation or extension of GR, or reveal the presence of an additional long range interaction in the physical law.  With this mission testing theory to several orders of magnitude higher precision, finding a violation of GR or discovering a new long range interaction could be one of this era's primary steps forward in fundamental physics. There are no analogs to the LATOR experiment; it is unique and is a natural culmination of solar system gravity experiments.  This paper organized as follows: Section \\ref{sec:sci_mot} discusses the theoretical framework and science motivation for the precision gravity tests that recently became available.  Section \\ref{sec:lator_description} provides an overview for the LATOR experiment including the current mission design.  Section \\ref{sec:error_bud} discusses a preliminary  error budget. In Section \\ref{sec:conc} we discuss the next steps that will taken in the development of the LATOR mission. ",
        "conclusions": "\\label{sec:conc}  LATOR mission is the 21st century version of Michelson-Morley experiment searching for a  cosmologically evolved scalar field in the solar system. This mission aims to carry out a test of the curvature of the solar system's gravity  field with an accuracy better than 1 part in 10$^{8}$. In spite of the previous space missions exploiting radio waves for tracking the spacecraft, this mission manifests an actual breakthrough in the relativistic gravity experiments as it allows to take full advantage of the optical techniques that recently became available.  LATOR has a number of advantages over techniques that use radio waves to measure gravitational light deflection. Thus, optical technologies allow low bandwidth telecommunications with the LATOR spacecraft. The use of the monochromatic light enables the observation of the spacecraft at the limb of the Sun. The use of narrowband filters, coronagraph optics and heterodyne detection will suppress background light to a level where the solar background is no longer the dominant noise source. The short wavelength allows much more efficient links with smaller apertures, thereby eliminating the need for a deployable antenna. Finally, the use of the ISS enables the test above the Earth's atmosphere -- the major source of astrometric noise for any ground based interferometer. This fact justifies LATOR as a space mission.  The LATOR experiment technologically is a very sound concept; all technologies that are needed for its success have been already demonstrated as a part of the JPL's interferometry program. The LATOR experiment does not need a drag-free system, but uses a geometric redundant optical truss to achieve a very precise determination of the interplanetary distances between the two micro-spacecraft and a beacon station on the ISS. The interest of the approach is to take advantage of the existing space-qualified optical technologies leading to an outstanding performance in a reasonable mission development time.  The  availability of the ISS makes this mission concept realizable in the very near future; the current mission concept calls for a launch as early as in 2011 at a cost of a NASA MIDEX mission.  Our next steps will be to perform studies of the optimal trajectory configuration and mission design including the launch vehicle choice trade studies. Our analysis will concentrate on the thermal design of the instrument, analysis of the launch options and configuration, estimates of on-board power and weight requirements as well as analysis of optics and vibration contamination for the interferometer.  We also plan to  develop an end-to-end mission simulation, including detailed mission error budget.  LATOR will lead to very robust advances in the tests of fundamental physics: it could discover a violation or extension of GR, or reveal the presence of an additional long range interaction in the physical law.  There are no analogs to the LATOR experiment; it is unique and is a natural culmination of solar system gravity experiments.  The work described here was carried out at the Jet Propulsion Laboratory, California Institute of Technology, under a contract with the National Aeronautics and Space Administration.  %**************************************"
    },
    "0311/astro-ph0311624_arXiv.txt": {
        "abstract": "{ We present observations of six isolated, pre-main-sequence, intermediate mass stars selected for shallow spectra at submillimeter wavelengths at 1.3, 2.6, 7.0, and 36 millimeters from the IRAM PdBI and the VLA. We analyze the new observations of these stars (HD34282, HD35187, HD142666, HD143006, HD150193, HD163296) together with similar observations of three additional stars from the literature (CQ Tau, UX Ori, TW Hya), in the context of self-consistent irradiated disk models. Our aim is to constrain the wavelength dependence of the dust opacity and the total dust mass in the disks. The shallow wavelength dependence of the opacity is confirmed and for a few stars extended to significantly longer wavelengths. For any plausible dust properties, this requires grain growth from interstellar sizes to maximum sizes of at least a few millimeters, and very likely to several centimeters or more. For four of the stars (HD34282, HD163296, CQ Tau, TW Hya), the millimeter emission has been spatially resolved, and the large disk radii ($>100$ AU) rule out that  high optical depths play a role. The mass of dust that has been processed into large grains is substantial, and in some cases implies a disk mass comparable to the mass of the central star. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311142_arXiv.txt": {
        "abstract": "We study the X-ray emission in a sample of galaxy clusters using the BeppoSAX PDS instrument in the 20 -- 80 keV energy band. We estimate the non-thermal hard X-ray cluster emission (HXR) by modeling the thermal contribution from the cluster gas and the non-thermal contamination from the unobscured AGN in the clusters. We also evaluate the systematic uncertainties due to the background fluctuations. Assuming negligible contamination from the obscured AGN, the resulting non-thermal component is detected at a 2$\\sigma$ level in $\\sim$50\\% of the non-significantly AGN-contaminated clusters: A2142, A2199, A2256, A3376, Coma, Ophiuchus and Virgo. The data are consistent with a scenario whereby relaxed clusters have no  hard X-ray component of non-thermal origin, whereas merger clusters do, with a 20 -- 80 keV luminosity of $\\sim 10^{43-44}$ h$_{50}^{-2}$ erg s$^{-1}$. The co-added spectrum of the above clusters indicates a power-law spectrum for the HXR with a photon index of 2.8$^{+0.3}_{-0.4}$ in the 12 -- 115 keV band, and we find indication that it has extended distribution. These indications argue against significant contamination from obscured AGN, which have harder spectra and centrally concentrated distribution. These results are supportive of the assumption of the merger shock acceleration of electrons in clusters, which has been proposed as a possible origin of the non-thermal hard X-ray emission models. Assuming that the Cosmic Microwave Background photons experience Inverse Compton scattering from the merger-accelerated relativistic electrons, and thus produce the observed HXR, the measured hard X-ray slope corresponds to a differential momentum spectra of the relativistic electrons with a slope of $\\mu$ = 3.8 -- 5.0. In presence of cluster magnetic fields this relativistic electron population produces synchrotron emission with a spectral index of 1.4 -- 2.1, consistent with radio halo observations of merger clusters. Thus both hard X-ray and radio observations of merger clusters are consistent with the Inverse Compton model. The observed slope of HXR is also consistent with that predicted by the non-thermal bremsstrahlung, which thus cannot be ruled by the fit to the current data, even though this model requires an extreme, untenable cluster energetics. Assuming centrally concentrated distribution of HXR, the data requires a harder slope for the HXR spectrum, which is consistent with secondary electron models, but this model yields a worse fit to the PDS data and thus seems to be disfavored over the primary electron Inverse Compton model. ",
        "introduction": "Non-thermal hard X-ray (HXR) emission has recently been observed in several clusters and groups of galaxies with the MECS and PDS instruments onboard BeppoSAX and with GIS instrument onboard ASCA. In the case of Coma (Fusco-Femiano et al. 1999), A2256 (Fusco-Femiano et al. 2000), HCG62 (Fukazawa et al. 2001) and A754 (Fusco-Femiano et al. 2003) detection of excess emission above the contribution from the hot ICM is statistically significant, while marginal evidence is provided for A3667 (Fusco-Femiano et al. 2001) and A2199 (Kaastra et al. 1999).   Most models of the HXR emission involve a population of electrons accelerated in the cluster medium. A natural source of acceleration in clusters is provided by merger shocks. In a strong cluster merger event, the electrons are accelerated to relativistic speeds (e.g. Bell, 1978a,b; Fujita \\& Sarazin 2001; Takizawa \\& Naito 2000). The Inverse Compton scattering of Cosmic Microwave Background photons from the relativistic electrons in clusters then can produce a non-thermal tail which exceeds the thermal bremsstrahlung emission at energies above 20 keV (e.g. Sarazin 1999, Blasi \\& Colafrancesco 1999). This model has been proposed as the simplest possible explanation of the HXR properties of galaxy clusters. If the acceleration is provided by a less energetic merger or turbulence (e.g. Ensslin et al. 1999), or if there is a high-energy cutoff in the electron velocity distribution, the resulting electron population is effectively transrelativistic. In this case, the dominating mechanism in producing hard X-rays has been proposed to be non-thermal bremsstrahlung (e.g. Sarazin \\& Kempner 2000). However, this solution faces some crucial problems mainly concerning the large energy injection required by such a mechanism and the resulting large heating expected (e.g., Petrosian 2001).  In the secondary electron population models the merger shocks and galaxy activity accelerate and inject large quantities of relativistic protons into the cluster atmosphere. Most of the relativistic protons can be confined and accumulated in the cluster medium for very long times, comparable with the cluster age $\\sim H_0^{-1}$, and can then produce secondary electrons via proton-proton collisions (e.g., Colafrancesco and Blasi 1998). The energy losses are balanced by the continuous refilling of the new electrons produced in situ (i.e. continuously in time and everywhere in space). Thus the resulting HXR spectrum in the secondary models reflects the electron spectrum right after the acceleration event, while in the primary models the more energetic electrons loose energy rapidly and the HXR spectrum steepens accordingly with time.  In the present work, we expand the database of cluster hard X-ray emission by studying a sample of clusters observed with the BeppoSAX PDS. We model the thermal and AGN contributions in the sample in order to obtain estimates for the non-thermal component. We propagate the modeling uncertainties, as well as the background fluctuation uncertainties in order to obtain reliable (and somewhat conservative) estimates for the non-thermal component. We furthermore study the co-added non-thermal hard X-ray spectrum of the sample, in order to investigate the origin of this emission.  We consider uncertainties and significances at 1 $\\sigma$ level, and use H = 50 $\\times$ h$_{50}$ km s$^{-1}$ Mpc$^{-1}$, unless stated otherwise. We define HXR in this paper as 20 -- 80 keV PDS net count rate, after removing the sky background, cluster thermal component and AGN contamination. ",
        "conclusions": "We have studied the hard X-ray emission in 20 -- 80 keV energy band in a sample of clusters using the BeppoSAX PDS instrument. After removing the contributions from the cluster thermal component and from unobscured AGN, in $\\sim$50\\% of the mildly AGN-contaminated clusters the non-thermal component is detected at 2 $\\sigma$ level, the clusters being  A2142, A2199, A2256, A3376, Coma, Ophiuchus and Virgo. All the clusters detected at 2$\\sigma$ level exhibit some degree of merger signatures, i.e. deviations from the azimuthally symmetric brightness and temperature distributions (except for A2199). Averaging the PDS 20 -- 80 keV count rates of the relaxed and merger clusters obtains a 2.5$\\sigma$ detection for the merger group, while the relaxed group count rate is consistent with zero. Assuming a power-law emission model with a photon index of 2.0 at the group average redshifts, the average count rates are consistent with a scenario whereby the relaxed clusters have no HXR component, while mergers do, with a 20 -- 80 keV luminosity of $\\sim 10^{43-44}$ h$_{50}^{-2}$ erg s$^{-1}$.  The co-added spectrum of our sample yields a best-fit photon index of 2.8$^{+0.3}_{-0.4}$ for the non-thermal emission in 12 -- 115 keV band, and we find indication that it has extended distribution. These indications argue against significant contamination from obscured AGN, which have harder spectra and centrally concentrated distribution.  The indicated connection between cluster mergers and the non-thermal hard X-ray emission is consistent with the Inverse Compton scattering of the Cosmic Microwave Background photons with merger--accelerated population of relativistic electrons. In this framework, the observed photon index is consistent with a scenario in which a strong acceleration event and consequent strong IC/CMB energy losses take place. In this scenario the measured hard X-ray slope corresponds to a differential momentum spectra of the relativistic electrons with a slope of $\\mu$ = 3.8 -- 5.0. The consequent synchrotron emission spectrum expected from the same electron population has a spectral index of 1.4 -- 2.1, consistent with radio halo observations of many merger clusters.  The observed slope of the HXR spectrum is also consistent with the predictions of the non-thermal bremsstrahlung model. Even though this fit cannot be ruled out by the current data, the bremsstrahlung model seems to face a strong energetics problem which does not make it a viable physical scenario.  Assuming that most of the non-thermal signal originates in the central regions of clusters, the HXR spectrum is forced to be harder, with a slope $\\sim$ 1.3--1.5, which turns out to be consistent with secondary electron models. However, this model provides a worse fit to PDS data and is thus disfavored by the statistical fit over the primary electron IC/CMB model.  In conclusion, spatially resolved hard X-ray spectroscopy is needed to disentangle between primary and secondary electron models for non-thermal hard X-ray emission in clusters of galaxies."
    },
    "0311/astro-ph0311468_arXiv.txt": {
        "abstract": "The  state of  the art  of optical studies  of  Isolated Neutron Stars (INSs) and their Pulsar Wind Nebulae (PWNe)  is reviewed. In addition, results  obtained   from recent  \\hst~   and  \\vlt~  observations  are presented and discussed. ",
        "introduction": "The Crab  pulsar (PSR B0531+21)  was the  first isolated  neutron star (INS) identified at  optical  wavelengths  (Cocke, Disney,  \\&  Taylor 1968).  Almost a decade went by till another INS, the Vela pulsar (PSR B0833-45), was detected  (Lasker 1976) and  confirmed  (Wallace et al. 1977).  The Vela pulsar, among the faintest objects known at that time ($V \\sim 23.6$), was also the last INS detected by photograhic plates. A  decade  later,  came  the first    CCD  detection of    the optical counterpart of another INS, Geminga  (Bignami et al.  1987).   Spurred by this  result and by  the  advent of  the new generation  telescopes (e.g.  the ESO NTT), optical observations of INSs were carried on with revived enthusiasm   and at the  beginning of  the 90s  yielded to the detections of the counterparts of PSR B0540-69 (Caraveo et al.  1992), readily confirmed by Shearer et al. (1994), PSR B0656+14 (Caraveo,   Bignami,  \\& Mereghetti 1994a)   PSR B1509-58 (Caraveo, Mereghetti,  \\& Bignami 1994b),  this  one later  revised by Wagner \\& Seifert (2000).  In the meantime, the identification of Geminga was confirmed through the proper  motion measurement of its counterpart  (Bignami, Caraveo,  \\& Mereghetti  1993).   The launch of \\hst~ in 1990 increased the chances of detecting new INSs by providing access to the near-UV window  through the short-wavelength sensitivity of the Faint Object Camera (FOC).  Indeed, soon came the detections of PSR B0950+08, PSR  B1929+10 (Pavlov, Stringfellow,  \\& Cord\\'ova 1996) and PSR B1055-52 (Mignani,  Caraveo, \\& Bignami 1997).  Likely optical counterparts were also detected for  some of the  INSs singled out  in the ROSAT data (a.k.a.  X-ray Dim INSs - XDINSs), namely RXJ 1856-3754 (Walter \\& Matthews 1997) and RXJ 0720-3125  (Kulkarni \\& van Kerkwijk 1998).  Thus, the number of INSs  observed in the optical/UV increased by a factor four in just a decade. ",
        "conclusions": "INSs with either a secured or a likely optical counterpart amounts now to 14, i.e., a  number comparable to those detected  in X-rays  in the pre-ROSAT  era.  Thus, the  optical  branch, originally  confined to a handful of representative cases,  is growing in importance occupying a larger  and larger  niche   in the  multiwavelength  studies  of INSs. Certainly, much work remains  to be done to reduce   the gap, both  in terms of quality  and  quantity, with the   X-ray  domain.  More  deep imaging observations are needed  to pinpoint new candidates, to  study correlations between the  emission  properties (luminosities, spectra) as a function  of the pulsars' parameters and   to search for  diffuse emission structures (PWNe and  bow-shocks).  More  timing observations are  required, both to secure faster  identifications and to provide a broader characterization  of the  pulsars' lightcurves.  Finally, more spectroscopy   is critical to better  define   the emission processes. These are tasks both for the current 10 m class telescopes and for the next generation  of extra-large telescopes   (Becker \\& Mignani, these proceedings)."
    },
    "0311/astro-ph0311004_arXiv.txt": {
        "abstract": "Kilometer-scale neutrino detectors such as IceCube are discovery instruments covering nuclear and particle physics, cosmology and astronomy. Examples of their multidisciplinary missions include the search for the particle nature of dark matter and for additional small dimensions of space. In the end, their conceptual design is very much anchored to the observational fact that Nature accelerates protons and photons to energies in excess of $10^{20}$ and $10^{13}$\\,eV, respectively. The cosmic ray connection sets the scale of cosmic neutrino fluxes. In this context, we discuss the first results of the completed AMANDA detector and the reach of its extension, IceCube. ",
        "introduction": "The flux of cosmic rays is summarized in Fig.\\,1a,b\\cite{gaisseramsterdam}. The  energy spectrum follows a broken power law. The two power laws are separated by a feature referred to as the ``knee\"; see Fig.\\,1a. There is evidence that cosmic rays, up to several EeV, originate in galactic sources. This correlation disappears in the vicinity of a second feature in the spectrum dubbed the ``ankle\". Above the ankle, the gyroradius of a proton exceeds the size of the galaxy and it is generally assumed that we are  witnessing the onset of an extragalactic component in the spectrum that extends to energies beyond 100\\,EeV. Experiments indicate that the highest energy cosmic rays are predominantly protons. Above a threshold of 50 EeV these protons interact with CMBR photons and therefore lose their energy to pions before reaching our detectors. This limits their sources to tens of Mpc, the so-called Greissen-Zatsepin-Kuzmin cutoff.   \\begin{figure}[ht] \\centering\\leavevmode \\includegraphics[width=5.5in]{halzen_fig1.eps} \\caption{At the energies of interest here, the cosmic ray spectrum consists of a sequence of 3 power laws. The first two are separated by the ``knee\" (left panel), the second and third by the ``ankle\". There is evidence that the cosmic rays beyond the ankle are a new population of particles produced in extragalactic sources; see right panel.} \\end{figure}   Models for the origin of the highest energy cosmic rays fall into two categories, top-down and bottom-up. In top-down models it is assumed that the cosmic rays are the decay products of cosmological remnants with Grand Unified energy scale $M_{GUT} \\sim 10^{24}\\rm\\,eV$. These models predict neutrino fluxes most likely within reach of AMANDA, and certainly IceCube.  In bottom-up scenarios it is assumed that cosmic rays originate in cosmic accelerators. Accelerating particles to TeV energy and above requires massive bulk flows of relativistic charged particles. These are likely to originate from the exceptional gravitational forces  in the vicinity of black holes. Examples include the dense cores of exploding stars, inflows onto supermassive black holes at the centers of active galaxies and annihilating black holes or neutron stars. Before leaving the source, accelerated particles pass through intense radiation fields or dense clouds of gas surrounding the black hole. This results in interactions producing pions decaying into secondary photons and neutrinos that accompany the primary cosmic ray beam as illustrated in Fig.\\,2.   \\begin{figure}[ht] \\centering\\leavevmode \\includegraphics[width=5in]{halzen_fig2.eps} \\caption{Diagram of cosmic ray accelerator producing photons and neutrinos.} \\end{figure}   How many neutrinos are produced in association with the cosmic ray beam? The answer to this question, among many others\\cite{PR},  provides the rationale for building kilometer-scale neutrino detectors. We first consider a neutrino beam produced at an accelerator laboratory; see Fig.\\,2. Here the target absorbs all parent protons as well as the muons, electrons and gamma rays produced. A pure neutrino beam exits the dump. If nature constructed such a ``hidden source\" in the heavens, conventional astronomy will not reveal it. It cannot be the source of the cosmic rays, however, because the dump would have to be partially transparent to protons. The extreme opposite case is a ``transparent source\" where the accelerated proton interacts once and escapes the dump after producing photons as well as neutrinos. Elementary particle physics is now sufficient to relate all particle fluxes because a fraction (1/6 to 1/2 depending on the energy) of the interacting proton goes into pion production. This energy is equally shared between gamma rays and neutrinos, of which one half are muon-neutrinos. Therefore, at most one quarter of the energy ends up in muon-neutrinos compared to cosmic rays. The flux of a transparent cosmic ray source is often referred to as the Waxman-Bahcall flux\\cite{wb1}. It is easy to derive and the derivation is revealing.  Fig.\\,1b shows a fit to the observed cosmic ray spectrum assuming an extragalactic component fitted above the ``ankle\". The energy content of this component is $\\sim 3 \\times 10^{-19}\\rm\\,erg\\ cm^{-3}$, assuming an $E^{-2}$ energy spectrum with a GZK cutoff. The power required to generate this energy density in the Hubble time of $10^{10}$\\,years is $\\sim 3 \\times 10^{37}\\rm\\,erg\\ s^{-1}$ per (Mpc)$^3$. This works out to\\cite{TKG} \\begin{itemize} \\item \\ $\\sim 3 \\times 10^{39}\\rm\\,erg\\ s^{-1}$ per galaxy, \\item \\ $\\sim 3 \\times 10^{42}\\rm\\,erg\\ s^{-1}$ per cluster of galaxies, \\item \\ $\\sim 2 \\times 10^{44}\\rm\\,erg\\ s^{-1}$ per active galaxy, or \\item  \\ $\\sim 2 \\times 10^{52}$\\,erg per cosmological gamma ray burst. \\end{itemize} The coincidence between these numbers and the observed electromagnetic energy output of these sources explains why they have emerged as the leading candidates for the cosmic ray accelerators. The coincidence is consistent with the relationship between cosmic rays and photons built into the ``transparent\" source previously introduced. The relationship can be extended to neutrinos.  Assuming the same energy density  of $\\rho_E \\sim 3 \\times 10^{-19}\\rm\\,erg\\ cm^{-3}$ in neutrinos with a spectrum $E_\\nu dN / dE_{\\nu}  \\sim E^{-\\gamma}\\rm\\, cm^{-2}\\, s^{-1}\\, sr^{-1}$ that continues up to a maximum energy $E_{max}$, the neutrino flux follows from $ \\int E_\\nu dN / dE_{\\nu}  =  c \\rho_E / 4\\pi  $. For $\\gamma = 1$ and $E_{max} = 10^8$\\,GeV, the generic source of the highest energy cosmic rays produces 50 detected muon neutrinos per km$^2$ per year\\cite{TKG}. [Here we have folded the predicted flux with the probability that the neutrino is actually detected given by\\cite{PR} the ratio of the muon and neutrino interaction lengths in ice, $\\lambda_\\mu / \\lambda_\\nu$.] The number depends weakly on $E_{max}$ and $\\gamma$. A similar analysis can be performed for galactic sources\\cite{PR}.  As previously stated, for one interaction and only one, the neutrino flux should be reduced by a factor $\\sim 4$. On the other hand,  there are more cosmic rays in the universe producing neutrinos than observed at earth because of the GZK-effect. The diffuse muon neutrino flux associated with the highest energy cosmic rays is estimated to be $ {E_\\nu}^2 dN / dE_{\\nu}  \\sim 5 \\times 10^{-8}\\rm\\, GeV \\,cm^{-2}\\, s^{-1}\\, sr^{-1} $, to be compared to the sensitivity achieved with the first 3 years of the completed AMANDA detector of $10^{-7}\\rm\\,GeV\\ cm^{-2}\\,s^{-1}\\,sr^{-1}$. The analysis has not been completed but a limit has been published that is 5 times larger obtained with data taken with the partially deployed detector in 1997\\cite{b10-diffuse}. On the other hand, after three years of operation IceCube will reach a diffuse flux limit of $E_{\\nu}^2 dN / dE_{\\nu} = 8.1 \\times 10^{-9}\\rm\\,GeV \\,cm^{-2}\\, s^{-1}\\, sr^{-1}$. ",
        "conclusions": ""
    },
    "0311/astro-ph0311232_arXiv.txt": {
        "abstract": "{We have determined the dynamical mass of the most luminous stellar cluster known to date, i.e. object W3 in the merger remnant galaxy NGC~7252. The dynamical mass is estimated from the velocity dispersion measured with the high-resolution spectrograph UVES on VLT.  Our result is the astonishingly high velocity dispersion of $\\sigma=45 \\pm 5~{\\rm km/s}$. Combined with the large cluster size $R_{\\rm eff}=17.5 \\pm 1.8~{\\rm pc}$, this translates into a dynamical virial mass for W3 of ~$ (8 \\pm 2) \\; 10^7 \\msun$. This mass is in excellent agreement with the value ($\\sim 7.2\\cdot10^7~\\msun$) we previously estimated from the cluster luminosity ($M_{\\rm V}=-16.2$) by means of stellar $M/L$~ratios predicted by Simple Stellar Population models (with a Salpeter IMF) and confirms the heavy-weight nature of this object. This results points out that the NGC~7252-type of mergers are able to form stellar systems with masses up to $\\sim~10^8~\\msun$. We find that W3, when evolved to $\\sim$~10 Gyr, lies far from the typical Milky Way globular clusters, but appears to be also separated from $\\omega$Cen in the Milky Way and G1 in M31, the most massive old stellar clusters of the Local Group, because it is too extended for a given mass, and from dwarf elliptical galaxies because it is much more compact for its mass. Instead the aged W3 is amazingly close to the compact objects named ultracompact dwarf galaxies (UCDGs) found in the Fornax cluster (Hilker et al.~1999; Drinkwater et al.~2000), and to a miniature version of the compact elliptical M32.  These objects start populating a previously deserted region of the fundamental plane. ",
        "introduction": "About a decade ago, HST observations (Holtzmann et al.~1992) confirmed the early ground-based discovery (Schweizer~1982) that mergers and merger remnant galaxies host luminous, compact objects. These are believed to be candidate young globular clusters (GCs) on the basis of their blue colours, high luminosities, and compactness as derived from the small radii. They were predicted to form during the (violent) interaction caused by the merger event (Schweizer~1987). The properties of these star clusters, such as their number, masses, ages and metallicities have attracted a number of studies (see review by Schweizer~1998) since they impact on the success of the scenario in which elliptical galaxies form via merging of spirals. Leaving aside pros and cons of this view, comprehensively discussed in Kissler-Patig~(2000), the formation of star clusters during mergers is nowadays an accepted fact. Besides helping in exploring the connection to the host galaxy, the properties of young stellar clusters aid in understanding the poorly known process of star cluster formation. Among the various properties, the present paper focuses on the cluster mass. A striking feature of the star clusters detected in galactic mergers and merger remnants are the high stellar masses reached by some of the members, as derived using $M/L$~ratios from stellar population models. For example, in the ``Antennae'' galaxies, the most luminous objects have masses up to a few $10^{6}~\\msun$ (e.g. Zhang \\& Fall~1999). Such {\\it luminous} masses are larger than the {\\it mean} mass of old globular clusters in the Milky Way ($M\\sim 2\\cdot 10^{5} \\msun$, e.g. Harris~1991), but are comparable to the mass of the most massive galactic GCs, like 47 Tuc ($\\sim 1.3\\cdot 10^{6}~\\msun$, Meylan \\& Mayor~1986) especially considering that (some) cluster mass will be lost. The percentage of mass loss due to stellar evolution between 30 Myr and 15 Gyr is only $\\sim~20 \\%$~(Maraston~1998), but dynamical processes such as evaporation and tidal disruption are also expected to reduce masses of GCs. However the amount is difficult to estimate, since it depends on the exact orbit of the cluster, its number of particles and the potential of the parent galaxy. Recent N-body simulations for potentials appropriate for spiral galaxies (Baumgardt \\& Makino~2003) indicate values from $\\sim ~50 \\%$~of dynamical mass-loss, to a negligible amount of dynamical mass loss when the cluster is populous enough ($N > 10^6$). For the potential of an early-type system similar predictions are not yet available.  This work focuses on the most extreme case of a superluminous star cluster known to date, i.e. object W3 in NGC~7252. Its luminosity-derived mass is in the range $\\sim10^{7}-10^{8} \\msun$~(Schweizer \\& Seitzer~1998; Maraston et al.~2001). These values are more than ten times larger than those of the star clusters in the ``Antennae'', in spite of W3 being significantly older ($t\\sim 300 - 500~{\\rm Myr}$, Schweizer \\& Seitzer~1998, Maraston et al.~2001). However, the mass determination via population synthesis models is affected by uncertainties in the age determination. As stellar $M/L$~ratios increase with age due to the fading of the light, an overestimation of the cluster age leads to an overestimation of its mass (other effects like metallicity and initial mass function playing a minor r\\^ole).  Given the potential impact of such high masses on the formation of stellar systems, in particular in violent interactions and extreme environments, the light-derived masses have to be checked by comparing them with the values determined dynamically. The first succesful attempt in this direction was by Ho \\& Filippenko~(1996) for the young star clusters of NGC~1569 (see Section~5), and more recently by Mengel et al.~(2002) for those in the ``Antennae''. Here we have employed the spectrograph UVES on VLT to obtain a high-resolution optical spectrum of W3 in order to measure its velocity dispersion, and hence its virial mass. This paper reports on the (spectacular) results for W3. We will devote a follow-up paper to present the results for the other star clusters measured during the same observing program. ",
        "conclusions": "\\subsection{Mass}  \\label{mass} Our main result is the perfect agreement between the dynamical mass determined here for W3, and the luminous mass as derived for it from model $M/L$~ratios (Maraston et al.~2001).  Recalling those results, a luminous mass $7.2\\times10^{7}~\\msun$~was obtained for a Salpeter IMF with lower and upper mass cutoffs of 0.1 and 100~\\msun, respectively; $4.\\times10^{7}~\\msun$~for a Gould IMF (Gould, Bahcall \\& Flynn~1997), i.e. an IMF the slope of which flattens with respect to Salpeter at the low mass-end ($M\\leq0.6~\\msun$). The mass derived using a Salpeter IMF is amazingly close to the dynamical estimate. In case one prefers the Gould-type IMF, roughly 45 \\%~of the total mass of W3 inside the half-mass radius should be dark, but it is clearly not possible to discriminate among the two options. It should be noted that the same agreement is found for the luminous masses as derived from the near-IR M/L~(Maraston et al.~2001).  The modeling of W3 as a simple stellar population of 300 Myr, $0.5~\\zsun$~appears very self-consistent from every side, colours, Balmer lines (Schweizer \\& Seitzer~1998), $M/L$. This favours the idea that W3 is a globular cluster, since all but one globular clusters in the Milky Way are simple stellar populations. However, as initially pointed out by Schweizer \\& Seitzer~(1998), the mass of the evolved W3 will still be $>100$~times larger than the mass of typical galactic GCs, having taken into account the stellar mass losses. Note that a flat IMF all over the mass range, as sometime advocated in the literature (e.g. Mengel et al.~2002) does not help in reducing significantly the total mass of the aged W3. Quantitatively, the decrease in the stellar mass between 300 Myr and 10 Gyr for a flat IMF with exponent 1.5 (in the notation in which the Salpeter's one is 2.35) is only a factor 1.64 larger than in case of Salpeter or Gould IMF (Maraston~1998).  Dynamical mass-loss, e.g. tidal stripping and evaporation, act in removing mass. However evaporation is not important for such a massive object, and we estimated that it can remove only 1\\% of mass (see Maraston et al. 2001). Tidal stripping might be more efficient, but as stated in Section~1, its impact is difficult to assess for our object since current N-body simulations (Baumgardt \\& Makino~2003) refer to clusters in spiral-like gravitational potential for the parent galaxy, while NGC~7252 has a de Vaucouleurs profile (Schweizer~1982). However, according to these simulations, the amount of dynamical mass-loss decreases strongly with an increasing number of particles, therefore it is probably negligible for W3 because of its very large number of stars (cfr. Section 4). \\begin{figure*} \\centering \\includegraphics[width=0.7\\textwidth]{AA4077Fig3.ps} \\caption {The $\\kappa$~space of stellar systems, in a remake of Fig.~11 of Burstein et al.~(1997). To easy the reading, the physical parameters mass, $M/L$~and surface brightness are also reported, according to the definition of $\\kappa$'s (Bender, Burstein \\& Faber~1992). Lines define average sequences for: B+E (bulges and ellipticals); dE (dwarf ellipticals); GC (globular clusters). The open circle highlights the position of the compact elliptical M32. For the sources of these data see Burstein et al.~(1997). The quantities refer to $H_{0}=75~{\\rm km/sec}$. The present position of W3 is displayed by means of the open star, while its predicted position at an age of 10 Gyr is indicated by the arrow. Crosses show G1 and $\\omega$Cen, the most massive stellar clusters known in the Local Group. The filled triangle indicates the very compact objects detected in the Fornax cluster (Hilker et al.~1999; Drinkwater et al.~2000) named ultra compact dwarf galaxies (UCDGs) (sources of the data of the latter three objects are given in the text).}  \\label{kappafig} \\end{figure*}  The very large mass and flattened structure of W3 (cf. Table~2) cast some doubt as to its nature as a GC. Indeed, in the Milky Way the maximum value of $\\sigma$~is $\\sim 18~{\\rm km/sec}$~for the old metal-rich GCs NGC~6441 and NGC 6388 (Pryor \\& Meylan~1993). The Magellanic Clouds GCs, many of which have ages/metallicities like W3, have normal luminosities ($\\sim~10^{4 - 5}~\\lsun$) and, when measured, normal $\\sigma$'s~(for the SMC GC NGC~419 Dubath et al.~(1997) give $\\sigma \\sim~9.5\\pm0.3~{\\rm km/s}$). The young (some tenths of Myr) star clusters of the nearby dwarf galaxy NGC~1569 have dynamical masses $\\sim 3\\cdot 10^{5}~\\msun$, and will evolve in perfectly normal Galactic-type GCs. Instead, the mass of some of the star clusters in the ongoing merger NGC~4038/4039 (``the Antennae''), although being coeval to those of NGC~1569, are at least ten times larger (see Section~1), as well as that of a very young ($\\sim~15~{\\rm Myr}$) star cluster in the nearby spiral NGC~6946 (Larsen et al.~(2001).  Worth noting is also the case of the star clusters in the 3 Gyr old merger remnant NGC~1316 in Fornax. As discussed by Goudfrooij et al.~(2001), the luminosity-derived mass of the brightest star cluster, being already 3 Gyr old is still $\\sim~1.4\\cdot 10^{7}~\\msun$. It would be interesting to verify this number with dynamical measurements.  It seems natural to conclude that the formation of abnormally massive star clusters is very much favoured by extreme environments like galaxy interactions/mergers.  With a mass of $\\sim~10^8~\\msun$~W3 is two orders of magnitudes more massive than any other GC with reliable dynamical mass estimates, and its classification does not appear straightforward. \\subsection{Clues from the $\\kappa$~space} \\label{kappa} In order to put the properties of W3 in context among other types of stellar systems, the most comprehensive way is to consider the fundamental plane of dynamically hot systems (Dressler et al.~1987; Djorgovski et al.~1987) in which the basic structural parameters: effective radius, mean surface brightness inside the effective radius, and central velocity dispersion, are considered {\\it simultaneously}. Particularly useful is the re-definition of the fundamental plane known as {$\\kappa$}-space (Bender, Burstein \\& Faber~1992) that combines the three variables mentioned above into more physically meaningful ones. The new variables called $\\kappa_1$, $\\kappa_2$ and $\\kappa_3$, are proportional to the mass, the product of $M/L$~and surface brightness (i.e. the compactness of the system) and the $M/L$, respectively.  The {$\\kappa$}-space is shown in Figure~\\ref{kappafig}, in a fashion similar to Fig.~11 of Burstein et al.~(1997). The sequences defined by the stellar systems: B+E (bulges plus ellipticals); dE (dwarf ellipticals); GC (globular clusters) have been drawn from their work.  The $\\kappa$-coordinates of the compact dwarf elliptical M32 (open circle) are from Bender et al. (1992).  In order to place W3 onto the $\\kappa$-space, the central velocity dispersion $\\sigma_{0}$~has been evaluated from its average value determined in this work, by means of the relation $\\sigma_0 = <\\sigma>/0.87=51.72~\\rm km/sec$~(derived from Djorgovski et al.~1997), where it has been assumed that all the light out to $3-5 R_{\\rm eff}$~fell into the UVES aperture. The arrow indicates the position of W3 when 10 Gyr old.  This implies a fading of 3.66 mag in B, according to the evolutionary synthesis models adopted here (Maraston~1998; Maraston et al.~2001).  The most massive star clusters (with dynamical mass measurements) known before W3 are G1 in M31, the most luminous stellar cluster of the Local Group, and $\\omega$Centauri in the Milky Way. As the aged W3 could be more related to these heavy-weight objects than to typical GCs, we have also plotted them into the $\\kappa$-space. For G1 a $\\sigma_{0}\\sim 27.8~{\\rm km/sec}$~is provided by Meylan et al.~(2001), and from their dynamical model we estimate $R_{\\rm eff}\\sim 4.47~{\\rm pc}$. For $\\omega$Cen $\\sigma_{0}\\sim 16~{\\rm km/sec}$~(from Pryor \\& Meylan 1993), and $R_{\\rm eff}\\sim 4.85~{\\rm pc}$~is from van den Bergh et al.~(1991).  It should be noted that both G1 and $\\omega$Cen have been argued to be the remnant nuclei of stripped galaxies (Meylan et al.~2001; Freeman 1993), because they show a metallicity spread, in other words they both are not simple stellar populations as all other galactic GCs.  The filled triangle in Figure~2 shows the position of the objects named Ultra Compact Dwarf galaxies (UCDGs) detected recently in a deep spectroscopic survey in Fornax (Hilker et al.~1999; Drinkwater et al.~2000). The nature of these objects is currently under debate, and as discussed by Phillipps et al.~(2001) they could be either extremely large star clusters or extremely compact small galaxies, perhaps the nuclei of stripped dE (see Drinkwater et al.~2003). This ambiguity pushes to a comparison with our object. To place the UCDGs in the $\\kappa$-space we evaluate as average magnitude $B\\sim~19.95$~using objects n. 1,2,4,5 from Table~3 of Karik et al.~(2003) that show consistent properties, i.e. we excluded object n.~3 that is significantly brighter. Mean values of $R_{\\rm eff}\\sim 20.75~\\rm pc$~and $\\sigma_{0}\\sim~27.54~\\rm km/sec$~are obtained by averaging the individual values of the 4 objects (M.~Hilker, {\\it private communications}, see also Drinkwater et al.~2003). The spectra of the UCDGs are consistent with old stellar populations (Hilker et al.~1999; Drinkwater et al.~2003), therefore they should be compared to W3 aged to 10 Gyr.  The following conclusions emerge from Figure~\\ref{kappafig}. It is confirmed that W3 lies far from the GC sequence because it is not compact enough given its mass compared to the GC sequence. This remains true when W3 is compared with G1 and $\\omega$Cen. W3 is also far from the position of dwarf ellipticals because it is far more concentrated at a mass matching the least massive dEs. Instead, the aged W3 shows properties that are remarkably similar to those of the mysterious UCDGs. Using the similarity in the opposite direction appears safer, since we presume to know how W3 has been formed, i.e. during the merger event. This suggests that the UCDGs might have formed in similar violent galaxy interactions at high redshift, since they are found to be old.  It should be noted however that while these objects are currently found in a galaxy cluster, NGC~7252 is rather isolated. This implies that the cluster environment is not an exclusive one for forming ultracompact massive objects. It would be interesting to search for these type of objects in the vicinity of other relatively isolated elliptical galaxies. Additionally, we note that other gas rich galaxy mergers, like the Antennae, seem not to have produced objects as massive as W3.  Finally, we briefly comment on the relation between W3 and the compact elliptical M32 (empty circle). Similar to W3, M32 has properties that place it in the realm of rare objects (see e.g. Bender, Burstein \\& Faber~1992), since for its mass is much more compact that the elliptical galaxies defining the sequence in Figure~2. In this context it is worth noting that, contrary to the compact objects of Fornax, M32 is the only close companion of the giant spiral M31. Similarly, in 10 Gyrs W3 will be the only visible close companion of NGC~7252.  It is unfortunate that the most straightforward difference between GCs and dwarf galaxies, i.e. the {\\it simple} (i.e. single, mono-metallic burst) versus the {\\it complex} (i.e. extended star formation) nature of their stellar populations is not assessable in case of distant objects, like W3 (and the UCDGs). In any case, we have learned that the galaxy interaction that has originated NGC~7252 roughly 1 Gyr ago was also capable to form a stellar system whose structural properties fall into a domain of rare objects. These objects at a given mass are more compact than dwarf galaxies, and at given compactness are more massive than GCs, and start populating a previously empty region of the fundamental plane.  It would be valuable to explore under which conditions current simulations of NGC~7252-type mergers (e.g. Mihos \\& Hernquist~1996) can produce W3-like objects."
    },
    "0311/astro-ph0311226_arXiv.txt": {
        "abstract": "The relationship between the star formation rate and surface density of neutral gas within the disk of M33 is examined with new imaging observations of \\co\\ J=1-0 emission gathered with the \\fcrao\\ 14m telescope and \\IRAS\\ \\HIRES\\ images of the 60\\micron\\ and 100\\micron\\ emission. The  Schmidt law, $\\Sigma_{SFR} \\sim \\Sigma_{gas}^n$, is constructed using radial profiles of the HI 21cm, CO, and far infrared emission. A strong correlation is identified between the star formation rate and molecular gas surface density. This suggests that the condensation of giant molecular clouds is the limiting step to star formation within the M33 disk. The corresponding molecular Schmidt index, $n_{mol}$, is 1.36$\\pm$0.08. The star formation rate has a steep dependence on total mass gas surface density, ($\\Sigma_{HI}+\\Sigma_{H_2}$), owing to the shallow radial profile of the atomic gas which dominates the total gas surface density for most radii. The disk pressure of the gas is shown to play a prominent role in regulating the molecular gas fraction in M33. ",
        "introduction": "The understanding of galaxy structure and evolution is ultimately linked to the understanding of the star formation process.  Star formation within galaxies is commonly described by two relationships which are oversimplifications of complex interstellar processes taking place over a large range of spatial scales (Elmegreen 2002). The first is an empirical relationship between the star formation rate per unit area, $\\Sigma_{SFR}$, and gas mass surface density, $\\Sigma_{gas}=\\Sigma_{HI}+\\Sigma_{H_2}$, $$ \\Sigma_{SFR} \\propto \\Sigma_{gas}^n  \\eqno(1) $$ often referred to as the Schmidt law (Schmidt 1959; Kennicutt 1989; Wong \\& Blitz 2002). The value of the Schmidt index, $n$, and its variation with galaxy type and environment provide valuable clues to the physical processes that regulate star formation in galaxies. The second relationship simply states that there is a critical gas surface density, $\\Sigma_c$, below which star formation is suppressed (Kennicutt 1989; Martin \\& Kennicutt 2001). Since stars form within the cold, dense, molecular gas phase of the interstellar medium, the critical surface density may simply reflect the necessary local conditions at which molecules are sufficiently self-shielded from dissociating UV radiation. However, several studies have shown that the critical surface density is dependent upon the galaxy rotation curve (Kennicutt 1989; Martin \\& Kennicutt 2001) and relate $\\Sigma_c$ to the surface density, $\\Sigma_{Q}$, which parameterizes the gravitational instability associated with the Coriolis force (Toomre 1964), $$\\Sigma_{Q}={ {c\\kappa}\\over{\\pi G} } \\eqno(2) $$ where $c$ is the sound speed and $\\kappa$ is the epicyclic frequency. Martin \\& Kennicutt (2001) demonstrate that $\\Sigma_{Q}$ provides an accurate gauge of the gas column density at which star formation is inhibited for most of the 32 galaxies in their sample.  Two galaxies in the sample of Martin \\& Kennicutt (2001) do not conform to the Toomre condition.  M33 and NGC 2403 are small, low luminosity, low mass galaxies.  Both systems are actively forming stars throughout their disks yet are sub-critical ($\\Sigma_{gas} << \\Sigma_{Q}$) over all radii (Kennicutt 1989; Martin \\& Kennicutt 2001). For both galaxies, this assessment is based upon limited imaging of the molecular gas component within the central region (Wilson \\& Scoville 1989; Thornley \\& Wilson 1995) or the major axis (Young \\etal 1995). Using wide field imagery of CO emission presented in this study, Corbelli (2003) has confirmed the sub-critical nature of the M33 disk.  Several studies have proposed that star formation in galaxies like M33, NGC 2403, and irregular galaxies with slowly rising rotation curves, may be regulated by gas instabilities modulated by the local shear rate (Hunter, Elmegreen, \\& Baker 1998; Martin\\& Kennicutt 2001; Corbelli 2003). The critical surface density may also be modulated by the additional gravity from either the stellar disk or the dark matter halo (Wang \\& Silk 1994; Hunter, Elmegreen, \\& Baker 1998; Corbelli 2003).  M33 is a nearby galaxy which shows no signs of interaction with other galaxies.  Therefore, it constitutes an excellent laboratory to evaluate the degree to which physical processes regulate star formation in disk galaxies. In this paper we present the first, unbiased, filled aperture census of the molecular gas component in M33 from an imaging survey of CO emission with the FCRAO 14m telescope.  A recent high resolution interferometric CO image of M33 is presented by Engargiola \\etal (2003). The FCRAO data have been previously used to derive the radial distributions of molecular hydrogen, the rotation curve, and the gravitationally unstable region in the disk of M33 (Corbelli 2003). Here, we examine the relationship between the star formation rate and gas surface density as parameterized by equation (1). It is important to assess whether the star formation properties (Schmidt Law, star formation efficiency) in sub-critical galaxies like M33 are modified due to different modes and processes which enable the condensation of cold, dense clouds from the diffuse ISM. Moreover, as a low luminosity, molecular poor galaxy, M33 provides an important complement to the sample of CO-bright spiral galaxies used by Wong \\& Blitz (2002) to investigate the relationship between gas content and star formation.  In $\\S$2, the CO data collection, calibration, and processing procedures are described. The spatial distribution of the molecular gas is presented in $\\S$3.  In $\\S$4, we derive the Schmidt law from HI, CO, and IRAS HiRes images and evaluate the role of pressure in regulating the molecular gas fraction. ",
        "conclusions": "In this paper we have presented an image of  $^{12}$CO J=1-0 emission from M33 and investigated the relationship between the star formation rate and the molecular and atomic neutral gas components. From our analyses, we derive the following conclusions. \\begin{enumerate} \\item The mass surface density of molecular gas is strongly correlated with the star formation rate per unit area even in regions that are dominated by atomic gas.  This proportionality suggests that the formation of molecular clouds is the limiting step to the formation of stars within the disk of M33. \\item The Schmidt indices for the total and molecular gas surface densities are $n=3.3\\pm0.1$ and $n_{mol}=1.36\\pm0.08$ respectively.  The large value for $n$ is a consequence of the variation of the star formation rate within the molecular gas substrate, ($n_{mol}$), and the condensation of molecular clouds from the atomic gas component ($f_{mol}$). \\item The pressure of the gaseous disk layer due to the weight of both gas and stars can account for the radial variations of the molecular gas fraction. \\item  The distributions of molecular gas and current star formation activity are extended with respect to the old, stellar population and suggests an outward progression of star formation within the M33 disk.  \\end{enumerate}"
    },
    "0311/astro-ph0311010_arXiv.txt": {
        "abstract": "We  present  an analysis  of  the  asymmetries  in the  population  of Galactic  M-giant  stars  present  in  the 2MASS  All  Sky  catalogue. Several large-scale asymmetries are  detected, the most significant of which is a strong elliptical-shaped stellar over-density, close to the Galactic plane at $(\\ell=240^\\circ, b=-8^\\circ)$, in the constellation of  Canis Major.   A small  grouping of  globular  clusters (NGC~1851, NGC~1904, NGC~2298,  and NGC~2808), coincident in  position and radial velocity, surround  this structure, as  do a number of  open clusters. The population  of M-giant  stars in this  over-density is  similar in number to that in the core  of the Sagittarius dwarf galaxy.  We argue that  this  object  is  the  likely dwarf  galaxy  progenitor  of  the ring-like structure  that has recently been  found at the  edge of the Galactic  disk.  A  numerical  study  of the  tidal  disruption of  an accreted dwarf  galaxy is presented.   The simulated debris  fits well the  extant  position,  distance   and  velocity  information  on  the ``Galactic  Ring'', as  well as  that of  the  M-giant over-densities, suggesting that all  these structures are the consequence  of a single accretion event.   The disrupted dwarf  galaxy stream orbits  close to the  Galactic Plane,  with  a pericentre  at  approximately the  Solar circle,  an orbital  eccentricity  similar  to that  of  stars in  the Galactic thick  disk, as  well as a  vertical scale height  similar to that of the thick disk.  This finding strongly suggests that the Canis Major dwarf  galaxy is  a building block  of the Galactic  thick disk, that the  thick disk  is continually growing,  even up to  the present time, and  that thick  disk globular clusters  were accreted  onto the Milky Way from dwarf galaxies in co-planar orbits. ",
        "introduction": "It  is   now  generally  accepted  that  galaxies   assembled  by  the hierarchical  merging of dark  matter halos  \\citep{white78, white91}, with the accretion of luminous (dwarf) galaxies playing a significant, or  probably  even a  major,  role in  the  formation  of the  visible structures.   In   the  Milky   Way,  this  picture   is  dramatically illustrated  by  the existence  of  the  Sagittarius dwarf  spheroidal galaxy, which  is in  an advanced state  of tidal disruption,  and has populated the  outer halo with  M giant stars  \\citep{ibata02}, carbon stars \\citep{ibata01a} and globular clusters \\citep{bellazzini03a}.  The accretion of  the Sagittarius galaxy is currently  the only strong evidence that the Milky Way is absorbing satellite galaxies.  However, the ring-like  structure that has recently been  discovered around the Galaxy  could  be  the  consequence  of  another  such  event.   First discovered as an  over-density in blue stars of  the Sloan Digital Sky Survey  \\citep{newberg}, this  ``Ring''  has since  been probed  using photometry  from   the  INT  Wide   Field  Camera  (\\citealt{ibata03}, hereafter  I03), SDSS  spectrometry  and photometry  (\\citealt{yanny}, hereafter Y03) and 2MASS M-giants (\\citealt{rochapinto}).  Surrounding the   Galactic   disk,  its   Galactocentric   distance  ranges   from $\\sim15\\kpc$ to  $\\sim20\\kpc$ in fields  taken within $30\\deg$  of the Galactic plane and for $122\\deg < \\ell < 227\\deg$.  It has  been proposed (I03  and Y03) that  the structure could  be the tidal stream  stripped away  by the Milky  Way tides from  a satellite galaxy that has an orbital plane  close to the plane of the Galaxy. In light  of this  interpretation,  it is  interesting  to reexamine  the \\citet{abadi} simulation  of the formation of a  disk galaxy assembled hierarchically in  a $\\Lambda{}$CDM  cosmology.  They showed  that the thin and thick disks of  their simulated galaxy, which were similar to those of the  Milky Way, do not share the  same origin.  The simulated thick  disk component  was  constructed mainly  through accretions  of satellites  with  an  orbit   close  to  the  Galactic  plane.   These satellites  are dense enough  for their  orbits to  circularize before they are disrupted.  \\citet{helmi03} also  used this simulation to follow  more closely the behavior of  an in-plane accretion.  They concluded  that it naturally leads  to the  formation of  a ring-like  structure around  the galaxy similar  to  the  one  found  by \\citet{newberg},  with  two  possible explanations. One possibility is that the ring is a tidal arc produced by  stars stripped  away from  the  parent satellite  during a  recent pericentric passage. This would  produce an asymmetric structure above and  below  the  disk,  limited  in  Galactic  longitude  and  with  a significant    velocity   gradient    in   the    Galactic   longitude direction. Alternatively, it could be a shell-like structure, produced by ancient minor mergers, similar to the shells observed in elliptical galaxies.  In this  case, the ring would be  rather symmetric and with no velocity gradient.  To discriminate between these  competing scenarios, it is essential to analyze a large-scale, homogeneous data  set.  Until now, this was the main limitation, but with the release  of the 2 Micron All Sky Survey, All Sky Release  (2MASS ASR) catalogue we now  have an invaluable tool to  determine the  characteristics of  this Galactic  ``Ring''.  Apart from  covering  all the  sky,  the  2MASS  ASR contains  precise  near infrared photometry, so  it is not much hampered  by extinction in the low Galactic  latitude regions with high  dust-contamination where the ``Ring'' is primarily located.  We use this  catalogue to probe the ring-like  structure and determine its dimensions and characteristics on the celestial hemisphere opposed to the Galactic centre. In section~2,  we discuss the data set we used for our  study, presented in  section~3. In section~4, we  present the results  of  our  search  of  clusters  in  the  ring-like  structure. Section~5 presents N-body simulations of the tidal disruption of the progenitor  of the  observed  structures. Finally,  in  section 6,  we discuss  our  results  and  the  consequences on  the  origin  of  the ``Ring''. Throughout this work, we 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{dehnen98b}). ",
        "conclusions": "We have shown  that the distribution of M-giant  stars at low latitude is highly  asymmetric.  The  most striking over-density  is due  to an elliptical object present  in the disk, in the  direction of the Canis Major   constellation   in  the   southern   Galactic  hemisphere   at $D_{GC}=12.0\\pm1.2\\kpc$. This structure contains as many M-giant stars as  the  remnant  of  the  Sagittarius dwarf  galaxy.   A  less  dense structure, but of huge proportions, the Northern Arc curves around the Galaxy in the  direction of the Galactic anticentre.   This arc ranges from the Galactic plane at $\\ell  \\sim 230\\deg$ to $b \\sim +30\\deg$ at $\\ell  \\sim 140\\deg$.   Between $140\\deg<\\ell<180\\deg$  it has  a mean displacement of $\\sim20\\deg$ above the Galactic Plane, with a distance of  $D_{GC}=18.1\\pm0.1\\kpc$, it  is  separated from  the  disk in  the radial direction along  the line of sight, and  has a Gaussian profile of width $11.6\\deg\\pm1.9\\deg$.  Several  other   structures  are  also  detected.    In  the  southern hemisphere,  the  ``Southern   Arc''  population  with  Galactocentric distance  $12.0\\kpc<D_{GC}<15.1\\kpc$ is  observed  between $110\\deg  < \\ell  <  210\\deg$;  this  may  be the  continuation  of  the  southern hemisphere  structure  reported  by  Y03 for  $R_{GC}=20\\pm2\\kpc$  and $180\\deg  < \\ell <  227\\deg$.  Two  other over-densities  are detected towards the  Galactic bulge, and close  to the Galactic  Plane: one is found between  $230\\deg < \\Phi < 270\\deg$  at $D_{GC}<14.5\\kpc$, while the other is between $90\\deg < \\Phi < 130\\deg$ at $D_{GC}=10.5\\kpc$.  Given these observations, we examine the idea of an in-plane accretion of a dwarf satellite galaxy  onto the Milky Way.  The realization that the four globular clusters NGC~1851, NGC~1904, NGC~2298, and NGC~2808, are close to each other in  phase space and are also close in position and distance to the Canis Major over-density also lends strong support to this possibility.  In this scenario the remnant of the satellite is likely the object  in the direction of Canis  Major.  During the tidal disruption  process, the  dwarf galaxy  loses stars  in a  long stream which loops  around the  Milky Way  and is now  seen as  the Monoceros ``Ring'' structure.  The two other  hypotheses presented by I03 do not seem to  be allowed  by our observations.   Indeed, the height  of the structure (generally more than $3\\kpc$) could hardly be explained by a spiral  arm.    Furthermore,  the  multiple  peaks   in  the  distance distribution  are more suggestive  of multiple  crossings of  a stream rather  than  of  a  ring-like  Galactic structure,  produced  by  the dynamical evolution of ancient stellar warps.  The simple N-body  models we presented reproduce many  of the observed spatial  features, including  the Canis  Major over-density,  the vast ring-like  arc  towards the  Galactic  anticentre, the  over-densities towards  $\\ell=70\\deg$ and  $\\ell=290\\deg$,  as well  as faint  debris scattered  out to  large  distance  below the  Galactic  Plane in  the anticentre  direction.   The   models  also  reproduce  the  available kinematic  data  and  point  to  the  possibility  that  several  star clusters,  apart from  those close  to the  Canis  Major over-density, could  be related  to the  Canis Major  Stream.  The  success  of this simple  model  strongly  suggests  that  many,  if  not  all,  of  the asymmetries in  the Galactic  M-giant population presented  above, are due to a single accretion  event. The accreted satellite must have had an orbital plane closely aligned to the plane of the Milky Way galaxy. Our analysis  shows that the ``shell'' model  of \\citet{helmi03} gives poor agreement with the observations,  but we confirm the tidal stream model for the ``Monoceros Ring'' proposed by Y03.  The observations and  model also suggest a connection  to the Galactic thick disk.  The orbital eccentricity of  the model is $e = (r_{apo} - r_{peri})/(r_{apo} + r_{peri}) \\sim 0.5$,  a value that is typical for Galactic  thick disk  stars of  metallicity ${\\rm  [Fe/H]  \\sim -1.2}$ \\citep{chiba}.  Furthermore, the  vertical exponential scale height of the debris $0.73\\pm0.05\\kpc$, is very close to the thick disk value of $0.8\\kpc$ determined  by \\citet{reyle}. It is  therefore possible that we have uncovered  one of the galactic building  blocks from which the thick  disk grows.   This analysis  suggests  that the  thick disk  is continually growing,  even up to the present  time, through accretions of dwarf galaxies in co-planar orbits. Although with a progenitor mass of $\\sim 10^9\\msun$,  as suggested by our analysis,  only a handful of accretions of  the magnitude of that  of the Canis  Major dwarf galaxy are  needed to populate  the thick  disk entirely.   Globular clusters with disk-like kinematics may also  have been brought to their present locations  in this  way. The  observations and  modeling  we presented strongly  support the  analysis  of \\citet{abadi},  in  which a  large proportion  of  thick  stars  have  their origin  in  accreted  galaxy fragments.  If the  scenario we have developed here  is correct, thick disks should  be a generic  component of giant spiral  galaxies, since accretions are a generic  feature of galaxy formation. Furthermore, at the  edge of the  disk where  the dynamical  times are  longest, their structure will  likely be a  complex mess where streaming  debris from the  most  recent  accretions  piles  up at  apocentre.   These  messy structures  may indeed  already have  been detected  in  the Andromeda galaxy \\citep{ibata01b,ferguson02}.  Further observational  work is needed  to constrain the  kinematics of the  structures  we report  here.   These  constraints  are needed  to determine  whether the structure  follows a  prograde or  a retrograde orbit, to verify  the kinematic predictions of our  models, as well as to  guide future simulations.   Detailed abundance  distributions will also be invaluable  in order to study the  chemical difference between an accreted  population, such as that  of the Canis  Major galaxy, and that  of the  normal Galactic  disk.  Future  simulations with  a live Milky   Way  will  provide   a  more   realistic  comparison   to  the observations. It  will be interesting  to investigate the  reaction of the Galactic disk stars and in  particular the reaction of the HI gas, to  this sizable  accretion, while  the back-reaction  onto  the Canis Major dwarf galaxy  will be essential in order  to model its dynamical evolution accurately."
    },
    "0311/astro-ph0311360_arXiv.txt": {
        "abstract": "Accretion onto compact objects plays a central role in high energy astrophysics. In these environments, both general relativistic and plasma effects may have significant impacts upon the spectral and polarimetric properties of the accretion flow.  In paper I we presented a fully general relativistic magnetoionic theory, capable of tracing rays in the geometric optics approximation through a magnetised plasma in the vicinity of a compact object.  In this paper we discuss how to perform polarised radiative transfer along these rays.  In addition we apply the formalism to a barotropic thick disk model, appropriate for low luminosity active galactic nuclei.  We find that it is possible to generate large fractional polarisations over the innermost portions of the accretion flow, even when the emission mechanism is unpolarised. This has implications for accreting systems ranging from pulsars and X-ray binaries to AGN. ",
        "introduction": "\\label{I} The spectral and polarimetric properties of astrophysical objects can provide significant insights into their structure and dynamics.  As a result, a number of theoretical investigations into the source of these properties have been undertaken.  Many of these have been primarily concerned with the spectral properties alone, typically comparing a physically motivated accretion flow to observations. However, with the measurement of polarisation in a number of sources, a significant fraction of the focus has been turned towards reproducing their polarimetric properties.  In the context of an accreting compact object, both general relativistic and plasma effects can play a role in determining these properties.  In \\citealp{Brod-Blan:03b} (hereafter Paper I) we demonstrated how to construct ray trajectories, in the geometric optics approximation, in a magnetoactive plasma in a relativistic environment.  In order to apply this to realistic accretion environments it is necessary to be able to perform radiative transfer along these rays.  Non-refractive, polarised radiative transfer through magnetised plasmas is flat space has been extensively studied.  A number of examples involving weak magnetic fields exist in the literature \\citep[see \\eg][]{Sazo-Tsyt:68,Sazo:69,Jone-ODel:77b,Jone-ODel:77a,Ginz:70}. More recently, investigations into the net effects of tangled magnetic fields (expected to be typical in magnetised accretion flows) have begun \\citep[see \\eg][]{Rusz-Bege:02}.  However, none of these deal with general relativistic environments.  The importance of refraction in the propagation of radio wavelengths has long been appreciated in the context of the ionosphere \\citep[see \\eg][]{Budd:64,Ginz:70}.  More recently, refraction has been studied in conjunction with pulsars \\citep[see \\eg][]{Welt-Stap-Horn-Edwa:03,Petr:02,Petr:00, Barn-Aron:86a,Aron-Barn:86b}.  Nonetheless, in all of these cases, the emission was assumed to originate from a region distinct from where the refraction occurred.  Refractive lensing of neutron stars was considered by \\citealp{Shav-Heyl-Lith:99}, but ignored general relativisitic effects.  General relativistic studies into the propagation of polarisation in vacuum have been done.  These have been primarily interested in the geometrical effects due to the parallel transport of the linear polarisation \\citep[see \\eg][]{Agol:97,Laor-Netz-Pira:90,Conn-Star-Pira:80}. Alternatively, in \\citealp{Brom-Meli-Liu:01}, polarised emission in a general relativistic environment is considered. However, none of the typical plasma transfer effects (\\eg~Faraday rotation) were included in these calculations. In \\citealp{Heyl-Shav-Lloy:03}, the vacuum birefringence due to strong magnetic fields was considered in the context of neutron star atmospheres. However, in both, refraction was completely ignored. There have been some attempts to study the problem of ray propagation in a covariant form \\citep[see \\eg][]{Brod-Blan:03b,Geda-Melr:01}, but in these the radiative transfer was not addressed.  As discussed in Paper I, refraction coupled with the presence of a horizon can be a source of significant polarisation when the observation frequency is near the plasma and cyclotron frequencies of emitting region.  The sense of the resulting net polarisation is determined by the plasma parameters at the surface at which the polarisation freezes out (when the modes cease to be adiabatic and must be treated as if they were in vacuum).  Typically, this will result in a net circular polarisation.  In a future paper we will discuss astrophysical environments in which this may be the case, including applications to \\Sg and high mass X-ray binaries.  We present a method for performing polarised radiative transfer through a strongly refractive magnetised plasma in a general relativistic environment. Additionally, we apply this to a model of a thick accretion disk. This is done in six sections with \\S\\ref{RP} briefly reviewing the formalism presented in Paper I, \\S\\ref{PRTiRP} discussing how to perform the radiative transfer in a magnetised plasma, \\S\\ref{LHSRiCP} presenting low harmonic synchrotron as a possible emission mechanism, \\S\\ref{R} presenting some results, and \\S\\ref{C} containing conclusions.  The details of constructing a magnetised, thick, barotropic disk are presented in the appendix. ",
        "conclusions": "\\label{C} We have presented refraction as a mechanism for the generation of polarisation when $\\omega_\\infty \\sim \\omega_P,\\,\\omega_B$.  That this will typically result in mostly circular polarisation is a result of the fact that the polarisation eigenmodes are significantly elliptical only when the wave-vector and the magnetic field are within $\\omega_B/\\omega$ of perpendicular, which is usually a small number near the surface where the polarisation freezes out. In addition to producing circular polarisation, this mechanism also significantly suppresses linear polarisation.  Because it does require significant refraction to take place, it is necessarily limited to approximately a decade in frequency, making it simple to identify.  As shown in section \\ref{R:PM:IP}, the resulting circular polarisation will be softer than the intensity.  However, because of optical depth effects, as the observation frequency increases the polarimetric properties of such a system will be dominated be increasingly smaller areas.  As a result, the fractional variability in such a system would be expected to increase with frequency.  Furthermore, even though the emission may arise from a large region, the polarimetric properties will continue to be determined by this compact area, making it possible to have variability on time scales short in comparison to those associated with the emission region. In addition, variability in the circular polarisation would be expected to be correlated with variability in the integrated intensity at frequencies where the emission is dominated by contributions from close to the horizon (\\eg~X-rays).  Possible applications to known astrophysical sources include the Galactic Centre (at submm wavelengths) and extinct high mass X-ray binaries (in the infrared).  These will be disscussed in further detail in an upcoming paper."
    },
    "0311/astro-ph0311156_arXiv.txt": {
        "abstract": "We present a high--resolution study of NGC 2264 IRS1 in CS 2$\\rightarrow$1 and in the 3-mm continuum using the IRAM Plateau de Bure Interferometer. We complement these radio data with images taken at 2.2 $\\mu$m, 4.6 $\\mu$m, and 11.9 $\\mu$m. The combined information allow a new interpretation of the closest environment of NGC 2264 IRS1. No disk around the B--type star IRS1 was found. IRS1 and its low--mass companions are located in a low--density cavity which is surrounded by the remaining dense cloud core which has a clumpy shell--like structure. Strong evidence for induced on--going star formation was found in the surroundings of IRS1. A deeply embedded very young stellar object 20$''$ to the north of IRS1 is powering a highly collimated bipolar outflow. The object 8 in the closer environment of IRS1 is a binary surrounded by dusty circumbinary material and powering two bipolar outflows. ",
        "introduction": "The near--infrared source NGC 2264 IRS1 (IRAS 06384+0932, Allen 1972) is a relatively isolated luminous infrared source (distance 800 pc) in the star--forming region NGC 2264. Previous observations \\citep{schwartz,Krugel,Phillips,Schreyer97} suggest that IRS1 is a young star located within a dense molecular core. This core is embedded in a more extended CO envelope. The 2--200 $\\mu$m luminosity of $\\approx$ 3500 L$_{\\odot}$ (Harvey,  Campbell, \\& Hoffman, 1977) implies an early B--type zero--age main--sequence star \\citep{schwartz}.    Previous ground--based near--infrared images (Scarrott \\& Warren--Smith 1989; Schreyer et al.\\ 1997) show a twisted  jet--like feature to the north of NGC 2264 IRS1, which could be a gas stream ``piercing'' the surrounding dark cloud. Two outflow systems were reported in this region \\citep{Schreyer97}, one at the position of IRS1 and a second one associated with a deeply embedded small star cluster $\\approx 47''$ to the southeast of IRS1. Ward--Thompson et al.\\ (2000) found a ridge (NGC 2264 MMS) of bright submm- and mm--emission around IRS1 wherein five emission peaks (MMS\\,1--5) are located. MMS\\,3 is associated with the gas clump at the position of the small star cluster \\citep{Schreyer97} MMS\\,4 and 5 are located on the eastern and southern side of IRS1, respectively. Wang et al.\\ (2002) found a number of bipolar jets in the region of NGC 2264 MM which point to complex star formation processes in the whole area (diameter $\\le 2'$). The high resolution H$^{13}$CO$^+$ 1\\pf0 map of Nakano, Sugitani, and Morita (2003) shows a shell of dense gas clumps around IRS1. Their corresponding continuum map displays four more compact millimeter sources, three in the closer surroundings of IRS1, and one is associated with the small star cluster to the southeast (MMS3).      In this paper\\footnote{Partly based on observations collected at the European Southern Observatory, La Silla, Chile, under programme IDs 62.I-0530 and 66.C-0209. }$^,$\\footnote{Based on observations carried out with the IRAM Plateau de Bure Interferometer. IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain).}, we focus on  the infrared source IRS1 and its nearest environment (diameter $\\le 51''$). We studied this region in the $K$(2.2$\\mu$m)--band with ESO's NTT as well as in the CS 2\\pf 1 line and the 3-mm continuum using the Plateau de Bure Interferometer (PdBI). Additional thermal infrared data came from observations with ESO's mid-infrared camera TIMMI2. The aim of this study was to answer the following questions: What is the reason for the displaced centre of the bipolar outflow seen with single--dish telescopes (e.g.\\ Schreyer et al.\\ 1997)? Does the presence of a ``wiggled jet'' possibly originating from IRS1 imply the existence of a disk around IRS1\\,? Is this twisted jet dense enough to be traced with the PdBI in order to obtain some information about the velocity structure? Is there an opposite ``jet--like'' feature\\, to confirm the suggestion of a (bipolar) jet? Finally, we wanted to learn what the nature of the central engine is. ",
        "conclusions": "We can conclude that IRS1 is in a more evolved evolutionary state than the young B--type objects AFGL 490 \\citep{Schreyer02} or G192.16--3.82 (Shepherd, Claussen, \\& Kurtz 2001) which show strong evidence for disks. These results showed that massive accretion disks orbiting B2--3 stars in the first 10$^3$--10$^5$ yrs of their main--sequence lifetime exist. However, these disks are unstable and might disappear, for instance, due to gravitational instabilities (e.g., Schreyer et al. 2002). IRS1 seems to be in this phase where the accretion disk already disappeared and a small cavity was created, although, the star is still embedded in the centre of a more extended cloud core \\citep{Krugel,Schreyer97}. The power of IRS1 leads to induced star formation in the surrounding denser cloud clumps, where a number of young stellar objects are embedded powering bipolar outflows. Our data show that the main source of the large-scale bipolar outflow is a deeply embedded young stellar object 20$''$ to the north of IRS1. In addition, the object 8 in the closer environment of IRS1 is a binary surrounded by  dusty circumbinary material and powering two bipolar outflows."
    },
    "0311/astro-ph0311630_arXiv.txt": {
        "abstract": "FREGATE, the gamma-ray detector of HETE-2 is entirely dedicated to the study of GRBs. Its main characteristic is its broad energy range, from 7 keV to 400 keV. This energy range can be further extended down to 2 keV using the data from the WXM, the X-ray detector of HETE-2. Such a large energy range allows to study in details the prompt emission of GRBs, determining with a high precision their spectral parameters. Moreover, because this energy range extends at low energies, the sample of GRBs detected by both  FREGATE and WXM contains a significant fraction of X-Ray Rich GRBs and X-Ray Flashes.  We present here the distributions of the spectral parameters mesured for the time integrated spectra of 50 GRBs. We  put emphasis on the distribution of the low energy spectral index $\\alpha$. Because FREGATE and WXM detected all classes of GRBs, we also discuss the connection between GRBs, X-Ray Rich GRBs and X-Ray Flashes. ",
        "introduction": "\\begin{figure} \\begin{tabular}{ccc} \\includegraphics[angle=-90,width=4cm]{GRB010213grbm.ps} &\\includegraphics[angle=-90,width=4cm]{GRB030329grbm.ps}&\\includegraphics[angle=-90,width=4cm]{GRB030328grbm.ps} \\\\ &&\\\\ GRB010213 & GRB030329 & GRB030328 \\\\ X-Ray Flash & X-Ray Rich GRB & GRB \\\\ $\\alpha = -1.2 , \\beta = -2.8,$ & $\\alpha = -1.2 , \\beta = -2.3, $ & $\\alpha = -1.1 , \\beta = -2.9, $\\\\ $E_o = 4.2$ keV, $E_p = 3.4$ keV & $E_o = 91.4$ keV, $E_p = 73.1$ keV & $E_o = 139.4$ keV, $E_p = 125.5$ keV \\\\ \\end{tabular} \\caption{Spectra of the different classes of GRBs: X-Ray Flashe, X-Ray Rich GRB and GRB using both FREGATE and WXM data.} \\label{spectre} \\end{figure} FREGATE is the gamma-ray detector of HETE-2 (see \\cite{atteia03} for a description of FREGATE). Its broad-energy range 7--400 keV which can be extended down to 2 keV using the WXM instrument (see \\cite{wxm} for a description of WXM) allows us to determine with high precision the spectral parameters of the prompt emission of the GRBs seen by both instruments. The two instruments also detected an important fraction of X-Ray Rich GRBs and X-Ray Flashes (see \\cite{heise01} for a  description of these new classes) and we are now able to discuss the differences and the similarities between these three populations.  We present in this paper an update of the results presented in \\cite{barraud03}: 'Spectral analysis of 35 GRBs/XRFs observed with  HETE-2/FRE\\-GATE'. This paper presented a first spectral analysis of 35 GRBs detected by HETE-2/FREGATE since its launch in October 2000 and which were well localized by either the instruments on-board HETE-2 (WXM or SXC, see \\cite{sxc} for a description of the Soft X-Ray Camera), or by the GRB InterPlanetary Network (IPN).  The update of the paper \\cite{barraud03} corresponds to an increase of the number of GRBs seen by both FREGATE and WXM which now reaches 50, all in the class of long GRBs. We didn't include the two short/hard bursts GRB020113 and GRB020531 detected by HETE-2. Another improvement is that the spectral parameters are now obtained from a joint fit of WXM and FREGATE data. We thus  obtain spectra ranging from 2 keV to 400 keV.  We focus here on the distribution of the spectral parameters: we show that the distribution of the low energy spectral index $\\alpha$ is compatible with the predictions of the synchrotron shock model and we show that a significant fraction of bursts have their peak energy $E_p$ lower than 50 keV. We also put emphasis on the hardness-intensity correlation. This correlation shows that the three populations, GRBs, X-Ray-Rich GRBs and X-Ray-Flashes form a continuum which strongly suggest that they are all produced by the same phenomenon. ",
        "conclusions": ""
    },
    "0311/astro-ph0311406_arXiv.txt": {
        "abstract": "I present a review of X-ray observations of supernova remnants with an emphasis on shell-type remnants. The topics discussed are the observation of fresh nucleosynthesis products, shock heating and cosmic ray acceleration. ",
        "introduction": "Although radio and optical observations have played an important role in identifying and characterizing supernova remnants (SNRs), most of the shock heated material is primarily visible in X-rays, having temperatures of typically $10^6-10^7$~K. The emission mechanisms are bremsstrahlung and line radiation, but a number of remnants are now known to emit X-ray synchrotron radiation from ultra-relativistic electrons, as well. This was already known for pulsar wind nebula dominated, or Crab-like remnants, which fall outside the scope of this review, but also shell-type remnants like SN 1006 emit synchrotron radiation from electrons accelerated by the blast wave \\cite{koyama95}.  The X-ray line emission of SNRs is in particular interesting, as it allows the study of the most abundant metals, the so called $\\alpha$-elements, which prominently feature in the X-ray spectra of young SNRs like \\casa, Tycho (SN 1576) and Kepler (SN 1064). This line emission, and mass estimates based on the X-ray emission, give important information on the nucleosynthesis yields of the explosions, which can then be compared to theoretical supernova (SN) models \\cite{ww95,nomoto97}.  The line and bremsstrahlung emission give us a window on the past of the SNR: its progenitor and the SN explosion. X-ray synchrotron radiation, on the other hand, is interesting as it provides a window on cosmic ray shock acceleration. SNRs are thought to be the primary source of cosmic rays up to, and possibly beyond, energies of $10^{15.5}$~eV, at which energy the spectrum has a break, usually referred to as ``the knee'', However, direct evidence for a SNR origin for cosmic ray particles other than electrons is lacking. For a long time the observational support for a SNR origin of cosmic rays was the detection of radio synchrotron emission from SNRs, revealing the presence of cosmic ray electrons up to energies of a few GeV. Recently the detection of X-ray synchrotron radiation from SNRs have substantially extended this to  $\\sim100$~TeV \\cite{koyama95,reynolds99}.  In this review I will discuss recent findings concerning both the study of fresh nucleosynthesis and cosmic ray acceleration in SNRs, emphasizing the role of X-ray observations, including results based on \\sax\\ observations.  \\begin{figure} \\psfig{figure=jvink1_f1.eps,width=0.5\\textwidth} \\caption{The supernova classification scheme, which is based on optical spectroscopy and light curve (LC) shape.\\label{scheme}} \\end{figure} ",
        "conclusions": "The last paragraph shows that to some extent the issue of  cosmic ray acceleration and explosive nucleosynthesis are intertwined. In order to assess abundance patterns an understanding of the radiation processes, bremsstrahlung, line emission, and synchrotron emission is needed. The possible presence of synchrotron emission in a line dominated spectrum has consequences for the mass estimates \\cite{favata97}. The discovery of X-ray synchrotron radiation from shell-type remnants have complicated the analysis of their X-ray emission, but has given unexpected new insights in the particle acceleration properties of supernova remnants.  These advances in our understanding have been made possible by the emergence of spatially resolved X-ray spectroscopy, first with \\asca\\ and \\sax, and now fruitfully continued with \\chandra\\ and \\xmm.  \\vskip 0.1cm I thank Kurt van der Heyden for generating Fig.~\\ref{typeIa}. This work was supported by NASA's Chandra Postdoctoral Fellowship Award No. PF0-10011 issued by the Chandra X-ray Observatory Center, which is operated by the SAO under NASA contract NAS8-39073."
    },
    "0311/astro-ph0311295_arXiv.txt": {
        "abstract": "We study the photometric properties of Lyman-break galaxies (LBGs) formed by redshift $z=3$ in a set of large cosmological smoothed-particle hydrodynamics (SPH) simulations of the $\\Lambda$ cold dark matter ($\\Lambda$CDM) model. Our numerical simulations include radiative cooling and heating with a uniform UV background, star formation, supernova feedback, and a phenomenological model for galactic winds.  Analysing a series of simulations of varying box size and particle number allows us to isolate the impact of numerical resolution on our results.  We compute spectra of simulated galaxies using a population synthesis model, and derive colours and luminosity functions of galaxies at $z=3$ after applying local dust extinction and absorption by the intergalactic medium (IGM). We find that the simulated galaxies have $\\U-G$ and $G-R$ colours consistent with observations, provided that intervening absorption by the IGM is applied.  The observed properties of LBGs, including their number density, colours, and luminosity functions, can be explained if LBGs are identified with the most massive galaxies at $z=3$, having typical stellar mass of $M_{\\star} \\sim 10^{10}\\himsun$, a conclusion broadly consistent with earlier studies based on hydrodynamic simulations of the $\\Lam$CDM model. We also find that most simulated LBGs were continuously forming stars at a high rate for more than one Gyr up until $z=3$, but with numerous starbursts lying on top of the continuous component. Interestingly, our simulations suggest that more than 50\\% of the total stellar mass and star formation rate in the Universe are accounted for by galaxies that are not detected in the current generation of LBG surveys. ",
        "introduction": "\\label{section:intro}  Based on surveys made with modern large telescopes, the number of observed Lyman-break galaxies \\citep[LBGs; e.g.][]{Steidel95} has reached about $\\sim 1000$ \\citep{Ste03} by now, providing a unique window for studying early galaxy formation. In these observations, a large sample of LBG-candidates is first selected as ``drop-outs'' in certain sets of colour-colour planes (e.g. $\\U-G$ vs. $G-R$), and then later spectroscopic follow-up observations are performed to obtain the redshifts of the identified candidates. Most surveys to date have been carried out at optical wavelengths, corresponding to the rest-frame far ultra-violet (UV) at redshift $z=3$. The LBGs found in this way show strong clustering \\citep{Ade98, Giavalisco98, Steidel98}, which has generally been interpreted as indirect evidence that LBGs reside in massive dark matter halos. Several studies of LBGs based on semi-analytic models for galaxy formation agree with the massive dark matter halo hypothesis \\citep[e.g.][]{Mo96, Bau98, Kau99, Mo99}, but a numerical study by \\citet{Jing} using collisionless N-body simulations expressed some concerns whether the clustering of dark matter halos of $<10^{12}\\himsun$ can explain the observed clustering strength of LBGs. However, a later investigation by \\citet{KHW99} argued that the CDM models have no difficulty in explaining the strong observed clustering of LBGs, although the substantial box size effects in their analysis introduced some uncertainty.  More recently, LBGs have also been studied in the near infrared (IR) \\citep*{Sawicki98, Papovich01, Pettini01, Sha01, Rudnick01}.  Near IR observations are less affected by dust obscuration and directly probe the rest-frame optical properties of LBGs at $z=3$.  Therefore, they allow a more reliable derivation of the stellar masses of LBGs.  Based on these observations, it has been suggested that the value of the extinction lies in the range $E(B-V)\\approx 0.0-0.3$ with a median value of 0.15, and that a significant old stellar population exists in LBGs at $z=3$, with a stellar mass of $\\approx 10^{10}\\Msun$ \\citep{Papovich01, Sha01}.  The existence of such a stellar component would again suggest that LBGs are embedded in massive dark matter halos which have continuously formed stars over an extended period of roughly one Gyr up to $z=3$. These systems would then most likely evolve into elliptical galaxies at the present day, or into the spheroidal components of massive spiral galaxies.  However, in a competing model, LBGs have been suggested to be merger-induced starbursting systems associated with low-mass halos \\citep*{Lowenthal97, Sawicki98, Som01}, and in some cases merger-induced starbursts are given a crucial role even when LBGs are the most massive galaxies at their time \\citep{Som01}. Given that the merger rate is expected to be quite high at $z\\simeq 3$, these scenarios provide an interesting alternative to the more conventional picture which associates LBGs with the most massive systems.  Self-consistent hydrodynamic simulations are an ideal tool for trying to distinguish between these different scenarios for the nature of LBGs.  \\citet{Dave99} and \\citet{Wei02} were the first to employ smoothed particle hydrodynamics (SPH) simulations to this end. However, their box size of $11.1\\,\\himpc$ represented an important limitation, because the space-density of LBGs is so low that only a few of them can be found in a volume of this size, as we will discuss further in Section~\\ref{section:color-color}. Therefore, simulations with a larger box size are desirable to obtain larger samples of simulated LBGs.  \\citet{Nag02} used an Eulerian hydrodynamic simulation with a box size of $\\Lbox=25\\,\\himpc$, tracing the merger history of galaxies from $z=5$ to $z=0$. The results of these earlier numerical studies were consistent with each other, and agreed reasonably well with the observations, within the uncertainties. In particular, the median stellar masses of LBGs were predicted to be $\\sim 10^{10}\\himsun$, and the simulated galaxies were experiencing significant star formation rates ($>30\\,\\Msun/\\yr$) for extended periods of time ($\\ge 1$ Gyr).   \\begin{table*} \\begin{center} \\begin{tabular}{cccccccc} \\hline Run & Box Size & ${N_{\\rm p}}$ & $m_{\\rm DM}$ & $m_{\\rm gas}$ & $\\epsilon$ & $z_{\\rm end}$ & wind \\\\ \\hline \\hline O3  & 10.00 & $2\\times 144^3$ &  $2.42\\times 10^7$ & $3.72\\times 10^6$ &2.78 & 2.75 & none \\cr P3  & 10.00 & $2\\times 144^3$ &  $2.42\\times 10^7$ & $3.72\\times 10^6$ &2.78 & 2.75 & weak \\cr Q3  & 10.00 & $2\\times 144^3$ &  $2.42\\times 10^7$ & $3.72\\times 10^6$ &2.78 & 2.75 & strong \\cr Q4  & 10.00 & $2\\times 216^3$ &  $7.16\\times 10^6$ & $1.10\\times 10^6$ &1.85 & 2.75 & strong \\cr Q5  & 10.00 & $2\\times 324^3$ &  $2.12\\times 10^6$ & $3.26\\times 10^5$ &1.23 & 2.75 & strong \\cr \\hline D4  & 33.75 & $2\\times 216^3$ &  $2.75\\times 10^8$ & $4.24\\times 10^7$ &6.25 & 1.00 & strong \\cr D5  & 33.75 & $2\\times 324^3$ &  $8.15\\times 10^7$ & $1.26\\times 10^7$ &4.17 & 1.00 & strong \\cr \\hline G5  & 100.0 & $2\\times 324^3$ &  $2.12\\times 10^9$ & $3.26\\times 10^8$ &8.00 & 0.00 & strong \\cr G6  & 100.0 & $2\\times 486^3$ &  $6.29\\times 10^8$ & $9.67\\times 10^7$ &5.00 & 0.00 & strong \\cr \\hline \\end{tabular} \\caption{Simulations employed in this study.  The box size is given in units of $\\himpc$, ${N_{\\rm p}}$ is the particle number of dark matter and gas (hence $\\times\\, 2$), $m_{\\rm DM}$ and $m_{\\rm gas}$ are the masses of dark matter and gas particles in units of $\\himsun$, respectively, $\\epsilon$ is the comoving gravitational softening length in units of $\\hikpc$, and $z_{\\rm end}$ is the ending redshift of the simulation. The value of $\\epsilon$ is a measure of spatial resolution.  From the top to the bottom row, we collectively call the first 5 runs (O3 to Q5) `Q-series', D4 \\& D5 are called `D-series', and G5 \\& G6 are called `G-series'. \\label{table:sim}} \\end{center} \\end{table*}   In this paper, we improve on the earlier numerical studies of LBGs by using a new set of high-resolution numerical simulations. These simulations are based on a novel model for the physics of star formation and feedback, and they use a more accurate implementation of SPH. For the first time, we also systematically study the effects of resolution and box size in the context of simulated LBG galaxies. The treatment of star formation and feedback we use is based on a sub-resolution multi-phase description of the dense, star-forming interstellar medium (ISM), and a phenomenological model for strong feedback by galactic winds, as recently proposed by \\citet{SH03a}. This model has been shown to provide converged star formation rates for well-resolved galaxies, with a cosmic star formation history consistent with recent observations \\citep{SH03b, Her03}.\\footnote[3]{We note an error in figure 12 of \\citet{SH03b} in which the observational estimates of the SFR were plotted too high by a factor of $h^{-1} = 1.4$.  When corrected, the observed points are in better agreement with the theoretical estimates; see astro-ph/0206395.} The inclusion of winds was motivated by the realisation that galactic outflows at high redshift \\citep{Pet02} likely play a key role in distributing metals into the intergalactic medium \\citep[e.g.][]{Aguirre01a, Aguirre01b}, as well as being important for the regulation of star formation activity. In fact, winds may also alter the distribution of neutral gas around galaxies \\citep{Ade03}, although the details of how this process may happen remain unclear \\citep[e.g.][]{Croft02, Kol03, Bruscoli}. In passing, we note that both \\citet{Des} and \\citet{Maselli} have found that the Lyman-$\\alpha$ transmissivity close to LBGs, as measured by \\citet{Ade03}, is better reproduced if LBGs are identified as dwarf starbursting galaxies as proposed in \\citet{Som01} and \\citet{WW}. We will discuss the work of \\citet{WW} in Section~\\ref{section:discussion}. Together with the increase in numerical resolution provided by our simulations, it is of interest to see how our refined physical modelling modifies the predictions for LBG properties within the $\\Lambda$CDM scenario.  This paper is organised as follows. In Section~\\ref{section:simulation}, we briefly introduce the numerical parameters of our simulation set. In Section~\\ref{section:method}, we then describe our method for computing spectra of simulated galaxies both in the rest-frame and the observed frame.  In Sections~\\ref{section:color-color} and \\ref{section:colour-magnitude}, we show the colour-colour diagrams and colour-magnitude diagrams of simulated galaxies, and we discuss the number density of colour-selected LBGs, as well as the stellar masses of LBGs at $z=3$. We then investigate the rest-frame $V$-band luminosity function and observed $R$-band luminosity function in Section~\\ref{section:lf}, followed by an analysis of the star formation histories of LBGs in Section~\\ref{section:sf}. Finally, we summarise and discuss the implications of our work in Section~\\ref{section:discussion}. ",
        "conclusions": "\\label{section:discussion}  We have used state-of-the-art hydrodynamic simulations of structure formation to study the properties of LBGs in a $\\Lam$CDM universe. Our simulations use a new `entropy conserving' SPH formulation that minimises systematic inaccuracies in simulations with cooling, as well as an improved model for the treatment of the multiphase structure of the ISM in the context of star formation and feedback \\citep{SH03a}. For the first time, our study uses a large series of simulations to investigate the properties of LBGs, allowing an exploration of an unprecedented large dynamic range in both mass and spatial scales, while simultaneously providing reliable estimates of systematic effects due to numerical resolution.  By comparing our results for LBG properties in simulations with different box size, we showed that a sufficiently large volume (at least $\\Lbox>50\\,\\himpc$) is crucial to faithfully sample the LBG population. In particular, the bright-end of the LBG luminosity function is invariably incomplete in simulations with small box size. Of equal importance is a proper treatment of effective feedback processes. By considering a series of runs with different strengths of galactic winds, we showed that without an effective feedback process, the number of LBGs is overpredicted. However, for our default model of strong winds, a reasonable space density of bright LBGs is obtained. In particular, we found that our G-series with a box size of $\\Lbox=100\\,\\himpc$ has a quite plausible LBG population, with luminosity functions in both rest-frame $V$-band and observed-frame $R$-band that match the observations reasonably well, at least at the bright-end.  Perhaps the most important conclusion of this paper is that the observed properties of LBGs, including their number density, colours, and luminosity functions, can be well explained if the LBGs are simply associated with the most massive galaxies at $z=3$, with median stellar mass of $M_* \\sim 10^{10}\\himsun$. This conclusion is consistent with earlier numerical studies based on hydrodynamic simulations of the $\\Lam$CDM model \\citep{Dave99, KHW99, Nag02, Wei02} as well as some semi-analytic models \\citep{Bau98, Kau99}, and does not provide direct support for alternative models which suggest that LBGs are star-bursting low-mass systems that later evolve into low-mass spheroids at $z=0$. This point is corroborated by the high rates of star formation with $>10 ~\\Msun~\\yr^{-1}$ seen over extended periods of time of order 1~Gyr in the simulated galaxies, leading to the build-up of typical stellar masses of $10^{10}\\himsun$ at $z=3$. These comparatively steady star formation histories are also consistent with observational studies by \\citet{Papovich01} and \\citet{Sha01}.  Note however that these same observations also find some evidence for starburst activity ($\\ge 100\\, \\Msun \\yr^{-1}$) within the last 500 Myr before $z=3$, lasting for periods of $\\sim 100$~Myr. Such long violent bursts are not seen in our SPH simulations, but they were present in the simulations analysed by \\citet{Nag02}. Instead, multiple shorter bursts with time-scale of $\\sim 10-20$~Myr are seen in our SPH simulations.  A recent study by \\citet{WW} however claims that the kinematic measurements of \\citet{Pettini01} and \\citet{Erb} favour the picture of LBGs being low-mass starbursting systems. But the measured `rotation curves' are strongly dependent on seeing, and it is very likely that these measurements underestimate the true rotation curves (M. Pettini, private communication), casting some doubt on these claims.  To settle this issue, we have to wait for the detection of a definite flattening of the rotation curves, which might become possible in the near future with near-IR spectrographs coupled to adaptive optics systems.  Another difficulty with the low-mass starbursting-galaxy picture for LBGs is that such systems would be expected to have relatively low chemical abundances, whereas at least the brighter LBGs have near-solar metallicity.   \\begin{figure} \\epsfig{file=f8.eps,height=3.5in,width=3.2in, angle=0} \\caption{Star formation histories of selected galaxies in the `G6'-run with the bin-size of 10 Myrs.  Galaxy ID, stellar mass in units of $\\himsun$, rest-frame $V$-band magnitude, and apparent $R$ magnitude for $E(B-V)=0.15$ are indicated in each panel for the corresponding galaxy. \\label{sf_LBG.eps}} \\end{figure}  The luminosity functions we measure for our simulated LBG galaxies are characterised by a steep faint end, with a slope close to $\\alpha \\sim -2$.  This is the slope of the dark matter halo mass function, so in our simulations, the galaxy luminosity function seems to largely follow the halo mass function at the low-mass end.  Observationally, we know that at $z=0$ the faint-end slope of galaxies is certainly much shallower than this, with a slope of order $\\alpha\\sim -1.2$ \\citep{Blanton}. At high-redshift, the faint-end of the galaxy luminosity function is not nearly as well constrained, however, and may be much steeper. In fact, \\citet{Sha01} report a slope of $\\alpha=-1.85$ for LBG galaxies at $z=3$ which is still consistent with our simulations.  However, as the authors explain, this should be taken with some caution because their observations only reach an absolute $V$-band magnitude of $-20.5$, which is not deep enough to constrain the faint-end slope reliably. If confirmed, a steep luminosity function at $z\\simeq 3$ would in any case have to evolve into a much flatter one at low redshift. Whether such a scenario is viable or not will be tested by current and future very deep surveys, such as DEEP2 \\citep{Davis, Madge, Coil}.  Barring confirmation of a rapid evolution of the faint-end slope, our results therefore hint that the simulations still overproduce the number of low-mass galaxies, despite the inclusion of strong feedback processes that are capable of accounting for the observed brightness of $L_{\\star}$-galaxies. The moderating effects of feedback on star formation do depend on galaxy mass in the physical model followed by the simulations in the sense that, for a given amount of star formation, the feedback by winds does comparatively more damage for small mass galaxies. This is because towards smaller galaxy mass scales, the winds find it easier to escape from their confining galactic potential wells. The winds also entrain more gas in the outflow for lower circular velocities, so that the net baryonic loss becomes larger towards smaller mass scales, making the total mass-to-light ratio of smaller halos in principle bigger.  But, this variation of the feedback efficiency with galaxy mass may not be strong enough in the present set of simulations to close the significant gap between the faint-end slopes of halo mass function and galaxy luminosity function.  Interestingly, the studies by \\citet{Chiu, Nag01b}, and \\citet{Nag02} based on Eulerian hydrodynamic simulations did find a flatter faint-end slope \\citep[see also][]{Harford}. Presently it is unclear whether this was due to resolution limitations in the low-mass end of the halo mass function, or due to genuine physical effects of the feedback model implemented in these Eulerian simulations. More work in the future will be needed to settle this very interesting question, which is of tremendous importance for the theoretical framework of hierarchical galaxy formation in CDM universes."
    },
    "0311/astro-ph0311180_arXiv.txt": {
        "abstract": "It is argued that the nanosecond giant subpulses detected recently in the Crab pulsar are generated by means of the coherent curvature radiation of charged relativistic solitons associated with sparking discharges of the inner gap potential drop above the polar cap. ",
        "introduction": "Although only four pulsars are known to emit giant pulses \\citep{lun95,cog96,rj01,jr03}, understanding the mechanism of their radiation can potentially lead to understanding the longstanding problem of a pulsar radio emission. Recently, the detection of extremely short and powerful, 2-nanosecond - 1000 Jy, subpulses within the radio giant pulses from the Crab pulsar has been reported by \\citet[][ HKWE03 hereafter]{hetal03}. This is an observational result of extraordinary importance, for it can shed a new light onto the mystery od coherent pulsar radio emission. HKWE03 argued that these nanosecond giant subpulses were true temporal modulations associated with an explosive collapse of nonlinear plasma turbulence \\citep{wetal98}. In this letter we argue that the time-scale of the Crab pulsar giant nanosecond subpulses is consistent with the angular beaming due to the curvature of dipolar field lines (different from the conventional angular beaming due to the pulsar rotation), thus with the coherent curvature radiation. Furthermore, we demonstrate that the observed fluxes of these giant subpulses are consistent with the emission of charged relativistic solitons, generated by means of the modulational instability of the strong Langmuir turbulence associated with a sparking discharge of the pulsar's polar gap (Melikidze, Gil \\& Pataraya 2000, MGP00 hereafter; Gil, Lyubarski \\& Melikidze 2004, GLM04 hereafter). ",
        "conclusions": "In this letter we argue that the 2-ns giant subpulses detected in the Crab pulsar are due to the curvature radiation of one or at most several solitons associated with a sparking discharges occurring near the local surface magnetic pole. This is the most elementary emission event that can be observationally resolved. The curvature radiation from the particular soliton(s) can be observationally distinguished by means of the kinematical boosting (eq.[\\ref{dtau}]) only if they are slightly more powerful than other solitons. This implies that the sparking event associated with giant subpulses is more energetic than an average one. Both luminosity and characteristic frequency of the soliton curvature radiation strongly depend on the Lorentz factor $\\gamma$, which in turn depends on the accelerating potential drop above the polar cap. Interestingly, pulsars exhibiting giant pulses are distinguished not only by their brightness, but also by extremely high values of the so-called complexity parameter $a=(r_p/h)$, where $r_p$ is the polar cap radius and $h$ is the polar gap height \\citep{gs00}. Since the accelerating potential drop $\\Delta V\\propto h^2$ and $h<r_p/\\sqrt{2}$ \\citep{rs75}, then pulsars with very high values of $a>>1$ must have a large reservoir of the maximum available potential drop over the actual potential drop $(\\Delta V_{max}/\\Delta V=a^2/2)$. This reservoir can be occasionally used to create exceptionally energetic spark(s).  \\begin{deluxetable}{lccccc} \\tabletypesize{\\scriptsize} \\tablecaption{Pulsar parameters\\label{table1}} \\tablecomments{Pulsars marked in boldface show giant pulses, others are candidates for giant pulse emission. Numbers in parenthesis denote ranking position in $B_{LC}=2850\\dot{P}^{0.5}P^{-2.5}$. The complexity parameter $a=425(\\dot{P})^{2/7}P^{-9/14}$.} \\tablewidth{3in} \\tablehead{\\colhead{PSR J} & \\colhead{$P$} & \\colhead{$\\dot{P}$} & \\colhead{$a$} &  \\multicolumn{2}{c}{$B_{LC}$} \\\\ & \\colhead{(msec)} & \\colhead{$(10^{-15})$} & & \\multicolumn{2}{c}{($10^5$G)}}  \\startdata {\\bf 0534+2200} & 34 & 420 & 249 & 8.7 & (2) \\\\ {\\bf 0540$-$6919} & 51 & 479 & 198 & 3.3 & (6) \\\\ 1513$-$5908 & 150 & 1540 & 138 & 0.4 & (21) \\\\ 1124$-$5916 & 135 & 745 & 120 & 0.4 & (22) \\\\ 1617$-$5055 & 69 & 137 & 114 & 0.8 & (12) \\\\ 1420$-$6048 & 68 & 83 & 99 & 0.7 & (14)\\\\ 1119$-$6127 & 408 & 4002 & 95 & 0.05 & (75)\\\\ {\\bf 0835$-$4510} & 89 & 125 & 94 & 0.4 & (20)\\\\ {\\bf 1824$-$2452} & 3 & 0.0016 & 33 & 7.3 & (4)\\\\ 1823$-$30A & 5.4 & 0.0034 & 28 & 2.4 & (8)\\\\ {\\bf 1939$+$2134} & 1.6 & 0.0001 & 23 & 8.8 & (1)\\\\ {\\bf 0218$+$4232} & 2.3 & 0.00007 & 16 & 3.1 & (7)\\\\ {\\bf 1959$+$2048} & 1.6 & 0.00002 & 14 & 3.6 & (5)\\\\ 2129$+$1210E & 4.6 & 0.00018 & 14 & 8.2 & (3)\\\\ \\enddata \\end{deluxetable}  The values of the complexity parameter $a$ can well exceed 100 in normal pulsars, while in millisecond pulsars they are limited to about 40 \\citep[see Fig.~1 in][]{gs00}. Therefore, when pulsars in both groups are ranked with respect to $a$, pulsars with giant pulses occupy the top of the lists (Table 1). Also, the Vela pulsar, which was reported to exhibit some kind of giant pulse behavior \\citep{joh01} is very high on the list. Moreover, sporadic large amplitude pulses (LAP) from two millisecond pulsars (J1959+2048 and J0218+4232) have been reported very recently by \\citet{jos03}. It should be emphasized that pulsars showing giant pulses have also high values of the magnetic field $B_{\\rm LC}\\gtrsim 10^5$~G at the light cylinder. This might, however, be a coincidence because roughly $B_{\\rm LC}\\sim a^2/P$, where $P$ is the pulsar period (Table 1). Although $B_{\\rm LC}$ is a good parameter to make a list of giant pulse candidates, only a value of $a^2$ has a physical meaning within our model. In Table 1 we propose a number of candidates for giant pulses with high values of the complexity parameter ($a>10$ for millisecond pulsars and $a>90$ for normal pulsars) and relatively low values of $B_{\\rm LC}$. The detection (non-detection) of giant pulses from these candidates would confirm (refute) our scenario. PSR J1119-6127 with low $B_{\\rm LC}\\approx 5\\times 10^3$ G seems the most interesting case.  The case of PSR J1959+2048 deserves some additional discussion in the light of the model discussed here. As emphasized by \\citet{jr03}, this pulsar is exceptional in the sense that despite the high value of $B_{\\rm LC}$ it does not show giant emission (although \\citet{jos03} found one significant LAP (out of million pulses) with an intensity $129$ times the mean intensity). This seems to be consistent with our model since this pulsar (which is fourth ranked with respect to $B_{\\rm LC }$) has a relatively small value of the complexity parameter $a$ as compared with two other millisecond pulsars showing giant pulses (Table 1). In fact, since $\\Delta V_{\\rm max}/\\Delta V=a^2/2$, the smaller the value $a$, the smaller the reservoir of the polar gap energy that can ignite giant emission.  Also the case of Vela pulsar is very interesting in view of the models of giant emission. \\citet{kjv02} reported giant micro-pulses in this pulsar, which are different from normal micro-pulses and are probably closely related to classical giant pulses. They have typically large amplitudes, appear to be narrower than normal micro-pulses and exhibit a power law in their cumulative probability distribution. The first two properties are quite natural within our model (a sequence of unresolved nano-pulses?) and we discuss the last issue bellow.  According to our scenario, both ordinary and giant pulses are related to the inner gap sparking activity and originate at relatively low altitudes, contrary to the suggestion that the latter arise in the outer gap region \\citep{rj01}. Generally, the following two steps are involved: first a corresponding spark should be relatively intense (which corresponds to relatively high Lorentz factors), so the associated solitons are distinguished from the background radiation (the soliton radiation intensity depends very strongly on the Lorentz factor, see eq.[43] in GLM04). Then, in the second step, the intensity of the curvature radiation of distinguished soliton(s) can be kinematically boosted by means of equation (\\ref{lum}). It is intuitive that the distribution of spark energy is quasi-gaussian, with a low energy cutoff corresponding to the pair formation threshold. Since the giant pulses should be associated with sparks corresponding to the high energy tail of this distribution, it is natural that they exhibit a power law in their cumulative probability distribution."
    },
    "0311/astro-ph0311149_arXiv.txt": {
        "abstract": "We present new results on the angular momentum evolution of dark matter halos. Halos, from N-body simulations, are classified according to their mass growth histories into two categories: the {\\it accretion category} contains halos whose mass has varied continuously and smoothly, while the {\\it merger category} contains halos which have undergone sudden and significant mass variations (greater than 1/3 of their initial mass per event). We find that the angular momentum  grows in both cases, well into the nonlinear regime. For individual halos we observe strong correlation between the angular momentum variation and the mass variation. The rate of growth of both mass and angular momentum has a characteristic transition time at around $z \\sim 1.5 - 1.8$, with an early fast phase followed by a late slow phase. Halos of the merger catalog acquire more angular momentum even when the scaling with mass is taken into account. The spin parameter has a different behavior for the two classes: there is a {\\it decrease} with time for halos in the accretion catalog whereas a small {\\it increase} is observed for the merger catalog. When the two catalogs are considered together, no significant variation of the spin parameter distribution with the redshift is obtained. We have also found that the spin parameter neither depends on the halo mass nor on the cosmological model. From our simple model developed for the formation of a disk galaxy similar to the Milky Way, we conclude that our own halo must have captured satellites in order to acquire the required angular momentum and to achieve most of the disk around $z \\sim 1.6$. The distribution of the angular momentum indicates that at $z \\sim 1.6$ only $22\\%$ of the halos have angular momentum of magnitude comparable to that of disk galaxies in the mass range $10^{10}  - 5\\times 10^{11} M_{\\odot}$, clearly insufficient to explain the present observed abundance of these objects. ",
        "introduction": "The origin of angular momentum remains a key factor in understanding the formation, evolution and particularly the morphological types of galaxies. In the gravitational instability paradigm of structure formation, angular momentum could arise from the tidal interaction of galaxies with their surroundings (Hoyle 1949; Peebles 1969; Doroshkevish 1970; White 1984). As a protogalaxy expands, its angular momentum grows linearly with time due to tidal interactions until it decouples from the background expansion, turns around and starts collapsing. After the turn around time the angular momentum essentially stops growing and becomes less sensitive to tidal torques (Peebles 1969), and hence this time marks the maximum angular momentum acquired by galaxies through this mechanism.  In a Universe dominated by dark matter, within the hierarchical scenario, galaxies are formed  when baryonic gas falls into the gravitational potential well of  dark matter halos. Therefore, many properties of galaxies, including their angular momentum, are expected to be closely related to those of their host halos. The acquisition of angular momentum by dark matter halos through tidal torques, can be formulated within the Lagrangian perturbation theory (White 1984; Catelan \\& Theuns 1996a, 1996b). At the first order, using Zel'dovich approximation, one can show that the angular momentum grows linearly with time. This result has been confirmed by N-body simulations. However, the tidal torque theory (TTT) is an oversimplified explanation of the complex evolution of the angular momentum of dark halos. Real progenitors of dark matter halos are not isolated, as assumed in TTT, but are continuously growing by accretion and merger. In the non-linear regime, no self-consistent theory, including  mass accretion and merger history of halos yet exists. Most of the investigations rely on numerical simulations and or on semi-analytic modelings. The transfer of orbital angular momentum to spin during a merger event has been verified in simulations, where one observes a correlation between a sharp increase in the halo mass and sudden variations in the dimensionless spin parameter $\\lambda ={J\\mid E \\mid ^{1/2}}/{GM^{5/2}}$ where $J$ is the angular momentum, $E$ is the total energy of the halo and $M$ is its mass. The spin parameter is essentially the ratio of the angular momentum of an object to that required for rotational support. There also seems to be a general tendency for $\\lambda$ to decrease during periods of slow accretion (Vitvitska et al. 2002). Effects of mergers are clearly seen in the probability distribution of the spin parameter. Halos that have undergone major merger events at $z \\leq 2-3$ have systematically larger spin parameters than those that have evolved only by accretion (Gardner 2001; Vitvitska et al. 2002; Maller, Dekel \\& Somerville 2002). Since on galactic scales ellipticals have lower angular momentum than spirals, these results could be troublesome if the former are merger remnants of the latter.  Thus, in spite of extensive works, many aspects of the evolution of the angular momentum in the nonlinear regime remain inconclusive. For instance, it remains unclear if the distribution of the spin  parameter depends on the redshift or not (Lemson \\& Kauffmann 1999; Maller, Dekel  \\& Somerville 2002). Concerning the variation of $\\lambda$ with the halo mass, the early investigations indicate a slight decrease with a mean slope ${d({\\rm log}<\\lambda>)}/{d({\\rm log}M)} = -0.17\\pm0.07$ (Barnes \\& Efstathiou 1987). Such a trend was also observed in more recent numerical simulations (Cole \\& Lacey 1996). Although this result seems to be compatible with the tidal torque theory, a small but opposite variation was obtained if halos were to acquire angular momentum by merging (Maller, Dekel \\& Somerville 2002). Recent tests of TTT by N-body simulations (Porciani, Dekel  \\& Hoffman 2002a, 2002b) have shown that the values of the spin amplitudes predicted by TTT are in agreement with those derived from simulations if the linear growth stops somewhat earlier than maximum expansion. After shell crossing, some numerical studies suggest that, in general, the angular momentum decays (Barnes \\& Efstathiou 1987; Sugerman, Summers \\& Kamionkowski 2000), which is expected due to spin conservation if the system remains isolated. On the other hand, erratic variations are expected due to the transfer of orbital angular momentum to spin during merger events or due to accretion (Vitvitska et al. 2002).  In the present work, considering that models based on TTT do not agree in detail with the results of N-body simulations in the nonlinear regime, we investigate an alternative approach in which  halos obtain most of their spin through the transfer of orbital angular momentum by continuous accretion and/or mergers. We identify the halos using FOF algorithm (Davis et al. 1985) and make two distinct catalogs: the ones which have never had a major merger event, and another catalog for halos which have undergone at least one merger episode, in which their mass increased by at least $1/3$ of their initial value during the event. We show that even after shell crossing, the mean angular momentum of dark halos still grows, but with a rate different from that predicted by TTT. Individual halos have erratic variations of the angular momentum, which are strongly correlated with the mass accretion rate, producing on the average a growth of the angular momentum. We recognize two scaling regimes for the growth of the mean halo masses and angular momenta. The corresponding exponents differ for the accretion and merger samples. We find that, when no distinction between the two catalogs are made, i.e., when all the halos are considered together, no significant variation of the spin parameter distribution with the redshift is observed. However, when the halos are considered separately, the evolution of spin parameter of the accretion sample indicates a small but significant {\\it decrease} with time, whereas for the merger catalog the spin parameter {\\it increases} with time. We show that this behavior is a consequence of the different time scaling laws followed by the dynamical variables defining the spin parameter.  Finally, we also discuss the formation of disk galaxies. Past cosmological N-body/SPH simulations have shown that the scale length and angular momenta of disks (baryonic component) are about one order of magnitude smaller than the observed values. This problem is twofold: firstly, in N-body simulations baryons loose a significant fraction of their angular momentum, resulting in disks that are too small and secondly the angular momentum distributions reveal too much low angular momentum material. If the specific angular momentum distribution is conserved, which is usually assumed to be so for disk formation, the resulting disk shows a more centrally concentrated mass distribution making it difficult to explain the exponential mass profile of spirals (Navarro, Frenk \\& White 1995; Navarro \\& Steinmetz 1997, 2000; van den Bosch 2001; van den Bosch, Burkert \\& Swaters 2001). We discuss a simple collapse model to form a baryonic proto-disk. We will show how the halo parameters, in particular the specific angular momentum distribution, affect the disk properties like its mass distribution, spin parameter and how the angular momentum acquired during the growth phases of the halo influences the evolution of the disk. The plan of this paper is as follows: in Section 2, we review briefly the basic principles of TTT and alternative theories, in Section 3 we describe our N-body simulations and our halo-finding algorithms and we discuss the angular momentum evolution by accretion and merger events in Section 4. In Section 5 we develop our scenario describing the formation of a disk galaxy and in Section 6, we summarize our main results and conclusions. ",
        "conclusions": "We have studied the effects of accretion and merger on the dynamical evolution of dark matter halos, specifically on their angular momentum. The difference between accretion and merger could essentially be semantical, since the criterion for separating the two cases is somehow arbitrary. In this work, accretion should  be understood as a \"continuous and almost smooth variation\" of the halo mass, while this variation in the merging case is  \"sudden and significant\". When a halo mass changes by more than a $1/3$ of its present value then we consider the event a merger and otherwise accretion. In this manner, we have built and studied two separate merger and accretion catalogs.  We have demonstrated that our results are only marginally modified if the mass-fraction threshold of $1/3$ is reduced by a factor of two.  We found that in both cases (accretion and merger), there are two distinct phases of mass growth: an early and a fast phase, followed by a late phase of slow mass growth, ({\\it e.g}.\\ in agreement with Zhao et al. 2003).  The transition occurs at $z \\sim 1.5-1.8$ (later than the value of $z \\sim 3$ found by Zhao et al. 2003). The  mass growth in both phases are well-represented by power laws. The exponents not only vary from one phase to another, but also change between the merger and the accretion catalogs. It is worth mentioning that in the hierarchical picture, if the power spectrum of the initial fluctuations is of the form $\\mid \\delta_k \\mid ^2 \\propto k^n$, in the linear regime the characteristic masses grow as $m \\propto t^{4/(n+3)}$ (Peebles 1980). In this case, for a Harrison-Zel'dovich spectrum, the masses at the early phases grow almost linearly with time, an evolution close to that found from our simulations for halos in the merger catalog. A similar exponent was also found by Toth \\& Ostriker (1992), i.e., $M \\propto t^{1.3}$, who considered the halo mass evolution in the  hierarchical picture, using a standard CDM power spectrum.  The algorithm used in this work is adequate to follow the evolution of the more massive halo progenitor either in the accretion or merger case. Therefore, our procedure is more suitable to test the evolution of the angular momentum, for example according to the random walk capture scenario (Vitvitska et al. 2002). It should be emphasized that unlike the present work, in that model no distinction between accretion and merger is made and essentially the effects of a cumulative capture of satellites by the progenitor is considered.  When considered individually,  halos show  erratic variations of the angular momentum strongly correlated (positively or negatively) with mass variations (Figs. 3). This correlation reinforces the fact that even after the first shell-crossing the angular momentum varies, and the reason is mainly the transfer of orbital angular momentum to the halo spin due to accretion and merger. Previously, a different evolution for individual halos in the simulations was found: in general the angular momentum decays after the first shell crossing (Sugerman, Summers \\& Kamionkowski 2000). It was then understood that such a trend is probably due to the redistribution of the angular momentum to bound particles outside the overdensity cut-off generated in the identification of the objects included in their catalog. Thus, unless all bound particles of a halo are identified, the apparent decay of the angular momentum at late times can arise as a numerical artifact.  The angular momentum distribution calculated at different redshifts indicate that the median and the mean value increase with time. Two regimes again are identified in both samples, with a transition at $z \\sim 1.5-1.8$, consistent with the behavior observed in the mass growth.  The two regimes of growth are again well represented by power laws (eq. 11) and are quite different from the linear variation expected from TTT. The present mass weighted distribution of the angular momentum (Fig. 5) and the distribution of the spin parameter (Fig.\\ref{fig6}) indicate that halos which have undergone merger episodes gain more angular momentum, confirming some previous investigations (Gardner 2001).  The distinction  between the two cases (accretion and merger) allows us to establish an important characteristic of the temporal behavior of the spin parameter. The median, the mean and $\\lambda_0$ increase with time for halos which have undergone major merger episodes, while a decrease is obtained for the accretion catalog. We have shown that such a behavior can be understood taking into account the different scaling laws followed by the dynamical variables defining the spin parameter. When no distinction is made, i.e., all halos are considered together, the statistical parameters of the $\\lambda$ distribution are practically independent of the redshift, apart from the mean which shows a slight increase with time. This result is consistent with studies in which no separation concerning the halo growth history is made (see, for instance, Lemson \\& Kauffmann 1999).  Concerning the possible dependence of the spin parameter on the halo mass (Barnes \\& Efstathiou 1987), we found that for $z < 3.5$ the angular momentum and total energy scale with mass approximately as $M^{5/3}$, i.e., the expected theoretical exponent. This results is valid both in $\\Lambda$CDM and in CDM cosmologies.  In this case, the spin parameter should be practically independent of the halo mass and of its previous growth history, since the same result was found for both catalogs.  We have also developed a simple model for disk formation based on considerations of the specific angular momentum distribution. The basic assumptions are the equality  between the specific angular momentum of dark matter and baryons, and rotational equilibrium after the collapse. However,  the distribution of the dark matter mass with the specific angular momentum  was calculated here in a cylindrical symmetry, centered along the spin axis. Two possible  progenitors of our own galactic halo were chosen based on their present mass. The first example concerns a halo evolved under accretion only, which is unable to acquire enough angular momentum to explain the observed value of the galactic disk. In the second example, the halo was taken from the merger catalog where it has captured a few satellites in the redshift interval $2 > z > 1$ and the required angular momentum is reached at $z \\sim 1.6$. The disk is built up gradually, the inner 5 kpc is formed around $z \\leq 5$, and the outskirts attain dimensions of about 30 kpc and 60\\% of the total mass at $z \\sim 1.6$, corresponding to an age of about 10.4 Gyr. Only 17\\% of the total mass was accreted in the last 5 Gyr. This evolutionary path is probably not unique, since other halos with different histories can also form disks with similar characteristics. However, the general picture agrees with the work by Helmi, White \\& Springel (2003), who also consider that the galactic  halo was formed gradually,  with more than 60\\% of the mass already present at around $11$ Gyr ago. The resulting projected mass density is exponential within the first 10 kpc, having a scale of length equal to 3.4 kpc (Fig.\\ref{fig12}).  The spin parameter of the disk is $\\lambda_b \\sim 0.49$, in agreement with the relation $\\lambda_{dm}/\\lambda_b \\approx f_b$, expected if the specific angular momenta of baryons and dark  matter are equal.  Finally, to conclude, it is important to notice that most disk galaxies  have masses in the range $10^{10} - 5 \\times 10^{11} M_{\\odot}$. Considering that the angular momentum of baryonic and dark matter scales as $M^{5/3}$, this translates into the fact that the halos associated to galaxies in that mass interval should be in the range $3 \\times 10^{66}$ up to $2 \\times 10^{69}  kg.m^2.s^{-1}$. If disks have ages comparable to our own galactic disk, i.e., have been formed around $z \\sim 1.6$, our simulations indicate that only 22\\% of the associated halos have acquired, at that redshift, the required angular momentum. If one neglects the eventual loss of spirals that have merged to form ellipticals, the estimated fraction is about a factor of three less than the present observed value. This is perhaps the ''true'' angular momentum problem !  \\vspace{1.0cm}  \\noindent {\\bf Acknowledgement}  \\noindent We thank the referee for the useful comments which have contributed to improve the text of this paper. S.\\ P.\\ acknowledges PhD fellowship from Universit\\'e de Nice Sophia-Antipolis (UNSA). R.\\ M.\\ is supported by Marie curie fellowship HPMF-CT 2002-01532."
    },
    "0311/astro-ph0311243_arXiv.txt": {
        "abstract": "{ From medium-resolution radio images, DR 16 was suspected to be a large cometary nebula. To test this suggestion we obtained a higher resolution (15\\arcsec) VLA continuum map. We also analyzed data from the Canadian Galactic Plane Survey in continuum, \\ion{H}{i} line, and IR.  These data were supplemented by published near-infrared (J, H, K) stellar photometric results and MSX 8.28 $\\mu $m data. We suggest that DR 16 is the diffuse \\ion{H}{ii} region of an ongoing star formation site at a distance of about 3 kpc. The complicated radio picture arises from the superposition of diffuse \\ion{H}{ii} with the remains of a giant bipolar outflow. The outflow was generated by a probable Herbig AeBe star, and the lobes are the remnants of its working surfaces. Additional ring-like features are discussed.  DR 16 is part of a larger volume of space in the local spiral arm where star formation is an ongoing process. \\vspace{-0.3cm} ",
        "introduction": "In the direction of Cygnus X our line of sight is tangential to the local spiral arm, and the superposition of many objects makes the region one of the most complex in the sky.  Although the radio emission in \\object{Cyg X} is mainly thermal in origin, individual sources reveal a surprisingly rich spectrum of properties, with objects in different stages of evolution occupying volumes of space with very different star formation histories.  Uncovering the nature of individual objects and establishing their distances demands extensive and thorough observations. \\begin{figure}[h] \\resizebox{\\hsize}{!}{\\includegraphics{0355.f1.eps}} \\caption{408 MHz map of the DR 16 area extracted and enlarged from the map in Paper XVIII. The coordinates are B1950. The contours are in steps of 40 K $T_\\mathrm{b}$ beginning at 140 K.} \\label{kar74} \\end{figure} The most recent radio continuum surveys (Wendker et al. \\cite{cx18}, Paper XVIII) showed that a resolution of about 1\\arcmin\\ begins to resolve many sources, so that observations of higher resolution will be effective in imaging their structure and uncovering their true nature. We have already studied the peculiar supernova remnant \\object{G76.9+1.0} (Landecker et al. \\cite{tll1}) and the ring \\object{G79.29+0.46} around a suspected LBV (Higgs et al. \\cite{lah2}). \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics{0355.f2.eps}} \\caption{VLA map of DR 16 at 1515 MHz. The gray scale is non-linear in order to emphasize the large-scale structure of the fainter emission and brightness increases from black to white.} \\label{vlaganz} \\end{figure*} Another interesting feature is \\object{DR 16} (Downes \\& Rinehart \\cite{dr}) -- \\object{G82.0+2.3} -- and its surroundings.  Images at 408 MHz (Paper XVIII) and 4.8 GHz (Wendker \\cite{cx15}, Paper XV) showed a cometary nebula. Recombination line observations at 6 cm (Piepenbrink \\& Wendker \\cite{cx17}, Paper XVII) gave nearly the same radial velocity at three positions, suggesting that the feature is one physical object, with a pronounced head and an impressive flaring tail about $0.5^{\\circ}$ in length. Fig.~\\ref{kar74} clearly conveys this impression. Nevertheless, in Paper XVIII we labelled the feature as a ridge (\\object{CXR 10}).  Optically only the northern part of the tail is visible in H$\\alpha $ (filaments \\object{DWB 155} and \\object{DWB 157} in the catalogue of Dickel et al. (\\cite{cx5}, Paper V)). Visual extinction in front of the head rises to more than $5^\\mathrm{m}$ (Dickel \\& Wendker \\cite{cx11}, Paper XI), totally obscuring it.  The physical picture commonly associated with the term `cometary nebula' is an object moving through its surroundings and producing some sort of wake. Supersonic motion produces a bowshock object, and is not implied in this case. Since sub-sonic or mildly supersonic velocities (less than 10 km/s) do not produce large extensions, a tail of length $0.5^{\\circ}$ immediately implies a very nearby object, or an exceptionally large one.  In this paper we present a new VLA 20-cm image with resolution 15\\arcsec\\, which reveals that the apparent head of the comet (which we now label DR 16) has a dumbbell structure. Furthermore, little evidence remains of a connection between the head and the other parts of the ridge (which we now refer to as the tail). We also present new observations from the database of the Canadian Galactic Plane Survey (the CGPS; see Taylor et al. \\cite{art}) and collect other published observations for a comprehensive study of the source. In Sect. 2 to 5 we describe new observations at radio and infrared wavelengths and present the results. A distance is suggested in Sect. 6. In Sect. 7 to 9 we discuss the source properties and suggest a source model and a source history. ",
        "conclusions": "Part of the following conclusions are based on two review papers, one on UCHIIR by Churchwell (\\cite{edrev}), and one on Herbig Ae/Be stars by Waters \\& Waelkens (\\cite{haebe}).  Massive stars produce UCHIIR which are dispersed after the stars reach the main sequence. The oldest, most massive stars in the DR 16 cluster have nearly all reached this state and can now provide the ionization of the dispersed diffuse gas.  The younger ones are far enough from the main sequence that their small UCHIIR are not observable.  As the lines corresponding to stellar evolution and to decreasing circumstellar extinction are nearly parallel in the CMD we have available (see Fig.~\\ref{cmd}), we cannot provide any more details of this picture.  The lower mass limit for young stars to which this picture applies is usually put around 10 $M_{\\sun }$ (or spectral type B0.5). Less massive pre-main-sequence stars are usually called HAeBe stars and develop more slowly towards the main sequence, but can shine through their natal cloud much sooner since their accretion disks are less massive. Probably as a consequence of this, outflows are less vigorous and hardly observed at stages near the main sequence.  Images of the thermal radio emission portray the average electron density along the line of sight, and this can smooth out small-scale structures. Remains of old UCHIIR and of bipolar outflows, which are probably ubiquitous in such a region, are thus not visible in our radio maps.  Our suggestion that the central star of the source 1515/5 is a HAeBe star and that the \\ion{H}{ii} structures 1515/3 and 1515/6 are the working surfaces of the lobes of a bipolar outflow impinging on neutral and/or molecular material implies the following specific evolutionary status.  The star is close enough to the main sequence that typical HAeBe properties are weak (this statement can be tested by spectroscopy).  Its mass is large enough that it also has properties usually ascribed to the massive star picture i.e. an earlier strong outflow which has now ceased.  On the other hand, the outflow is sufficiently recent that the density fluctuations it produced are still visible, although its ionization may by now be maintained by the diffuse radiation rather than by the originating star. This transition status makes the star quite interesting.  On the grander scale, DR 16 can now be added to the list of sources in Cyg X which comprise an \\ion{H}{ii} region and a small stellar cluster and lie at about a distance of 3 kpc. This is beyond the zone of influence of the Cyg OB2 association, often regarded as all dominating.  The Cyg OB2 association was suggested to be a young globular cluster rather than a classical open cluster (Kn\\\"odlseder \\cite{knoe}).  Distances around 3 kpc have been suggested for \\object{DR 21} and \\object{DR 23} (Papers XVII and XI, respectively). The projected distance of DR 16 from these is only 100 pc or so, which may indicate that this part of the local spiral arm contains an extended complex of small star forming regions. DR 16 itself is quite bright, about one third of the Orion nebula.  Compared to DR 15, itself a bright and conspicuous source in Cyg X, it is about 15 times brighter. On the other hand, it has only about 1/7th of the radio power of DR 21, usually regarded as a prominent cluster of just forming stars."
    },
    "0311/astro-ph0311075_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{intro}   Type Ia supernovae (SNe Ia) are believed to be the destructive thermonuclear explosions of white dwarfs. Their light curves are during the first couple of years powered by the radioactive decay of freshly synthesized $^{56}$Ni, releasing $\\gamma$-rays and positrons in the ejecta. At phases later than about 200 days, virtually all gamma-rays escape freely from the ejecta, and the luminosity is then provided by the kinetic energy deposited by the positrons.  Whether or not the positrons are able to slip out of the ejecta depends on the strength and geometry of the magnetic field (e.g., \\cite{ps98,m99,m01}). While a weak and radially combed magnetic field might allow positron escape, thus providing a steep light curve, a strong, tangled magnetic field would efficiently trap all the positrons, and drive the light curve to the radioactive decay rate.  As a first step to investigate whether or not observations of the positron phase can establish conclusions about positron escape, we have conducted a photometric study at late phases of the SN Ia 2000cx, and modeled its light curve in detail. Here we present some preliminary results from that study. The final analysis will be reported elsewhere. ",
        "conclusions": ""
    },
    "0311/astro-ph0311596_arXiv.txt": {
        "abstract": "We present an empirical infrared spectral library of medium resolution (R$\\approx$2000-3000) $H$ ($1.6\\,\\mu$m) and $K$ ($2.2\\,\\mu$m) band spectra of 218 red stars, spanning a range of [Fe/H] from $\\sim$$-$2.2 to $\\sim$+0.3. The sample includes Galactic disk stars, bulge stars from Baade's window, and red giants from Galactic globular clusters. We report the values of 19 indices covering 12 spectral features measured from the spectra in the library. Finally, we derive calibrations to estimate the effective temperature, and diagnostic relationships to determine the luminosity classes of individual stars from near-infrared spectra.  This paper is part of a larger effort aimed at building a near-IR spectral library to be incorporated in population synthesis models, as well as, at testing synthetic stellar spectra. ",
        "introduction": "Over the last two decades, evolutionary synthesis modeling has become a common tool to study unresolved stellar populations of galaxies in the optical and UV passbands \\citep{bru93,wor94a,vaz99,lei99}. The heavy obscuration in starburst galaxies -- often with A$_V$$\\geq$5 mag~ \\citep{eng97} -- requires an expansion of these models into the near infrared (IR) domain, because the extinction there is reduced tenfold in comparison with the optical region: A$_K$/A$_V$=0.112 \\citep{rie85}. Therefore, it is not surprising that many studies of embedded stellar populations in galaxies have been conducted in the near-IR \\citep{rie80,rie93}.  IR evolutionary synthesis models based on synthetic spectra \\citep{ori93,kur94} generally have difficulties reproducing broad-band colors because of the complicated opacity calculations in the near-IR. Empirical libraries based predominantly on bright and nearby Milky Way stars \\citep{kle86,lan92,wal96,wal97,mey98,wal00} are limited to near solar metallicity. At the same time, realistic galaxy modeling requires metallicities ranging from [Fe/H]=$-$1.76 for the most metal-poor galaxy known I\\,Zw\\,18 \\citep{alo03}, to supersolar values, measured for at least for some giant elliptical galaxies \\citep{cas96}.  The properties of the most prominent near-IR spectral libraries available in the literature are summarized in Table\\,\\ref{TblLibRev}. For a review of the earlier work see \\citet{mer79}. This summary shows that despite the sizable quantity of available stellar spectra, up until now there was no uniform near-IR dataset with high signal-to-noise and resolution, covering the entire range of spectral classes, luminosity, and metallicity necessary to carry out evolutionary population synthesis in the near-IR. The deficit of some types of stars such as metal-poor super giants is understandable in the context of the Milky Way star formation history, but the lack of many other types is rectifiable. We concentrate on metal features that have equivalent widths in a typical starburst galaxy larger than 1\\AA\\, -- as they can be measured reliably in  spectra with signal-to-noise ratios of $\\sim$30-50 \\citep{eng97}.  A second application of our library, albeit no less important, is the analysis of individual stars hidden behind A$_V$$\\geq$10 mag of visual extinction. A typical case for such a population is represented by the Arches cluster member starss in the Galactic Center region \\citep{nag93}.  Here we describe the observations and the sample of an empirical near-IR stellar library that is designed to meet the requirements of population synthesis. The main advantage in comparison with previous work is the expanded metallicity coverage. We also present a few diagnostics methods to derive parameters of individual stars. The evolutionary population model will be published in a subsequent paper. The next section describes the observations and the data reduction technique. Sec.\\,3 summarizes the sample selection. Spectral indices are defined in Sec.\\,4, and some diagnostic applications for parameters of individual stars are considered in Sec.\\,5. In Sec.\\,6. we give the summary. ",
        "conclusions": "We have assembled a library of moderately high-resolution ($\\approx$\\,2000-3000) $H$ (1.6\\,$\\mu$m) and $K$ (2.2\\,$\\mu$m) spectra of 218 red stars, mostly supergiants and giants. The majority of these stars have well-known photospheric parameters from high-resolution optical spectroscopy. These stars dominate the near-IR emission in both starburst and elliptical galaxies. Our library covers a range of effective temperatures, and metal abundances from [Fe/H]$\\sim$$-$2.2 to +0.3. This library will offer a unique opportunity to study directly the most obscured stellar populations in starburst galaxies, as well as in the center of the Milky Way.  Although the main motivation behind the creation of this library is the study of unresolved extragalactic stellar populations, the obtained spectra can be used to derive parameters of individual stars. We calibrated some line ratios as indicators of the stellar effective temperature. Finally, we demonstrated how some diagnostic relationships can distinguish (super)giants from dwarf stars."
    },
    "0311/astro-ph0311305_arXiv.txt": {
        "abstract": "We present a new method for multi-component power spectra estimation in multi-frequency observations of the CMB. Our method is based on matching the cross and auto power spectra of observation maps to their expected values. All the component power spectra are estimated, as well as their mixing matrix. Noise power spectra are also estimated. The method has been applied to full-sky Planck simulations containing five components and white noise. The beam smoothing effect is taken into account. ",
        "introduction": "\\label{}  Precise measurements of Cosmic Microwave Background (CMB) angular power spectrum, is one of the main objectives of modern Cosmology. Several experiments, such as ACBAR \\cite{ACBAR}, Archeops\\cite{Archeops}, Boomerang \\cite{Boomerang}, CBI \\cite{CBI}, DASI \\cite{DASI}, MAXIMA \\cite{MAXIMA}, VSA \\cite{VSA} and more recently WMAP \\cite{WMAP} already give precise results over a large range of angular scales. An unbiased estimate of the CMB angular power spectrum requires important efforts in data processing because of the presence of several sources of contamination in the observations. Foregrounds, in particular (Galactic dust, synchrotron, extragalactic point sources and Sunyaev-Zel'dovich effects), will be the main contaminants for CMB measurements, such as that of the future Planck mission. Fortunately, the avaibility of several detectors operating at several frequencies, ranging from 30 GHz to 857 GHz for Planck, permit the separation of most astrophysical components due to their different emission laws.\\\\ Starting from the different observation maps, current methods for CMB power spectrum estimation consist first in making the cleanest CMB anisotropy map possible, using component separation techniques, like Wiener filtering \\cite{Wiener,wiener96tegmark,allskywienerPrunet} or MEM \\cite{MEM,allskyMEM}. Spectral estimation is then performed on the individual component maps. This approach is not fully satisfactory for two reasons: First, component separation requires prior knowledge of the electromagnetic spectra of the components, which are not necessarily very well known. Second, after component separation, in the subsequent spectral estimation step, it is necessary to remove the power spectrum of the residual noise in the extracted CMB map. A small error in the noise evaluation introduces a bias in CMB power spectrum estimation.\\\\ A new approach has been considered in paper \\cite{MDMC}. This approach is based on the analysis of cross- and auto-power spectra of different observation maps by different detectors. Considering the simple case where only CMB anisotropies are present with noise, the cross-power spectra give a direct measurement of the CMB power spectrum, assuming that the noise is uncorrelated between detectors. In this limited context, the technique has been used by the WMAP team. Our approach deals simultanously with multi-component data. By jointly analysing the different observation maps, we estimate the power spectra of all the components, without making any priors on their emission laws. Thus, the contribution of each component to each detector, given by the mixing matrix $A$ (e.g. (1)), is estimated directly from the data. The noise power spectra are also estimated. The method is based on likelihood maximization in the Whittle approximation assuming that the observations are a linear mixing of independant components and independant noise. It can be viewed as a blind multi-detector multi-component (MDMC) spectral matching, in which the spectral diversity of the various components is used. The method has been implemented on harmonic coefficients of all sky maps. We account for the finite spatial resolution of the detectors. In this paper, we describe the basics of the spectral matching method and we present its application on full-sky simulated Planck observations. ",
        "conclusions": "\\label{}  We have presented our blind spectral matching method for all-sky multi-detector CMB observations. The main objective is to measure the power spectra of all the components, including the CMB, and their contribution levels at each observation frequency. The method exploits in particular the cross-power spectra of observation maps, giving an unbiased estimate of component power spectra. The method has been applied on full-sky Planck simulations containing five components and white noise. The power spectrum of the CMB is accurately estimated up to $l \\simeq 2000$ in bins of size $\\Delta l = 10$.   {\\bf{Acknowledgments.}} The author would like to thank the Planck collaboration and in particular M.A.J. Ashdown, V. Stolyarov and R. Kneissl for the full sky simulated maps, and also J. Bartlett for a critical reading of the manuscript. The HEALPix package (see http://www.eso.org/ science/healpix/) was used for the spherical harmonics decomposition of the input maps."
    },
    "0311/astro-ph0311419_arXiv.txt": {
        "abstract": "Radio and X-ray studies of young supernovae probe the interaction between the supernova shock waves and the surrounding medium and give clues to the nature and past of the progenitor star. Here we discuss the early emission from type Ic SN 2002ap and argue that repeated Compton boosting of optical photons by hot electrons presents the most natural explanation of the prompt X-ray emission. We describe the radio spectrum of another type Ic SN 2003dh (GRB030329) obtained with combined GMRT and VLA data. We report on the low frequency radio monitoring of SN 1995N and our objectives of distinguishing between competing models of X-ray emission from this SN and the nature of its progenitor by X-ray spectroscopy. Radio studies on SN 2001gd, SN 2001ig and SN 2002hh are mentioned. ",
        "introduction": "\\label{sec:introduction}  The association of long duration Gamma Ray Bursts with core collapse supernovae is by now established with SN 2003dh (GRB030329) \\cite{Hjo03}. Observational evidence supports the viewpoint that GRBs are energetic explosions like supernovae occuring in star formation regions and a large fraction of their energy is directed in relativisitic jets. Supernovae are known to be explosions of massive (and intermediate mass) stars, although the nature of the progenitor stars for varied supernova types remains an open subject.  A supernova explosion from the core collapse of a massive star drives a powerful shock wave into the circumstellar medium (CSM) of the progenitor. The shock wave with a speed approximately 1000 times larger than the speed of the progenitor star's wind quickly probes the circumstellar medium established by the wind lost over many thousands of years before the explosion. Relativistic electrons and magnetic fields in the interaction region give rise to nonthermal radio emission. The interaction of a supernova shock with its surrounding medium produces X-ray emission probing the environment. Multiwaveband studies from radio and X-ray bands thus provide a handle on the past history of the parent star and the nature of the interaction region. Fingerprinting the exploded ejecta composition through X-ray spectroscopy \\cite{Zim03} can give clues to the mass of the exploding star (e.g. in SN 1998S \\cite{Poo02}).  In this paper, we discuss several supernovae: SN 2002ap -- a `hypernova' of type Ic associated with GRBs {\\it that had no GRB counterpart}; another type Ic `hypernova' SN 2003dh which was associated with GRB030329; and several other supernovae such as: SN 2001gd, SN 2001ig and SN 2002hh. We also mention how upcoming X-ray spectroscopic studies of SN 1995N will help discriminate between the sites of observed X-ray emission. ",
        "conclusions": ""
    },
    "0311/astro-ph0311133_arXiv.txt": {
        "abstract": "We present the first linear multicolor polarization observations for a sample of 35 stars in the direction of the Galactic cluster NGC 6231. We have found a complex pattern in the angles of the polarimetric vectors. Near the core of this cluster the structure shows a semi-circular pattern that we have interpreted as a re-orientation of the dust particles, showing the morphology of the magnetic field. We propose that a supernova event has occurred some time ago and produced a shock on the local ISM. We discuss in this paper independent confirmations of this event, both from the studies on the diffuse interstellar absorptions \\citep{cro01} and the results of the pre-main sequence stars (PMS) given by \\cite{sung98}. We also show that a supernova is supported by the evolutionary status of the cluster. ",
        "introduction": "The polarimetric techniques are a very useful tool to find important information \\break ($P_{\\lambda max}$,$\\lambda_{max}$, magnetic field direction, etc) from the dust located in front of a star (or a luminous object), or from dust located inside a stellar cluster. Young open clusters are very good candidates to carry out polarimetric observations, because previous photometric and spectroscopic studies of these clusters give detailed information of the main sequence, so we can characterize the physical parameters of the member stars (age, distance, memberships, etc), in order to study the extinction of the dust in the direction to the cluster and within the cluster. For example, from the study of the $P_{\\lambda max}$ versus the $E_{B-V}$ of each star of a cluster it is possible to find the polarimetric components of the extinction produced by Galactic dust located between the Sun and the cluster (e.g. Tr 27, \\citet{fein00}). By subtracting the effect of this dust on the line of sight from the data it is possible to study the component associated with the internal extinction of the cluster. In addition to this valuable information, stars with intrinsic polarimetric properties can be isolated. The angle of the polarization vector is an estimator of the direction of the Galactic magnetic field. By observing members and nonmembers stars (foreground and background) the direction of the magnetic field towards the cluster at several distances can be inferred.  In a typical open cluster the polarimetric data show, for ``normal'' stars (that is, not considering stars with intrinsic polarization), a distribution of angles that can be described by a mean fixed value which characterizes the cluster. The polarization angles of the cluster members show some scatter ($\\sim 10 ^{\\circ}$ or less) over this mean angle, mainly because of density inhomogeneities in the dust along the line of sight to the cluster, inhomogeneities of the dust inside the cluster, to the intrinsic polarimetric properties of the stars and also due to more complex magnetic field structures. For example, in Tr 27 the polarimetric data show a mean angle of $\\theta\\sim 30^{\\circ}$ for the cluster with a FWHM of $\\sim 10^{\\circ}$ \\citep{fein00}, in Ara OB1 the polarimetric observations give a FWHM of $\\sim 9^{\\circ}.9$ for NGC 6167, $\\sim 13^{\\circ}.1$ for NGC 6193 and $\\sim 9^{\\circ}.5$ for NGC 6204 \\citep{wald99}. In the case of Stock 16 \\citep{fein03} the FWHM is  6$^{\\circ}$.7, which seems to be slightly smaller than the other clusters.   NGC 6231 (l= 345$^\\circ$.5 , b=1$^\\circ$.2) is known in the literature to be one of the youngest open clusters in the Galaxy and one of the brightest in the Southern Hemisphere. The cluster is embedded in an extended system of O and early-B stars, namely the very young association Sco OB1, that contains the ring-shaped HII region IC4678, centered on NGC 6231. Usually, NGC 6231 is considered as the nucleus of Sco OB1 \\citep{per91}. Its age  determinations range between $3 \\times  10^{6}$ years \\citep{mer86} and $ 6.9\\times 10^6 $ years \\citep{per91}. It is known as one of the richest open clusters, containing WR stars, $\\beta$ Cepheid variables and  more than 100 O and B stars \\citep{sho83}. According to photometric CCD observations of \\citet{sung98} and of \\citet{bau99} there is  evidence of a population of pre-main sequence stars (PMS). The distance obtained by deep CCD photometry by \\cite{bau99} locates the cluster NGC 6231 at $1990 \\pm 200$ pc from the Sun.      The average visual extinction toward the cluster is about $A_{V} \\sim 1.4 $ \\citep{balo95,bau99}. According to \\citet{rab97} in the northern part of the cluster the absorption is constant, but it increases from the central part towards the South. A similar trend, but less pronounced, is observed from West to East in the southern part. Part of the absorption suffered by the light from the cluster stars is due to matter within the distance interval 100-1300 pc from the Sun \\citep{per91}, but the field shows a relatively good transparency between about 400 and 1000 pc at a constant value averaging of $A_{V}=0.75$ mag (Raboud et al. 1997). ",
        "conclusions": "\\subsection{Serkowski law} To analyze the data, the polarimetric observations were fitted using  the Serkowski's law of interstellar polarization  \\citep{serk75}. This is:  $P_{\\lambda}/P_{\\lambda max}=e^{-Kln^2(\\lambda_{max}/\\lambda)} \\  \\  \\ \\ [1]$   If the polarization is produced by aligned interstellar dust particles, the observed data (in terms of wavelength, UBVRI) will follow [1], where each star is characterized by  $P_{\\lambda max}$ and  $\\lambda_{max}$.  Adopting $K = 1.66  \\lambda_{max}+ 0.01$  \\citep{whittet92}, we have fitted our observations and computed the $\\sigma_{1}$ parameter (the unit weight error of the fit) in order to quantify the departure of our data from the ``theoretical curve'' of  Serkowski's law. A $\\sigma_{1}$ value larger than 1.5   is considered as an indication of the presence of a component of intrinsic stellar polarization. Another criterion of intrinsic stellar polarization is to compute the dispersion of the position angle for each star normalized by the average of the position angle errors ($\\bar{\\epsilon}$, Marraco et al. 1993).  The  $\\lambda_{max}$ values can also be used to test the origin of the polarization. In fact, since the average value  of $\\lambda_{max}$ for the interstellar medium is 0.545 $\\mu m$ \\citep{serk75}, objects showing  $\\lambda_{max}$ rather lower than this value are also candidates for having an intrinsic component of polarization \\citep{orsa98}. The values we have obtained for $P_{\\lambda_{max}}$, the $\\sigma_{1}$ parameter, $\\lambda_{max}$, and $\\bar{\\epsilon}$ together with the identification of stars are listed in Table 3.  Figure 2 shows the observed $P_{\\lambda}$ and $\\theta_{\\lambda}$ vs ${\\lambda}$ of those stars which are candidates for having an intrinsic component of polarization (stars \\#70,73 and 254). For comparison purposes, the best Serkowski's law fit has been plotted as a continuous line. Star \\#220 also shows an intrinsic component of polarization on Table 2, but it is already known as polarimetric variable. First \\citet{lun} and then \\citet{st87} have made an extensive polarimetric study of this double-line binary (WC7 + O5-8) system.   \\subsection{Polarization efficiency}   Fig. 3 shows the plot of  $P_{\\lambda max}$ vs.  $E_{B-V}$ \\citep{sung98,rab97,bau99} for each of the three groups. This plot is very useful for study the polarization efficiency, defined as the ratio $P_{\\lambda max}/E_{B-V}$ which indicates how much polarization is obtained for a certain amount of extinction.   To understand the behavior of the polarimetry and the extinction a detailed discussion of the distribution of the dust on the line of sight to the cluster is necessary. Neckel \\& Klare (1980) have studied the extinction values of the ISM and distances in our Galaxy using more that 11.000 stars. Their Figure 6j(211) (which is the plot in the direction of NGC 6231) shows two steeps in the absorption increase: the first one is very near to the Sun and associated with the Lupus cloud \\citep{cro00} at 150 pcs ($A_{V}=0.8$) and the second step is about 700 to 1000 pc and is also detected in the diffuse interstellar lines \\citep{cro01}. This second step of absorption is responsible for $A_{V}\\sim0.6$. \\citet{bau99} have reconstructed Neckel \\& Klare's plot with new CCD data, and found a more complex structure obtaining a similar result for the distribution of dust between the Sun and the cluster. Nevertheless, Figure 7 of \\citet{bau99} shows that NGC 6231 has also intracluster dust.   Most of the studies show that the cluster has  differential reddening. For example, \\citet{sung98} have made a bidimensional map of the distribution of $E_{B-V}$ over the NGC 6231 cluster (see Fig. 5 of their paper). They show that the core of the cluster is less reddened and seems to lie in a hole in the reddening material ($E_{B-V} \\sim 0.45$). To the South they find that the extinction increases to $E_{B-V} \\sim 0.65$. The differential reddening of the cluster is extended discussed in \\citet{rab97} and \\citet{bau99}.      Fig. 3 is the plot of  $P_{\\lambda max}$ vs $E_{B-V}$ for each group of stars. The upper panel are stars of the first group, the middle panel is for second group of stars and the bottom panel are data of the third group of stars. The solid line in the first panel of Fig. 3 is the empirical upper limit relation for the polarization efficiency by the interstellar polarization $P_{\\lambda max} = R\\ A_{v} = 9 E_{B-V}$ (for normal dust, $R=3.2$) given by Serkowski et al. (1975). The dashed line ($P_{\\lambda max}/ E_{B-V} = 5 $) represents the observed normal efficiency of the polarizing properties of the dust (Serkowski et al., 1975).   Nevertheless, the $E_{B-V}$ data were collected from different sources and probably some of these values have large errors, some qualitative properties emerge from the figure. It is noticiable that all the groups have a polarizing efficiency lower than the  $P_{\\lambda max}/E_{B-V} \\sim 5$ expected. Because for an extinction of $E_{B-V}=0.43$ \\citep{bau99} and the average polarization efficiency, the typical value should be around $P_{\\lambda max} \\sim 2\\%$. This value is  much higher  than the one observed in the cluster (see Table 3). Therefore, in all the groups we have an important depolarization and this efect is notorious for the stars of the semi-circular structure (group 2) that have a much lower polarization in comparison with the first and third groups (see Fig. 3 middle panel). There are not stars with polarization greater than 0.95\\% in this group.           There is a question that remains to be answered. Why we can see this semicircular structure behind two layers of dust which orientate the polarization vectors in different directions. It is pretty obvious that just having only one dust layer on the line of sight would compose and change the direction of the polarization vectors, washing out this semicircular structure. But on the other hand, the observed semicircular structure seems quite perfect. As we have discussed in the previous paragraph, in the particular case of NGC 6231 there are two layer of dust, a first one at Lupus (hereafter the first component of extinction) and a second one in front of the cluster locate at 1 kpc, \\citep{rab97,bau99,nec80}, hereafter the second component. It is very important to notice that the two components are orthogonal and one of them can compensate or diminish the other. The first component is dominated by dust in the Galactic plane, which has an angle of polarization of $\\sim 45$ degrees, the second component seems to be at right angles to  the  Galactic plane at $\\sim$ 135 degrees. Fig. 3 shows a huge depolarization for the second group of stars. We think that in the northern part the first component dominates the observed polarization, while in the South the second component is the strong one. But in the core, where the semicircular structure lays, both components (first and second) cancel each other, therefore making a polarization ``window'' where we can observe inside the cluster. This interpretation also explains the depolarization of the stars locate in the semi-circular structure (second group).  We tested this interpretation of the extinction toward NGC 6231, by analyzing the polarization of the nonmembers stars. For example, in the first group, \\#102 is considered a F0III \\citep{per91} star so it must be located at about 379 pc from the Sun and star \\# 189 is a B9-8V \\citep{rab97} being located at 1 kpc. Both stars are behind the Lupus cloud and have  polarization angles similar to the direction of the Galactic disk. Therefore, we can deduce that the Lupus cloud is responsible for the  polarization of the light in the direction of the Galactic disk. This cloud is reddening the stars behind by $E_{B-V}=0.25$ in average \\citep{nec80}. For larger distances \\citet{rab97} have found that the interstellar medium shows a relatively good transparency between 400 and 1000 pc.  Star \\# 105 is also a nonmember star but its polarization (and location in the cluster) is related with the third group of stars.  Assuming that this star is a main sequence object reddened by normal dust it is possible to found an unique solution from the color-color diagram that gives an $E_{B-V}=0.52$. This value is larger than the estimate for the Lupus cloud and shows that this object is located behind another layer of dust which is polarizing the stellar light at right angles to the direction of the Galactic disk. Therefore, this star is probably located behind the second cloud at a distance of more 1 kpc from the Sun.          \\subsection{ The semicircular polarization pattern }  As despicted in Fig. 1, the polarization vectors of the stars belonging to our second group (at the core of the cluster) displayed a  semicircular pattern. This is very unlikely to be the result of a distribution of polarization angles at random, because the pattern is well organized. We think we are in the presence of an energetic phenomenon that has changed the magnetic field orientation, resulting in a realignment of the dust grains, which manifests itself in a rotation of the polarization vectors. We propose that a supernova event has occurred some time ago and produced a shock over the ISM. It has pushed the gas and dust (and the associated magnetic field), causing the dust grains to re-oriented to the new direction of the magnetic field. We will discuss in the next paragraphs the relevant questions associated with the proposed supernova event.  Where did this phenomenon occur ?.  Tracing  lines perpendicular to the most conspicuous polarimetric vectors it is possible to obtain an approximate  center at the crossing point where these lines overlap each other. This center seems to be at $\\alpha =16^{h}\\ 54^{m}\\ 22^{s}$, $\\delta = -41^{\\circ} \\ 46^{'}\\ 30^{\"}  $, J2000.0 (Fig. 4). This semi-circle is very regular and the vectors converge clearly to a center. We think this is a solid argument about the existence of center of this pattern, which means that the energetic phenomenon was produced at a certain fixed location in the cluster. This supports  the idea that a supernova exploded in the near past of the cluster, located at the center of the crossing lines as shown in Fig. 4.   Could  a supernova event have happened in NGC 6231 ? The earliest star on the main sequence, not considering binaries, is at present of spectral type O9 (\\#292, HD 326329), and there is evidence of more massive stars that had evolve out of the main sequence \\citep{bau99}. Therefore, it is very probable that a star more massive than HD 326329 had evolved to a supernova stage in the past. Also, \\cite{bau99} have computed the initial mass function of NGC 6231 (hereafter IMF) with binary correction and they have obtained a slope of $x=1.14\\pm0.10$, for $x=log(dN/\\delta M)/log(M)$. This slope of the distribution allows the existence of a former star of about 70 $M_{o}$ (an O3-O4) without any change in the parameters of the observed IMF. This means that one or few (because we are dealing with small number statistics) very massive stars could have lived in the past with no contradiction with the actual IMF. So, the evolutionary status of the cluster population in NGC 6231 indicates that some massive star have evolved and probably became supernovae.              Is there  other evidence on the ISM of this supernova event ? Crawford and collaborators (e.g. \\cite{cro89a}, \\cite{cro89b,cro92,cro01}) have made a very detailed study of the interstellar lines (NaI, CaII, KI) towards Sco OB1. Although these studies are related to the whole Sco OB1 association, these series of papers include in the analysis some of the brightest stars of NGC 6231. They have found several components and classified them into three major groups \\citep{cro01}.  The first one consists of several components with radial velocities in the range $-20 \\leq v_{helio} \\leq 6\\ km\\ s^{-1}$ which are consistent with the galactic rotation along the 1900 pc to Sco OB1. All the clouds in this component are presumably foreground, and  unrelated with Sco OB1. For example components in the velocity range $-6 \\leq v_{helio} \\leq 6\\ km\\ s^{-1}$ are probably associated with the Lupus molecular cloud complex at a distance of 150 pc.  The second group of interstellar lines found by Crawford (2001, and references therein) is the ``shell component'' which has velocities more negative than -20 $km\\ s^{-1}$ (-22 $km\\ s^{-1}$ is the velocity expected for the galactic rotation model at a distance of 1900 pc) and these clouds of interstellar matter seem to be related with the shells surrounding the Sco OB1 association. In a previous work \\cite{cro89a} found a very low NaI/CaII relation in this component. This finding was explained by the removal of the Ca atoms from the grain surfaces by shock velocities of $\\geq 20\\ km\\ s^{-1}$. New data  \\citep{cro01} indicate that the observed shell component has an upper limit for the temperature of the order of $T=450^{\\circ}K$ (as observed in the spectrum of HD 152235, a member of NGC 6231), which is low for a post-shock temperature ($T\\geq 10^{4}$). But as \\cite{cro01} noted, the cooling time of a shock for the regular interestellar medium ($v_{s}\\ = 20\\ km\\ s^{-1}, n_{H}\\ = 10\\ cm^{-3}$) gives $t_{cool}=6300\\ yr$, meanwhile the depletion time scale for the Ca to become part of the interestellar grains is $t_{dep}\\ =\\ 1.6\\times 10^{7}\\ yr$. Therefore, shocks in the near past can account for the low NaI/CaII observed in the interstellar lines.   The third component of interestellar line has velocities more positive than those predicted by the galactic rotation models. In NGC 6231, HD 152249 exhibits an interstellar absorption component at a velocity of +19 $km\\ s^{-1}$. Note that HD 152249 is located in the core of the NGC 6231, and 1' to the South of the semicircular polarization pattern. This component has a low NaI/CaII ratios ($\\leq 0.4$), which (as described above) is an indication of shocks. This component could be located near the sun in the ``Upper-Centarus Lupus'' or it could be an expanding interestellar structure related to NGC 6231.  Both second and third interstellar components exhibit the low NaI/CaII ratio and in our opinion it is a strong support for shocks in the region and gives clues over a turbulent scenario in the near past of NGC 6231. These interstellar line components are probably located in the cluster, as discussed by Crawford in his series of papers, and were produced by shocks in the region. This point of view of Crawford is consistent with the interpretation of our polarimetric observations.       \\cite{sung98}, studying the Pre Main Sequence (PMS) stars in NGC 6231, have found a discontinuity in the number of stars near the MS band with ages between 30 Myr and 12 Myr. They have suggested that probably some PMS are not $H\\alpha$ emitters and they are not detected by the typical $R - H\\alpha$ vs. V-I diagram. They explained the lack of $H\\alpha$ emission suggesting that all the material surrounding the PMS stars, material that would normally produce strong H$\\alpha$ emission, could be swept away by stellar winds from massive stars, Wolf-Rayet stars or past supernova explosions. Their result is also compatible and a strong support to our polarimetric findings.    Other interesting result showing some clues to support the idea of energetic events on the cluster NGC 6231 came from \\citet{man96}. To explain the COS B diffuse emission of gamma-rays observations, \\citet{man96} propose that strong winds in the cluster interact with the interstellar medium creating the suitable condition for cosmic ray enhancement.      We have observed multicolor (UBVRI) linear polarization ($P$ and $\\theta$) for 35 of stars on the direction of NGC 6231. We have found a very unusual distribution of the angles of the polarization vectors. To the North of the cluster the polarization angles are dominated by the direction of the Galactic disk; to the South are perpendicular to this orientation. Close to the core of the cluster, the polarization vectors are well organized in a regular pattern, showing a semicircular morphology which  we interpreted as a re-orientation of the dust particles by the magnetic field.  We suppose that an energetic event (very probably a supernova) has occurred some time ago and produced a shock pattern on the local ISM. A supernova is compatible with the evolutionary status of the cluster. We discuss in  this paper independent confirmations of this event from  studies of the interstellar lines \\citep{cro01} to the photometric results of the PMS stars given by \\cite{sung98}."
    },
    "0311/hep-ph0311281_arXiv.txt": {
        "abstract": "The ultrahigh energy neutrino-nucleon cross sections are  computed taking into account different phenomenological implementations of the non-linear QCD dynamics. Based on the color dipole framework, the results for the saturation model supplemented by DGLAP evolution as well as for the BFKL formalism in the geometric scaling regime are presented. They are contrasted with recent calculations using NLO DGLAP  and unified BFKL-DGLAP formalisms. ",
        "introduction": "The inelastic interaction of neutrinos with nucleons is in general described by the QCD-improved parton model. In ultrahigh energy neutrino interactions one probes a kinematical region which is not accessible in the current collider experiments.  It is known so far that these reactions can be sensitive to the domain for the usual Bjorken variable $x\\simeq m_{W,Z}^2/E_{\\nu} \\sim 10^{-8}$ at $E_{\\nu}\\sim 10^{12}$ GeV and electroweak boson virtualities $Q^2\\sim m_{W,Z}^2\\approx 10^{3}$ GeV$^2$, where $m_{W,Z}$ are the boson masses. In contrast, in the current electron-proton collider DESY HERA the interactions are probed in the typical kinematical range $x\\simge 10^{-5}$ and $0 \\leq Q^2\\simle 10^3$ GeV$^2$ (the variables $x$ and photon virtuality $Q^2$ are correlated, with small $x$ data associated to low virtualities). There, the measurements of the deep inelastic $ep$ scattering structure functions have put important constraints on parton distributions, mainly sea-quarks and gluons, in the small $x$ region. In perturbative QCD these distributions are expected to grow as $x$ decreases, which has been strongly confirmed by HERA data. Therefore, in order to provide accurate predictions for the ultrahigh energy neutrino-nucleon cross sections, a precise extrapolation of the structure functions to the small $x$ and large $Q^2$ region probed in these interactions is needed.  In the usual QCD-improved parton model, the power growth of the parton distributions functions in very high neutrino energies lead consequently to a power increasing of the neutrino-nucleon total cross section, which in turn implies in a violation of unitarity at high energies. This issue is an outstanding theoretical challenge, which has produced intense work towards a complete understanding of the non-linear (higher twist) QCD corrections in the high parton density regime at fixed impact parameter (for a recent review, see  \\cite{armestoreview} and references therein). At present, one has reliable estimates for the kinematical domain where non-linear (saturation) effects should  play an important role. Namely, below a typical transverse  momentum scale $Q_{\\mathrm{sat}}^2(x)\\sim x^{-\\lambda}$, the so-called saturation scale \\cite{QSATDEF}, the growth of parton (gluon) density is tamed towards the black-disk limit of the target. This slow down of the gluon distribution function prevents the total cross section violating unitarity. An important prediction of the non-linear QCD approaches is the property of geometric scaling \\cite{SGK} below saturation scale, where the gluon density (and cross sections) scales with the saturation momentum, which grows with $1/x$ as a power. It has also been shown that such scaling survives even above the saturation regime \\cite{IANCUGEO,STASTOGEO,MUNIERWALLON}: the kinematical region has been estimated to be $Q_{\\mathrm{sat}}^2 (x)<Q^2<Q_{\\mathrm{sat}}^4 (x)/\\Lambda^2_{\\mathrm{QCD}}  $ relying on the BFKL equation \\cite{BFKL} in the presence of saturation \\cite{IANCUGEO}. This feature has important consequences in the underlying QCD dynamics at the nucleus-nucleus accelerators RHIC and LHC and should pose strong constraints to the cross sections beyond current accelerator energies. Concerning high energy neutrino cross sections, there are estimates constrained by HERA data based on DGLAP \\cite{DGLAP} and effects of BFKL \\cite{BFKL} linear evolutions \\cite{DKRS,GQRS,RSSSV,GRVNEUT,BASU,KMSNEUT} as well as implementations of non-linear QCD corrections \\cite{KUTAK}. The linear NLO DGLAP and unified BFKL-DGLAP results are found to be  consistent with each other presenting somewhat small  deviations \\cite{KMSNEUT}. An interesting analytical calculation using an approximate DGLAP solution with initial conditions from a soft non-perturbative model has been recently reported in Ref. \\cite{FIORE}. Saturation effects are shown to be quite small in recent DGLAP analysis \\cite{BASU}, whereas it was found a reduction in the cross section by a factor 2 at $E_{\\nu}\\sim 10^{12}$ GeV  considering the saturation model or screening effects via Balitsky-Kovchegov equation \\cite{KUTAK}. Recently, it has been claimed in Ref. \\cite{JAMAL} that the geometric scaling property can lead to an enhancement of the neutrino-nucleon total cross section by an order of magnitude compared to the leading twist cross section and would lead it further to its unitarization.   In this work one presents estimates for the ultrahigh energy neutrino-nucleon cross sections, focusing on the saturation models which encode the physics of non-linear QCD evolution in a phenomenological way. The main point is those models provide a suitable dynamical extrapolation to very high energies/virtualities in very contrast with usual DGLAP based approaches, which are not a priori valid in that regime and are dependent on quark and gluon distributions also unknown (no constraint from accelerator data) in that kinematical limit.  Based on the color dipole framework, first one presents the result for the saturation model \\cite{GBW} supplemented by QCD evolution \\cite{BGBK}. The original saturation model \\cite{GBW} was proposed for either low values of $Q^2 \\simle 150$ GeV$^2$ , whereas high energy neutrino interactions are sensitive to large virtualities $Q^2\\sim m_{W,Z}^2$. Therefore, it is necessary to  perform calculations with complete QCD evolution. Furthermore, one computes  estimates of the ultrahigh neutrino-nucleon cross section through an implementation of the dipole cross section considering BFKL formalism in the geometric scaling region \\cite{IIM}.In the last section, the numerical results for both cases are shown and contrasted with recent NLO DGLAP estimates and with unified BFKL-DGLAP formalism without saturation effects. ",
        "conclusions": "\\begin{figure}[t] \\begin{tabular}{cc} \\epsfig{file=nuccsatmod.eps,width=43mm,height=55mm} & \\epsfig{file=nuncsatmod.eps,width=43mm,height=55mm} \\\\ (a) & (b) \\end{tabular} \\caption{The $\\nu \\,N$  charged current (a) and  neutral current (b) as a function of neutrino energy. The curves correspond to the results for  the GBW saturation model (dot-dashed curves) and the BGBK model (solid lines). The LO and NLO DGLAP results are represented by dashed and long-dashed curves, respectively.} \\label{fig:1} \\end{figure}  In what follows, one  computes the ultrahigh energy CC and NC neutrino-nucleon cross sections using the different models containing saturation effects. In Fig. \\ref{fig:1} one presents the results  for the saturation model (GBW), the BGBK model and also   one includes a DGLAP  calculation. The result for the saturation model (dot-dashed lines)  has been recently computed in Ref. \\cite{KUTAK}, which we reproduce here. The introduction of the DGLAP evolution in the saturation model enhances the cross section by a factor 3 or even more at $E_{\\nu} \\sim 10^{12}$ GeV. This is shown by the BGBK curves (solid lines). It should be noticed that the power growth for saturation model is milder than the BGBK model and it gives an upper limit for the role played by saturation effects in the cross section. It should be emphasized this assertion concerns to a comparison between phenomenological saturation model and its  version including DGLAP evolution. Namely, other non-linear QCD approaches can give a larger amount of saturation in the same kinenatical regime and the phenomenological saturation model probably is inaccurate in this region. For sake of completeness, the DGLAP results (LO and NLO) are also presented. For this purpose, we follow the calculations in Ref. \\cite{BASU}, where a detailed analysis of the NLO corrections to neutrino-nucleon cross sections was performed. The BGBK model has an energy dependence at ultrahigh energy neutrinos similar to the NLO DGLAP (long-dashed lines) for $E_{\\nu} \\simge 10^{9}$ GeV.   In order to investigate the BFKL physics in the geometric scaling region, we use the recent fit to the dipole cross section presented in Ref. \\cite{IIM} (the CGC model), which encodes the BFKL formalism in the geometric scaling regime and an interpolation to the saturation domain, as presented in the last section. It should be noticed that an analysis including DGLAP evolution in the fit would be timely. However, as the CGC model produces an ``effective'' anomalous dimension, i.e. $\\gamma_{\\mathrm{eff}}\\,(x,\\rr Q_{\\mathrm{sat}}) \\equiv \\gamma_{sat} + (\\kappa \\lambda Y)^{-1}\\ln \\,(4/\\rr^2Q_{\\mathrm{sat}}^2)$ \\cite{IIM} far from the geometric scaling region $\\rr^2 Q_{\\mathrm{sat}}^2(x) \\ll 1$, which is closer to the DGLAP one,  the estimate presented here would be not strongly spoiled by QCD evolution. The results for the CGC model (solid lines) is presented in Fig. \\ref{fig:2}, contrasted with  the  NLO DGLAP calculation (dashed lines)  and also with the unified BFKL-DGLAP formalism (dot-dashed lines), which embodies non-leading $\\ln (1/x)$ contributions \\cite{KMSNEUT}.  The latter incorporates both the $\\ln(1/x)$ BFKL resummation and the complete LO DGLAP evolution, with the dominant non-leading $\\ln(1/x)$ contributions resummed to all orders. The result found for CGC model is quite similar to  both DGLAP and unified BFKL-DGLAP formalisms in the range of energy considered here. It is not observed any enhancement in the neutrino-nucleon cross section in comparison with the leading twist DGLAP calculation, even at very high energies of order $E_{\\nu}\\sim 10^{12}$ GeV. This fact is in disagreement with Ref. \\cite{JAMAL}, where it has been shown that the geometric scaling lead to an enhancement of the neutrino-nucleon total cross section by an order of magnitude versus the  leading twist cross section. We also verify that  at $E_{\\nu}\\simge  10^{11}$ GeV the power on energy starts to slow down and it should be increasingly smoother (logarithmic growth) at higher energies. This expected change of behavior on energy of the cross section is estimated to be reached at $E_{\\nu}\\sim 10^{18}$ GeV in Ref. \\cite{JAMAL}, suppressing  the cross section compared to the leading twist result.  \\begin{figure}[t] \\begin{tabular}{cc} \\epsfig{file=nuccbfkl.eps,width=43mm,height=55mm} & \\epsfig{file=nuncbfkl.eps,width=43mm,height=55mm} \\\\ (a) & (b) \\end{tabular} \\caption{The $\\nu \\,N$  charged current (a) and  neutral current (b) as a function of neutrino energy. The solid curves correspond to the results for  the CGC model. The NLO DGLAP result is represented by the dashed and the unified BFKL-DGLAP formalism by the dot-dashed curves, respectively.} \\label{fig:2} \\end{figure}  We can qualitatively understand the result above analyzing the ratio between the dipole cross section for  BFKL  and DGLAP approaches in the geometric scaling regime. As discussed before, the ultrahigh energy neutrino cross section is sensitive to the kinematical region of small $x\\sim m_{W,Z}^2/E_{\\nu} $ and high $Q^2\\sim m_{W,Z}^2$. Concerning DGLAP formalism, in such a region we are in the  double logarithmic approximation (DLA) limit. Concerning the non-linear QCD approaches, at the vicinity of saturation, the DLA saddle point is no longer equal to one, but it was found to be $\\gamma_{\\mathrm{sat}}^{\\mathrm{DLA}}\\simeq 1/2$ \\cite{IANCUGEO,STASTOGEO,GBMS}. The ratio of the cross section in the geometric scaling region is given by, \\begin{eqnarray} R_{\\mathrm{scal}} \\sim \\frac{ \\left[\\rr^2 Q_{\\mathrm{sat}}^2 \\right]^{ \\gamma_{\\mathrm{sat}}^{\\mathrm{BFKL}}}}{\\left[\\rr^2 Q_{\\mathrm{sat}}^2\\right]^{\\gamma_{\\mathrm{sat}}^{\\mathrm{DLA}}}}\\simeq \\left[\\frac{Q_{\\mathrm{sat}}^2}{m_{W,Z}^2} \\right]^{\\left(\\gamma_{\\mathrm{sat}}^{\\mathrm{BFKL}} - \\gamma_{\\mathrm{sat}}^{\\mathrm{DLA}}\\right)} \\sim {\\mathcal O}\\,(1)\\,,\\nonumber \\end{eqnarray} where we have disregarded the diffusion factor in the BFKL expression and consider $\\rr^2\\simeq 1/Q^2 \\sim 1/m_{W,Z}^2$. The result $R_{\\mathrm{scal}}\\simeq 1$ leads to similar results between the all twist and leading twist calculation in the geometric scaling region.  In fact, there is room for a slight  suppression since $(\\gamma_{\\mathrm{sat}}^{\\mathrm{BFKL}} - \\gamma_{\\mathrm{sat}}^{\\mathrm{DLA}}) \\simeq 0.13$ and $Q_{\\mathrm{sat}}^2/m_{W,Z}^2 < 1$. On the other hand, the argument in Ref. \\cite{JAMAL}  is exactly true in the comparison between CGC and GBW models. In this case, as we will discuss in what follows, $\\gamma_{\\mathrm{sat}}^{\\mathrm{GBW}}=1$ in the geometric scaling region and a sizeable enhancement between CGC and GBW is really observed.  The qualitative analysis presented above can also help us to understand the enhancement of the BGBK model (with DGLAP evolution) in relation to GBW saturation model. We can think in an effective ``anomalous dimension'' for the saturation model in the geometric scaling regime: from Eq. (\\ref{gbwdip}), $\\sigma_{dip}^{\\mathrm{GBW}} \\simeq \\rr^2 Q_{\\mathrm{sat}}^2/4$ and then one  verifies that $\\gamma_{\\mathrm{sat}}^{\\mathrm{GBW}}=1$. In this case, the ratio now reads as, \\begin{eqnarray} R_{\\mathrm{scal}} = \\frac{\\sigma_{dip}\\,(\\mathrm{BGBK})}{\\sigma_{dip}\\,(\\mathrm{GBW})} \\sim \\left[\\frac{Q_{\\mathrm{sat}}^2}{m_{W,Z}^2} \\right]^{\\left(\\gamma_{\\mathrm{sat}}^{\\mathrm{DLA}} - \\gamma_{\\mathrm{sat}}^{\\mathrm{GBW}}\\right)} \\gg 1\\,, \\end{eqnarray} whereas at lower energies (beyond the geometric scaling regime) the DGLAP anomalous dimension takes the usual value and the ratio is closer to one, as expected.  As a summary, we have presented the results for the ultrahigh energy neutrino cross section considering the current phenomenology on non-linear QCD dynamics, which play an important role in the correct extrapolation to the regime of very high gluon density. The saturation model with QCD evolution was studied as well as the physics of BFKL dynamics at geometric scaling region. The latter presents suppression in the cross section at $E_{\\nu}\\simge 10^{12}$ GeV, whereas the BGBK model is consistent with NLO DGLAP calculations. An outstanding advantage in the color dipole framework is the possibility to provide analytical expressions for the extrapolation to ultrahigh energies, towards the black limit of the nucleon target. This cannot be taken into account in a leading twist approach as the QCD-improved  parton model. Then these saturation (all twist) approaches  play an important role, since a precise knowledge about the neutrino  interactions and production rates is essential for estimating background, expected fluxes and detection probabilities \\cite{KMSNEUT,STASTO}. In particular, reliable estimates for UHE neutrino cross section are strongly needed for the planned experiments \\cite{KUSENKO} to detect them via nearly horizontal air showers in the Earth atmosphere, as discussed in \\cite{FIORE}."
    },
    "0311/astro-ph0311127_arXiv.txt": {
        "abstract": "We obtained spectra of 60 red, starlike objects ($E<$ 18.8) identified with FIRST radio sources, $S_{1.4GHz} >$ 1 mJy. Eight are QSOs with redshift $z>$ 3.6. Combined with our pilot search (Benn et al 2002), our sample of 121 candidates yields a total of 18 $z >$ 3.6 QSOs (10 of these with $z >$ 4.0). 8\\% of candidates with $S_{1.4GHz}<$ 10 mJy, and 37\\% of candidates with $S_{1.4GHz}>$ 10 mJy are QSOs with $z >$ 3.6. The surface density of $E <$ 18.8, $S_{1.4GHz} >$ 1mJy, $z>$ 4 QSOs is 0.003 deg$^{-2}$. This is currently the {\\it only} well-defined sample of radio-loud QSOs at $z \\approx$ 4 selected independently of radio spectral index. The QSOs are highly luminous in the optical (8 have $M_B <$ -28, $q_0$ = 0.5, $H_0$ = 50 kms$^{-1}$Mpc$^{-1}$). The SEDs are as varied as those seen in optical searches for high-redshift QSOs, but the fraction of objects with weak (strongly self-absorbed) Ly$\\alpha$ emission is marginally higher (3 out of 18) than for high-redshift QSOs from SDSS (5 out of 96). ",
        "introduction": "The evolution with redshift of the space density of QSOs places strong constraints on the abundance of massive objects at the earliest cosmological epochs. High-redshift QSOs can be found efficiently by colour selection from very large samples (e.g. the Sloan Digital Sky Survey, Anderson et al 2001, see Benn et al 2002, hereafter Paper I, for a summary of previous searches), but space-density measurements based on such selection could be biased by e.g. redshift-dependent dust extinction.  Any such bias is reduced when selecting in the radio, but searches for high-redshift radio QSOs have so far concentrated on radio-bright objects with flat radio spectra, which yields small samples with relatively low surface density on the sky. E.g. Snellen et al (2001) and Hook et al (2002) sought red starlike optical counterparts of $S_{5GHz} >$ 30 mJy flat-spectrum radio sources, and identified a total of 8 $z >$ 4 QSOs. Here, we aim to identify high-redshift radio quasars without any bias in radio spectral index, and with higher surface density on the sky, through spectroscopy of red starlike counterparts of FIRST radio sources, with $S_{1.4GHz} >$ 1 mJy. Combined with our pilot search (Paper I), this yields 18 $z >$ 3.6 radio-loud\\footnote{All 18 QSOs are radio-loud according to the criterion of Gregg et al 1996, i.e. P$_{1.4 GHz}$ $>$ 25.5; see Table 1). } QSOs (9 previously known, mainly from multi-colour optical searches). This includes the largest sample to date of radio-selected QSOs at z $>$ 4.   \\begin{table*} \\vbox to220mm{\\vfill Landscape table to go here. \\caption{} \\vfill} \\label{landtab} \\end{table*}  \\begin{table*} \\begin{minipage}{170mm} \\caption{Candidates which are not $z >$ 3.6 QSOs} \\begin{tabular}{r r rrr rrl ll} \\hline \\multicolumn{1}{c}{RA}   & \\multicolumn{1}{c}{Dec}   & \\multicolumn{1}{c}{S$_{1.4}$}  & \\multicolumn{1}{c}{R-O}    & \\multicolumn{1}{c}{$E$} & \\multicolumn{1}{c}{$O-E$}& \\multicolumn{1}{c}{ID}& \\multicolumn{1}{c}{z}& \\multicolumn{1}{c}{$\\sigma_z$}  & \\multicolumn{1}{c}{Notes} \\\\  \\multicolumn{1}{c}{J2000}   & \\multicolumn{1}{c}{J2000}   & mJy     & \\multicolumn{1}{c}{$\\prime\\prime$} &     &     &   &  & &     \\\\ \\multicolumn{1}{c}{(1)}  & \\multicolumn{1}{c}{(2)}   & \\multicolumn{1}{c}{(3)}& \\multicolumn{1}{c}{(4)}& \\multicolumn{1}{c}{(5)}& \\multicolumn{1}{c}{(6)}& \\multicolumn{1}{c}{(7)}& \\multicolumn{1}{c}{(8)}& \\multicolumn{1}{c}{(9)}& \\multicolumn{1}{c}{(10)} \\\\ \\hline \\multicolumn{10}{l}{\\bf Low-redshift QSOs and galaxies (`?' = probable galaxy)} \\\\ \\rule{0mm}{3.5mm} 08 53 36.66& 17 43 47.9&  40.4&0.4&18.4&    3.2&Q&     &     &Jaunsen et al (1995), $z$ = 3.21 \\\\ 09 29  00.44& 16 28 06.3&   4.8&0.9&17.0&$>$ 4.7&?&     &     &                                  \\\\ 09 41 39.71& 35 32 33.5&   4.6&0.6&18.8&$>$ 2.9&?&     &     &                                  \\\\ 09 44 28.41&  06 41 44.7&   3.2&0.8&18.6&$>$ 3.1&?&     &     &                                  \\\\ 10  01 42.41& 20 48 18.2&   2.2&0.5&18.6&$>$ 3.1&Q&     &     &Stocke et al (1991), BL Lac, $z$ = 0.346 \\\\ 10  08 34.99& 35 51 23.4&   6.5&1.1&18.6&$>$ 3.1&?&     &     &                                  \\\\ 10 13 08.38&  08 27 15.7&  28.5&0.4&18.7&$>$ 3.0&?&     &     &                                  \\\\ 10 24 42.51& 56 11 25.1&   0.8&0.3&18.7&$>$ 3.0&?&     &     &                                  \\\\ 10 25  06.45& 19 15 44.6&   2.0&0.7&18.6&$>$ 3.1&?&     &     &                                  \\\\ 11 34 58.18& 12 32 21.2&   6.7&0.5&18.0&    3.3&?&     &     &                                  \\\\ 12  02  09.51& 31 30 40.1&   5.7&1.0&18.7&$>$ 3.0&?&     &     &                                  \\\\ 12  05  01.33& 14 28 30.7&  15.4&1.4&18.7&    3.2&?&     &     &                                  \\\\ 12 20 16.87& 11 26 28.2&   2.3&0.4&18.1&    3.6&Q&1.893&0.010&CIII 1909, CIII 2326, MgII 2798 \\AA  \\\\ 12 39 32.76& -00 38 38.1&   1.3&0.6&18.4&$>$ 3.3&Q&     &     &Boyle et al (1990), BAL, $z$ = 2.18\\\\ 12 39 56.76& 42 40 59.8&   3.2&0.4&18.8&$>$ 2.9&?&     &     &                                  \\\\ 12 53 56.32&  00 52 42.9&  12.7&1.4&18.7&$>$ 3.0&?&     &     &                                  \\\\ 12 55 04.57& 43 10 38.1&   3.2&0.2&18.8&$>$ 2.9&?&     &     &                                  \\\\ 12 55 27.65& 55 18 19.0&   1.3&0.8&18.6&$>$ 3.1&?&     &     &                                  \\\\ 13  00 38.85& -03 35 16.8&   1.4&0.6&17.4&    3.1&G&     &     & Extended on JKT image, no spectrum obtained \\\\ 13  00 53.46& -03 11 28.2&   7.5&1.0&18.0&    3.0&G&     &     &                  \\\\ 13  06 21.99&  03 36  06.3&   2.0&0.8&18.4&$>$ 3.3&?&     &     &                                  \\\\ 13 10 43.44& 43 39  04.8&   4.7&0.7&18.6&$>$ 3.1&?&     &     &                                  \\\\ 13 18 51.57& -00 53 23.4&   4.1&1.5&18.5&    3.6&?&     &     &                                  \\\\ 13 32 29.08& 26 34 33.2&   1.3&0.3&18.8&$>$ 2.9&?&     &     &                                  \\\\ 13 39 43.42& 40 34 25.4&   1.0&1.0&18.7&$>$ 3.0&?&     &     &                                  \\\\ 13 43 14.20& 22  08 45.7&   1.5&0.4&18.4&$>$ 3.3&G&     &     & Extended on JKT image, no spectrum obtained \\\\ 13 49 07.12& -02 23 24.9&   3.1&0.4&18.0&$>$ 3.7&?&     &     &                                  \\\\ 14 13 59.17&  02 55 27.2&   1.1&0.3&17.9&    3.1&G&     &     &                                  \\\\ 14 20 54.98& 21 10 29.6&   1.9&0.5&18.1&$>$ 3.6&G&     &     &                                  \\\\ 14 28  01.12& 25 45 40.2&   3.3&0.7&18.7&$>$ 3.0&?&     &     &                                  \\\\ 14 32 22.27& 48 34 42.6&  16.7&0.7&18.8&$>$ 2.9&S&0.1906    &0.0005    & OIII 4959/5007, H$_\\alpha$                                  \\\\ 14 52 43.77&  01 27 33.1&  11.9&0.5&18.5&$>$ 3.2&?&     &     &                                  \\\\ 14 53 46.71& 35 53 11.6&   1.3&0.9&18.8&$>$ 3.0&G&     &     &                                  \\\\ 15 11  07.20& 11 45 58.8&   1.0&0.4&18.8&$>$ 2.9&G&     &     & Extended on JKT image, no spectrum obtained \\\\ 15 34 22.24& 18 51 25.2&   1.5&0.3&18.4&    3.2&?&     &     &                                  \\\\ 15 45 20.00& 10 28 55.7&   5.6&1.3&18.8&$>$ 3.0&G&     &     &                                  \\\\ 16 40 06.14& 17 38 49.4&  15.6&1.0&18.2&    3.0&?&     &     &                                  \\\\ 17 00 01.29& 19 29 31.6&   2.7&0.8&18.7&$>$ 3.0&G&     &     & Extended on JKT image, no spectrum obtained \\\\ \\multicolumn{10}{l}{\\bf Stars} \\\\ 09 01 46.15& 15 02 21.8&   3.7&1.3&18.6&$>$ 3.1&*&     &     &                                  \\\\ 09 09 53.20& 03 01 35.8&   1.5&0.5&18.6&$>$ 3.1&*&     &     &                                  \\\\ 09 12 20.95& 14 18 48.3&   1.8&1.5&18.4&$>$ 3.3&*&     &     &                                  \\\\ 09 17 33.46& 08 56 30.7&   2.1&1.4&17.9&    3.7&*&     &     &                                  \\\\ 09 27 20.11& 14 25 49.8&  11.4&0.5&18.7&$>$ 3.0&*&     &     &                                  \\\\ 09 47 34.27&-01 51 31.7&   7.5&0.8&18.4&    3.2&*&     &     &                                  \\\\ 09 51 53.79& 00 52 05.3&   7.2&0.8&18.6&    3.3&*&     &     &                                  \\\\ 11 07 19.77& 15 16 20.1&  14.3&1.2&18.0&    3.0&*&     &     &                                  \\\\ 11 09 25.01& 12 10 30.9&  11.1&1.0&18.3&    3.1&*&     &     &                                  \\\\ 11 22 09.69& 02 13 36.8&   2.3&1.2&17.9&$>$ 3.8&*&     &     &                                  \\\\ 12 13 06.97& 13 43 07.3&  62.7&0.5&18.4&    3.1&*&     &     &                                  \\\\ 13 26 58.60& 17 35 14.5&  14.8&1.4&17.2&    3.2&*&     &     &                                  \\\\ 14 33 23.85& 37 39 08.9&  15.8&1.2&18.7&$>$ 3.0&*&     &     &                                  \\\\ 14 49 00.23& 07 37 36.2&   1.3&1.3&18.5&$>$ 3.2&*&     &     &                                  \\\\ 15 14 50.17& 06 36 05.3&   2.4&1.2&17.2&    3.3&*&     &     &                                  \\\\ 15 16 43.79& 22 20 46.3&   4.0&1.5&17.8&    3.7&*&     &     &                                  \\\\ 15 39 59.30&-02 15 09.7&  14.0&1.3&18.3&$>$ 3.4&*&     &     &                                  \\\\ 16 03 05.11& 14 43 42.9& 285.3&1.4&17.5&    3.1&*&     &     &                                  \\\\ 16 23 57.72& 23 50 12.0&   2.5&1.0&18.8&$>$ 3.0&*&     &     &                                  \\\\ 16 41 53.13& 12 55 12.7&   1.8&0.6&17.8&    3.4&*&     &     &                                  \\\\ 16 48 20.27& 50 39 04.0&  62.6&1.2&18.8&$>$ 2.9&*&     &     &                                  \\\\ \\hline \\end{tabular} Columns 1-6 and 8-10 are as in Table 1. Column 7 indicates spectral type (see caption of Fig. 3). For 3 of the QSOs, redshifts are available from the literature (indicated); spectra were not obtained of these. \\end{minipage} \\end{table*} ",
        "conclusions": "We identified 194 high-redshift QSO candidates which: \\begin{enumerate} \\item coincide with FIRST radio sources $S_{1.4GHz} >$ 1 mJy; \\item are classified starlike on APM (POSS-I) and on POSS-II; \\item have APM $E <$ 18.8 and $O-E \\geq$ 3.0, or are invisible on the $O$ plate. \\end{enumerate} The sample covers an area of 7030 deg$^2$. We have observed 121 of the candidates and find 18 to be QSOs with $z >$ 3.6, 10 of these with $z >$ 4 (Benn et al 2002, and this paper). We estimate that we have found 75\\% of the high-redshift QSOs present amongst the 194 candidates. The surface density of $z >$ 4 QSOs with $E <$ 18.8, $S_{1.4GHz} >$ 1 mJy, is 0.0031 deg$^{-2}$.  This is currently the {\\it only} well-defined sample of $z \\approx $ 4 radio QSOs selected independently of radio spectral index. Vigotti et al (2003) use a subsample of 13 of these QSOs, with $z>$ 3.8, $M_{AB}$(1450 \\AA) $>$ -26.9 and log $P_{1.4 GHz}$(W Hz$^{-1}$) $>$ 25.7, to measure the  space density of QSOs at $z$ = 4, and determine the change in space density between $z$ = 2 and $z$ = 4.  These QSOs are highly luminous in the optical (8 have $M_B <$ -28, $q_0$ = 0.5, $H_0$ = 50 kms$^{-1}$Mpc$^{-1}$). The SEDs are remarkably varied. They include: 3 QSOs with very weak Ly$\\alpha$ (0831+52, 0918+06, 1639+40), one with an unusually-high density of CIV absorption systems (0747+27, Richards et al 2002), one with a probable DLA (0941+51), the luminous and much-studied 0831+52 (Irwin et al 1998), and a QSO with narrow Ly$\\alpha$ of high equivalent width (1309+57). 1624+37 (outside the sample) is an unusual BAL QSO.    \\vspace{3mm} {\\bf Acknowledgments}\\\\ JH was a 1-year placement student at ING during 2000-2001. RC and JIGS acknowledge financial support from DGES project PB98-0409-C02-02. CRB, MV, RC and JIGS acknowledge financial support from the Spanish Ministerio de Ciencia y Tecnologia under project AYA2002-03326. KHM was supported through a European Community Marie Curie fellowship. The Isaac Newton and the Jacobus Kapteyn Telescopes are operated on the island of La Palma by the Isaac Newton Group in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias. The 100-m Effelsberg radio telescope is operated by the Max-Planck-Institut f\\\"ur Radioastronomie. The APS Catalog of POSS I (http://aps.umn.edu) is supported by NASA and the University of Minnesota. We are grateful to the anonymous referee for helpful suggestions."
    },
    "0311/astro-ph0311311_arXiv.txt": {
        "abstract": "We introduce a Markov Chain Monte Carlo simulation and data analysis package that extends the \\textsc{Cmbeasy} software.  We have taken special care in implementing an adaptive step algorithm for the Markov Chain Monte Carlo in order to improve convergence. Data analysis routines are provided which allow to test models of the Universe against measurements of the cosmic microwave background, supernovae Ia and large scale structure. We present constraints on cosmological parameters derived from these measurements for a $\\Lambda$CDM cosmology and discuss the impact of the different observational data sets on the parameters. The package is publicly available as  part of the {\\sc Cmbeasy} software at www.cmbeasy.org. ",
        "introduction": "The wealth of recent  precision measurements in cosmology \\cite{Bennett:2003bz,Hinshaw:2003ex,Kogut:2003et,Kuo:2002ua,Readhead:2004gy,Dickinson:2004yr,Riess:2004nr,Tonry:2003zg,Knop:2003iy,Barris:2003dq,Percival:2001hw,Verde:2002,Peacock:2001gs,Tegmark:2003ud,Tegmark:2003uf} places stringent constraints on any model of the Universe. Typically, such a model is given in terms of a number of cosmological parameters. Numerical tools, such as {\\sc Cmbfast} \\cite{Seljak:1996is}, \\textsc{Camb}  \\cite{Lewis:1999bs} and \\mbox{\\sc Cmbeasy}  \\cite{Doran:2003sy}, permit to calculate the prediction of a given model  for the observational data. While these tools are  comparatively fast, scanning the parameter space for the most likely model and confidence regions can become a matter of time and computing power.  The cost of evaluating models on a n-dimensional grid in parameter space increases exponentially with the number of parameters. In contrast, the Markov Chain Monte Carlo (MCMC) method scales roughly linearly with the number of parameters \\cite{Christensen:2000ji,Christensen:2001gj,Lewis:2002ah}.  The MCMC method has already been used to constrain  various models \\cite{Knox:2001fz,Kosowsky:2002zt,Verde:2003ey,Caldwell:2003hz}. A popular tool for setting up MCMC simulations is the \\textsc{Cosmo-Mc} package \\cite{Lewis:2002ah} for the  \\textsc{Camb} code, an improved proposal distribution for the local Metropolis algorithm has been proposed in \\cite{Slosar:2003da}.  In this paper, we  introduce the \\class{AnalyzeThis} package\\footnote{It is part of the cmbeasy v2.0 release.} for {\\sc Cmbeasy}. It includes a parallel MCMC driver, as well as routines to calculate the likelihood of a model with respect to various data sets. We took special care in designing a step-proposal strategy that leads to fast convergence and mixing of the chains. This strategy is applied during the early stages of the simulation. As soon as the likelihood contour has been roughly explored, the adaptive step proposal freezes in. This ensures that the MCMC results are not contaminated by the adaptive steps. At the same time, the automated step optimization considerably improves performance and is rather convenient. The raw data files can be processed from within a graphical user interface (gui). Using the gui, one can marginalize, visualize and print  one and two dimensional likelihood surfaces (see figure \\ref{fig::gui}).   The plan of this paper is as follows: we describe the MCMC method and our implementation in section \\ref{sec::simulation}. A brief introduction to the software is given in section \\ref{sec::user}. We analyze the constraints on cosmological parameters from observational data sets in section  \\ref{sec::constraints}. In section \\ref{sec::conclusion}  we present our conclusions, while the format of the MCMC data files is defined in \\ref{appendix::format}. ",
        "conclusions": "\\label{sec::conclusion}  We have introduced the {\\sc AnalyzeThis} package, which can be used to constrain cosmological models using observational data sets. The \\class{AnalyzeThis}  class provides many functions to compute the likelihood of a model with respect to  measurements of the CMB, SNe Ia and Large Scale Structure.\\footnote[8]{We would be happy to include contributions of routines for recent and future measurements. So if the readers favorite measurement is not included yet, please send in a few lines of code.}  In order to constrain models of the Universe with a substantial number of parameters, we include a Markov Chain Monte Carlo driver. As the MCMC step strategy determines the convergence speed of the chains, we implemented a multivariate Gaussian sampler with an additional dynamical scaling. We stop the  adaptive improvement of the step proposal density as soon as a good level of convergence is reached and discard all data calculated before. Using this approach, we combine the advantage of a single-run automatic optimization of the step strategy with the demand of a static step proposal density.  The output of the Monte Carlo chains is in a human readable format and may be processed by any software, even during the run. For convenience however, one may use {\\sc Cmbeasy}'s  gui to marginalize, plot and print two and one dimensional likelihood surfaces.  Finally, we discussed the impact of the different data sets on the parameters for the case of a $\\Lambda$CDM cosmology with five parameters.   \\ack We would like to thank Robert R.~Caldwell, Robert A.~Knop,  Havard B.~Sandvik, Max Tegmark, Adam G.~Riess and Timothy J.~Pearson for helpful discussions. M. Doran is supported by NSF grant PHY-0099543.  C.M. M{\\\"{u}}ller is supported by GRK grant 216/3-02.   \\appendix"
    },
    "0311/astro-ph0311582_arXiv.txt": {
        "abstract": "% The analysis of Planetary Nebulae (PNe) provides a tool to investigate the properties of their exciting central stars (CSPN) at the moment of the PN ejection as well as on the properties of the ambient interstellar medium (ISM). The spectral analysis of the CSPN is a prerequisite to calculate the ionizing flux which is a crucial input for reliable PN modeling. In the framework of a systematic study of PNe interacting with the ISM (Kerber et al\\@. 2000), we present preliminary results of ongoing NLTE spectral analyses of ten of their CS based on new optical medium-resolution spectra. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311061_arXiv.txt": {
        "abstract": "We have obtained \\emph{FUSE} and \\emph{HST} STIS spectra of HR 4796A, a nearby 8 Myr old main sequence star that possesses a dusty circumstellar disk whose inclination has been constrained from high resolution near-infrared observations to be $\\sim$17$\\arcdeg$ from edge-on. We searched for circumstellar absorption in the ground states of \\ion{C}{2} $\\lambda$1036.3, \\ion{O}{1} $\\lambda$1039.2, \\ion{Zn}{2} $\\lambda$2026.1, Lyman series H$_{2}$, and CO (A-X) and failed to detect any of these species. We place upper limits on the column densities and infer upper limits on the gas masses assuming that the gas is in hydrostatic equilibrium, is well-mixed, and has a temperature, $T_{gas}$ $\\sim$ 65 K. Our measurements suggest that this system possesses very little molecular gas. Therefore, we infer an upper limit for the gas:dust ratio ($\\leq$4.0) assuming that the gas is atomic. We measure less gas in this system than is required to form the envelope of Jupiter. ",
        "introduction": "Planets, asteroids, and comets are believed to form in circumstellar disks around stars with ages $<$few 100 Myr. Current (uncertain) models suggest that Jovian planets form either via rapid gravitational collapse through disk instability within a few hundred years (Boss 2000) or via coagulation of dust into solid cores within the first $\\sim$1 Myr and accretion of gas into thick hydrogen atmospheres within the first $\\sim$10 Myr (Ruden 1999). At present, the  timescales on which giant planets form and accrete their atmospheres have not been well constrained observationally. Studies of objects with ages $\\sim$10 Myr afford the opportunity to constrain the timescale for gas dissipation around young stars and therefore the timescale for giant planet formation.  Although the circumstellar disks around pre-main sequence stars possess both gas and dust, the circumstellar disks around main sequence stars may be gas poor. Millimeter searches for CO around stars with ages 1 - 10 Myr suggest that within a few Myr, less than a Jupiter mass of molecular gas remains in the circumstellar environment (Zuckerman, Forveille, \\& Kastner 1995). However, CO may be depleted in the outer regions of disks because it freezes out onto dust grains at temperatures, T$_{gas}$ $\\sim$ 20 K. Since gas with interstellar or solar abundance is mainly composed of hydrogen and H$_{2}$ does not freeze onto grains, recent searches for circumstellar gas have focused on H$_{2}$. \\emph{ISO} searches for emission from the S(0) and S(1) transitions of H$_{2}$ at 28 $\\mu$m and 17 $\\mu$m respectively, suggest that 50 - 2000 $M_{\\earth}$ H$_{2}$ exists around $\\beta$ Pic, 49 Cet, and HD 135344 (Thi et al. 2001). However, \\emph{FUSE} searches for H$_{2}$ absorption in the Lyman band toward $\\beta$ Pic place upper limits on the column density of H$_{2}$, N(H$_{2}$) $\\leq$10$^{18}$ cm$^{-2}$, which is 3 orders or magnitude lower than that inferred from \\emph{ISO} (Lecavelier des Etangs et al. 2001). The Thi et al. (2001) results not only apparently conflict with the ultraviolet H$_{2}$ measurements but also with millimeter emission line studies which suggest that $\\beta$ Pic possesses $<$10$^{-5}$ M$_{\\earth}$ CO (Liseau \\& Artymowicz 1998). Ground-based follow-up of \\emph{ISO} H$_{2}$ detections toward Herbig Ae and T-Tauri stars yielded upper limits which were smaller than \\emph{ISO} detections (Richter et al. 2002), suggesting that the H$_{2}$ must lie in an extended cloud. However, an H$_{2}$ cloud around $\\beta$ Pic should have produced a detectable ultraviolet absorption line (Jura 2003). Recently, fluorescent H$_{2}$ Lyman-band emission has been detected from the surface layers of the circumstellar disk around the $\\sim$10 Myr old star TW Hydrae (Herczeg et al. 2002) using \\emph{HST} STIS. We have carried out a \\emph{FUSE} and STIS search for circumstellar molecular and atomic gases around the $\\sim$8 Myr (Stauffer et al. 1995) debris disk object HR 4796A, another member of the TW Hydrae Association, to help constrain the gas dissipation timescale.  HR 4796A is a main sequence A0V star with a \\emph{Hipparcos} distance of 67.1 pc (see Table 1 for a summary of stellar properties) and a fractional dust luminosity L$_{IR}$/L$_{*}$ = 5$\\times$10$^{-3}$ (Jura et al. 1998). Submillimeter observations suggest that it possesses $\\geq$0.25 $M_{\\earth}$ dust (Greaves et al. 2000), more circumstellar dust than any other main sequence A-type star in the Yale Bright Star Catalog (Jura et al. 1991). High resolution scattered light imaging has resolved a narrow, dusty ring at 70 AU from HR 4796A with an inclination of $\\sim$17$\\arcdeg$ from edge-on (Schneider et al. 1999). High resolution mid-infrared imaging has resolved the disk in thermal emission (Koerner et al. 1998; Jayawardhana et al. 1998). The presence of a central clearing (Jura et al. 1998) and of a brightness asymmetry from the two lobes of the disk (Telesco et al. 2000; Weinberger, Becklin, \\& Schneider 2000) have led to speculation that the circumstellar disk possesses a planetary companion at $\\leq$70 AU (Wyatt et al. 1999). Blackbody fitting to the \\emph{IRAS} fluxes suggests that the grain temperature, $T_{dust}$ = 110 K (Jura et al. 1995).  Previous studies of circumstellar gas around HR 4796A have not been very successful. Searches for millimeter and submillimeter emission from CO (J=3$\\rightarrow$2, J=2$\\rightarrow$1, J=1$\\rightarrow$0), used as a tracer for H$_{2}$, have failed (Zuckerman et al. 1995, Liseau 1999, \\& Greaves et al. 2000), perhaps because the angular size of the source as measured in the infrared is significantly smaller than the beam of available telescopes. Another possibility is that circumstellar molecular gas is not present because it is rapidly photodissociated under the intense stellar ultraviolet field of the central A0V star. Holweger, Hempel, \\& Kamp (1999) have detected two narrow, circumstellar \\ion{Ca}{2} K absorption features, one component close to the stellar velocity and another weaker one blueshifted by $\\sim$10 km/sec which they attribute to possible circumstellar gas. ",
        "conclusions": "We have obtained ultraviolet spectra (905 \\AA \\ - 1187 \\AA, 1425 \\AA \\ - 1620 \\AA, 1880 \\AA \\ - 2155 \\AA) of the nearby, debris disk object HR 4796A using \\emph{FUSE} and STIS on \\emph{HST}. We argue the following:  1. The X-ray flux from the M-dwarf companion HR 4796B does not significantly heat circumstellar gas around HR 4796A. The estimated X-ray heating rate is at least $\\sim$40 times smaller than the estimated photoelectric heating rate.  2. Circumstellar gas around HR 4796A may be warmed to a temperature $T_{gas}$ $\\sim$ 65 K by photoelectric heating if the gas:dust ratio is 4, suggesting a molecular gas disk thickness parameter $H$ = 6.1 AU and an atomic gas disk thickness parameter $H$ = 8.1 AU at 70 AU from the star. This is approximately consistent with detailed heating and cooling models which suggest a maximum gas temperature, $T_{gas}$ 61 K, for a gas:dust ratio of 2, assuming a 1-D approximation for the escape probability of line photons.  3. The circumstellar disk around HR 4796A possesses very little $H_{2}$. We estimate that $\\leq$1.1 $\\times$ $10^{-3}$ $M_{\\earth}$ H$_{2}$ exists around HR 4796A assuming that gas is in hydrostatic equilibrium, is well-mixed, and has a temperature T$_{gas}$ = 65 K.  4. We estimate that $\\leq$1.0 $M_{\\earth}$ atomic hydrogen exists around HR 4796A. This corresponds to a gas:dust mass ratio $\\leq$4.0 assuming that gas is in hydrostatic equilibrium, is well-mixed, and has a temperature T$_{gas}$ = 65 K.  5. Measurements of \\ion{C}{2} column density may provide the strongest constraint on the gas mass in the HR 4796A circumstellar disk. While the measured H$_{2}$, CO, and \\ion{Zn}{2} column densities are consistent with the column densities predicted for the gas:dust ratio = 100 model, the measured \\ion{C}{2} column denisty N(\\ion{C}{2}) = $1.4 \\times 10^{14}$ cm$^{-2}$ is most consistent with the gas:dust ratio = 2 model.  6. The circumstellar disk around HR 4796A is too gas depleted to support the formation of giant planets. Any Jovian planets in this system must have formed on a timescale $\\leq$8 Myr."
    },
    "0311/astro-ph0311194_arXiv.txt": {
        "abstract": "{We present high-contrast imaging data on the disk of the classical T Tauri star TW Hya. The images were obtained through the polarimetric differential imaging technique with the adaptive optics system NACO. Our commissioning data show the presence of polarized disk emission between 0.1\\arcsec{} and 1.4\\arcsec{} from the star. We derive the first Ks-band radial polarized intensity distribution. We show that the polarized intensity compares well to shorter wavelengths surface brightness observations and confirm the previously reported gradual slope change around 0.8\\arcsec{}. These results show the potential of the new polarimetric differential imaging technique at 8m-class telescopes to map the inner regions of protoplanetary disks. ",
        "introduction": "The simulations of viscous disks  with an embedded Jupiter-mass planet (see, e.g. D'Angelo et al. \\cite{dan03}) predict dramatic changes in the disk structure: gaps are being opened and shock waves follow the path of the orbiting body. Detecting these signatures is an important step in confirming that our models for planet formation and disk evolution are right and would provide a unique possibility to study a newly formed exoplanet and its interaction with the circumstellar material.  However, the investigations of protoplanetary disks require high-contrast, high-spatial-resolution imaging very close to the bright central star. Although in recent years numerous groups studied disks around nearby stars at their planet-forming age with the HST, instrumental limitations do not allow to probe the disk structure at scales similar to our inner Solar System. The technique of polarimetric differential imaging (PDI) as described by Kuhn et al.~(\\cite{kuh01}) is capable of tracing scattered (i.e. polarized) light from a circumstellar disk very close to the central star. This promising method enhances the contrast between the disk and the star and, thereby, might be important to detect the detailed inner structure of protoplanetary disks.  The basic idea of the PDI is to take the difference of two orthogonally polarized, simultaeneously acquired images of the same object in order to remove all non-polarized light components. The non-polarized light consists mainly of the central star's light and the speckle noise, which is -- close to the star -- the dominant noise source of ground-based observations at optical and near-infrared wavelengths (Racine et al.~\\cite{rac99}). After this subtraction only the polarized light, such as scattered light from the disk, is left over.  To explore the potential of this technique for planet formation studies as well as to demonstrate the NACO/VLT platform capabilities we targeted the 8 Myr old (Webb et al.~\\cite{web99})  and very close ($56 \\pm 7$ pc, Wichmann et al.~\\cite{wic98}) classical T Tauri star TW Hya (RA: $11^{\\rm h} 01^{\\rm m} 51.9^{\\rm s}$  Dec:$-34\\degr 42' 17''$ J2000). At the visual wavelength regime TW Hya shows a polarization variability between 0-3\\% (Mekkaden \\cite{mek98}). Being an excellent target for planet formation studies, during the past few years several attempts have been made to image its disk as close as possible to the star and to identify possible disk structures (Krist et al.~\\cite{kri00}, Weinberger et al.~\\cite{wei02}). In addition, the presence of a giant planet was predicted (Calvet et al. \\cite{cal02}) from modeling the spectral energy distribution (SED) (see also Steinacker \\& Henning~\\cite{ste03}). The first application of the PDI to TW Hya by Potter (2003\\cite{pot03a}, 2003\\cite{pot03b}) made use of the 36 element curvature sensing Adaptive Optics (AO) system Hokupa'a (Graves et al.~\\cite{gra00}), of the Gemini North telescope. This H-band data set probed the circumstellar material closer than 0.5\\arcsec to the star.  In the following we describe the first results obtained with the PDI technique and the 188 element Shack-Hartmann AO system NACO, which is attached to the Nasmyth B port of the VLT UT4. ",
        "conclusions": "\\subsection{Scattering}  The presence of the extended polarization signal across our FOV confirms the existence of  light scattering material. The sinusoidal modulation of the Q- and U-images confirms the centrosymmetric polarization pattern. The fact that we detected such emission even as close as $\\sim$0.1\\arcsec{} ($\\sim$6 AU) to TW Hya is the first direct evidence that its dust disk extends so far in. This detection does not support the model of Krist et al. (\\cite{kri00}) who proposed the presence of an inner 'dark' zone ($r\\simeq 18$ AU = 0.3\\arcsec{}) with suppressed flaring angle, necessary to explain the exceptional brightness of the TW Hya disk.  Based on the modeling of the SED, Calvet et al. (\\cite{cal02}) predicted the presence of a mostly dust-evacuated inner gap at $r<3-4$  AU ($r<0.07''$). Such a gap could be a tell-tale signpost of a massive planet in close orbit around TW Hya and would be of high importance. However, our commissioning data do not allow to draw any definite conclusions on the existence of this feature.  \\subsection{Radial Profile and Slope Change}  In Fig.~\\ref{profile} we show the comparison of the radial polarized intensity derived from our PDI observation and (arbitrarily shifted) radial surface brightness profiles measured by the HST (Weinberger \\cite{wei02}). Although these measurements probe different quantities at different wavelengths (Ks and J, H,  respectively)  a remarkable similarity is evident between the radial slopes. Based on this similarity, we conclude that the polarization degree is nearly independent of the radius and thus the polarized intensity characterizes the surface brightness. Thus, the PDI method can lead to contrast enhancement without the loss of information.  The radial surface brightness profile between $0.9\\arcsec{}$ and $1.4\\arcsec{}$ can be well fitted by a power-law function $I_{\\rm Ks,1} \\propto r^{-3.1 \\pm 0.3}$. This is slightly steeper than the $I_{\\rm J,H} \\propto r^{-2.6 \\pm 0.1}$ behaviour found by Weinberger et al. (\\cite{wei02}) at 1.1 and 1.6 $\\mu$m, for the region between 0.8\\arcsec{} and 2.7\\arcsec{}.  To place our results in context, we compare them to surface brightness slopes predicted by simple analytical approximations (e.g. Whitney \\& Hartmann \\cite{whi92}). These approximations consider only isotropic scattering and simple geometric effects and do not include wavelength-dependent terms. The predicted surface brightness power-law exponent for an optically thick, geometrically thin disk is $-3$, while for a flaring disk an exponent of -2 is expected. Our Ks-band results in the $0.9\\arcsec{} < r < 1.4\\arcsec{}$ region support the flat disk geometry or a very small flaring. However, the surface brightness slope is determined by an intricate interplay between many disk and dust parameters and is not a conclusive measure of the disk geometry.  As to underline this statement, a strong, gradual change in the brightness profile around 0.8\\arcsec{} from TW Hya is evident. The annulus between $0.5\\arcsec{} < r < 0.7\\arcsec{}$ is characterized by the very gentle surface brightness behaviour of $I_{\\rm Ks,2} \\sim r^{-0.9\\pm0.2}$. The slope of the curve is even less steep than that expected from a flared disk (Whitney \\& Hartmann~\\cite{whi92}).  Although the radial brightness slope is strongly influenced by the disk geometry, slope changes can have different origins. Even though a detailed analysis of the reasons of the slope changes is beyond the scope of the current paper, we refer to some of them. The most obvious one is an actual change in the surface density slope of the disk, translating into a slope change in the surface brightness. Such clear-cut changes might be the results of external or internal perturbation of the disk by a companion or a massive planet as seen in the case of HR 4796A (see, e.g. Boccaletti et al.~\\cite{boc03}).  A second cause might be the change of the flaring angle and through this the change in the illumination of the disk's surface. Alternatively, a radial dependence of the scattering cross section of the dust grains might introduce modulations in the surface brightness profile. Disentangling the effects of these processes requires high-angular resolution observations at longer wavelengths.  We note, that this change in the slope at 0.8\\arcsec{} has been detected previously at shorter wavelengths by HST direct imaging and coronographic studies (Krist et al.~\\cite{kri00}, Weinberger et al.~\\cite{wei02}), together with other slope changes in the outer parts ($>$1.4\\arcsec) of the disk.  Our data were acquired during the commissioning of the NACO instrument and suffer from high detector noise and limited AO performance. The current performance of NACO is expected to result in roughly five times better SNR.  The first results of a PDI survey of the TW Hya Association  are presented by Hu\\'elamo \\& Brandner (\\cite{hue03}).  For future PDI observers it might be important to point out some of the limiting factors in our data set. First, the dynamic range of the data set has to be increased by applying series of different exposure times. Furthermore, the flat field calibration plays a critical role and should be carried out with great care. Finally, the neccessarily saturated star will pose limitations to the accurate alignment of the individual frames, which will in practice determine the inner radius until which the polarization signal can be detected."
    },
    "0311/astro-ph0311477_arXiv.txt": {
        "abstract": "{ We present the results of our investigation on the phenomenon of mass segregation in young star clusters in the Magellanic Clouds.  HST/WFPC2 observations on NGC 1818, NGC 2004 \\& NGC 2100 in the Large Magellanic Cloud and NGC 330 in the Small Magellanic Cloud have been used for the application of diagnostic tools for mass segregation: i)  the radial density profiles of the clusters for various mass groups and ii) their mass functions (MFs) at various radii around their centres. All four clusters are found to be mass segregated, but each one in a different manner. Specifically not all the clusters in the sample show the same dependence of their density profiles on the selected magnitude range, with NGC 1818 giving evidence of a strong relation and NGC 330 showing only a hint of the phenomenon. NGC 2004 did not show any significant signature of mass segregation in its density profiles either. The MFs radial dependence provides clear proof of the phenomenon for NGC 1818, NGC 2100 and NGC 2004, while for NGC 330 it gives only indications. An investigation of the constraints introduced by the application of both diagnostic tools is presented.  We also discuss the problems related to the construction of a reliable MF for a cluster and their impact on the investigation of the phenomenon of mass segregation. We find that the MFs of these clusters as they were constructed with two methods are comparable to Salpeter's IMF. A discussion is given on the dynamical status of the clusters and a test is applied on the equipartition among several mass groups in them. Both showed that the observed mass segregation in the clusters is of primordial nature. ",
        "introduction": "Mass segregation nowadays is a well known phenomenon, which manifests itself as a strong concentration of the heaviest stars towards the centre of the cluster. This phenomenon has been extensively observed during the last decade in a plethora of young open clusters (Jones \\& Stauffer 1991; Pandey et al. 1992; Shu et al. 1997; Raboud \\& Mermilliod 1998; Hillenbrand \\& Hartmann 1998), as well as old globular clusters in the Galaxy (Cote et al. 1991; Pandey et al. 1992; Paresce et al. 1995; King et al.  1995; de Marchi \\& Paresce 1996; Ferraro et al. 1997; Sosin 1997; Andreuzzi et al. 2000; Howell et al. 2000). Young star clusters in the Magellanic Clouds (MCs) have also recently been investigated for mass segregation (Malumuth \\& Heap 1994; Brandl et al. 1996; Fischer et al. 1998;  Kontizas et al. 1998;  Santiago et al. 2001; de Grijs et al. 2002a,b,c; Sirianni et al. 2002). Last, but not least noticeable, is the work by Stolte et al. (2002), who found mass segregation in the Arches cluster, a young, massive star cluster and one of the densest in the Milky Way.  Galactic globulars show mass segregation due to their dynamical evolution, through which equipartition of kinetic energy among different mass groups in the cluster can be achieved. According to Spitzer (1969) this leads to stars of different masses being found in different ``layers'' with the heavier stars gravitating toward the centre, and thus mass stratification occurs. This phenomenon is mostly known as {\\em Dynamical Mass Segregation} (for reviews see Lightman \\& Shapiro 1978; Meylan \\& Heggie 1997). Earlier studies (King 1966; Da Costa 1982; Bolte \\& Marleau 1989) showed that low-mass stars predominate the galactic globulars. This implies a time-scale for energy equipartition of the order of the relaxation time of the clusters (the time the system needs to reach a quasi-Maxwellian equilibrium in its interior). For typical globular clusters the relaxation time (at the half-mass radius) is significantly shorter than the cluster age. In consequence, since the cluster had the time to relax, it exhibits dynamical mass segregation.  The discovery of mass segregation in young clusters (see review by Clarke et al. 2000) in the MCs introduced new problems concerning its interpretation, since it can be either the signature of the initial conditions and loci, where massive stars are formed (e.g. Hillenbrand \\& Hartmann 1998; Bonnell \\& Davies 1998), or it can be the result -- at least partially -- of early dynamics (e.g. Malumuth \\& Heap 1994; Brandl et al.  1996). This phenomenon is usually referred to as {\\em Primordial Mass Segregation}. Observations and theoretical predictions on the formation of massive stars emphasise the role of early dynamics. An example has been given by Fischer et al. (1998), who suggest that the steepening of the mass function they observed at larger radii from the centre of NGC 2157 is most likely an initial condition of the cluster formation, since interactions among the protostellar cloudlets play a crucial role in the massive-star-forming process.  \\begin{table*}[t] \\begin{center} \\caption{Ages, metallicities and structural parameters for the sample of our clusters, as found in the literature. The structural parameters (last four columns) have been estimated by Mackey \\& Gilmore (2003a,b). The core radii of the clusters are given in col. 6, while $L_{\\rm m}$ and $M_{\\rm m}$ are the integrated luminosity and mass to the distance from the centre of each cluster where a maximum measured extent of its profile is reached ($\\simeq 1\\farcm3$ for all clusters). The central densities of the clusters as estimated by Mackey \\& Gilmore are given in the last column.} \\begin{tabular*}{\\textwidth}[]{@{\\extracolsep{\\fill}}lcccccccc} \\hline Cluster & $\\log{\\tau}$&Age& Metallicity& Met.&$r_{\\rm c}$&$\\log{L_{\\rm m}}$& $\\log{M_{\\rm m}}$&$\\log{\\rho_{\\rm 0}}$\\\\ Name    & (yr)&ref.& [Fe/H]& ref.&(arcmin)&(L{\\solar})&(M{\\solar})& (M{\\solar} pc$^{-3}$)\\\\ \\hline \\hline NGC 330 &7.50$^{+0.10}_{-0.50}$& 1,2,3& $-0.82 \\pm 0.11$&4  &0.15 $\\pm$ 0.01& 5.46$^{+0.12}_{-0.13}$&4.41$^{+0.12}_{-0.13}$&1.84$^{+0.07}_{-0.07}$\\\\ NGC 2004&7.20$^{+0.30}_{-0.10}$& 1,5  & $-0.56 \\pm 0.20$&6  &0.11 $\\pm$ 0.01& 5.37$^{+0.16}_{-0.17}$&4.27$^{+0.16}_{-0.17}$&2.32$^{+0.10}_{-0.09}$\\\\ NGC 1818&7.40$^{+0.30}_{-0.10}$& 1,7    & 0.00 to $-$0.40 &7,8&0.17 $\\pm$ 0.01& 5.11$^{+0.10}_{-0.10}$&4.01$^{+0.10}_{-0.10}$&1.55$^{+0.06}_{-0.06}$\\\\ NGC 2100&7.20$^{+0.20}_{-0.20}$& 1,5  & $-0.32 \\pm 0.20$&6  &0.08 $\\pm$ 0.01& 5.46$^{+0.22}_{-0.23}$&4.31$^{+0.22}_{-0.23}$&2.64$^{+0.13}_{-0.12}$\\\\ \\hline \\end{tabular*} \\parbox[l]{\\textwidth}{References: (1) Keller et al. (2000); (2) Chiosi et al. (1995); (3) Da Costa \\& Hatzidimitriou (1998); (4) Hill (1999); (5) Elson (1991); (6) Jasniewicz \\& Th\\'{e}venin (1994); (7) de Grijs et al. (2002b); (8) Johnson et al. (2001)} \\end{center} \\label{params} \\end{table*}  In addition, according to the cluster formation model proposed by Murray \\& Lin (1996) cloudlets fall toward the central region of the protocluster cloud, and cohesive collisions cause their masses to grow. More massive stars are subject to many dissipative mergers (Larson 1991) and so they are preferably formed in the cluster centre (Bonnell \\& Davies 1998), leading to an initial mass segregation. Models suggesting enhanced accretion rates, which grow with the mass of the accreting protostar (Behrend \\& Maeder 2001), or high collision probabilities in dense cluster centres (Bonnell et al. 1998) predict the formation of stars up to 100 M$_{\\odot}$ in the densest central regions of a rich star cluster. Furthermore the total mass and the maximum stellar mass of a cluster strongly depends on the star formation rate and local density (Elmegreen 2001), which is enhanced by gravitational collapse or cloud collisions in high-density environments (Elmegreen 1999, 2001). In conclusion, protoclusters with high central density are more prone to exhibit primordial mass segregation.  We investigate the phenomenon of mass segregation in four young star clusters in the Magellanic Clouds: NGC 1818, NGC 2004 and NGC 2100 in the Large Magellanic Cloud (LMC) and NGC 330 in the Small Magellanic Cloud (SMC) with the use of HST/WFPC2 observations (Keller et al. 2000). These clusters, which are among the younger in the MCs with ages between 10 and 50 Myr, cover a range of metallicities (see Table 1).  This study is based on a consistent set of HST observations, which allowed us to search for mass segregation in a uniform manner and to investigate the differences among the clusters concerning the phenomenon. No such investigation seems to be available in the literature except for the case presented by de Grijs et al. (2002a,b,c). In the sample of their clusters these authors include NGC 1818, while Sirianni et al. (2002) recently presented their result on the low-mass IMF and mass segregation in NGC 330. The results of both these studies are based on different data sets than our HST observations on these clusters. Under these circumstances it would be interesting to compare our results, which have been preliminarily presented by Keller et al. (2001) with theirs.  This article is laid out as follows. In Sect. 3 we use our observations to examine the radial density profiles of the clusters and their dependence on the selected magnitude range, as an indication of stellar stratification. Our data were also used in Sect. 4 to derive the mass functions of the clusters and examine their radial dependence, which is a widely used method for the detection of mass segregation. The results of both investigations are given in Sect. 5 and a discussion on the nature of the observed mass segregation concerning its dynamical origin is presented in Sect. 6. Finally, Sect. 7 highlights our conclusions.  \\begin{figure} \\centerline{\\hbox{ \\psfig{figure=h4664f2.ps,height=8.0truecm,width=9.truecm,angle=0} }} \\caption{Completeness of the WFPC2/F555W photometry for various radial distances for the observed star clusters.} \\label{compl} \\end{figure} ",
        "conclusions": "The radial distribution of the cluster density in various magnitude ranges and its MF within different radial distances are two useful diagnostic tools for mass segregation. Both require 1) very good resolution towards the cluster centre in order to have accurate radial density distributions for various magnitude ranges and 2) very good photometry for accurate MFs in rings for various radial distances in order to find at which distance mass segregation leads to a change in the slope of the MF. Our HST/WFPC2 observations on young clusters in the MCs provide a suitable data set for the application of these tools. The conclusions of this investigation on the phenomenon of mass segregation in these clusters are the following.  We searched for mass segregation in a sample of four young star clusters in the MCs, NGC 1818, NGC 2004 and NGC 2100 in the LMC and NGC 330 in the SMC. These clusters are among the younger MC clusters and they vary in age, metallicity and other parameters. Mass segregation was exhibited in all four of them, but in a variety of related parameters, as well as of degrees of significance. From the dependence of their density profiles on the selected magnitude range it was found that NGC 1818 exhibits strongly the phenomenon of mass segregation with stars segregated almost in the whole of the observed magnitude range, while NGC 330 and NGC 2100 show a rather weak dependence and it is important only for stars around $V\\simeq 18$ mag ($\\sim$ 4 M{\\solar}).  For NGC 2004 there is an even weaker dependence also for $V \\sim 18$ mag.  We constructed the mass functions (MFs) of the clusters for the observed range of magnitudes in order to check for any radial dependence of their slopes as the signature of mass segregation. This construction was performed with two methods. First by directly counting stars between evolutionary tracks according to their positions in the HRD and second by translating their luminosities into masses using mass-luminosity relations and then constructing the distribution of the estimated masses. The $first$ $method$ which is more accurate concerning the shape of the MF but less statistically significant due to the lower numbers, was used to test the results of the $second$ $method$, which is more statistically complete but more model dependent.  We constructed the MFs of the clusters considering the dependence of their shape to the chosen binning for counting stars, the correction for incompleteness and the field contamination. The use of a mass-luminosity relation can affect the shape of the resulting MF (which tends to be flatter for older population models). Still, it was found that the models used for the $second$ $method$ were sufficient, since the overall MFs do not differ significantly from one method to the other and they have slopes clustered around a Salpeter IMF for a mass range of $\\sim$ 3 - 9 M{\\solar} for NGC 330, 3 - 14 M{\\solar} for NGC 2004, 4 - 11 M{\\solar} for NGC 1818 and 4 - 14 M{\\solar} for NGC 2100. It should be noted that not all of them may be represented by a single-power law for the whole mass range mentioned.  The construction of the radial MFs for the investigation of their slopes proved to be a rather complicated procedure due to several factors. Counting stars for the construction of the MF of a cluster in a specific ring around the centre can produce statistical problems due to the low numbers. Thus, the $first$ $method$ was excluded from the investigation of the radial dependence of the MF slope of the clusters, since by definition, it can rely on a smaller sample of stars.  NGC 1818 and NGC 2004 as was found from the radial dependence of their MFs are mass-segregated at radial distances of about 0\\farcm6 and 0\\farcm4 respectively. NGC 2100 was also found to be mass-segregated probably throughout its observed extent, but a steep radial dependence of its MF in the inner 0\\farcm2 was also observed and this can be assigned to a local central mass segregation of the most massive group of stars ($\\sim$ 13 M{\\solar}) as was found from the relative numbers of stars inside and outside this radial limit. This was also the case for NGC 1818 where 99\\% of the stars with masses around 10 M{\\solar} were found concentrated within about 0\\farcm2 from its centre. NGC 330 was also found to be segregated from the radial dependence of its MF, but in a less pronounced way up to a distance of $\\sim$ 0\\farcm5, with the best segregated mass group to be around 4 M{\\solar}. This cluster could also be considered as mass segregated throughout its extent.  We estimated the median relaxation time of the clusters in their core and observed half-mass radii, as well as in their ``radii of segregation'' as derived from the radial dependence of their MF slopes. This estimation was based on assumptions on the slope, and on the low-mass limit, and we found that the clusters in this sample are not dynamically relaxed within these radii, except for NGC 2004 and NGC 2100, which were found to be very close to relaxation only in their cores. Furthermore, the test of conditions for equipartition among three massive groups (Inagaki \\& Saslaw 1985) in our clusters showed that {\\em no equipartition has taken place} in any cluster at any of the above radii for stars with masses larger than 5 M{\\solar}. This clearly suggests that the observed mass segregation is of primordial nature. We present a discussion where we highlight the most important aspects of theories of star formation predicting the phenomenon of initial mass segregation (Sect. 6.2).  In general it seems that the diagnostic tools applied here, which are those most commonly used for the investigation of mass segregation, are rather sensitive to selection effects, such as the selection of the radial distances for which the MF slopes are computed, the binning selected for the construction of the MF, the selection of the mass range where the slopes are estimated, etc. While the presence of mass segregation in our sample of clusters is certainly a robust result, the description of the phenomenon can be only qualitative.  Specifically, we observe that the clusters of our sample exhibit mass segregation in different ways, but to describe these differences we would need to define a new parameter. This parameter may be called ``degree of mass segregation'', and it should be a function of parameters involved in the magnitude dependence of the density profile and the radial dependence of the MF slope for each cluster. It is not yet possible to define such a parameter with the use of the available tools, due to the limitations mentioned above.  Gouliermis \\& de Boer (2003) recently proposed a new method for the detection of stellar stratification in star clusters. They present a self-consistent, model-free tool for the investigation of the phenomenon, which can also be used to quantify the degree of mass segregation in star clusters in a uniform manner. This tool would probably provide a coherent scheme for the description of stellar stratification in star clusters, so that global and comparable results can be achieved."
    },
    "0311/astro-ph0311566.txt": {
        "abstract": "The GraF instrument using a Fabry-Perot interferometer cross-dispersed with a grating was one of the first integral-field and long-slit spectrographs built for and used with an adaptive optics system. We describe its concept, design, optimal observational procedures and the measured performances. The instrument was used in 1997-2001 at the ESO 3.6\\,m telescope equipped with ADONIS adaptive optics and SHARPII+ camera. The operating spectral range was 1.2\\,-\\,2.5 $\\mu m$. We used the spectral resolution from 500 to 10\\,000 combined with the angular resolution of 0.1{\\arcsec}\\,-\\,0.2{\\arcsec}. The quality of GraF data is illustrated by the integral field spectroscopy of the complex $0.9{\\arcsec} \\times 0.9{\\arcsec}$ central region of $\\eta$ Car in the 1.7 $\\mu m$ spectral range at the limit of spectral and angular resolutions. ",
        "introduction": "%________________________________________________________________ High angular resolution observations at the diffraction limit of the ground based large telescopes were pioneered using the speckle analysis of images \\cite{labeyrie}, however their sensitivity was severely limited by the necessity to keep exposures short in order to ``freeze'' the turbulence induced patterns. The adaptive optics overcame this obstacle allowing long exposures (COME-ON and ADONIS at the ESO 3.6\\,m telescope, \\opencite{adonis97}, PUEO at the CFHT, \\opencite{pueo98}, NAOS at the ESO VLT, \\opencite{lagrange03}, and others).  The limiting magnitude on the telescopes equipped with the adaptive optics (hereafter AO) is now defined by the detector and optics efficiency as for any observations, while the angular resolution is close, at least in the IR, to the telescope diffraction limit, about 0.1{\\arcsec} for a 4\\,m aperture at $\\lambda = 2$\\,$\\mu m$.  The achievements of the AO were firstly exploited for imaging (see reviews by \\opencite{close00}, \\opencite{lai00}, \\opencite{menard00}).  The next step was to use the adaptive optics for spectroscopy.  Indeed, the combination of high angular and high spectral resolutions is a key issue for a number of observational programmes such as physics and evolution of multiple stellar systems, morphology and dynamics of circumstellar gas and jets in young and evolved stars, etc.  Driven by this ideas, the GraF project was started in 1995 with the aim to build an imaging spectrograph optimized for the use with the ADONIS at the ESO 3.6\\,m telescope. The instrument is based on the integral field spectroscopic properties of the Fabry-Perot interferometer (hereafter FPI) used in cross-dispersion with a grating (le Coarer et al., 1992, 1993). It was tested at the telescope in 1997 and successfully used for a number of astronomical programmes in 1998-2001 (for preliminary results see Chalabaev et al., \\citeyear{jena99}, \\citeyear{graf99}, \\citeyear{etacar99}, \\opencite{trouboul99}).  In the present article we describe in details the optical concept and the instrument built to fit the constraints imposed by ADONIS \\cite{adonis97} and the SHARPII+ camera \\cite{sharp95}. We give the account of the observational procedures emerged from our experience as optimal, present the measured performances, and illustrate them by a sample of reduced data obtained at the limit of the instrument possibilities in terms of the angular and spectral resolutions.  The discussion will be restricted to the applications of a high spectral resolution ($R \\simeq 5000$\\,-\\,$10000$) in a moderate passband ($\\Delta \\lambda \\leq 50$) combined with a high angular resolution ($\\simeq$ 0.1{\\arcsec}\\,-\\,0.2{\\arcsec}) in a small field of view ($FOV \\leq 15{\\arcsec}$). The operating spectral range of the instrument was $1.2\\,-\\,2.5$ $\\mu m$.  We hope that this one of the first experiences of the imaging spectroscopy with the adaptive optics (see also \\opencite{tiger95}, \\opencite{lavalley97}) will be useful for designers and users of future spectro-imaging instruments combining the high angular and high spectral resolutions. % %__________________________________________________________________ ",
        "conclusions": "%__________________________________________________________________ % The GraF instrument for imaging spectroscopy based on the Fabry-Perot interferometer in the cross-dispersion with a grating was successfully tested and operated at the 3.6\\,m telescope with the ADONIS adaptive optics, allowing to combine the spectral resolution up to 10\\,000 with the high angular resolution of 0.1{\\arcsec}\\,-\\,0.2{\\arcsec} provided by ADONIS. The ability of simultaneous imaging in several spectral passbands proved to allow in certain cases the simultaneous calibration of the spatial point-spread function ${\\cal{P}}$.  The subsequent deconvolution of data in the crowded field of $\\eta$ Car is shown to be of a high quality, with the final angular resolution close to 0.1{\\arcsec} at $\\lambda = 1.7~\\mu$m, close to the telescope diffraction limit, in spite of the fact that at this rather short wavelength the adaptive optics correction of the atmosphere induced wavefront perturbations was only partial, and the point-spread function was highly variable in time.  The simultaneous mapping of $\\eta$ Car with resolutions R up to 10\\,000 and $\\delta \\phi \\approx$ 0.1{\\arcsec} achieved from the ground with ADONIS/GraF in the 1.7 $\\mu m $ spectral range compares favorably with R up to 5\\,000 and $\\delta \\phi \\approx$ 0.1{\\arcsec} obtained on the same object with the space-born instrument HST/STIS in the visible and far-red spectral range \\cite{hillier01}.  In its practical aspects, the optical concept of the FPI cross-dispersed with a grating proved to be compact and versatile, allowing to switch to the long-slit spectroscopy, or to direct imaging within a minute.  The potential of the GraF concept has been further proven by the construction of new instruments either adopting its principles as the GriF scanning integral-field spectrograph for the CFHT (\\opencite{grif00}, \\opencite{grif02}), or developing them further as the Tunable Echelle Imager (\\opencite{baldry}).  High quality data were obtained during the period of the scientific use of the instrument at the telescope in 1998-2001 on different programmes of stellar physics.  They will be a subject of forthcoming papers. % \\begin {acknowledgements} % It is a pleasure to thank people at Laboratoire d'Astrophysique de l'Observatoire de Grenoble (LAOG) and elsewhere for discussions and help.  F.\\ Lacombe (Paris-Meudon Observatory) wrote additional code of the ADOCAM control software making possible GraF and SHARPII+ interaction; E.\\ Stadler helped with the mechanical design; C.\\ Tran-Thiet, and G.\\ Duvert contributed to the data reduction software, R.\\ Conan and E.\\ Zubia to the numerical simulations, H.\\ Lanteri and C.\\ Aime advised on the data deconvolution, C.\\ Perrier, N.\\ Hubin, J.-L.\\ Beuzit, A.\\ Tokovinin, R.\\ Gredel, S.\\ Gilloteau, R.\\ Petrov, P.\\ L\\'{e}na, J.\\ Melnick, G.\\ Monnet have been advising all along the project advancement; J.-L.\\ Beuzit and A.C.\\ Danks helped during the tests at the telescope, J.\\ Bouvier, T.\\ B\\\"{o}hm, J.\\ Eisl\\\"{o}ffel, C.\\ Dougados contributed to the scientific case of the GraF proposal document, the ESO La Silla staff, in particular J.\\ Roucher, E.\\ Barrios, J.\\ Fluxa, A.\\ Gilliotte, F.\\ Marchis, M.\\ Maugis, V.\\ Merino, P.\\ Prado, A.\\ Sanchez, R.\\ Tighe, was extremely helpful during the GraF tests at the 3.6\\,m telescope.  The funding was provided by CNRS/INSU/MENSR/Programme National de Haute R\u008esolution Angulaire en Astronomie, R\\'{e}gion Rh\\^{o}ne-Alpes, LAOG and Observatoire de Grenoble. Scientific collaborations budget was partially supported by the grants from GdR ``Milieux circumstellaires''.  The telescope technical time was provided by ESO. DLM is thankful to ``Soci\\'{e}t\\'{e} des Amis des Sciences'' for the studentship which made possible his stay at Grenoble. \\end {acknowledgements} % %__________________________________________________________________"
    },
    "0311/astro-ph0311200.txt": {
        "abstract": "We have used the {\\it Hubble Space Telescope} and Wide Field Planetary Camera~2 to image the putative tidal dwarf galaxy located at the tip of the Southern tidal tail of NGC 4038/9, the Antennae.  We resolve individual stars, and identify two stellar populations.  Hundreds of massive stars are present, concentrated into tight OB associations on scales of $200$~pc, with ages ranging from $2$-$100$~Myr.  An older stellar population is distributed roughly following the outer contours of the neutral hydrogen in the tidal tail; we associate these stars with material ejected from the outer disks of the two spirals.  The older stellar population has a red giant branch tip at $I=26.5\\pm 0.2$ from which we derive a distance modulus $(m-M)_0=\\mMo\\pm\\errmMo$.  The implied distance of $\\DMpc\\pm\\errDMpc$ Mpc is significantly smaller than commonly quoted distances for NGC 4038/9. In contrast to the previously studied core of the merger, we find no super star clusters.  One might conclude that SSCs require the higher pressures found in the central regions in order to form, while spontaneous star formation in the tail produces the kind of O-B star associations seen in dwarf irregular galaxies.  The youngest population in the putative tidal dwarf has a total stellar mass of $\\approx \\mYoung M_\\odot$, while the old population has a stellar mass of $\\approx \\mOld ~M_\\odot$. If our smaller distance modulus is correct, it has far-reaching consequences for this proto-typical merger. Specifically, the luminous to dynamical mass limits for the tidal dwarf candidates are significantly less than $1$, the central super star clusters have sizes typical of galactic globular clusters rather than being $1.5$ times as large, and the unusually luminous X-ray population becomes both less luminous and less populous. ",
        "introduction": "\\label{sec:intro}     Interactions and mergers of galaxies are becoming increasingly recognized as perhaps the most important physical process affecting galaxy evolution.  In the context of hierarchical clustering models of galaxy formation, luminous galaxies such as the Milky Way could have been formed as the product of early mergers of less massive galaxies (Kauffmann \\& White \\citeyear{kauffmannWhite93}).  Considering the crucial role that galaxy merging plays in the evolution of massive galaxies in CDM cosmology, it is important to consider also the derivative effects that this merger activity has on the evolution of galaxy populations.  Since the pioneering work of Toomre \\& Toomre (\\citeyear{TT}) it has been known that dramatic tidal tails are among the most striking signs of a merger in progress.  It has been suggested that the material in tidal tails might detach from the parent galaxies and collapse into gravitationally bound clumps, leading to independent objects (so-called Tidal Dwarf Galaxies or TDGs) that might form new stars and be later recognized as dwarf galaxies (Zwicky \\citeyear{zwicky56}; Schweizer \\citeyear{schweizer78}, hereafter S78; Barnes \\& Hernquist \\citeyear{barnes-hernquist92}; Duc et al. \\citeyear{duc_etal97}, Duc \\& Mirabel \\citeyear{duc_mirabel98}).  This process could be more significant at high redshift where the merger rate is dramatically higher.      One of the best cases for a TDG actually in the process of formation is near the end of the long, curving southern tidal tail of the well known nearby merging system NGC 4038/9, ``The Antennae''.  Schweizer (\\citeyear{schweizer78}) was  the first to discuss the tip of the Southern tail in some detail. Making use of stacked photographic images obtained using the CTIO 4-m telescope (reproduced in Fig.~\\ref{fig:dss}), Schweizer noted the presence of a low surface brightness object (\\( V>25 \\) magnitudes~arcsec\\( ^{-2} \\)) slightly \\emph{beyond} the end of the tail (indicated by an arrow in the right panel of Fig.~\\ref{fig:dss}).  He noted that at the adopted distance of $\\sim 28$~Mpc, this feature would have an absolute magnitude \\( M_{V}\\simeq -16.5 \\) making it roughly twice as luminous as the local group dwarf irregular galaxy IC~1613. Furthermore, the then-current WSRT H{\\sc i} maps of this system by van der Hulst (1979) showed a strong concentration of atomic gas coincident with this extension, and Schweizer hypothesized that it was a distinct dwarf stellar system, possibly created during the interaction as envisioned by Zwicky (1956).   \\begin{figure*} \\plotone{Saviane.fig1.ps} \\caption{{\\em Left panel}: The position of the \\noun{wfpc2} field is represented by the black outline over a reproduction from the DSS field (North is up and East to the left).  The bar at the lower left of the figure is \\protect\\( 2\\protect \\) arcmin, and the \\noun{wfpc2} field is \\protect\\( \\sim 10\\protect \\)~arcmin SW from the nuclei of the two galaxies. NGC~4038 is the northern object (plus the southern tidal tail), and NGC~4039 is the southern one (plus the northern tidal tail). The arrow points to the object identified by MDL92 as a TDG. {\\em Right panel}: Reproduction of the CTIO plate in Schweizer (\\citeyear{schweizer78}), where the arrow points to the low surface brightness object he identified as a candidate TDG beyond the tip of the tail. The white outline represents the \\noun{wfpc2} field. \\label{fig:dss} } \\end{figure*}    Besides this possible TDG, S78 observed star formation taking place \\emph{within} the tip of the tail (indicated by an arrow in the left panel of Fig.~\\ref{fig:dss}).  This region of the tail was subsequently studied by Mirabel et al. (\\citeyear{mirabel}; hereafter MDL92). MDL92 described a chain of H{\\sc ii} regions delineating a twisting stellar bar embedded in an envelope of diffuse optical emission.  Optical spectroscopy of three of the H{\\sc ii} regions (denoted as Regions I, II, and III by MDL92) showed them to be ionized by stars as young as 2 Myr (Region I) and 6 Myr (Regions II and III), and to have an oxygen abundance of \\( 12+\\log (\\rm O/H)=8.4 \\) (i.e. $\\simeq 1/3$ solar). The total luminosity of the three H{\\sc ii} regions was found to be \\( \\sum L(H\\alpha ) = 3.9\\times 10^{39}\\)~erg\\, s\\(^{-1}\\), and the apparent magnitude of the entire region is \\( V=17.3\\pm 0.3 \\), corresponding to \\( M_{V}=-15.5 \\) (all values referenced to their assumed distance of \\( 33.2 \\) Mpc, i.e. \\( m-M=32.6 \\)). These properties, along with the atomic gas mass of \\( M(HI)= 10^{9} \\) \\( M_{\\odot } \\) and derived dynamical mass of \\( 8.3\\times 10^{9}\\, M_{\\odot } \\), are similar to those of dwarf irregular galaxies, and MDL92 concluded that this regions was in fact a dwarf irregular galaxy forming out of the tidal debris.  More recent radio observations have mapped the total extent of the H{\\sc i} within the tidal tails (Gordon et al. \\citeyear{gordon_etal01}, Hibbard et al. \\citeyear{hibbard_etal01}). These observations show that, while the density of the atomic gas is enhanced in the vicinity of the putative TDGs identified by S78 and MDL92, the entire tail is a continuous gas rich structure. The gas kinematics suggest that the tail bends back along our line of sight just in the vicinity of the TDG candidates, giving rise to the appearance of a distinct clump of gas associated with the putative tidal dwarfs. No clear kinematic signature of a self-gravitating entity of the mass suggested by MDL92 was seen, but the geometric velocity gradients caused by the tail bending away from our line of sight may mask the kinematic signature of smaller self-gravitating condensations (Hibbard et al. \\citeyear{hibbard_etal01}). Even though the dynamical nature of the TDG candidates is still in question, the H{\\sc i} observations show that the atomic gas in the vicinity of star forming regions I, II, and III of MDL92 is denser than anywhere else in the system.  To study further the star forming regions identified by S78 and studied by MDL92, and to examine the distribution and ages of the underlying old and young stellar populations of the TDG candidate identified by MDL92, we targeted this region for broadband UBVI observations with Wide Field and Planetary Camera~2 (\\noun{wfpc2}) on board the Hubble Space Telescope (HST). Figure~\\ref{fig:dss} illustrates the location of our single \\noun{wfpc2} field relative to NGC 4038/9. In addition to investigating the population of a TDG candidate, these observations also afford us the chance to investigate star formation taking place in an environment relatively isolated from the large scale structure of a disk, in the far outskirts of a galactic potential well. We stress that a study of the S78 object with the recently installed ACS would be valuable as well, as the population of this object may well sample some of the original disk population of NGC 4038 and additionally give an improved red giant branch tip distance test of our proposed short modulus.  \\begin{figure*} \\plotone{Saviane.fig2.ps}  \\caption{True-color image of our HST field, where red has been assigned to the F814W image, blue to the F555W, and green to an image made from a combination of both filters. The extent and orientation of the field is outlined in Fig.~\\ref{fig:dss}. The three star formation regions studied by Mirabel et al. (\\citeyear{mirabel}) are marked along the bar-like structure that extends westward from the two bright stars. \\label{fig:hst-mosaic}} \\end{figure*}  A color image of the HST mosaic is shown in Fig.~\\ref{fig:hst-mosaic}, and a visual inspection of the image gives a clear impression of the population, even without the formality of a color-magnitude diagram (CMD). The light of the tidal tail region imaged is dominated by bright blue stars indicative of a young stellar population (which contrast strongly with the reddish background galaxies).  Fainter and redder stars indicative of an older population can also be seen.  The youngest and brightest blue stars are located along a bar-like structure that extends from the two bright foreground stars in the upper right corner of the WF3, crosses the upper left corner of the WF4, and ends in the lower right corner of the PC. These stars form groups that look like OB associations, and the three most prominent ones coincide with the location of the three H{\\sc ii} regions studied by MDL92. Following that study, we refer to these regions as Regions {\\sc I, II, III}, and have labeled them as such in Fig.~\\ref{fig:hst-mosaic}. Besides these active sites of star formation, another concentration of young stars can be seen in the WF3 chip, East of the two foreground stars.   The paper is organized as follows. The data reduction and calibration are presented in \\S\\ref{sec:obs-reduc}, and the resulting CMD is used in \\S\\ref{sec:distance} to find the distance to the Antennae. We interpret the CMD with the aid of that of NGC~625 (a star forming dwarf irregular galaxy) and theoretical isochrones, and the distance modulus is then found using the luminosity of the tip of the red giant branch of the old population. The stellar content of the NGC~4038 tail is discussed in \\S\\ref{sec:star-content}, and in particular, we discuss the spatial distribution of populations as a function of age in \\S\\ref{sec:space-age}. A number of stellar associations are defined, and their morphology is compared to the central star clusters in \\S\\ref{sec:recent-sf}. Their CMD is interpreted in \\S\\ref{sec:comp-isos}, and an estimate of the mass in young and old stars is obtained in \\S\\ref{sec:youngstars} (plus Appendix~\\ref{sec:estimating_young}) and \\S\\ref{sec:fornax}, respectively.  In order to interpret the luminosity function of the young population, a model of constant and continuous star formation is developed in Appendix~\\ref{sec:simple-sf}, and the luminosity function of the old population is compared to that of a Virgo dwarf elliptical in Appendix~\\ref{sec:vcc1104}. Our conclusions are given in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  In this paper we have analyzed the stellar content of a region of linear dimension $\\sim 8$~kpc, in the southern tail of the Antennae system of interacting galaxies. According to dynamical simulations (e.g. Barnes 1988), the material of the tail has been pulled out of NGC~4038 by the tidal field of NGC~4039 $\\sim$ 450 Myr ago (for Milky Way sized progenitors), and previous groundbased studies of this area have shown that stars are now actively forming in a few regions (S78, MDL92).  The \\noun{wfpc2} camera on board of the HST has now been used to obtain, for the first time, broadband $F555W$ and $F814W$ photometry and astrometry of more than $4000$ stars within the region studied by  MDL92.   A population of faint red stars is seen in the CMD. This population shows no clumping in its spatial distribution, and it was presumably ejected from the outer disks of the interacting galaxies. We have isolated a subsample of this population by selecting only stars from regions far from the active star forming areas. The selected sample shows a sharp cutoff in the LF at $I_0<26.5$. The LF is almost identical to that of the low surface brightness component of the star forming dwarf irregular galaxy NGC~625, and based on this comparison we interpreted the cutoff as the RGB tip of a few Gyr old population. This yields a distance to the Antennae of $\\DMpc \\pm \\errDMpc$~Mpc, much shorter than the typically quoted $\\sim 20$~Mpc. The choice of a shorter distance is reinforced by comparing the LF of this population to that of VCC1104, a dwarf elliptical galaxy in the Virgo cluster (see Appendix~\\ref{sec:vcc1104}), and to that of the Fornax dwarf spheroidal galaxy in the Local Group.  Our imaging with \\noun{wfpc2} along with the VLA neutral hydrogen map of Hibbard et al. (\\citeyear{hibbard_etal01}) gives a better picture of the stellar populations at the end of the Southern tail of the Antennae.  We find that the regions of most active star formation coincide with the H{\\sc i} areas of highest density ($> 8\\times 10^{20}~\\rm cm^{-2}$), and that the brightest young star forming associations are similar to those seen in dwarf irregular galaxies. In particular, there are two elongated structures that contain most of the associations, and one of them is the ``bar'' identified by MDL92. Our color-magnitude diagrams confirm the intuitive picture that the smooth structure of the tail is composed of older stars from the disks of NGC 4038/9 that were launched along with the gas, due to the interaction.  The youngest stars are mostly concentrated in OB associations, but can also be found scattered over the whole region.     In contrast to the inner regions, there are no ``super star clusters'' seen in this region of the tail. The young tidal associations are more diffuse and have  orders of magnitude fewer stars than the central star clusters. The eight most prominent associations have been used to study the recent star formation in the area. We are able to measure the luminosity function of their main sequence to reach stars as faint as \\( M_{V}=-4.2 \\) corresponding (roughly) to ages of $20$~Myr and masses of \\( 14\\, M_{\\odot } \\) (in the hypothesis of a short distance modulus).  In this range, it would appear that the star formation has been continuous.  Considering the stellar mass budget of the whole area that we studied, by comparing the LF of the red-low surface brightness component to that of the Fornax dSph galaxy, we find that most of the mass ($\\approx 6.7\\times 10^7~M_\\odot$) is contained in an old ($> 2 $~Gyr) stellar population, of order the stellar mass of Fornax itself.  With a simple model of constant star formation we show there is a further $\\approx 2\\times 10^5 M_\\odot$ of stars formed {\\it in situ} (including young stars outside the eight associations), within the last $20$~Myr.   Detailed investigations of the star formation process occurring well outside the boundaries of a disk or the potential well of a galaxy are rare.  Mould et al. (1999) images the radio jet-induced star formation near NGC 5128, and finds a population similar to ours, with young OB associations but no massive star clusters.  Clearly, star formation requires only the presence of gas of sufficient density, and not the large scale organizing gravitational potential of a disk, spiral arm, or galaxy.  One of the most important results of this work is the downward revision of the distance to the Antennae, which follows the first determination of the distance modulus based on a standard candle. This revision has consequences for the the nature of the TDG candidates, for the stellar populations within the inner regions, for the dynamical nature of the system, and for the dynamical models of the local Universe.    With respect to the TDG candidates, Hibbard et al. (\\citeyear{hibbard_etal01}) used the H{\\sc i} kinematics to explore the dynamical nature of the regions studied by both MDL92 and S78. Specifically, they evaluate the ratio of the luminous mass, as inferred from the H{\\sc i} and optical luminosity, and the dynamical mass inferred from the H{\\sc i} line width and physical radius. They derive ratios of $M_{lum}/M_{dyn} \\sim 0.3-0.7$ for a distance of 19.2 Mpc, where the range indicates stellar mass-to-light ratios ranging from $M_*/L_B =$ 0--5 (the analysis in \\S\\ref{sec:mass-ssp} suggests a mass-to-light ratio near the bottom of this range). Since the baryonic mass scales as the square of the distance while the dynamical mass scales linearly with distance, a distance modulus of \\mMo\\ requires revising these estimates downward by a factor of 1.4 (i.e. $M_{lum}/M_{dyn} \\sim 0.2-0.5$). If either of these regions are self-gravitating entities, i.e. {\\it bona fide} tidal dwarf galaxies, then most of their mass must be in the form of dark matter. This is quite hard to understand if they formed from disk material. Rather than witnessing the formation of a single large dwarf galaxy, we find it more likely that we are seeing several smaller scale regions forming stars. The mass scales of these regions are more typical of dwarf spheroidal galaxies than dwarf irregulars, and they may have more relevance to remnant streams around galaxies than the formation of dwarf irregular galaxies (Kroupa 1998).   With respect to the populations believed to result from the merger induced star formation taking place in the inner regions, all linear dimensions quoted in W99 would be reduced by a factor of $1.4$, so the effective radii of the youngest clusters created in the merger would now be comparable to those of Galactic globulars (W99 find a median effective radius $\\sim 1.5$ times larger than that of GGCs). The young clusters' LF would of course still be a power law, however, the bend found by W99 at $M_V\\approx -10.4$ would now be a factor $2$ fainter in luminosity and the corresponding mass would be smaller, assuming the same $M/L$ value. W99 identified the bend with the characteristic mass of Milky Way globular clusters, although their value of $1\\times 10^5~M_\\odot$ is smaller than that of GGCs ($\\approx 2\\times 10^5~M_\\odot$). With the shorter distance modulus, the mass of the bend would now be $\\approx 5\\times 10^4~M_\\odot$, making this identification weaker.    With a smaller distance, the ``Ultra Luminous'' X-ray source population discovered by Fabbiano, Zezas \\& Murray (\\citeyear{fabbiano_etal2001}; see also Zezas et al. \\citeyear{zezas_etal2002}, Zezas \\& Fabbiano \\citeyear{zezas_fabbiano2002}), while still extreme in its properties, becomes both less luminous (peak X-ray luminosity of $L_X=5\\times10^{39}$ ergs s$^{-1}$ vs. $2\\times10^{40}$ ergs s$^{-1}$) and less populous (6 instead of 18 with $L_X>10^{39}$ ergs s$^{-1}$). The X-ray sources in the Antennae are still luminous, but the new distance makes their luminosity function comparable to those found in other starburst galaxies (cf. Zezas \\& Fabbiano \\citeyear{zezas_fabbiano2002}).  Finally, with respect to the dynamical models of the local Universe, we compare the predicted and observed radial velocity of the Antennae at the new distance, using the Tonry et al. (\\citeyear{tonry_etal00}) galaxy flow model of the local Universe, and find that the model predicts a velocity $500 \\kms$ smaller than that of NGC~4038/9. The discrepancy would hold even if we used the W99 distance (it would then be $>300 \\kms$).  Considering the importance of the distance estimate for a physical understanding of the Antennae, we would urge that the distance modulus be confirmed with deeper HST/ACS observations.   \\begin{figure}  \\plotone{Saviane.fig17.ps}  \\caption{Interpreting the LF of a composite stellar population (heavy solid line). We assume that it is made by the convolution of several LFs of simple stellar populations (thin inclined lines). The single populations are generated at a constant rate $\\dot n$, and stars brighter than the termination of the main sequence are removed from the LFs (dotted lines).  The figure has been obtained for $\\dot{n} =0.21\\, \\, \\, \\rm Myr^{-1} $, from $100$~Myr to $4$~Myr ago.  The labels identify parameters that are used in the text to develop our model. \\label{fig:sf-labels} } \\end{figure}"
    },
    "0311/astro-ph0311106_arXiv.txt": {
        "abstract": "We present the results of a search for clusters in dwarf irregular galaxies in the Virgo and Fornax Cluster using HST WFPC2 snapshot data.  The galaxy sample includes 28 galaxies, 11 of which are confirmed members of the Virgo and Fornax clusters.  In the 11 confirmed members, we detect 237 cluster candidates and determine their V magnitudes, V-I colors and core radii.  After statistical subtraction of background galaxies and foreground stars, most of the cluster candidates have V-I colors of -0.2 and 1.4, V magnitudes lying between 20 and 25th magnitude and core radii between 0 and 6 pc. Using H$\\alpha$ observations, we find that 26\\% of the blue cluster candidates are most likely HII regions.  The rest of the cluster candidates are most likely massive ($>$10$^4$ M$_\\odot$) young and old clusters.  A comparison between the red cluster candidates in our sample and the Milky Way globular clusters shows that they have similar luminosity distributions, but that the red cluster candidates typically have larger core radii.  Assuming that the red cluster candidates are in fact globular clusters, we derive specific frequencies (S$_N$) ranging from $\\sim$0-9 for the galaxies.  Although the values are uncertain, seven of the galaxies appear to have specific frequencies greater than 2.  These values are more typical of ellipticals and nucleated dwarf ellipticals than they are of spirals or Local Group dwarf irregulars. ",
        "introduction": "The study of extragalactic globular cluster systems in galaxies of a variety of Hubble types has led to the development of a number of scenarios that place globular clusters within the context of galaxy formation and evolution.  \\citet{carney01} and \\citet{vandenbergh00b} reviewed the myriad pieces of evidence derived from globular cluster systems and the scenarios designed to explain them.  The properties of globular clusters in giant elliptical galaxies, in particular, have led to suggestions that globular cluster systems form 1) {\\it in situ}, perhaps in the hierarchical collapse of the halo \\citep{cote00} and possibly in multiple phases \\citep{forbes97}; 2) in galaxy mergers, both through the formation of new clusters in the colliding gas \\citep[e.g.][]{schweizer96} and through the combination of the pre-existing globular cluster systems \\citep[e.g.][]{forbes00}; 3) by accretion of the clusters of smaller galaxies \\citep{cote98}, as exemplified by the $\\sim$4 GCs being contributed to the Milky Way \\citep{dacosta95} by the Sagittarius dwarf galaxy \\citep{ibata94}.  More recently the second two of these ideas have been incorporated into the current cosmological scenario using simulations \\citep[e.g.][]{kravtsov03} and semi-analytic methods \\citep[e.g.][]{beasley02}.  The generally low globular cluster specific frequencies \\citep[$S_N$, the number of clusters normalized by the galaxy magnitude;][]{harris81} observed in nearby dwarf and spiral galaxies pose a difficulty for the third of the above possibilities, a point which was raised by \\citet{vandenbergh82} in argument against spiral-spiral mergers explaining the high $S_N$ of cD galaxies. However, it may be misleading to compare the specific frequences of spiral and dwarf galaxies residing in nearby low-density galaxy groups with those of giant elliptical galaxies found in more distant rich clusters of galaxies, since the role of the environment in the assembly of globular cluster systems is not understood.  A census of the globular cluster systems of a wide variety of galaxies drawn from the same environment is needed to properly assess the roles that each of the formation mechanisms may play.  \\citet{miller98} embarked on such a census through the study of the globular clusters of dwarf elliptical (dE) galaxies in the Virgo, Fornax, and Leo clusters.  The results of that survey revealed a distinction between the globular cluster populations of nucleated dE (dE,N) and nonnucleated dE galaxies; in particular, dE,N galaxies were found to have a factor of 3 higher $S_N$ than nonnucelated dEs.  The high $S_N$ of dE,N align them more closely with giant elliptical galaxies than with dwarf irregular (dIrr) and spiral galaxies, which from measurements in {\\it local} galaxies have $S_N\\sim0.5$ \\citep{harris91}; the nonnucleated dE's, on the other hand, {\\it do} have $S_N$ that resemble those of dIrrs and spiral galaxies, once fading of the young stellar populations is taken into account. The difference in $S_N$ between dE and dE,N galaxies thus suggests globular cluster formation processes in dE,N and giant ellipticals that are different from those of spirals and dIrrs.  Moreover, it argues that not all dE galaxies represent the stripped or harassed remnants of dIrr galaxies, processes described by e.g. \\citet{lin83} and \\citet{moore98}.  These conclusions depend, however, to large degree on the validity of comparing globular cluster populations in rich-cluster dE galaxies with those of local dIrr and spiral galaxies.  In this paper, we discuss the results of a {\\it Hubble Space Telescope} (HST) snapshot survey of dIrr galaxies in the Virgo and Fornax galaxy clusters, the same environments from which many of the dE galaxies of \\citet{miller98} were drawn.  We describe the observational sample and data reduction in \\S2.  In \\S3, we describe our technique for detecting candidate globular clusters.  We present the properties of our sample of candidate globular clusters, derive specific frequencies, and examine individual galaxies in \\S4.  Some of the more interesting results are discussed in \\S5, followed by our conclusions in \\S6. ",
        "conclusions": "We have studied the cluster systems of dwarf irregular galaxies using HST/WFPC2 observations of 28 galaxies in the Virgo and Fornax clusters, 11 of which are confirmed members.  In these 11 confirmed members we have found 237 cluster candidates and determined their magnitudes, core radii, and V-I colors.   Based on source counts of neighboring, empty, WFPC2 fields, we expect a contamination of 51 background and foreground sources, which were statistically subtracted from the data.  The cluster candidates were divided by color in order to separate red (with age $\\gtrsim$ 1 Gyr) and blue (age $\\lesssim$ 1 Gyr) cluster candidates.  We consider the red cluster candidates to be globular cluster candidates.  The main results presented here are: \\begin{enumerate} \\item The red cluster candidates have the same magnitudes as expected for globular clusters.  More specifically, the luminosity function of Galactic globular clusters \\citep{harris96} is well matched by the luminosity function of our red cluster candidates down to our 90\\% completeness limit at M$_{\\rm V} \\sim$ -7. \\item The median color (V-I = 1.0) of the red cluster candidates is bluer than the median colors of globular cluster systems in elliptical galaxies \\citep{kundu01}, even without a correction for internal reddening.  This is consistent with the globular cluster systems in these dwarf irregular galaxies being metal-poor. \\item The distribution of core radii for the red candidate clusters is significantly larger than for Galactic globular clusters.  Larger radii are also seen in the globular clusters of the LMC and other local group dwarves \\citep{vandenbergh91} suggesting that this might be a general property of globular clusters in dwarf galaxies. \\item The blue cluster candidates are in general more compact and fainter than the red cluster candidates.  Based on single-stellar population models, they have masses ranging between 10$^4$ and 10$^5$ M$_\\odot$.  H$\\alpha$ data on 7 Virgo cluster galaxies suggests that 26\\% of the blue cluster candidates are in fact HII regions. \\item The specific frequencies for the galaxies exhibit a large scatter, but are surprisingly large, with 7 of the 11 galaxies having specific frequency values greater than two.  These values are surprising when compared to Local Group dwarf irregulars and other nearby late-type galaxies, which have specific frequencies less than two. We note, however, that our specific frequency values are uncertain due to the possible presence of reddened young clusters, massive intermediate-age clusters and incompleteness.  \\end{enumerate}  We plan to follow up this work with spectroscopic observations or infrared imaging of some of these cluster systems.  These observations of the cluster candidates will test the accuracy of our photometric results by enabling accurate age and metallicity determination.  This work was supported by the National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.  It also has been supported by the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., on behalf of the international Gemini partnership of Argentina, Australia, Brazil, Canada, Chile, the United Kingdom, and the United States of America.  The authors would like to thank: Paul Hodge, for his guidance and help observing, Peter Stetson for help with DAOPHOT II, Eduardo Bica, for providing us with his tables, Nick Suntzeff for his helpful comments, Ricardo Covarrubias, and Armin Rest.  \\clearpage"
    },
    "0311/astro-ph0311276_arXiv.txt": {
        "abstract": "We address the issue of the production of Lyman continuum photons by T Tauri stars, in an attempt to provide constraints on theoretical models of disc photoionisation.  By treating the accretion shock as a hotspot on the stellar surface we show that Lyman continuum photons are produced at a rate approximately three orders of magnitude lower than that produced by a corresponding black body, and that a strong Lyman continuum is only emitted for high mass accretion rates.  When our models are extended to include a column of material accreting on to the hotspot we find that the accretion column is extremely optically thick to Lyman continuum photons.  Further, we find that radiative recombination of hydrogen atoms within the column is not an efficient means of producing photons with energies greater than 13.6eV, and find that an accretion column of any conceivable height suppresses the emission of Lyman continuum photons to a level below or comparable to that expected from the stellar photosphere.  The photospheric Lyman continuum is itself much too weak to affect disc evolution significantly, and we find that the Lyman continuum emitted by an accretion shock is similarly unable to influence disc evolution significantly.  This result has important consequences for models which use photoionisation as a mechanism to drive the dispersal of circumstellar discs, essentially proving that an additional source of Lyman continuum photons must exist if disc photoionisation is to be significant. ",
        "introduction": "The evolution and eventual dispersal of discs around young stars is an important area of study, as it provides constraints on theories of both star and planet formation.  It is now well established that the majority of young stars at ages of $\\sim 10^6$yr have circumstellar discs which are optically thick at optical and infrared wavelengths \\citep*{strom89,kh95,haisch01}.  These discs are relatively massive, with masses of the order of a few percent of a solar mass \\citep*{beckwith90,ec03}.  However by an age of $\\sim 10^7$yr most stars are no longer seen to have such massive discs, although low-mass ``debris discs'' may remain (eg.~\\citealt*{ms97,wyatt03}).  The mechanism by which these discs are dispersed remains an important unsolved question.  One fact which is clear, however, is that a large fraction of the mass from the disc will eventually be accreted on to the central (proto)star. The inner edge of the accretion disc is typically truncated by the magnetosphere at a radius of around 5$R_*$ \\citep*{mch97,cg98} and so material falling from the disc on to the central star can attain an extremely high velocity, resulting in a so-called ``accretion shock'' when this material impacts upon the stellar surface.  Existing models of the accretion shock \\citep*{cg98,gul98,lamzin98,gul00} have paid a great deal of attention to the emission in the ultraviolet (1000-3000\\AA) and visible (3500-7000\\AA) wavebands, comparing theoretical predictions to observed spectra.  Extremely good models have been constructed, and these emission spectra are now well understood.  However, very little attention has been paid to the emission shortward of the Lyman break ($<$ 912\\AA), primarily because absorption by interstellar H{\\sc i} makes it impossible to observe young stars in this wavelength regime.  Recent theoretical studies \\citep*{cc01,mjh03,acp03} have suggested that photoionisation by the central object may play an important role in disc dispersal, and so the emission shortward of the Lyman break has become important.  Currently the origin of photoionising radiation from young stars such as T Tauri stars is unclear, and even the magnitude of such emission is poorly constrained \\citep{gahm79,ia87}.  To date, models of disc photoionisation have used either a constant ionising flux (eg.~\\citealt{holl94,cc01}), assumed to be chromospheric in origin, or modelled the accretion-driven flux simply as a constant temperature hotspot on the stellar surface, emitting as a blackbody \\citep{mjh03}.  The latter produces an ionising flux which is proportional to the mass accretion rate, thus decreasing dramatically with time as the accretion rate falls, and so naturally the models of \\citet{cc01} and \\citet{mjh03} have produced markedly different results.  These models \\citep*{holl94,ry97,hollppiv} also find that the mass-loss rate scales approximately with the square root of the ionising flux, and so very high ionising fluxes are required to influence disc evolution significantly: an ionising flux of $10^{41}$photons s$^{-1}$ will drive mass-loss from the disc at around the $10^{-10}$$M_{\\odot}\\mathrm{yr}^{-1}$ level.  In this paper we study the issue of the ionising flux generated by an accretion shock.  The simplified model of the accretion shock adopted by \\citet{mjh03} neglects two key points, which we address in turn.  Firstly, it seems likely that the Lyman continuum emission from such a hotspot would resemble a stellar atmosphere rather than a blackbody; whilst the two are almost identical at longer wavelengths, photo-absorption by H{\\sc i} provides a strong suppression of the flux shortward of 912\\AA.  Secondly, the photons emitted by the accretion shock must pass {\\it through} the column of material accreting on to it in order to interact with material in the disc, and again we would expect photo-absorption by H{\\sc i} in the column to supress the Lyman continuum significantly.  In order to address these issues we have modelled this process in some detail.  In section \\ref{sec:atmos} we investigate the effect of replacing the blackbody hotspot with a more realistic stellar atmosphere.  In section \\ref{sec:columns} we investigate the effect of passing these photons, from both the blackbody and the stellar atmosphere, through a column of accreting material.  In section \\ref{sec:dis} we discuss our results and the limitations of our analysis, and in section \\ref{sec:summary} we summarise our conclusions. ",
        "conclusions": "\\label{sec:dis} There are obvious caveats to these models.  The first, and most significant, is that the {\\sc cloudy} code is only completely reliable at densities less than $10^{13}$cm$^{-3}$; it is prone to numerical problems at higher densities, and for cases of high $\\dot{M}_{\\mathrm {d}}$ and small $f$ the density in our models can exceed this value.  However, the main uncertainties at these densities are regarding the treatment of heavy elements.  The dominant effects in the regime in which we are interested are photoionisation and recombination of atomic hydrogen, and these processes {\\it are} treated reliably by the code.  However in order to compensate for this numerical problem we were forced to limit the density artificially to have a maximum value of $5\\times10^{13}$cm$^{-3}$.  This affected the 3 models with the highest densities ($\\dot{M}_{\\mathrm {d}} = 1\\times10^{-7}$$M_{\\odot}\\mathrm{yr}^{-1}$, $f=0.01,0.05$ and $\\dot{M}_{\\mathrm {d}} = 5\\times10^{-8}$$M_{\\odot}\\mathrm{yr}^{-1}$, $f=0.01$), and the reduction in density results in these models under-estimating the ionising flux somewhat.  This does introduce some uncertainty into our models, but we consider the effect to be small relative to the gross effects which dominate the calculations.  A further caveat regards the issue of non-radial emission, both from the hotspot and from the ``sides'' of the column, as mentioned in section \\ref{sec:columns}.  The ionising photon rates from our models are those emitted from the top of the accretion columns only, and neglect emission from the sides of the columns.  Further, the incident flux from the hotspot is assumed to be purely radial.  The Lyman continuum emission in which we are interested arises from the radiative recombination of hydrogen atoms, an intrinsically isotropic emission process, and so emission from the sides of the column could be significant.  However, as the recombination process is isotropic it is reasonable to assume that the Lyman continuum emitted from the sides of a column will be comparable to that emitted from the top of a column that is truncated at a height equal to its diameter.  Fig.\\ref{fig:phi} shows that the diffuse Lyman continuum from the top of the column decreases dramatically as the column height increases (and the density decreases).  Consequently the Lyman continuum emission from the sides of the column will only be significant over a distance comparable to the hotspot diameter: the sides of the column only produce a significant Lyman continuum near to the stellar surface.  As noted in section \\ref{sec:columns}, Lyman continuum photons incident on the columns will be attenuated by a factor of $10^{13}$ over a distance of approximately 10$^{-5}$$R_{\\odot}$.  As a result of this, the majority of the non-radial Lyman continuum photons emitted by an isotropically emitting hotspot will be absorbed by the column.  The only such photons not absorbed will be those emitted within a fraction of an attenuation length of the edge of the hotspot, in a direction away from the centre of the hotspot.  A hotspot covering 1\\% of the stellar surface has a radius of 0.2$R_{\\odot}$, so less than $10^{-5}$ of the hotspot photons are emitted within 10$^{-5}$$R_{\\odot}$ of the hotspot edge.  Given this fact, and also the behaviour of the ``raw'' hotspot emission described in Fig.\\ref{fig:atmos}, we neglect this effect.  The net result of these two simplifications is that in reality the emission from the bottom of any accretion column will dominate the Lyman continuum emitted by the column, and accretion columns of any height will emit ionising photons at a rate comparable to that provided by the lower part of the column (the left-hand end of the curves in Fig.\\ref{fig:phi}).  This will increase the largest calculated ionising photon rates by a small geometric factor but still cannot increase their flux to significantly greater than 10$^{31}$photons s$^{-1}$; the photospheric emission will still dominate the overall Lyman continuum.  More importantly, much higher ionising fluxes (of the order of 10$^{41}$photons s$^{-1}$) are required to have a significant effect on disc evolution \\citep{cc01,mjh03}, and the uncertainties caused by the approximations we have made are negligible compared to this 10 orders-of-magnitude difference.  Further simplifications used in our models regard the geometry of the accretion column.  Our models use a column with a constant covering factor - essentially a truncated radial cone - and so the area of the outer surface at a given radius $R$ is proportional to $R^2$.  However in reality the column is channelled by the magnetosphere and, as explained in section \\ref{sec:columns}, has an area proportional to $R^3$; our model accretion columns are somewhat less flared than we would expect to see in reality.  However the difference between the two is only significant at large radii; small column heights, which provide the highest ionising fluxes, will not show significant deviation between the two cases.  Again, the net result of this is that we probably under-estimate the ionising flux by a small factor, but not by enough to alter the results significantly.  As our accretion columns are radial cones they do not bend to follow the magnetosphere, as expected in more realistic magnetospheric accretion models (eg.~\\citealt{gl78}).  In such a model the column would bend over to meet the accretion disc, with a curvature dependent on both the latitude of the hotspot and the strength of the magnetosphere.  However, as discussed above, the emission from the bottom part of the accretion column dominates over that from the upper parts (those affected by this curvature), so this simplification will not affect our results significantly.  There is also the issue of an infall velocity, which our models do not address.  In reality the accretion column will be falling towards the stellar surface at close to the free-fall velocity, which can be several hundred kms$^{-1}$, and so this could modify the absorbing effect of the cloud.  However the infall velocity is much less than the speed of light, and there are no strong emission lines near to the Lyman break, and so we consider the impact of this effect on the emitted Lyman continuum to be negligible.  It should be noted that our model makes no predictions as to the behaviour of the photons emitted at wavelengths longward of the Lyman break.  As shown in Fig.\\ref{fig:spectra}, the emission from the top of the column longward of the Lyman break is essentially identical to that from the hotspot at the base of the column.  We approximate the accretion shock crudely, and so are not able to fit our models to observed spectra in a manner similar to \\citet{cg98} or \\citet*{jkvl00}.  However the Lyman continuum emitted by our columns is insensitive to variations in the hotspot spectrum, due to the high optical depth of the columns to Lyman continuum photons.  Consequently we find that any reasonable accretion-shock model will produce a similar Lyman continuum, and that this Lyman continuum is independent of the emission at longer wavelengths.  We have adopted stellar parameters of $R=1R_{\\odot}$ and $M=1M_{\\odot}$ to provide direct comparisons with the models of \\citet{cc01} and \\citet{mjh03}.  However in the case of T Tauri stars a radius of $2R_{\\odot}$ and a mass of $0.5M_{\\odot}$ would be more realistic \\citep{gul98}.  The result of this will be that our models over-estimate the ionising flux somewhat, due to both reduction in the energy released by accretion on to the stellar surface and also due to a reduction in the star's surface gravity.  Once again, however, it is unlikely that these factors are significant in comparison to the gross effects we have already considered. Similarly, the use of the ``constant temperature'' hotspot to heat the column probably over-estimates both the ionising flux and heating provided by the hotspot, and thus over-estimates the diffuse Lyman continuum.  In effect we have constructed a ``best-case'' model, designed to produce the maximum ionising flux, and still found the ionising flux to be less than that emitted by the stellar photosphere.  It seems extremely unlikely that any conceivable accretion column could produce ionising photons at a rate significantly greater than this.    In an attempt to provide some constraints on the nature and magnitude of the ionising continuum emitted by T Tauri stars, we have constructed models which treat the accretion shock as a hotspot on the stellar surface beneath a column of accreting material.  We have modelled these columns for a variety of different accretion rates, hotspot sizes and column heights, and have found that: \\begin{itemize} \\item A hotspot radiating like a stellar atmosphere radiates ionising photons at a rate some 3 orders of magnitude less than the corresponding blackbody. \\item A constant area hotspot radiating like a stellar atmosphere can only emit ionising photons at greater than photospheric rates for mass accretion rates greater than $10^{-7}$$M_{\\odot}$yr$^{-1}$.  Such accretion rates are near the upper limit of the rates derived from observations \\citep{hcga98,jkvl00}. \\item Photoionisation of neutral hydrogen in the accretion column attenuates the Lyman continuum from any hotspot to zero over a very short length scale.  The ionising photons which do emerge are due to radiative recombination of hydrogen atoms in the column, and the rate of ionising photon emission is less than the photospheric level for all of the accretion columns we have modelled. \\end{itemize} In short, we find that accretion shocks and columns are extremely unlikely to produce Lyman continuum photons at a rate significantly greater than that expected from the stellar photosphere.  The photospheric level itself is some 10 orders of magnitude below the rates required for photoionisation to affect disc evolution significantly, and so it seems that the Lyman continuum emitted by an accretion shock will not be large enough to be significant in disc photoionisation models.  These models have provided an attractive explanation of some observed disc properties (eg.~\\citealt{cc01,acp03}) but we have shown, as suggested by \\citet{cc01}, that they must be powered by something other than the accretion-shock emission.  In order to produce ionising photons at a rate large enough to enable photoionisation models to provide a realistic means of disc dispersal, the central objects must sustain a source of ionising photons that is not driven by accretion from the disc."
    },
    "0311/astro-ph0311040_arXiv.txt": {
        "abstract": "{ We combine recent ISO observations of the vibrational ground state lines of H$_2$ towards Photon-Dominated Regions (PDRs) with observations of vibrationally excited states made with ground-based telescopes in order to constrain the formation rate of H$_2$ on grain surfaces under the physical conditions in the layers responsible for H$_2$ emission. We use steady state PDR models in order to examine the sensitivity of different H$_2$ line ratios to the H$_2$ formation rate $R_{f}$. We show that the ratio of the 0-0 S(3) to the 1-0 S(1) line increases with $R_{f}$ but that one requires independent estimates of the radiation field incident upon the PDR and the density in order to infer $R_{f}$ from the H$_2$ line data. We confirm the earlier result of \\cite{habart2003a} that the H$_2$ formation rate in regions of moderate excitation such as Oph W, S140 and IC 63 is a factor of 5 times larger than the standard rate inferred from UV observations of diffuse clouds. On the other hand, towards regions of higher radiation field such as the Orion Bar and NGC 2023, we derive H$_2$ formation rates consistent with the standard value. We find also a correlation between the H$_2$ 1-0 S(1) line and PAH emission suggesting that $R_f$ scales with the PAH abundance. With the aim of explaining these results, we consider some empirical models of the H$_2$ formation process.  Here we consider both formation on big ($a\\sim$0.1 $\\mu$m) and small ($a\\sim$10 \\AA) grains by either direct recombination from the gas phase or recombination of physisorbed H atoms with atoms in a chemisorbed site. We conclude that indirect chemisorption is most promising in PDRs. Moreover small grains which dominate the total grain surface and spend most of their time at relatively low temperatures may be the most promising surface for forming H$_2$ in PDRs. ",
        "introduction": "The formation of molecular hydrogen is a key process affecting the thermal and density structure and the chemical evolution of the interstellar medium \\cite[see, for example,][]{combes00}. Although there is a consensus that H$_2$ forms on the surface of dust grains \\cite[]{gould63,Hollenbach71,jura75,duley84}, the mechanism is not yet understood. This is partly due to our ignorance concerning interstellar grain composition, form, structure and physico-chemical state. It is also caused by our lack of understanding of surface reactions in the interstellar context.  Numerous theoretical and experimental studies have thus been dedicated to the study of the H$_2$ formation process \\cite[]{sandford93,duley96,parneix98,pirronello97,pirronello99a,takahashi99,katz99,williams00,sidis2000,biham2001,joblin01,cazaux2002,cazaux2002a}. Another approach to this issue is to examine the H$_2$ formation rate in different regions of the ISM in order to see how it depends upon the local physical parameters. One can for example study the correlation of H$_2$-related quantities (abundance, rotational excitation, vibrational excitation) with the local dust properties. ISO observations of H$_2$ lines and small carbonaceous grains emission in several PDRs, and new H$_2$ UV absorption observations by FUSE of diffuse clouds, open new perspectives on our understanding of the H$_2$ formation process. In particular, the confrontation between observations and theoretical predictions provides strong constraints upon the H$_2$ formation rate \\cite[]{habart2003a,gry2002,tumlinson2002}.  In this paper, we consider PDRs and investigate the H$_2$ formation process using observations of H$_2$ emission obtained by ISO and ground-based telescopes. PDRs, where stellar radiation plays a dominant role in determining the chemical and thermal state of the gas \\cite[for a recent review see][]{hollenbach99}, are privileged objects for the study of physical and chemical processes of the interstellar medium. In Sect. \\ref{objects}, we review the data available for five nearby PDRs and summarize their physical conditions. In Sect. \\ref{observations}, we estimate for each PDR of our sample the H$_2$ formation rate using the H$_2$ line intensity ratios as a diagnostic. Also, in order to probe the influence of PAHs on the H$_2$ formation, we compare in Sect. \\ref{SGs_diagnostic} the intensities of emission measured in H$_2$ fluorescent lines with those in the aromatic bands. Then, with the aim of explaining these results, we consider in Sect. \\ref{mechanisms} empirical models of the H$_{2}$ formation mechansim by recombination of an H atom in a bound site with a second H atom which is either in the gas phase or on a neighbouring physisorbed site. In Sect. \\ref{comparison}, we compare the observational constraints on the H$_2$ formation rate with different model predictions. Our conclusions are summarized in Sect. \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}   The main aim of this study has been to provide estimates of the molecular hydrogen formation rate in a sample of nearby PDRs using results from both ISO and ground--based telescopes.  The physical conditions in the PDR layers from which H$_2$ emission is observed ($n_H\\simeq 10^3-10^5$ cm$^{-3}$, $T_{gas}\\ge$ 300 K)  differ considerably from those in the diffuse clouds where one can observe the UV lines of H$_2$ ($n_H \\simeq$100-1000 cm$^{-3}$, $T_{gas}\\simeq$50-100 K) and upon which most estimates of the molecular hydrogen formation rate have been based. Thus  the results from PDRs allow important constraints to be placed upon the mechanisms for forming molecular hydrogen on grain surfaces. We confirm the earlier result of \\cite{habart2003a} that the H$_2$ formation rate in regions of moderate excitation ($\\chi \\le$1000) such as Oph W, S140 and IC 63 is a factor of $\\sim$5 times larger than the standard rate estimated by Jura for diffuse clouds (and confirmed by recent FUSE data). On the other hand,  towards regions of higher radiation field such as the Orion Bar and NGC 2023, we derive H$_2$ formation rates consistent with the standard value.  Thus, the higher grain  and gas temperatures in PDRs do not seem to impede the formation of H$_2$.  We have attempted to interpret these results with simple empirical models of the formation of H$_2$ on grain surfaces. From these, we conclude that an ``indirect chemisorption'' model analogous to that discussed by \\cite{cazaux2002,cazaux2002a} is capable of explaining the data. This result requires an activation barrier energy against the recombination of a physisorbed H atom and a neighbouring chemisorbed H atom $E_a\\lesssim$0.1 eV. Another condition appears to be that one needs an appreciable fraction of surface sites occupied (few percent at least) with a binding energy of order 1 eV relative to the total number of surface binding sites (presumably mainly physisorbed with binding energies of order 0.05 eV). Moreover, we suggest that small (size $<$ 100 \\AA) grains may be the most promising surface for forming H$_2$ in PDRs. There is in fact enough grain surface in small grains to allow the formation rate to be larger than the standard H$_2$ formation rate and small grains spend most of their time at low temperature \\cite[]{guhathakurta89}. H$_2$ formation by indirect chemisorption upon small grains should be effective even in highly excited regions where large grains are quite warm.  Our results show that formation of molecular hydrogen in PDRs is likely to take place with a different mechanism than in the diffuse interstellar medium where the formation by physisorbed H atoms (Langmuir-Hinshelwood) probably dominates. In fact, in cold diffuse clouds the surface density of strongly bound H atoms should be low and consequently the formation by chemisorption would not be efficient.  There are several fundamental uncertainties in  our present estimates of the H$_{2}$ formation rates in PDRs which future work should try to eliminate. One is due to the fact that we have used steady--state PDR models. This assumption can cause appreciable errors and it would be useful to calculate the expected H$_2$ line intensities for models with lower H$_{2}$ formation rates but where advection has been taken into account. However, \\cite{stoerzer98} have modelled PDR structure assuming an ionization front moving into the PDR and found that non-equilibrium effects are probably minor in objects similar to the Orion Bar.  Other uncertainties are of an observational nature. It would be useful to have reliable estimates of the inclination and density distribution in the PDRs of interest in order to better constrain the models. There are indications in the Orion Bar for instance \\cite[]{walmsley2000} that the column density is higher in the molecular layers of the PDR than in the region (discussed in this paper) where the H$_2$ lines are formed. Density gradients perpendicular to the PDR  photodissociation front clearly need to be taken into account when considering the spatial distribution of the various H$_2$ lines.  The theoretical models of H$_2$ formation discussed here are clearly very preliminary.   More detailed models need to explicitly consider the degree of occupation of chemisorbed sites by H-atoms as well as the mobility of H atoms on various types of grain surfaces. Our results however do suggest that more detailed consideration of H$_2$ formation on the surface of small grains would be worthwhile.  In this case, one should ideally follow the thermal fluctuations of these small particles and take into account the dependence on grain size in order to examine properly their contribution to H$_2$ formation. Clearly also our present estimates for processes such as tunneling are very crude.  Additional information on the H$_2$ formation process could probably be obtained using high quality data for H$_2$ in excited vibrational states which gives constraints on the excitation state of the newly formed H$_2$. It also would be useful to obtain estimates of the H$_2$ formation rate under a variety of different conditions. One such condition might be in the Magellanic clouds where the different metallicity, extinction curve, and radiation field  potentially may  influence both the available grain surface area and the efficiency of H$_2$ formation. Another is in the thin high excitation clouds   such as that found by \\cite{meyer2001} towards the exciting star of NGC 2023.   It is clear that given our uncertainty about grain compositions and size distributions in different ISM locations, one is forced to some extent to use the astrophysical data to guide our estimates of processes such as H$_2$ formation.  This is perhaps philosophically not as satisfactory  as the traditional approach of employing experimentally determined or theoretically calculated rates but it is likely that nature does not give us the choice."
    },
    "0311/astro-ph0311089_arXiv.txt": {
        "abstract": "There is increasing observational evidence that hot, highly ionized interstellar and intergalactic gas plays a significant role in the evolution of galaxies in the local universe. The primary spectral diagnostics of the warm-hot interstellar/intergalactic medium are ultraviolet and X-ray absorption lines of O\\,{\\sc vi} and O\\,{\\sc vii}. In this paper, I summarize some of the recent highlights of spectroscopic studies of hot gas in the Local Group and low-redshift universe. These highlights include investigations of the baryonic content of low-$z$ O\\,{\\sc vi} absorbers, evidence for a hot Galactic corona or Local Group medium, and the discovery of a highly ionized high velocity cloud system around the Milky Way. ",
        "introduction": "We live in a wonderful age of discovery and exploration of the universe. As we peer farther and farther back in time, it is becoming ever more important to make sure that we observe the local universe as well as possible.  Observations of galactic systems and the intergalactic medium (IGM) in the low-redshift universe are required to study the universe as it has evolved over the last $\\sim5$ billion years.   They are essential for the interpretation of higher redshift systems, and they form a framework for studies of such key topics as galactic evolution, ``missing mass'', and the distribution of dark matter.  Studies of hot gas and its relationship to galaxies are shedding new light on these and other astronomical topics of interest today. In this review, I summarize some basic information about the elemental species and types of observations that can be used to study hot gas. I also  provide overviews of recent spectroscopic observations of hot gas in the low-redshift universe (\\S2), the Local Group (\\S3), and the high velocity cloud (HVC) system that surrounds the Milky Way (\\S4). Some concluding remarks on future observations can be found in \\S5.  Table~1 contains a summary of some of the most important diagnostics of hot gas in the local universe.  The species listed are detectable in either the ultraviolet or X-ray bandpasses accessible with spectrographs aboard the Hubble Space Telescope (HST), the Far Ultraviolet Spectroscopic Explorer (FUSE), the Chandra X-ray Observatory, and XMM-Newton.  Most of the diagnostics listed can be observed at their rest wavelength, which means that they can be used to study hot gas in the Local Group.  Others must be redshifted into one of the observable bandpasses; the observed wavelength at $z=0.5$ is listed for for comparison in the table.   \\begin{table} \\begin{center} \\caption{Diagnostics of Hot Gas at Low Redshift ($z < 0.5)^a$} \\begin{tabular}{lccccc} \\hline Ion, $\\lambda_{\\rm rest}$ & $f$  & $\\log f\\lambda$  & $\\lambda_{z=0.5}$ & $T_{CIE}$$^b$ & b$_{th}$$^c$ \\\\ & & & (\\AA) & (K) & (km~s$^{-1}$)\\\\ \\hline \\multicolumn{6}{c}{FUSE, HST}\\\\ \\hline H\\,{\\sc i} Ly-series & \\ldots& \\ldots & 1369-1824  & \\ldots & 13--129$^d$\\\\ C\\,{\\sc iv} 1548.195  & 0.1908 & 2.470 & 2322.292 & $1.0\\times10^5$ & 11.8 \\\\ C\\,{\\sc iv} 1550.770  & 0.0952 & 2.169 & 2326.155 & $1.0\\times10^5$ & 11.8 \\\\ N\\,{\\sc v} 1238.821 &  0.1570 & 2.289 & 1858.232 & $1.8\\times10^5$ & 14.6 \\\\ N\\,{\\sc v} 1242.804 & 0.0782 & 1.988 & 1864.206 & $1.8\\times10^5$ & 14.6 \\\\ O\\,{\\sc iv} 787.711 & 0.111 & 1.942 & 1181.567 & $1.6\\times10^5$ & 12.9 \\\\ O\\,{\\sc v} 629.730 & 0.515 & 2.511 & \\phantom{0}944.595 & $2.5\\times10^5$ & 16.1 \\\\ O\\,{\\sc vi} 1031.926 & 0.1329 & 2.137 & 1547.889 & $2.8\\times10^5$ & 17.1 \\\\ O\\,{\\sc vi} 1037.617 & 0.0661 & 1.836 & 1556.425 & $2.8\\times10^5$ & 17.1 \\\\ Ne\\,{\\sc viii} 770.409 & 0.103 & 1.900 & 1155.614 & $5.6\\times10^5$ & 21.5 \\\\ Ne\\,{\\sc viii} 780.324 & 0.0505 & 1.596 & 1170.486 & $5.6\\times10^5$ & 21.5 \\\\ \\hline \\multicolumn{6}{c}{Chandra, XMM-Newton} \\\\ \\hline O\\,{\\sc vii} 21.602 & 0.696 & 1.177 & \\phantom{0}\\phantom{0}32.403 & $8.0\\times10^5$ & 28.8 \\\\ O\\,{\\sc viii} 18.967 & 0.277 & 0.720 & \\phantom{0}\\phantom{0}28.450 & $2.2\\times10^6$ & 47.8 \\\\ O\\,{\\sc viii} 18.972 & 0.139 & 0.421 & \\phantom{0}\\phantom{0}28.459 & $2.2\\times10^6$ & 47.8 \\\\ Ne\\,{\\sc ix} 13.447 & 0.0724 & 0.988 & \\phantom{0}\\phantom{0}20.170 & $1.5\\times10^6$& 35.2 \\\\ \\hline \\end{tabular} \\end{center} $^a${$f$-values and wavelengths (in \\AA) are from Morton (1991), Verner, Barthel, \\& Tytler (1994), and Verner, Verner, \\& Ferland (1996).}  $^b${Temperature of maximum ionization fraction in collisional ionization equilibrium (Sutherland \\& Dopita 1993).}  $^c${Thermal line width, b = $(2kT/m)^{1/2}$, at $T=T_{CIE}$ unless indicated otherwise.}  $^d${Value of b for $T=10^4-10^6$\\,K.}  \\end{table}   The O\\,{\\sc vi} $\\lambda\\lambda1031.926, 1037.617$ resonance doublet lines are the best lines to use for kinematical investigations of hot ($T \\sim 10^5-10^6$\\,K) gas in the low-redshift universe.  O\\,{\\sc vi} has a higher ionization potential than other species observable by FUSE and HST, and oxygen has the highest cosmic abundance of all elements other than hydrogen and helium. X-ray spectroscopy of the interstellar or intergalactic gas in higher ionization lines (e.g., O\\,{\\sc vii}, O\\,{\\sc viii}) is possible with XMM-Newton and the Chandra X-ray Observatory for a small number of sight lines toward AGNs and QSOs, but the spectral resolution (R~$\\equiv \\lambda/\\Delta\\lambda \\la 400$) is modest compared to that afforded by FUSE (R~$\\sim 15,000$) or HST/STIS ($R\\sim45,000$).  While the X-ray lines provide extremely useful information about the amount of gas at temperatures greater than $10^6$\\,K, the interpretation of where that gas is located, or how it is related to the $10^5-10^6$\\,K gas traced by O\\,{\\sc vi}, is hampered at low redshift by the kinematical complexity of the hot interstellar medium (ISM) and IGM along the sight lines observed, as discussed below.  Nevertheless, the X-ray diagnostics are of fundamental importance for determining the ionization state of the nearby hot gas, and they provide information about hotter gas that is not traced by species in the ultraviolet wavelength region of the electromagnetic spectrum.  Table~2 contains a high-level summary of key considerations for absorption and emission-line spectroscopy of hot interstellar and intergalactic plasmas.  Both types of observations have their strengths, and the combination of ultraviolet and X-ray information holds great promise for studies of the ISM and IGM.  \\begin{table} \\begin{center} \\caption{Spectroscopy of Hot Interstellar/Intergalactic Gas} \\begin{tabular}{cccc} \\hline & {\\it UV Absorption} & {\\it X-ray Absorption} & {\\it X-ray Emission} \\\\ \\hline {\\it Observatory}           & FUSE, HST    & $\\begin{array}{c}{\\rm Chandra} \\\\ {\\rm XMM-Newton} \\end{array}$              & $\\begin{array}{c}{\\rm Chandra, ROSAT} \\\\ {\\rm XMM-Newton} \\end{array}$\\\\ \\vspace{0.2cm} $\\begin{array}{c}{T_{CIE}\\,{(K)}} \\\\ {Range} \\end{array}$ & $10^5-10^6$  & $10^6-10^8$          & $10^6-10^8$\\\\ \\vspace{0.2cm} $\\begin{array}{c}{Density} \\\\ {Dependence} \\end{array}$ & N$_{\\rm ion} \\propto n_e$  & N$_{\\rm ion} \\propto n_e$ & I$_{\\rm xray} \\propto n_e^2$\\\\ \\vspace{0.2cm} $\\begin{array}{c}{Limiting} \\\\ {Sensitivity}^a \\end{array}$ & $\\begin{array}{c}{\\rm log N(O^{+5}) \\sim 13} \\\\ {\\Rightarrow \\rm log N(H^+) \\sim 18} \\end{array}$ & $\\begin{array}{c}{\\rm log N(O^{+6}) \\sim 16} \\\\ {\\Rightarrow \\rm log N(H^+) \\sim 20.3} \\end{array}$ & \\ldots \\\\ \\vspace{0.2cm} {\\it Spatial Info.}  & Point source & Point source & Extended source\\\\ \\vspace{0.2cm} $\\begin{array}{c}{Spectral} \\\\ {Resolution} \\end{array}$ & $\\begin{array}{c}{\\rm R > 15,000} \\\\ {\\rm Detailed\\,kinematics} \\end{array}$ & $\\begin{array}{c}{\\rm R < 1000} \\\\ {\\rm General\\,kinematics} \\end{array}$ & $\\begin{array}{c}{\\rm Low\\,(broadband)} \\\\ {\\rm R < 1000} \\end{array}$ \\\\ \\hline \\end{tabular} \\end{center} $^a${Assuming a metallicity of $\\sim0.1$ solar and peak ionization fractions (see text).}  \\end{table} ",
        "conclusions": "The detection of hot gas locally and at low redshift has altered our perspective on the presence of highly ionized gas outside of galaxies and has led to refined  estimates of the baryonic content of the present-day universe.  Clearly, there is much work yet to be done in determining the spatial distribution of the hot gas, its physical conditions, and its association with galaxies and the larger-scale gaseous structures from which galaxies form.  A deeper understanding of the spatial distribution of the hot gas would be possible with ultra-sensitive maps of the diffuse O\\,{\\sc vi} emission associated with the cosmic web.  Emission maps would enable estimates of the filling factor of the hot gas, and would set the context for ongoing absorption-line studies of the hot gas.  Experiments to produce O\\,{\\sc vi} emission maps over large regions of the sky are well-suited to small and medium Explorer-class missions and have been proposed to NASA.  The Hubble Space Telescope Cosmic Origins Spectrograph, which is to be installed as part of the next HST servicing mission, will greatly enhance studies of O\\,{\\sc vi} systems at redshifts $0.14 <  z <0.5$ by increasing the number of accessible background sources to use for the absorption-line measurements. Larger redshift paths will be probed, and more sight lines will be examined.  In addition,  the high sensitivity of the spectrograph will permit the acquisition of very high signal-to-noise spectra for bright extragalactic sources. In the longer term, larger (6-20\\,m) ultraviolet/optical telescopes in space will make it possible to select many closely spaced lines of sight for detailed spectroscopic and imaging investigations of galactic and intergalactic structures associated with the hot gas.  A key scientific motivation for such investigations will be to determine feedback mechanisms operating on the hot gas and the degree to which these processes affect the formation and evolution of galaxies."
    },
    "0311/astro-ph0311383_arXiv.txt": {
        "abstract": "The variation of the mass-to-light ratios M/L of early type galaxies as function of their luminosities L is investigated. It is shown that the tilt $\\beta=0.27$ (in the B--band) of the fundamental plane relation  $M/L\\propto L^{\\beta}$ can be understood as a combination of two effects: about one-quarter  (i.e. $\\Delta \\beta =0.07$) is a result of systematic variations of the stellar population properties with increasing luminosity. The remaining three-quarters (i.e. $\\Delta \\beta =0.2$) can be completely  attributed to nonhomology effects that lead to a systematic change of the surface brightness profiles with increasing luminosity. Consequently, the observed tilt in the K--band ($\\beta=0.17$) where stellar population effects are negligible, is explained by nonhomology effects alone. After correcting for nonhomology, the mean  value of the mass-to-light ratio of elliptical galaxies  (M/L$_B$) is 7.1$\\pm$2.8 (1$\\sigma$ scatter). ",
        "introduction": "Traditionally, the global properties of elliptical galaxies have been characterized by  three observables: the central velocity dispersion $\\sigma_0$, the effective or half-light radius $r_e$, and the mean surface brightness within the effective radius $I_e$. It is generally assumed that elliptical galaxies are relaxed particle systems. From a theoretical point of view, the virial theorem then predicts a relationship between the gravitational mass, the velocity dispersion and the radius. This can be written as a relationship between $I_e$, $\\sigma_0$ and $r_e$. These observables should then be constrained to a two-dimensional plane in the above three-space. Elliptical galaxies  are indeed found to lie along a plane; however, the observed plane (popularly known as the fundamental plane [FP]; Djorgovski \\& Davis 1987, Faber et al. 1987, Dressler et al. 1987)  deviates from the prediction using a constant $M/L$. This deviation is known as the ``tilt'' of the FP.  Since the above prediction is based on the assumption that the mass--to--light ratio and the global structure of the elliptical galaxies are not dependent on the luminosity, the explanations for the tilt have been attributed to deviations from one or both of these hypotheses.  Following on that,  Faber et al. (1987) suggested that the  $M/L$ values increase systematically with increasing luminosity. Using this interpretation, the observed averaged value of the FP tilt over a large group of galaxy samples in the B--band is written as $M/L\\propto L^{\\beta}$, with $\\beta$=0.27$\\pm$0.08 (Prugniel \\& Simien 1996). Since brighter ellipticals are systematically redder, at least part of the tilt must be understood as stellar population variations along the different luminosities of elliptical galaxies. However, the presence of a ``remaining'' tilt ($\\beta$=0.17$\\pm$0.01)  in the K--band FP (Pahre, Djorgovski \\& de Carvalho 1998), where population variations are weak, implies that this solution alone cannot explain completely the observed slope. In fact, the variation in the properties of the  stellar population as deduced from broad band colors and the Mg$_2$ index explain only one-third of the slope, i.e., $\\Delta \\beta\\sim0.10$ in the B-band (Tinsley 1978, Dressler et al. 1987, Prugniel \\& Simien 1996). We show this more explicitly in Figure \\ 1a, where the luminosity dependence of estimated stellar M/L ratios in $B$ (open circles) and $K$-bands (filled circles) is plotted for a sample of 911 morphologically selected early-type galaxies from the Sloan Digital Sky Survey (SDSS) Early Data Release and the Two Mass All Sky  Surveys (2MASS) (Bell et al.\\ 2003). The dependence of stellar M/L ratios on galaxy luminosity is weak in the $B$-band, in line with previous studies, and is negligible in the $K$-band.  Taken together, the above is consistent with the conclusion that the dominant part of the tilt of the FP ($\\Delta \\beta \\sim 0.2$) is not a population effect.  Changes in the mixture of dark-to-visible matter as function of luminosity would  be a plausible explanation for the remaining tilt. However, observations (Romanowsky et al. 2003) demonstrate  that the amount of dark matter within an effective radius is in general negligibly small. Furthermore, systematic variations in kinematical structure (van der Marel 1991, Bender, Burstein and Faber 1992, hereafter BBF)  and radial anisotropy (Ciotti, Lanzoni \\& Renzini 1996, Ciotti \\& Lanzoni 1997)  cannot produce the observed tilt. The role played by rotation is also small (Busarello et al. 1997).  In this Letter, we argue that the remaining tilt of the FP is due to a systematic variation of the structural (and its associated dynamical) nonhomology of the elliptical galaxies as function of their luminosity. Following an analogy with the FP of the globular clusters, Djorgovski (1995) suggested a systematic change in the galaxy concentration as a viable hypothesis. Hjorth \\& Madsen (1995) explained part of the tilt of the FP ($\\Delta\\beta=0.12$) using a galaxy model with a variable central gravitational potential correlated with the luminosity. Following a similar idea but using a model directly related with the observations (the S\\'ersic (1968) r$^{1/n}$ law), Ciotti et al. (1996) suggested that a suitable variation (by a factor of $\\sim$ 2--4) of the shape parameter $n$  (i.e. a different degree of concentration) along the FP would be able to explain the tilt. However, Prugniel \\& Simien (1997) (and see also Graham \\& Colless (1997) in the V--band) showed that  the structural nonhomology  contributes only $\\Delta\\beta\\sim0.06-0.08$ and  to fully account for the tilt of the FP they needed to include some contributions from stellar population and rotational support. These analysis could be  affected by a limited sample  which had a small range in S\\'ersic indices 2$\\lesssim$n$\\lesssim$5 which is far from providing the full range of nonhomology observed by other authors with 1$\\lesssim$n$\\lesssim$10 in the same absolute magnitude range -22$\\lesssim$M$_B$$\\lesssim$-15 (see e.g., Graham \\& Guzm\\'an 2003).  Taking into account the full range of structural nonhomology observed, we quantify the contribution to the tilt from the structural nonhomology  in the B band. We show that the structural nonhomology is able to explain $\\approx$3/4 of the tilt in the B--band (i.e.\\ $\\Delta\\beta$=0.2), which, in combination with the tilt expected from population variations (i.e.\\ $\\Delta\\beta$$\\sim$0.1) can explain completely the observed tilt of the FP. ",
        "conclusions": "It is encouraging that the tilt $\\beta\\sim0.1$ of the FP after correction for to nonhomology  is in very good agreement with the value expected due to systematic variations of stellar population properties  of elliptical galaxies ($\\Delta\\beta\\sim0.07$). The observed tilt of $\\beta=0.29(\\pm0.04)$ in the B--band FP therefore can be understood as a combination of three-quarters structural nonhomology  and one-quarter stellar population effects. The tilt of the K--band FP where the change in the stellar population is negligible is approximately $\\beta=0.2$. This is indeed in good agreement with our expectation if the tilt were due only to the nonhomology and if color gradients inside galaxies do not change their S\\'ersic indices. Indeed, radial V--K gradients in the surface brightness profiles of ellipticals are on average only 0.16 mag dex$^{-1}$ (Peletier, Valentijn \\& Jameson 1990) which is small compared to the variations in surface brightness within an effective radius. The homologous FP of our sample has a scatter in log(M/L) of order 25\\% which is consistent with previous results (BBF; Pahre, Djorgovski \\& Carvalho  1998). The scatter is not significantly reduced in the non-homologous FP which indicates that it is mainly a result of population effects and not of structural variations within a given luminosity bin.   Our results rule out that even very massive ellipticals are strongly dominated by dark matter in their inner regions, which is in agreement with Rix et al. (1997) and Gerhard et al. (2001). After correcting for nonhomology, the mean value of the mass-to-light ratio (M/L$_B$) of early-type galaxies in our luminosity range is 7.1$\\pm$2.8 (1$\\sigma$ scatter). This value is in good agreement with the value expected for an old stellar populations of M/L$_B$=7.8$\\pm$2.7 (Gerhard et al 2001)."
    },
    "0311/astro-ph0311456_arXiv.txt": {
        "abstract": "We have observed the SiO $J$=5\\too 4 line towards a sample of 25 low-luminosity (L$_* < 10^3$ L$_\\odot$) protostellar outflow systems.  The line was detected towards 7 of the 25 sources, a detection rate of 28 per cent. The majority (5 out of 7) of sources detected were of class 0 status. We detected a higher fraction of class 0 sources compared with the class I and II sources, although given the small numbers involved the significance of this result should be regarded as tentative.  Most of the detected sources showed emission either at or close to the central position, coincident with the protostar. In four cases (HH\\,211, HH\\,25MMS, V-380 Ori\\,NE and HH\\,212) emission was also detected at positions away from the center, and was stronger than that observed at the centre position.  SiO abundances of 10$^{-8}$ to 8$\\times$10$^{-7}$ are derived from LTE analysis. For 2 sources we have additional transitions which we use to conduct statistical equilibrium modeling to estimate the gas density in the SiO-emitting regions. For HH\\,25MMS these results suggest that the SiO emission arises in a higher-density region than the methanol previously observed. We find that the most likely explanation for the preferential detection of SiO emission towards class 0 sources is the greater density of those environments, reinforced by higher shock velocities. We conclude that while not all class 0 sources exhibit SiO emission, SiO emission is a good signpost for the presence of class 0 sources. ",
        "introduction": "In principle SiO has great potential as an efficient shock tracer, a millimetre-wavelength counterpart to the near-infrared lines of H$_2$. Its abundance in quiescent gas within cloud cores is $< 10^{-12}$ relative to molecular hydrogen (Irvine, Goldsmith \\& Hjalmarson 1987; Ziurys, Friberg \\& Irvine 1989), which can rise to as high as $10^{-6}$ in outflows (Mart\\'{\\i}n-Pintado et al. 1992). The formation of SiO in post-shock gas probably takes place as a consequence of grain sputtering or grain-grain collisions releasing Si-bearing material into the gas phase which reacts rapidly with O-bearing species (principally O$_2$ and OH) to form SiO (Schilke et al. 1997; Caselli, Hartquist \\& Havnes 1997).  A number of outflow sources have been detected in SiO emission (e.g. Mart\\'{\\i}n-Pintado, Bachiller \\& Fuente 1992) but it is only recently that extensive surveys of outflows have been carried out and published. Harju et al. (1998) observed the locations of water and/or OH masers, with the result that their sample is dominated by massive young stellar objects (YSOs). Furthermore they only observed at the maser position in each source although they were able to obtain high sensitivity spectra. The sample of Codella, Bachiller \\& Reipurth (1999) includes YSOs with luminosities in the range 1 to 10$^4$ L$_\\odot$, and they made the surprise discovery that a number of sources show evidence for high- and low-velocity SiO components with different spatial distributions and linewidths. Furthermore, the abundance of SiO was different in each component, with the high-velocity SiO two orders of magnitude more abundant than the low-velocity SiO.  In addition, the low-velocity component was typically only seen in the lower-lying transitions of SiO, probably indicating that the low-velocity gas is also of lower density and/or temperature.  In an effort to understand the conditions under which SiO emission is observed, we have conducted a survey of the $J$=5\\too 4 transition of SiO towards a sample of 25 low-luminosity YSOs, made at the James Clerk Maxwell Telescope (JCMT) on Mauna Kea, Hawaii. The following section describes the observations, with subsequent sections presenting the results, analysis and discussion of the data. ",
        "conclusions": "We have observed the SiO $J$=5\\too 4 line towards a sample of 25 low-luminosity (L$_* < 10^3$ L$_\\odot$) protostellar outflow systems. The line was clearly detected towards 7 of the 25 sources, a detection rate of 28 per cent. The majority (5 out of 7) of sources detected are of class 0 status. We detect a greater fraction of class 0 sources compared with other classes (45 per cent compared with 18 and 0 per cent for class I and II). However, given the small numbers involved, the significance of this result should be regarded as tentative.  In four cases (HH\\,25MMS, HH\\,212, V-380 Ori\\,NE and HH\\,211) emission was detected at positions offset from the center (in two other cases it is known that SiO emission also occurs away from the central source position). LTE analysis yields SiO abundances ranging from 10$^{-8}$ to as high as 8$\\times$10$^{-7}$. In two cases (HH\\,25MMS and HH\\,212) we have conducted LVG modeling of the emission, and find that the SiO emission arises in regions with temperatures of typically 30--100 K, densities of 10$^5$ to 10$^6$ cm$^{-3}$ and SiO abundances of 10$^{-11}$ to 10$^{-9}$ relative to H$_2$. These abundances are lower than those derived from LTE analysis, a result consistent with the notion that the SiO emission arises in clumps which are much smaller than the 22-arcsec beam.  Higher resolution observations has shown this is true, and we feel that further high-resolution studies are desireable in order to determine the conditions of the post-shock gas.  We have explored a number of possible reasons for the detection of SiO in some sources but not others, and conclude that the shock velocity and density of the ambient medium play an important role.  We can therefore account for the fact that most of the detected sources have class 0 status. We conclude that while not all class 0 sources exhibit SiO emission, SiO emission is a good signpost for class 0 sources."
    },
    "0311/astro-ph0311168_arXiv.txt": {
        "abstract": "I consider three processes which enhance mass loss rate from a common envelope of a giant star with a main sequence or a white dwarf companion spiraling-in inside its envelope. I consider deposition of orbital energy and orbital angular momentum to the giant's envelope, and in more detail the formation of jets by an accreting companion and their propagation in the envelope. I find that in many cases the deposition of orbital angular momentum to the envelope may be more important to the mass loss process than the deposition of orbital energy. Jets blown by an accreting companion, in particular a white dwarf, orbiting inside the outer regions of the giant's envelope may also dominate over orbital energy deposition at early stage of the common envelope evolution. These imply that studies which ignore the deposition of angular momentum to the envelope and the effects of the accreting companion may reach wrong conclusions. ",
        "introduction": "As a star in a binary system swells to become a giant it engulfs its companion if the orbital separation is smaller than some critical value and if the companion is not too massive; a common envelope (CE) phase commences. (for a review see Iben \\& Livio 1993, and Taam \\& Sandquist 2000 ). Because of tidal interaction and friction, the orbit shrinks. Several parameters can be defined to characterized the CE evolution (e.g. Livio \\& Soker 1988), but the most commonly used parameter is the ratio of the binding energy of the ejected envelope $\\Delta E_{\\rm bind}$ to the orbital energy that is released during the CE phase $\\Delta E_{\\rm orb}$: $\\alpha_{\\rm CE} \\equiv \\Delta E_{\\rm bind}/\\Delta E_{\\rm orb}$. Note that different definitions for the binding energy exist (e.g., O'Brien, Bond, \\& Sion 2001). Since the orbital energy that is released depends mostly on the final orbital separation, the value of $\\alpha_{\\rm CE}$ can be in principle calculated for systems whose final orbital separation is known, assuming the giant structure at the onset of the CE is known (O'Brien et al.\\ 2001; Maxted et al.\\ 2002). The use of the $\\alpha _{\\rm CE}$ is common also in numerical simulations of the CE phase (e.g. Sandquist, Taam \\& Burkert 2000 for a recent paper). However, numerical simulations can't include the effect of enhanced mass loss rate from giant stars that have a very high  mass loss rate. The spun-up envelope of red giant branch (RGB) and asymptotic giant branch (AGB) stars may have a much higher mass loss rate, with the energy source being the giant's luminosity rather than the orbital energy (Soker \\& Harpaz 2003).  In some systems the usage of the above expression in a simple manner yields $\\alpha_{\\rm CE} >1$. For example, Maxted et al.\\ (2002) assume that negligible mass has been lost prior to the onset of the CE phase in PG1115+166, and find $\\alpha_{\\rm CE} >1$. This led some researchers to argue that the energy stored in the envelope, and in particular the ionization energy, i.e., the energy released when the envelope material recombines, is the extra energy needed to expel the CE (e.g., Han, Podsiadlowski, \\& Eggleton 1994; Dewi \\& Tauris 2000; Maxted et al.\\ 2002). This proposed mechanism was criticized in previous papers (Harpaz 1998; Soker 2002; Soker \\& Harpaz 2003). In Soker (2002) I criticized the paper by Maxted et al.\\ (2002) for not considering the mass lost from the envelope prior to the onset of the CE, when the system is synchronized, i.e., the giant's rotation period equals the binary orbital period, and the binary orbital shrinkage proceeds very slowly. Eggleton (2002), for example, notes that a close companion may substantially enhance mass loss rate prior to the onset of a Roch lobe overflow (RLOF), with the possibility of preventing a CE phase altogether.  Soker \\& Harpaz (2003) criticize Han et al.\\ (2002) for claiming that the ionization energy in the envelope is a significant factor in the CE evolution. Soker \\& Harpaz (2003) consider the mass lost by RGB stars as they expand by a relatively large factor from the moment of synchronization to the RLOF. Soker \\& Harpaz then argue that Han et al.\\ (2002) include a mass loss rate prior to the onset of the CE that is too low, and do not include the energy radiated by the accreting white dwarf companion, as well as that emitted by the core of the giant star. In a later paper Han et al.\\ (2003) briefly refer to Soker \\& Harpaz criticism, keeping the dispute alive. Since the applicability of the $\\alpha_{\\rm CE}$ parameter is a fundamental question in the CE process, and the CE evolution is the channel for the formation of many close binary systems, I elaborate on some questions regarding energy and angular momentum budget in the CE phase. In the first several sections I study the way by which an accreting companion can deposits energy via jet formation. I then (section 7) put all into a coherence picture. A short summary is in section 8. ",
        "conclusions": "The main goal of the present paper is to point to the caution one must take in using the $\\alpha_{\\rm CE}$ parameter when studying CE evolution. Namely, the orbital energy deposited to the giant's envelope is not always the main effect leading, directly or indirectly, to the removal of the envelope. For that I considered here the deposition of energy from the accreting companion and the deposition of orbital angular momentum to the giant's envelope. The main results can be summarized as follows. \\begin{enumerate} \\item When inside the envelope of a giant, a main sequence companion is unlikely to blow jets, or a collimated fast wind (CFW, i.e., less collimated jets), or it will marginally do so only when in the outer parts of the envelope. \\item A WD companion is more likely to blow jets or a CFW \\item These jets, even if exist, whether from a WD or a MS companion, are not likely to exit the envelope at a high speed during the CE phase. Hence, they are not likely to play a major role in shaping the circumbinary matter. Jets might be blown by the companion before entering the CE (Soker \\& Rappaport 2000), or one or two of the stars after the CE ends (Soker \\& Livio 1994). This explains the observations that PNs with binary nuclei are not bipolar PNs, i.e., have no lobes, beside NGC 2346, with the longest known orbital period. I do expect that some binary progenitors of bipolar PNs entered the CE phase at late stages, and that now the orbital separation is $\\sim 0.1-1 \\AU$. These systems are hard to detect (Bond 2000). To obtain a quantitative result, the CE population synthesis calculations of Yungelson, Tutukov, \\& Livio (1993) should be repeated but with enhanced mass loss rate from rotating AGB stars included. \\item The CFW or jets, if exist, may inflate a bubble (with a complicated structure because of the orbital motion), hence playing a significant role in expelling the outer layers of the envelope when the companion is still orbiting in the outer envelope region. \\item In many cases the effects due to angular momentum deposition into the envelope seem more influential in removing the envelope than orbital energy deposition, assuming that fast rotating envelopes have high mass loss rates. This is true for stellar as well as substellar companions. The energy source is the giant luminosity due to nuclear energy production in the core. The Eddington luminosity of an accreting stellar companion is of the order of the giant's luminosity, and can farther increase the mass loss rate (Iben \\& Livio 1993; Armitage \\& Livio 2000). \\item My results here iterate earlier claims (Soker 2002; Soker \\& Harapz 2003) that a high degree of cautious should be taken when applying the $\\alpha_{\\rm CE}$ parameter for the removal of CEs. For example, the conclusions of some papers that another energy source, e.g., ionization energy of the envelope, is required to remove the envelope (see criticism in Soker \\& Harpaz 2003) are questionable. \\end{enumerate}"
    },
    "0311/astro-ph0311397_arXiv.txt": {
        "abstract": "We present results from observations of 29 BL Lacertae objects, taken with the Whipple Observatory 10\\,m Gamma-Ray Telescope between 1995 and 2000.  The observed objects are mostly at low redshift ($z < 0.2$) but observations of objects of $z$ up to 0.444 are also reported. Five of the objects are EGRET sources and two are unconfirmed TeV sources. Three of the confirmed sources of extragalactic TeV gamma rays were originally observed as part of this survey and have been reported elsewhere. No significant excesses are detected from any of the other objects observed, on time scales of days, months or years. We report 99.9\\% confidence level flux upper limits for the objects for each observing season. The flux upper limits are typically 20\\% of the Crab flux although, for some sources, limits as sensitive as 6\\% of the Crab flux were derived. The results are consistent with the synchrotron-self-Compton (SSC) model predictions considered in this work. ",
        "introduction": "\\label{intro}  BL Lacertae objects (BL Lacs) are members of the blazar class of active galactic nuclei (AGN). Like all blazars, they exhibit rapid, large amplitude variability at all wavelengths, high optical and radio polarization, and in some cases, apparent superluminal motion and/or gamma-ray emission. All of these observational properties lead to the widely held belief that blazars are AGN with jets oriented nearly along our line of sight \\citep{Urry95}. The broadband spectral energy distribution (SED) of blazars, when plotted as $\\nu$F$_\\nu$ versus frequency, shows a double peaked shape, with a smooth extension from radio to between IR and X-ray frequencies (depending on the specific blazar type), followed by a distribution that typically starts in the X-ray band and can peak in the gamma-ray band, at energies as high as several hundred GeV. EGRET has detected more than 65 blazars \\citep{Hartman99}, 14 of which have been identified as BL Lacs \\citep{Dermer00}.  The low energy part of the blazar SED is believed to be incoherent synchrotron radiation from a relativistic electron-positron plasma in the blazar jet. The origin of the high energy emission is still a matter of considerable debate (e.g., \\citealp{Buckley98}; \\citealp{Mannheim98}).  Leptonic models are the most popular models used to explain the observed gamma-ray emission from blazars. In synchrotron-self-Compton (SSC) models, the gamma rays are produced through inverse Compton scattering of low energy photons by the same electrons that produce the synchrotron emission at lower energies \\citep{Konigl81,Maraschi92,Dermer92,Sikora94,Bloom96,Sikora01}. In external Compton models, the dominant source of seed photons for upscattering in the inverse Compton process are ambient photons from the central accretion flow, the accretion disk, the broad line region, the torus, the local infra-red background, or some combination of these \\citep{Sikora94, Dermer93, Blandford95, Ghisellini96, Wagner95}. Another set of models proposes that the gamma rays are the result of ultra-high energy (E $\\gtrsim 10^{19}$ eV) protons producing TeV gamma rays as proton synchrotron radiation \\citep{Aharonian00}, in proton induced electromagnetic cascades \\citep{Mannheim98}, or by a combination of both processes \\citep{Mucke03}.  Among blazars, BL Lacs are believed to be the best candidates for TeV emission. They have weak or absent optical emission lines, indicating that they may have less TeV-absorbing material near the emission region \\citep{Dermer94}. Although they show relatively low luminosities compared to Flat Spectrum Radio Quasars (FSRQs) and Optically Violent Variables (OVVs), their SEDs peak at higher energies \\citep{Fossati98}. In blazar unification models in which electrons are assumed to be the progenitors of the gamma rays (e.g. \\citealp{Ghisellini98}), the lower luminosity of the BL Lacs, relative to the FSRQs implies that in BL Lacs, electrons are cooled less efficiently. This means that the electrons reach higher energies and that the emitted gamma rays will therefore also be of higher energy.  Consistent with these expectations, BL Lacs are the only type of blazars detected at very high energies (VHE, E $\\gtrsim$ 250\\,GeV). Currently, six BL Lacs are confirmed sources of VHE gamma rays, Mrk\\,421 \\citep{Punch92}, Mrk\\,501 \\citep{Quinn96}, H\\,1426+428 \\citep{Horan02} and 1ES\\,1959+650 \\citep{Nishiyama00}; the original detection of the nearby BL Lac, 1ES\\,2344+514 \\citep{Catanese98} was confirmed by the HEGRA group \\citep{Tluczykont03} while that of the Southern Hemisphere BL Lac, PKS\\,2155-304 \\citep{Chadwick99} was confirmed recently by H.E.S.S. \\citep{Djannati03}. Unconfirmed detections have been reported for the BL Lac objects 3C\\,66A \\citep{Neshpor98} and BL Lacertae \\citep{Neshpor01}. The TeV blazars, 1ES\\,2344+514, H\\,1426+428 and 1ES\\,1959+650 were originally observed at Whipple as part of the BL Lac survey described here. Since they were detected during these observations, the results were described in detail elsewhere \\citep{Catanese98,Horan02,Holder03} and are summarized here.  \\citet{Padovani95a} introduced the terminology ``Low-frequency peak BL Lacs'' (LBLs) to describe those BL Lacs in which the lower energy SED peak occurs in the radio band and ``High-frequency peaked BL Lacs'' (HBLs) for those whose lower energy peak occurs in the X-ray band. Recently, deeper BL Lacs surveys (e.g. \\citealp{Perlman98}; \\citealp{Laurent99}; \\citealp{Caccianiga99}) have revealed evidence for the existence of BL Lacs with properties intermediate to those in the LBL and HBL classes. Indeed, W\\,Comae has recently been classified as an intermediate BL Lac \\citep{Tagliaferri00}.  This suggests that, rather than being separate subclasses, LBLs and HBLs represent the edges of a sequence of progressively different BL Lacs. The term ``extreme blazars'' was introduced by \\citet{Ghisellini99} to describe those BL Lacs whose first peak extends into the hard X-ray band. Such objects, which lie at the end of the ``blazar sequence'' proposed by \\citet{Fossati97}, are good candidates for TeV emission since the second peak in their SEDs also lies at higher energies meaning that they can be powerful at TeV energies.  With the exception of 3C\\,66A and BL Lacertae, all of the claimed and confirmed TeV gamma-ray emitting BL Lacs are HBLs. H\\,1426+428, Mrk\\,501, 1ES\\,2344+514 and, to a lesser extent, 1ES\\,1959+650, all have very hard X-ray energy spectra and fall into the class of extreme blazars. All of the TeV blazars but these four were listed as detections in the Third EGRET catalog \\citep{Hartman99}. The second peak in the SED of these and other extreme objects can only be studied in detail at TeV energies as it lies above 100 GeV, where EGRET and GLAST are less sensitive. Ground-based gamma-ray telescopes therefore offer a unique opportunity to study this class of higher-peaked blazars. Indeed, the majority of blazars detected by EGRET were FSRQs whose second SED peak falls in the MeV to GeV band.  In order to improve our understanding of the gamma-ray emission from BL Lacs, we need to detect more of these objects at very high energies. Several groups have published upper limits on VHE emission from BL Lacs in the last several years \\citep{Roberts98,Roberts99,Chadwick99,Aharonian00}, including the VERITAS collaboration \\citep{Kerrick95}, but these efforts have been on smaller groups of objects, or not directed specifically at BL Lacs. In this paper we present the results of our BL Lac observing program from January 1995 to July 2000. The database comprises observations of 29 objects totaling 143 hours. We present the results of searches for emission spanning time-scales of 30 minutes to six years. No statistically significant excess emission above the background is found and we discuss the implications of the non-detections.  These results supersede preliminary analyses presented in conference proceedings \\citep{Catanese97a, Horan00, Horan03}. Five of the objects presented here were later observed more intensively as part of a HBL observation program carried out by \\citet{delaCalle03}. In that survey, eight HBLs, selected from a list of TeV candidate objects derived by \\citet{Costamante02}, were subjected to intensive Whipple observations during 2001 and 2002. These objects were all predicted to be VHE emitters based upon the location of their synchrotron peak and on their high density seed photons for the inverse Compton process. No evidence for TeV emission from these objects was found during this intensive monitoring campaign \\citep{delaCalle03}. ",
        "conclusions": "\\label{res}  Table~\\ref{results-tab} summarizes the results of the observations of the BL Lacs observed but not detected between January 1995 and July 2000 while Table~\\ref{detected} summarizes the results of the observations taken on H1426+428, 1ES1959+650 and 1ES2344+514, the three BL Lacs that were detected during this survey. For each target object, we list the observation exposure during each season, the significance of the excess or deficit of these observations, the maximum significances for a night or month of observations and the flux upper limits expressed as fractions of the Crab count rate and in integral flux units (assuming a Crab-like spectrum). Many of the objects were observed over a number of different observing seasons resulting in a range of upper limits above different values of E$_p$. No evidence for a statistically significant excess or deficit was seen in the detected count rate from any of the objects for any of the time periods examined. The distribution of the significances for each object for each observing season is shown in Figure~\\ref{sigmas}. This distribution has a mean of 0.005 with a standard deviation of 0.976. The black curve shows the expected shape if the significances were drawn from a Gaussian distribution. The Kolmogorov-Smirnov test returns a 95\\% probability that the data are normally distributed.  Table~\\ref{detected2} summarizes the detections of H1426+428, 1ES1959+650 and 1ES2344+514, the three target BL Lacs that have been subsequently confirmed as TeV emitters.  In order to investigate how much the flux upper limits change if the spectral index is not the same as that of the Crab Nebula, estimates of the flux upper limits for the BL Lacs were made assuming source spectral indices of -2.2 and -2.8. The integral flux from the Crab Nebula above 300 GeV was used to scale the previously calculated upper limits. This integral flux was assumed to remain constant when the spectral index was changed. E$_p$ was not adjusted to account for the response of the telescope to the different input spectra but this effect should be very small. Table~\\ref{newindices} lists the flux upper limits for each BL Lac for each observing season when these different source spectral indices were assumed.  \\citet{Costamante02} have made predictions for the TeV flux from fourteen of the BL Lacs included in this paper using two different methods. In the first, an SSC model was used to fit the multiwavelength data gathered on each of the objects while in the second approach, a phenomenological description of the average SED of the blazars was derived based upon their observed bolometric luminosity \\citep{Fossati98}. The resulting flux predictions from each of the two methods were given above 300 GeV and above 1 TeV. In order to compare the upper limits presented here with these predictions, our upper limits that were not already derived above 300 GeV, were extrapolated to this energy. This was done by first expressing the flux upper limit as a fraction of the Crab flux at that energy, F$_{BLLac}$($>E_p$). This was then scaled to 300 GeV assuming a Crab-like spectrum. Thus, an upper limit on the integral flux for each BL Lac above 300 GeV for each observing season, F$_{BLLac}$($>$300GeV), was calculated:  \\begin{equation}\\label{flux} F_{BLLac}(>300GeV)=F_{BLLac}(>E_p)\\,F_{Crab}(>E_p)\\left({\\displaystyle\\frac{300}{E_p}}\\right)^{-1.5} \\end{equation}  \\noindent F$_{Crab}$($>$E$_p$) is the integral flux from the Crab Nebula above E$_p$ in units of photons cm$^{-2}$ s$^{-1}$, assuming an integral spectral index of -1.5. F$_{BLLac}$($>E_p$) is the upper limit on the flux from the BL Lac above E$_p$ expressed as a fraction of the Crab Nebula integral flux at this energy.  When deriving their flux predictions, \\citet{Costamante02} did not take into account absorption of the gamma rays by the infra-red background. The predictions above 300 GeV could therefore change by factors on the order of 5 for objects at redshifts above 0.2. The upper limits presented here were compared with these predictions and, those of four BL Lacs, shown in Table~\\ref{foscos}, were found, during all seasons in which they were observed, to be lower than the predicted fluxes according to the Fossati approach \\citep{Fossati98} adapted in \\citep{Costamante02}.  Those of two more BL Lacs, also listed in Table~\\ref{foscos}, were found to be lower during some of the observing periods. All of the upper limits calculated were higher than the fluxes predicted using the one-zone, SSC model \\citep{Costamante02}. It should be noted however, that the upper limits quoted here pertain only to the specific period during which the observations were made. As demonstrated by \\citet{Bottcher02}, spectral fitting of blazars is subject to very large uncertainties when non-simultaneous multiwavelength data are used. Indeed, it is also shown that, even with the best currently available simultaneous optical - X-ray data, there is a very wide range in the predicted fluxes above 40 GeV. Given an observed X-ray flux, the predicted gamma-ray flux depends very sensitively on the model parameters and, even for simultaneous data, can vary by large factors due to the uncertainty in these parameters.  In the absence of dedicated simultaneous multiwavelength data, it is difficult to use these data to constrain emission models. However, many of the objects surveyed here are monitored by the All Sky Monitor (ASM) on board the Rossi X-ray Timing Explorer. For many other TeV blazars, the soft X-ray flux has been seen to rise during periods when the gamma-ray flux was also high. For example, Mrk\\,501 was observed to have a higher than average flux in the ASM during 1997, the same year during which the greatest TeV flaring activity detected to date was observed \\citep{Catanese97b,Quinn99}. In order to see if any of the BL Lacs surveyed here were particularly active in the X-ray band during these observations, the X-ray curves for the objects presented here that are monitored by ASM were analyzed. If heightened X-ray activity was detected in the absence of corresponding gamma-ray activity, this could have interesting consequences for emission models. Out of the 29 objects presented here, 25 of them are monitored on a regular basis by the ASM. The nightly average light curves for each of these were generated. The flux from each object was found to be, on average, very low ($\\lesssim$ 0.01 Crab units ) with no evidence for dramatic flaring or for any sustained period of X-ray activity.  Current and future observing campaigns at the Whipple Observatory make use of BL Lac monitoring at X-ray wavelengths to try to predict when an object might be in a higher flux state and thus detectable in the VHE band. Elevated gamma-ray fluxes are often accompanied by a corresponding increase in the X-ray flux (e.g. \\citealp{Maraschi99}; \\citealp{Jordan01}). Thus, by monitoring the X-ray activity from blazars, we can identify periods of increased activity during which the VHE flux may also be stronger. However, the relationship between the X-ray and gamma-ray flux has been shown to be complicated with gamma-ray flares being detected in the absence of X-ray flares \\citep{Holder03, Krawczynski03}, and vice-versa \\citep{Rebillot03}.  Finally, it should be noted that, although none of the objects presented here were found to have a statistically significant TeV flux during these observations, three of the confirmed TeV blazars were detected during this survey. When 1ES2344+514, H1426+428 and 1ES1959+650 were initially observed at Whipple, it was as part of this BL Lac campaign and, like the objects listed here, they were not detected. In subsequent years however, continued monitoring with deeper exposures revealed these objects to be TeV emitters when in more active states, although not always detectable when in their quiescent state. Continued VHE observations of the BL Lacs presented here, in particular those shown to have extreme properties \\citep{delaCalle03}, accompanied by monitoring of their X-ray flux level, may reveal many more of them to be TeV emitters. Indeed, since extreme BL Lacs, the best blazar candidates for TeV emission, have lower luminosity at all wavelengths than their lower energy peaked counterparts, their flux level often lies below the detection threshold of the current generation of imaging atmospheric Cherenkov telescopes when they are in quiescent state. As the next generation of ground-based gamma-ray telescopes comes online with their increased flux sensitivity, they will offer a unique opportunity to study this low luminosity class of blazar and should detect many more of these objects.  Future, X-ray all-sky monitor experiments like LOBSTER \\citep{Black03} and EXIST \\citep{Grindlay02}, with their improved flux sensitivity and increased bandwidth, will allow for more detailed monitoring of the X-ray emission from blazars thus providing valuable information which can be used to trigger observations at gamma-ray energies. These X-ray missions, coupled with the next generation of higher sensitivity VHE observatories such as VERITAS, HESS, MAGIC and CANGAROO, should allow both lower power extreme BL Lacs and lower frequency peaked blazars (LBLs and FSRQs) to be detected."
    },
    "0311/astro-ph0311504_arXiv.txt": {
        "abstract": "If the transverse jet {\\bf B} fields observed in BL~Lac objects on parsec scales represent the toroidal component of the jet {\\bf B} fields, this should give rise to rotation-measure gradients {\\em across} the jets, due to the systematic change in the line-of-sight component of the jet {\\bf B} field. We have found evidence for such rotation-measure gradients in one-third to one-half of the 1~Jy BL~Lac objects we have analyzed. We present and discuss these new results, together with some of their implications for our understanding of the parsec-scale jets of AGN. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311218_arXiv.txt": {
        "abstract": "{ The shape of the gamma-ray line from radioactive \\Aln, at 1808.7~keV energy in the frame of the decaying isotope, is determined by its kinematics when it decays, typically 10$^6$~y after its ejection into the interstellar medium from its nucleosynthesis source.  Three measurements of the line width exist: HEAO-C's 1982 value of $(0+3)$~keV FWHM, the GRIS 1996 value of $(5.4\\pm 1.3)$~keV FWHM, and the recent RHESSI value of $(2.0\\pm 0.8)$~keV FWHM, suggesting either ``cold'', ``hot'', or ``warm'' \\Al in the ISM.  We model the line width as expected from Galactic rotation, expanding supernova ejecta, and/or Wolf-Rayet winds, and predict a value below 1~keV (FWHM) with plausible assumptions about \\Al initial velocities and expansion history.  Even though the recent RHESSI measurement reduces the need to explain a broad line corresponding to 540~km~s$^{-1}$ mean \\Al velocity through extreme assumptions about grain transport of \\Al or huge interstellar cavities, our results suggest that standard \\Al ejection models produce a line on the narrow side of what is observed by RHESSI and INTEGRAL. Improved INTEGRAL and RHESSI spatially-resolved line width measurements should help to disentangle the effects of Galactic rotation from the ISM trajectories of \\Aln. ",
        "introduction": "Current Galactic nucleosynthesis reveals itself through the decay of \\Aln, one of its radioactive by-products with a mean lifetime of $1.04\\times 10^6$~y.  \\Al undergoes $\\beta^+$-decay into an excited state of $^{26}$Mg, which de-excites through emission of a gamma-ray photon at 1808.7 keV. This gamma-ray line has been observed and imaged throughout the Galaxy \\citep{mahoney82,diehl95,oberlack97,knoedlseder99,plueschke01}. Sources of \\Al may be AGB stars and novae, but massive stars (via core-collapse supernovae and winds from Wolf-Rayet stars) have been found the most plausible and probably dominating sources \\citep{prantzos96}. The rather irregular \\Al emission along the plane of the Galaxy, and its consistency with the patterns of tracers of massive-star activity, is the main argument for favouring massive stars as the sources \\citep{diehl96,knoedlseder99,plueschke01}.  Flux measurements have been employed to study the nature of the sources, comparing with predicted \\Al yields from models of the source types. The amount of \\Al present in the ISM of the Galaxy has been estimated at $\\approx 2~\\mathrm{M}_\\odot$, and used to argue for the roles of different source types. But the uncertainties about the spatial distribution and total number of nucleosynthesis events add to source yield uncertainties, providing only qualitative arguments for the nature of \\Al sources \\citep{prantzos96}. Therefore, locally constrained candidate source populations have been studied, such as in the Cygnus region \\citep{knoedlseder00, plueschke01}. In such a case, the distance to the sources is constrained to a smaller interval, along with the radial velocity due to Galactic rotation.  Measurements of Galactic \\Al with high-resolution spectrometers have produced somewhat controversial results: The initial discovery with the space-borne \\mbox{HEAO-C} Ge spectrometer had reported a narrow line \\citep[intrinsic width (FWHM) $(0+3)$~keV;][]{mahoney84}. But the GRIS balloon-borne Ge spectrometer found the line to be significantly broadened (intrinsic FWHM $(5.4\\pm1.4)$~keV; \\citet{naya96}).  Such a broad line, however, cannot be explained easily \\citep[see][]{chen97}. If the origin of the line broadening was thermal, the \\Al decay region would need to be at a temperature of $\\approx 4.5\\times 10^8$~K. Alternatively, \\Al isotopes would have to maintain a mean velocity around 540 km~s$^{-1}$ over their 1 My decay time, travelling kpc distances at those speeds. One may either assume that a substantial fraction of \\Al is injected into such rather large interstellar cavities, or that a substantial fraction of \\Al condenses onto dust grains before deceleration, so it maintains its momentum throughout passages of supernova remnant shells or other obstacles.  The velocity could also be a result of re-acceleration of dust grains by interstellar shocks in the neighbourhood of the source, allowing them to maintain a high velocity over the \\Al decay time scale \\citep{sturner99}.  None of these explanations is straightforward or without problems.  A firm measurement of the \\Al gamma-ray line width is desirable, before questioning our understanding of \\Al fate from its production sites until decay in interstellar space.  New measurements have been obtained recently with Ge spectrometers aboard the high energy solar spectroscopy imager RHESSI \\citep{smith03} and the INTEGRAL observatory. The RHESSI result was derived through Earth occultation analysis of data while pointing at the sun. They obtain an intermediate intrinsic line width of about $(2.0 \\pm 0.8)$~keV FWHM. The preliminary INTEGRAL measurement also suggests a narrow line, but systematic uncertainties are still large \\citep{diehl03}. Within the given uncertainties, the current set of measurements is mildly inconsistent. But more measurements have been recorded with INTEGRAL's SPI Ge spectrometer already, so that longitude-resolved line width results are at the horizon.  We explore the potential of \\Al decay line shape measurements for diagnostics on the sources of \\Al and their interstellar environment. We follow up on earlier analysis by \\citet{gehrels96}, who showed that the structure of the Galaxy should be reflected in the position of the line due to Doppler shift from Galactic rotation. In this paper we employ an updated spatial distribution model for the plausible sources of \\Al in the Galaxy.  We also account for the fate of ejected \\Al in interstellar space through different model variants which reflect ejections by winds and/or supernovae into cavities of plausible sizes for the massive-star environment of the \\Al sources. Our aim is to illustrate the diagnostic power of measuring the \\Al gamma-ray line's position and width; this should be feasible with current instrumentation through imaging spectroscopy, even if a detailed decomposition of the line shape may still be beyond reach. ",
        "conclusions": "We model \\Al sources in the Galaxy, adopting massive stars as the dominating sources. Using Galactic rotation and plausible assumptions for the fate of \\Al from ejection by the supernova or Wolf Rayet star into until decay in interstellar space, we derive an expected profile for the 1809~keV gamma-ray line from the decay of Galactic \\Aln. For this, we integrate along the lines of sight over the Doppler-shifted source regions of \\Al emission with their intrinsic state of dynamical evolution.  Measurements of the 1809~keV emission from galactic \\Al play in principle a role similar to the \\ion{H}{i} 21~cm radiation, with the benefit that the Galaxy is always optically thin to MeV gamma radiation. Current instruments are unable to improve the knowledge of galactic rotation over \\ion{H}{i} results, the basis for the inner galaxy rotation curve by \\citet{olling00}, but the knowledge of rotation can be used to test the position of \\Al sources.  Our result demonstrates that the gamma-ray line profile reflects the kinematics of decaying \\Al in the ISM. However, current Ge detectors will probably be unable to detect line shape departures from a simple Gaussian shape for spectra such as predicted by our model. Nevertheless, centroid and width of the line will provide a diagnostic for the \\Al sources in the Galaxy.  Detection of the expected amount of line shift would put a limit on the contribution of local emission to the \\Al flux in the direction of the inner galaxy. By their massive star origin, \\Al and \\ion{H}{ii}-regions are connected; the latter having also been used to measure galactic rotation \\citep{brand93}. \\ion{H}{ii} regions are created by the ionising radiation from O stars, therefore probe the first 2~My of star formation. The peak of \\Al emission from a coeval group of massive stars only begins at 2~My, continuing for another 10~My \\citep{plueschke01}. This demonstrates that \\Al and the other analyses of galactic rotation may complement each other.  Our model parameters, chosen to represent supernova-produced \\Al with ejection into an typical ISM environment, predict a line width of $\\approx$~1~keV for the inner Galaxy. This value is on the low side of the recent RHESSI measurement \\citep{smith03}, and significantly below the ``broad'' line reported from the GRIS balloon measurement \\citep{naya96}.  We suggest, therefore, that \\Al decay in the interstellar medium is likely to teach us more about the ejection and expansion characteristics of nucleosynthesis ejecta, and thus about the morphology of the interstellar medium in the vicinity of massive stars, our assumed sources of \\Al in the Galaxy. For example, if \\Al sources typically are surrounded by cavities from earlier massive-star action, this could lead to some additional broadening of the observed \\Al gamma-ray line. Further imaging spectroscopy analysis of RHESSI data and of measurements with the SPI Ge spectrometer on INTEGRAL \\citep[launched in October 2002;][]{winkler03} promise to provide the data for such studies."
    },
    "0311/gr-qc0311013_arXiv.txt": {
        "abstract": "We present a complete treatment in the strong field limit of gravitational retro-lensing by a static spherically symmetric compact object having a photon sphere. The results are compared with those corresponding to ordinary lensing in similar strong field situations. As examples of application of the formalism, a supermassive black hole at the galactic center and a stellar mass black hole in the galactic halo are studied as retro-lenses, in both cases using Schwarzschild and Reissner-Nordstr\\\"{o}m geometries. ",
        "introduction": "Gravitational lensing is an important tool in astrophysics. Observable phenomena of lensing by stars and galaxies can be explained in the weak field approximation, that is to say, keeping only the first order term in the expansion of the deflection angle \\cite{schne}. But when the lens is a black hole, a full general relativistic treatment is in order.\\\\  Recently, several articles studying strong field lensing appeared in the literature. Virbhadra \\& Ellis \\cite{virbha1} numerically analyzed  the situation where the lens is a Schwarzschild black hole placed in the center of the Galaxy. They obtained the lens equation with an asymptotically flat metric and found that, besides the primary and secondary images, two infinite sets of faint relativistic images are formed by photons that make complete turns (in both directions of rotation) around the black hole before reaching the observer. Fritelli, Kling \\& Newman \\cite{fritelli} found an exact lens equation without any reference to a background metric and compared their results with those of Virbhadra \\& Ellis. Bozza et al \\cite{bozza1} developed an approximate analytical method (generally called, as well, the strong field limit) by which they obtained the positions and magnifications of the relativistic images for the Schwarzschild black hole. Subsequently, Eiroa, Romero \\& Torres \\cite{eiroto} applied the strong field limit to the Reissner-Nordstr\\\"{o}m black hole and discussed the possibility of detection of the relativistic images in the next few years. Bozza \\cite{bozza2} extended the strong field limit to any static spherically symmetric lens and analyzed the case of a charged Brans-Dicke black hole. Bhadra \\cite{bhadra} studied the Gibbons-Maeda-Grafinkle-Horowitz-Strominger charged black hole of string theory in the strong field limit. Virbhadra \\& Ellis numerically investigated the lensing by naked singularities \\cite{virbha2}. Bozza \\& Mancini \\cite{bozza3} used the strong field limit to study the time delay between different relativistic images, showing that different types of black holes are characterized by different time delays. Spinning black holes are more difficult to tackle. Bozza \\cite{bozza4} analyzed the case of quasi equatorial lensing by rotating black holes and V\\'{a}zquez \\& Esteban \\cite{vazquez} studied  the Kerr black hole where the observer and source have arbitrary inclinations. Many of these recent advances (if not all) deal with theoretical implications of General Relativity in a regime which is well beyond our current technological capabilities for detection. Henceforth, the emphasis of these studies is not as much in predicting new phenomena that could be tested in a short period of time as it is in gaining insight on how gravitational lensing in particular, and gravity in general, behave in the strong field regimes. \\\\  In ordinary lensing situations, the lens is placed between the source and the observer. But if the lens is a compact object with a photon sphere and the observer is placed  between the source and the lens, it leads to the formation of images with deflection angles closer to odd multiples of $\\pi$. The observer sees the images in front and the source behind. This situation is called retro-lensing and was studied for the first time by Holtz \\& Wheeler \\cite{holtz}, who analyzed only the two stronger images for a black hole placed in the galactic bulge with the sun as source. They also proposed retro-lensing as a new mechanism for searching black holes. A similar idea was suggested in the context of defocusing gravitational lensing by Capozziello et al \\cite{capoz}. De Paolis et al \\cite{depaolis} recently analyzed the retro-lensing scenario for a bright star close to the massive black hole at the galactic center.\\\\  In this paper we give a complete treatment of retro-lensing by a static, spherically symmetric lens, using the strong field limit. In Section 2 we obtain the lens equation for retro-lensing and the deflection angle in the strong field limit. In Section 3, the position and magnification of the images are calculated. In Section 4, the results are compared with those for ordinary lensing and astrophysical examples are given in Section 5. Final concluding remarks are given in Section 6. Throughout the paper we use units such that $G=c=1$. ",
        "conclusions": "We have presented the strong field limit formalism for gravitational retrolensing. We have found the positions and magnifications of the relativistic images within the strong field limit, and studied some cases which might present particular interest and be a benchmark point for comparison with the normal lensing situation. As in the case of ordinary lensing, the relativistic images are faint. However, if we compare the magnification of  retro-lensed relativistic images with those of the ordinary strong field lensing, we find that the former are significanlty stronger. Additionally, as no brighter weak field images appear in retro-lensing, it would be easier to detect them. The angular separation of the relativistic images is of course beyond the angular resolution of current or foreseeable technologies, and not only angular resolution but also sensitivity improvements are needed to detect them (see the discussion in Eiroa, Romero \\& Torres \\cite{eiroto}). The image separations and magnifications present  values that make the observational identification of these relativistic  images a challenge for the future: we are not aware of any current research, in any wavelength, which would allow to resolve the retro-lensing images in the next years. In the case of charged black holes, it is only for extreme values of charge that important changes are obtained."
    },
    "0311/astro-ph0311054_arXiv.txt": {
        "abstract": "The arrival times of gravitational waves and optical light from orbiting binaries provide a mechanism to understand the propagation speed of gravity when compared to that of light or electromagnetic radiation. This is achieved with a measurement of any offset between optically derived orbital phase related to that derived from gravitational wave data, at a specified location of one binary component with respect to the other. Using a sample of close white dwarf binaries (CWDBs) detectable with the Laser Interferometer Space Antenna (LISA) and optical light curve data related to binary eclipses from meter-class telescopes for the same sample, we determine the accuracy to which orbital phase differences can be extracted. We consider an  application of these measurements involving a  variation to the speed of gravity, when compared to the speed of light, due to a massive graviton. For a subsample of $\\sim$ 400 CWDBs with high signal-to-noise gravitational wave and optical data with magnitudes brighter than 25, the combined upper limit on the graviton mass is at the level of $\\sim 6 \\times 10^{-24}$ eV. This limit is two orders of magnitude better than the present limit derived by Yukawa-correction arguments related to the Newtonian potential and applied to the Solar-system. ",
        "introduction": "In a recent paper, we discussed the optical followup study of close white dwarf binaries (CWDBs) that will be detectable with the Laser Interferometer Space Antenna (LISA) mission \\cite{Cooetal03}. Recent estimates suggest that close to 3000 binaries will be detected via gravitational waves (GWs) at frequencies above 3 mHz, with most of them restricted to the frequency range between 3 mHz and 6 mHz \\cite{webb,Neletal02,Set02}. These binaries will only be localized to the accuracy of $\\sim$ one sqr. degree with a three-year LISA observation based on the Doppler effect related to the orbit around the Sun and the amplitude modulation due to the rotation of LISA detectors. For a precise location of an individual binary, optical observations are required; an easy identification of the binary is facilitated by the presence of eclipses in the optical light curve with a period twice the GW period. Considering the impact on astronomy and astrophysics, in general, one can certainly expect detailed optical followup observations of LISA binaries \\cite{Cooetal03}. In addition to CWDB related physics, such as tidal heating, using both the optical light curve and gravitational wave data, we can also extract important information related to certain fundamental parameters in physics. We discuss such a possibility in the present paper by considering observational prospects in the LISA era.  In addition to a localization of the direction from which GWs are emitted, with gravitational waves, LISA also allow a measurement of the binary distance, the chirp mass, the binary orbital inclination angle, and the binary orbital phase (modulus $\\pi$). The optical followup observations also allow a measurement of the orbital phase. In general, one expects the orbital phase measured with gravity wave data to agree precisely with that of the optical light curve. A difference in the two is only expected if optical light travel at a different speed than that of gravity, which determines the propagation of gravitational waves. While a difference is not expected, a potential non-zero mass related to the graviton particle can lead to a slight difference. Thus, a comparison of the orbital phase related to an optical eclipsing light curve of CWDBs with the phase determined from LISA data will allow a constraint on the graviton mass.  Note that previous studies have considered such a possibility, again with respect to LIGO and LISA data \\cite{Wil98,Will:um,LarHis00,Cuetal02}. While one of the considered methods involve the dispersion of the gravity waves alone, as a function of the frequency \\cite{Wil98,Will:um}, the comparison between optical or electromagnetic data and gravity waves is expected to provide an improved constraint. Previous analyses, however, only considered a handful of objects which are already known to be eclipsing binaries in optical data and are expected to be GW sources detectable with LISA (e.g., 4U1820-30) \\cite{Cuetal02}. The present study considers a sample of CWDBs detectable and locatable with LISA above its confusion noise and is bright enough optically for followup studies with a few meter-class telescope.  The paper is organized as follows.  In \\S~\\ref{sec:lisa}, we briefly discuss observations with LISA, requirements for an optical followup study and, then, consider the extent to which orbital phases of individual binaries can be determined. In \\S~\\ref{sec:results}, we put these measurements in the context of an improved limit on the graviton mass. We conclude with a summary in \\S~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}   The arrival times of gravitational waves and optical light from orbiting binaries provide a mechanism to understand the propagation effects of gravity, when compared to light. This is easily achieved via binary orbital phase measurement and by looking for an phase offset in the optically derived orbital phase related to that derived from gravity wave data. Using a sample of close white dwarf binaries (CWDBs) detectable with the Laser Interferometer Space Antenna (LISA) and the associated optical light curve data related to binary eclipses, we determine the accuracy to which orbital phase difference can be extracted. We consider an application of these measurements, determining   an upper limit on the graviton mass.  For a subsample of $\\sim$ 400 CWDBs with high signal-to-noise gravity wave and optical data with magnitudes brighter than 25, the combined upper limit on the graviton mass is at the level of $\\sim 6 \\times 10^{-24}$ eV, which is two orders of magnitude better than the limit derived by Yukawa-correction arguments related to the Newtonian potential and applied to the Solar-system."
    },
    "0311/astro-ph0311262_arXiv.txt": {
        "abstract": "We model the spectral energy distribution of the type 1 Seyfert galaxies, fitting data from simultaneous optical, UV, and X-ray monitoring observations.  We assume a geometry consisting of a hot central Comptonizing region surrounded by a thin accretion disk.  The properties of the disk and the hot central region are determined by the feedback occurring between the hot Comptonizing region and thermal reprocessing in the disk that, along with viscous dissipation, provides the seed photons for the Comptonization process.  The constraints imposed upon this model by existing multiwavelength data allow us to derive limits on the central black hole mass, the accretion rate, and the radius of the transition region between the thin outer disk and the geometrically thick, hot inner region in these objects. ",
        "introduction": "In a series of three papers (Chiang \\& Blaes 2001, Chiang 2002, and Chiang \\& Blaes 2003), we showed that the optical/UV (OUV) and X-ray light curves obtained from multiwavelength monitoring campaigns of bright type 1 Seyfert galaxies could be readily explained by thermal reprocessing of radiation from a central X-ray source by a thin accretion disk.  These observations had been puzzling since there is no apparent correlation between the OUV and X-ray light curves (Nandra et al. 1998; Edelson et al. 2000) as predicted by early models of thermal reprocessing (Krolik et al.\\ 1991; Courvoisier \\& Clavel 1991; Collin-Souffrin 1991).  The resolution of these difficulties were pointed to by Nandra et al. (2000) who noted that for the NGC 7469 observations the OUV variations {\\em are} correlated with the X-ray spectral index as well as with an extrapolation of the measured 2--10 keV fluxes up to 100 keV.  This latter result leads to the natural conclusion that the disk flux from thermal reprocessing is driven by the bolometric X-ray luminosity, which can differ substantially from that seen in narrow band observations (e.g., 2--10 keV) owing to variations in the X-ray spectral shape.  If the hard X-ray emission in type 1 Seyfert galaxies is produced by thermal Comptonization in a hot plasma, then the X-ray spectral index is determined by the plasma temperature and the optical depth to Compton scattering.  Furthermore, if the thermal reprocessed emission comprises a significant fraction of the seed photons for the Comptonization process, then the feedback between the plasma and the reprocessing material acts as a thermostat, regulating the temperature of the plasma (Stern, et al. 1995). For a given optical depth, the spectral index of the emission is then set by the strength of the feedback, and that, in turn, is determined by the geometry of the hot plasma and the reprocessing material and by the amount of seed photons that are not due to reprocessing (Poutanen, Krolik, \\& Ryde 1997; Zdziarski, Lubi\\'nski, \\& Smith 1999). ",
        "conclusions": ""
    },
    "0311/astro-ph0311112_arXiv.txt": {
        "abstract": "We present X-ray spectra of four ultraluminous infrared galaxies (ULIRGs) with detectable broad near-infrared emission lines produced by active galactic nuclei (AGNs): Mrk 463, PKS 1345+12, IRAS 05189$-$2524, and IRAS 07598+6508. With the exception of IRAS 07598+6508, high quality X-ray spectra obtained with {\\it XMM} or {\\it Chandra} show modest (30--340 eV) equivalent widths of the 6.4 keV Fe K$\\alpha$ emission line and clear signatures for absorption at a level of 4--33 $\\times$ 10$^{22}$ cm$^{-2}$ for the main power-law components from the AGNs. These spectral properties are typical of Compton-thin AGNs, and so we estimate absorption-corrected 2--10 keV X-ray luminosities for the AGNs, Lx(2--10keV), using the Compton-thin assumption. We compare the Lx(2--10 keV) values with broad optical/near-infrared emission-line luminosities, and confirm a previous finding by Imanishi \\& Ueno that the Lx(2--10 keV) to broad-emission-line luminosity ratios in ULIRGs are systematically lower than those of moderately infrared-luminous type-1 AGNs. A comparison of independent energy diagnostic methods suggests that the AGNs are underluminous in the 2--10 keV band with respect to their overall spectral energy distributions, as opposed to the broad emission lines being overluminous. This X-ray under-luminosity should be taken into account when using 2--10 keV X-ray data to investigate the energetic contribution from AGNs to the large infrared luminosities of ULIRGs. ",
        "introduction": "The star-formation rate is one of the most important indicators of the nature of a galaxy; it is usually estimated from continuum or emission-line luminosities \\citep{ken98}. In dusty objects, the infrared luminosity is particularly important, because the correction for dust extinction is generally not significant. An observed infrared luminosity can be converted into a star-formation rate, assuming that the luminosity is powered primarily by star-formation activity. Although this assumption seems to be true in the majority of moderately infrared-luminous galaxies with L$_{\\rm IR}$ $<$ 10$^{12}$L$_{\\odot}$, it has been suggested that active galactic nuclei (AGNs) provide an important energetic contribution in ultraluminous infrared galaxies (ULIRGs) with L$_{\\rm IR}$ $>$ 10$^{12}$L$_{\\odot}$ (Veilleux, Kim, \\& Sanders 1999a). Therefore, the AGN contribution to the observed infrared luminosity of ULIRGs needs to be properly estimated.  To estimate this contribution, an indicator which can trace AGN power, by disentangling from star-formation activity, is required. Two useful indicators are currently known. The first is the absorption-corrected intrinsic 2--10 keV X-ray luminosity. In an AGN, strong 2--10 keV emission is produced by the inverse-Compton process in the close vicinity of an accretion disk around a central supermassive black hole, whereas the 2--10 keV emission from star-formation is usually much weaker. The second indicator is the broad hydrogen emission lines in the optical and near-infrared domains, whose line widths are larger than $\\sim$1500 km s$^{-1}$ in full-width at half maximum (FWHM). Such broad emission lines are believed to originate in high-velocity gas (the so-called ``broad-line regions'') around a central supermassive black hole in an AGN, rather than being associated with phenomena related to star formation (Veilleux, Sanders, \\& Kim 1997b). In many dust-unobscured type-1 AGNs with moderate infrared luminosity, both intrinsic 2--10 keV and broad emission-line luminosities have been derived, and a correlation between these luminosities has been found \\citep{war88}. If this correlation holds for all types of AGNs, then both of these indicators can be used to estimate AGN's power in a galaxy.  However, results using {\\it ASCA} data have shown that the absorption-corrected 2--10 keV to broad emission-line luminosity ratios in some ULIRGs are systematically lower than those of dust-unobscured, moderately infrared-luminous, type-1 AGNs \\citep{ima99a}. One possible explanation is the large time variability in X-ray luminosity. Some AGNs that were initially classified as X-ray underluminous based on X-ray data taken at one epoch showed normal X-ray luminosities when observed at a later epoch \\citep{ris03}. A second possibility is the misidentification of the origin of the detected 2--10 keV emission. Based on modest equivalent widths of the Fe K$\\alpha$ line at 6.4 keV (EW$_{6.4}$ $<$ several $\\times$ 100 eV) and signatures for Compton-thin absorption (N$_{\\rm H}$ $<$ 10$^{24}$ cm$^{-2}$) from {\\it ASCA} spectra, absorption-corrected 2--10 keV luminosities, Lx(2--10 keV), were derived, assuming that the detected X-ray emission is a direct component from an AGN behind Compton-thin absorbing material \\citep{awa91,kro94,ghi94}. If, instead, the detected X-ray emission is a scattered/reflected component of AGN emission behind Compton-thick absorbing material (N$_{\\rm H}$ $>$ 10$^{24}$ cm$^{-2}$), the intrinsic 2--10 keV luminosity will be larger than the estimate based on the Compton-thin assumption. In the ULIRG, Superantennae, although no significant Fe K$\\alpha$ emission was seen in an {\\it ASCA} spectrum due to limited photon statistics \\citep{ima99b,pap00}, a strong Fe K$\\alpha$ line (EW$_{6.4}$ $\\sim$ 1.4 keV) was reported in a higher quality spectrum obtained later with {\\it XMM} \\citep{bra03a}. This demonstrated that the detected 2--10 keV emission is produced by a scattered/reflected component, which was difficult to be recognized based on {\\it ASCA} data. Higher quality {\\it XMM} and {\\it Chandra} data obtained at a different epoch to the {\\it ASCA} era are very useful to (1) investigate whether the X-ray underluminosity of AGNs in ULIRGs found by \\citet{ima99a} is caused by X-ray time variability and (2) better understand the origin of detected 2--10 keV emission using higher photon statistics at 2--10 keV and, in particular, near 6.4 keV.  In this paper, we present {\\it XMM} and {\\it Chandra} X-ray spectra of ULIRGs that have detectable broad optical/near-infrared emission lines, in order to investigate the 2--10 keV to broad-line luminosity ratios. Throughout this paper, we adopt H$_{0}$ $=$ 75 km s$^{-1}$ Mpc$^{-1}$. Luminosities derived in previous literature using different values of H$_{0}$ were re-calculated. Lx(2--10 keV) denotes the absorption-corrected 2--10 keV luminosity of an AGN, unless otherwise stated. ",
        "conclusions": "\\subsection{The 2--10 keV to Broad Line Luminosity Ratio}  For Mrk 463, PKS 1345+12, and IRAS 05189$-$2524, the modest EW$_{6.4}$ ($<$ several $\\times$ 100 eV) and clear signatures for absorption with N$_{\\rm H}$ $<$ several $\\times$ 10$^{23}$ cm$^{-2}$ suggest that the detected hard component is direct emission from an AGN behind Compton-thin absorbing material \\citep{awa91,kro94,ghi94}, and so the absorption-corrected 2--10 keV luminosities, Lx(2--10 keV), estimated in $\\S$4, can be taken as the intrinsic X-ray luminosity of the AGN. Therefore, the 2--10 keV to broad-line luminosity ratios for the AGNs in these ULIRGs can be compared with those for moderately infrared-luminous type-1 AGNs \\citep{war88}. For IRAS 07598+6508, however, the constraint on EW$_{6.4}$ ($<$13 keV) is so loose that it is unclear whether the observed 2--10 keV flux is a direct component behind Compton-thin absorbing material (EW$_{6.4}$ $<$ several $\\times$ 100 eV). The extremely small Lx(2--10 keV)/L$_{\\rm IR}$ ratio and small N$_{\\rm H}$ ($\\sim$8 $\\times$ 10$^{20}$ cm$^{-2}$) suggest that the observed 2--10 keV emission is a scattered/reflected component behind Compton-thick absorbing material. This source will be excluded from the following discussions.  The correlation between 2--10 keV and broad emission lines established for type-1 AGNs \\citep{war88} is for the broad component of an optical H$\\alpha$ ($\\lambda_{\\rm rest}$ = 0.65 $\\mu$m) emission line. Thus, we need to convert the luminosity of a broad near-infrared emission line (Pa$\\alpha$ for PKS 1345+12 and IRAS 05189$-$2524 \\footnote{ In IRAS 05189$-$2524, broad components were detected at H$\\alpha$, Pa$\\beta$, and Pa$\\alpha$ (see $\\S$ 2). We use the Pa$\\alpha$ line, because dust extinction is the lowest. } ; Veilleux et al. 1997b,1999b, and Pa$\\beta$ for Mrk 463; Veilleux et al. 1997a) to that of broad H$\\alpha$. For emission lines from the broad-line regions in an AGN, the case-B assumption is not applicable to low-level transition lines, such as H$\\alpha$, due to collisional effects \\citep{ost89}. For the broad H$\\alpha$ to Pa$\\alpha$ luminosity ratio, the value of 20 observed in dust-unobscured type-1 AGNs \\citep{hil96} was adopted. Since the deviation of this ratio from case-B ($\\sim$10) is only a factor of $\\sim$2, the choice of this ratio will not affect our main conclusions significantly. For high-transition broad emission lines, such as Pa$\\alpha$ and Pa$\\beta$, case-B is applicable \\citep{rhe00}. Thus, a broad Pa$\\alpha$ to Pa$\\beta$ luminosity ratio of 1.9 (case-B) was adopted.  Absorption-{\\it corrected} 2--10 keV and extinction-{\\it uncorrected} broad emission-line luminosities for the three ULIRGs are plotted in Figure~\\ref{fig3} (solid squares), where the estimated broad H$\\alpha$ emission-line luminosities are $\\sim$5, $\\sim$3, and $\\sim$1 $\\times$ 10$^{43}$ ergs s$^{-1}$ for Mrk 463, PKS 1345+12, and IRAS 05189$-$2524, respectively. The type-1 AGNs studied by \\citet{war88} are plotted for comparison (open circles). In addition to the three ULIRGs, the other two Seyfert-2 ULIRGs with detectable broad near-infrared Pa$\\alpha$ emission lines, IRAS 20460+1925 and IRAS 23060+0505 \\citep{vei97b}, have available {\\it ASCA} 2--10 keV spectra \\citep{oga97,bra97}. Since their observed 2--10 keV fluxes are sufficiently high ($>$10$^{-12}$ ergs s$^{-1}$ cm$^{-2}$), strong constraints were obtained for EW$_{6.4}$ (260$^{+145}_{-137}$ eV for IRAS 20460+1925; Ogasaka et al. 1997, $<$290 eV for IRAS 23060+0505; Brandt et al. 1997). The detected 2--10 keV fluxes show signatures for absorption with N$_{\\rm H}$ $=$ 2--8 $\\times$ 10$^{22}$ cm$^{-2}$, so that the detected 2--10 keV flux is taken as a direct component from an AGN behind Compton-thin absorbing material. For IRAS 20460+1925, Lx(2--10 keV) is $\\sim$1 $\\times$ 10$^{44}$ ergs s$^{-1}$ \\citep{oga97} and Lx(2--10 keV)/L$_{\\rm IR}$ is $\\sim$1 $\\times$ 10$^{-2}$. For IRAS 23060+0505, Lx(2--10 keV) is $\\sim$2 $\\times$ 10$^{44}$ ergs s$^{-1}$ \\citep{bra97} and Lx(2--10 keV)/L$_{\\rm IR}$ is $\\sim$2 $\\times$ 10$^{-2}$. Broad H$\\alpha$ emission-line luminosities were derived from Pa$\\alpha$ data \\citep{vei97b} using the method described above and estimated to be $\\sim$2 and $\\sim$1 $\\times$ 10$^{44}$ ergs s$^{-1}$ for IRAS 20460+1925 and IRAS 23060+0505, respectively. These two ULIRGs are also plotted in Fig.~\\ref{fig3} (solid squares). The five Seyfert-2 ULIRGs show absorption-{\\it corrected} 2--10 keV to extinction-{\\it uncorrected} broad line luminosity ratios roughly an order of magnitude smaller than those in dust-unobscured type-1 AGNs with modest infrared luminosities. While no significant dust extinction is expected for broad emission lines in type-1 AGNs, some degree of dust extinction may be present for the five Seyfert-2 ULIRGs. For these five ULIRGs, since near-infrared Pa$\\alpha$ or Pa$\\beta$ emission lines were used to estimate the broad H$\\alpha$ luminosities, the effects of dust extinction are expected to be reduced. If some amount of dust extinction is present for the broad near-infrared emission lines and its correction is applied, these five ULIRGs will move to the right in Fig. 3, making the deviation of these ULIRGs even larger. Even using new {\\it XMM} and {\\it Chandra} data, we confirmed systematically smaller 2--10 keV to broad line luminosity ratios in ULIRGs, originally suggested by \\citet{ima99a}.  It is unlikely that the systematic difference is caused by the uncertainty of the luminosity conversion from broad Pa$\\alpha$ or Pa$\\beta$ to H$\\alpha$ for the following two reasons. First, in a number of moderately infrared-luminous type-2 AGNs, the 2--10 keV to broad-H$\\alpha$ luminosity ratios derived from observed broad near-infrared line luminosities agree with those for unobscured type-1 AGNs \\citep{ima99a}. Second, for IRAS 05189$-$2524, the intrinsic broad H$\\alpha$ luminosity estimated from the observed broad Pa$\\alpha$ luminosity (L$_{\\rm H \\alpha}$ $\\sim$ 1 $\\times$ 10$^{43}$ ergs s$^{-1}$; see previous paragraph) is in good agreement with that derived directly from the broad H$\\alpha$ luminosity \\citep{you96} after correction for dust extinction estimated by \\citet{sev01} is applied (L$_{\\rm H \\alpha}$ $\\sim$ 1 $\\times$ 10$^{43}$ ergs s$^{-1}$). Both of these results suggest that (1) the luminosity conversion from broad near-infrared emission lines to broad H$\\alpha$ is reasonable, and that (2) at least for some sources with detectable broad near-infrared emission lines, dust extinction is not significant in the near-infrared domain. Thus, the significant deviation of ULIRGs in Fig. 3 is believed to be real.  \\subsection{X-ray Underluminous or Broad Line Overluminous?}  Estimating the fractional AGN contribution to the large infrared luminosities of ULIRGs is of particular importance to the study of their energy sources. It is not clear which of the two AGN indicators, 2--10 keV emission or broad optical/near-infrared emission lines, is a better tracer of this fraction. The 2--10 keV luminosity depends primarily on the number of hot electrons in the close vicinity of an accretion disk around a central supermassive blackhole in an AGN and the number of incident UV photons to be Compton up-scattered. The broad emission-line luminosity is roughly proportional to the product of the number of UV photons from the central AGN and the covering factor of broad-line gas clouds. The infrared luminosity of a ULIRG powered by an AGN is determined mainly by the number of UV photons from the AGN and by the covering factor of dust. In the nuclear region of a ULIRG, a large amount of gas and dust is available \\citep{san96}, which may increase the covering factors both of dust and broad-line gas clouds around the AGN. If this is the case, broad emission lines are a better indicator of the AGN's contribution to the infrared luminosity, and AGNs in ULIRGs are underluminous in 2--10 keV X-rays.  In fact, in a number of ULIRGs, although an AGN is consistently suggested to be the primary energy source by some independent methods, only energy diagnostics based on the 2--10 keV luminosity suggest that the AGN is energetically insufficient. Examples include IRAS 05189$-$2524 \\citep{vei99b,soi00,imd00,sev01} and Mrk 463 \\citep{uen96,gen98,ima02}. Thus, it is more likely that the AGNs in ULIRGs are underluminous in X-rays with respect to their overall spectral energy distributions, than that their broad lines are overluminous.  If X-ray underluminosity (relative to UV) is an intrinsic property of ULIRGs, one possible explanation is a high mass accretion rate onto a central supermassive black hole \\citep{bec03}. This is plausible in ULIRGs, given the high nuclear concentration of gas and dust \\citep{san96}.  Alternatively, the apparent X-ray underluminosity of AGNs in ULIRGs may be caused by our misidentification of the origin of the detected 2--10 keV emission from ULIRGs. Using reasonable assumptions for an AGN, the scattered/reflected X-ray component of AGN emission behind Compton-thick absorbing material is expected to show a very large EW$_{6.4}$ with $>$1 keV and only weak absorption \\citep{mat03}. A modest EW$_{6.4}$ and clear Compton-thin absorption signature, as were found in the five ULIRGs, make it most plausible that we are observing direct X-ray emission from an AGN behind Compton-thin absorbing material \\citep{awa91,kro94,ghi94}. However, Thomson scattering by highly ionized gas produces a small EW$_{6.4}$. If the scattered component suffers some degree of absorption, both the modest EW$_{6.4}$ and clear signature for Compton-thin absorption can be explained \\citep{mal00}. In this scenario, the intrinsic 2--10 keV luminosity of an AGN could be substantially larger, depending on the scattering efficiency. Since such X-ray spectra are usually interpreted as Compton-thin, and the absorption-corrected 2--10 keV luminosity is determined based on this assumption, we will underestimate the true Lx(2--10 keV).  A third possibility which could explain the observed X-ray underluminosity is partial Compton-thick X-ray absorption. If the X-ray-absorbing material in front of an AGN consists of small clumps of Compton-thick and thin clouds, whose size scales are substantially smaller than the size of the X-ray-emitting region of the AGN, then the 2--10 keV spectrum will suggest Compton-thin absorption. The actual Lx(2--10 keV) could be significantly larger than a simple estimate assuming uniform Compton-thin X-ray absorption. The covering factor of Compton-thick clouds in front of an AGN along our line-of-sight determines the degree of our underestimate for Lx(2--10 keV), but an X-ray spectrum at $>$10 keV is required to determine this \\citep{bra03b}.  Whether the X-ray underluminosity of AGNs in ULIRGs is intrinsic or caused by our underestimate of Lx(2--10 keV), the use of 2--10 keV observations to estimate AGNs' contribution to the infrared luminosities of ULIRGs will be quite misleading, unless the X-ray underluminosity is properly taken into account.  The discovery of underluminous X-ray emission relative to broad emission lines from AGNs is only for ULIRGs with modest dust obscuration along our sightlines, because in order for broad optical/near-infrared emission lines to be detected, dust obscuration toward AGNs must be relatively low. The putative AGNs in the majority of ULIRGs are so highly dust-obscured that broad emission lines are not detectable \\citep{vei99b}, and so we have no information on the broad line luminosities. It is unclear whether X-ray underluminosity is a general property of AGNs in the majority of ULIRGs, and there is currently no way to determine this. However, when using 2--10 keV data to determine the energy contribution of AGNs in ULIRGs, it is important to consider that the AGN contribution may be significantly underestimated."
    },
    "0311/astro-ph0311438_arXiv.txt": {
        "abstract": "It is suggested that a Poynting flux-dominated outflow with a sufficiently strong magnetic field can pick up hard X-ray photons when the magnetic field is of sufficient strength. The zeroth generation  X-rays pair produce, and the pairs radiate extremely energetic first generation photons that could be detected by extensive air shower arrays and/or MILAGRO if they escape the production site intact. Giant flares from magnetars may thus yield bursts of UHE photons. GRB-associated Poynting flows may be unstable to pair production near their source, and their energy rapidly converted to pairs. ",
        "introduction": "There is growing evidence that soft gamma ray repeaters (SGR's) and anomalous X-ray pulsars are  objects with ultrastrong magnetic fields, $10^{14.5}$ G or more \\cite{DT,TD,{WO}}. Giant flares from SGR's appear to be caused by a reconnection or reconfiguration of an  ultrastrong magnetic field, in a manner similar to the mechanism that causes solar flares. These flares have two stages: a) the bright phase, lasting $\\sim 10^{-1}$s, when the energy is presumably released, b) a period lasting several minutes during which time the remaining plasma is apparently trapped on closed field lines.  More than $10^{44}$ ergs may be released in giant flares, while smaller, more frequent bursts field between $10^{38}$ and $10^{42}$ ergs.  A rapidly rotating neutron star with a magnetic field of at least $3 \\times 10^{14}$ G  was proposed by Usov \\cite{US} as being the source of gamma ray bursts (GRB's). One indication (of sorts) is that synchrotron models of afterglow frequently invoke a rather large fraction of the total fireball energy in Poynting flux to obtain good fits (e.g. \\cite{LW}). Other models \\cite{T94,LPB} explicitly invoke Poynting flux as being the original form in which GRB energy is expelled by the compact object. The rate at which  magnetars form can be estimated by dividing the number in the Galaxy ($\\sim10$) by their age ($\\sim 10^4$ yr), and the rate so estimated,  once per $\\sim 10^3$ years per galaxy, is consistent with what is needed to produce GRB's. Perhaps magnetars are the fossil remnants of GRB's.  It would be useful to identify a qualitative consequence of ultrastrong magnetic fields, one that would not exist to any degree were the fields much smaller.  In this letter, it is suggested that pair production by photons in a magnetic field could be just such a consequence and could result in the emission of ultrahigh energy (UHE) photons.   The basic point is that such coupling is possible in a magnetic field only when the product of the photon energy $\\epsilon_{\\gamma}$ and field strength $B_{\\bot}$ normal to the photon path  satisfies  \\begin{equation} \\epsilon_{\\gamma}B_{\\bot} \\ge 2 \\times 10^{18}eV-G. \\label{criterion} \\end{equation} Thus, when the field strength exceeds $10^{14}$G, a photon above 20 KeV (routine for typical astrophysical thermal X-ray emitters)  suffices. (Hereafter the subscript in $B_{\\bot}$ will be dropped.)  We will assume  in this letter  without rigorous proof that, in a flare-like event,  plasma can be ejected from a magnetar or similar object to infinity, much like plasma is ejected from the Sun by solar flares, but with a very large Lorentz factor ($\\Gamma\\ge 10^4$) characteristic of the large Alfven Lorentz factor associated with tenuous plasma in a strong magnetic field. At this $\\Gamma$, the plasma moves essentially at the speed of light, to within one part in $10^8$.  There is little relative motion between the photon and the plasma over a hydrodynamic timescale, which, in the frame of the fluid, is much shorter than the crossing time  measured in the lab frame. ",
        "conclusions": ""
    },
    "0311/astro-ph0311324_arXiv.txt": {
        "abstract": "We isolate samples of $\\nbhbbr$ bright ($g < 18$) and $\\nbhbfaint$ faint ($g > 18$) high-Galactic latitude blue horizontal branch stars with photometry and spectroscopy in the Sloan Digital Sky Survey (SDSS).  Comparison of independent photometric and spectroscopic selection criteria indicates that contamination from F and blue-straggler stars is less than $10\\%$ for bright stars ($g<18$) and about $25\\%$ for faint stars ($g>18$), and this is qualitatively confirmed by proper motions based on the USNO-A catalog as first epoch.  Analysis of repeated observations shows that the errors in radial velocity are $\\approx 26 \\kms$.  A relation between absolute magnitude and color is established using the horizontal branches of halo globular clusters observed by SDSS. Bolometric corrections and colors are synthesized in the SDSS filters from model spectra.  The redder stars agree well in absolute magitude with accepted values for RR Lyrae stars.  The resulting photometric distances are accurate to about $0.2$ magnitudes, with a median of about 25 kpc. Modest clumps in phase space exist and are consistent with the previously reported tidal stream of the Sagittarius dwarf galaxy.  The sample is tabulated in electronic form in the online version of this article, or by request to the authors. ",
        "introduction": "Blue horizontal branch (BHB) stars have long been important objects of study as tracers of the Galactic potential \\citep{pier_1983,sommerlarsen_christensen_carter_1989,clewley_etal_2002}. Their constant absolute magnitude allows accurate space positions to be determined.  Historically, many researchers have used samples of BHB stars, or other standard-candle stars such as RR Lyraes, to study the local halo, i.e.~the vicinity of the solar neighborhood where typical distances are of order the solar Galactocentric radius $R_0 = 8 \\kpc$.  These samples can be used to constrain the kinematics of the local halo \\citep{pier_1984,layden_etal_1996}. The distant halo (i.e.~typical distances $\\gg R_0$), however, is less well constrained (but see \\citet{sommerlarsen_etal_1997}). Isolating a sample of distant BHB stars would enable study of the structure and kinematics of the Galactic halo on larger scales \\citep[][hereafter Paper~II]{halo2}, and estimates of the mass of the Galaxy.  The Sloan Digital Sky Survey (SDSS), currently halfway through a photometric and spectroscopic survey of a quarter of the sky, offers an excellent opportunity to isolate and study BHBs. The SDSS spectroscopic survey targets objects between about $g$ = 15.5 to 19 (the bright limit is set to avoid contamination in adjacent spectroscopic fibers at the slit end of the fiber bundle), although a significant number of stars outside this range is also observed. Thus BHB stars at distances from $\\sim 7$ to $\\sim 60 \\kpc$ from the sun appear in the spectroscopic survey.  This paper presents $\\nbhb$ BHB stars selected from SDSS that probe the outer halo. The plan of the paper is as follows.  Pertinent characteristics of the Sloan Digital Sky Survey are summarized in \\S\\ref{s:sdss}. \\S\\ref{s:color_cut} describes the preliminary color cut. \\S\\ref{s:balmer_cut} discusses spectroscopic BHB classification criteria. For the faint stars, the spectroscopic classification breaks down so we resort to a more stringent color cut described in \\S\\ref{s:fainter_stars}. Radial velocity errors are evaluated in \\S\\ref{s:duplicates}. \\S\\ref{s:absmag} discusses the calculation of bolometric corrections for accurate distances, and evaluates magnitude errors. Some sanity checks are presented in \\S\\ref{s:sanity_checks}, including an analysis of the limited proper motion data available, and a map of the Galactic halo which shows the conspicuous Sagittarius tidal stream (figure~\\ref{f:map_3d}). The results are summarized in \\S\\ref{s:summary}. The sample is available in Table~\\ref{t:data} in the electronic version of this paper. ",
        "conclusions": "\\label{s:summary}  Using broadband colors and spectroscopic surface-gravity indicators, we have selected $\\nbhb$ blue horizontal branch stars from the SDSS.  For the $\\nbhbbr$ bright stars ($g<18$) the contamination is probably $\\lesssim 10\\%$, and for the $\\nbhbfaint$ faint BHB stars ($g>18$) the contamination is probably $\\sim 25\\%$. The absolute magnitude of BHB stars was determined from globular cluster color-magnitude diagrams, and an absolute-magnitude-color relation established. This was used to estimate photometric distances, with errors $\\approx 0.2$ mag. The stars also have measured radial velocities with errors $\\approx 26 \\kms$. Proper motions determined by matching against the digitally scanned Palomar Sky Survey \\citep[USNO-A,][]{monet_etal_1998} also indicate that contamination from the disk is small. We emphasize, however, that no kinematic information was used to select the sample. The data can be downloaded as Table~\\ref{t:data} in the electronic version of this paper."
    },
    "0311/astro-ph0311374.txt": {
        "abstract": "We introduce a procedure to identify very soft X-ray sources (VSSs) in external galaxies. Our immediate goal was to formulate a systematic procedure to identify luminous supersoft X-ray sources (SSSs), %A well-defined selection algorithm allows so as to allow comparisons among galaxies and to study environmental effects. The focus of this paper is on the design of the selection algorithm and on its application to simulated data. In the companion paper we test it by applying it to sources discovered through {\\it Chandra} observations of $4$ galaxies. We find that, in its application to both simulated and real data, our procedure also selects somewhat harder sources, which we call quasisoft. Whereas values of $k\\, T$ for SSSs are typically tens of eV, some quasisoft sources (QSSs) may have $k\\, T$ as high as $\\sim 250-300$ eV. The dominant spectral component of other QSSs may be as soft as SSS spectra, but the spectra may also include a low-luminosity harder component. We sketch physical models for both supersoft and quasisoft sources. Some SSSs are likely to be accreting white dwarfs; some of these may be progenitors of Type Ia supernovae. Most QSSs may be too hot to be white dwarfs. They, together with a subset of SSSs, may be neutron stars or, perhaps most likely, accreting intermediate-mass black holes. ",
        "introduction": "\\subsection{Goals}  First identified as a class of X-ray sources approximately 11 years ago, luminous supersoft X-ray sources (SSSs) were characterized by $k\\, T$ on the order of tens of eV, and $L$ in the range $10^{37}-10^{38}$ erg s$^{-1}$.\\footnote{See e.g., RX J0019.8+2156: Reinsch, Beuermann, \\& Thomas 1993, Beuermann et al. 1995; RX J0925.7-4758: Motch, Hasinger, \\& Pietsch 1994; GQ Mus: Ogelman et al. 1993; 1E 1339.8+2837: Hertz, Grindlay, \\& Bailyn 1993; AG Dra: Greiner et al. 1996; RR Tel: Jordan, M\u00fcrset, \\& Werner 1994; 1E 0035.4-7230: Seward \\& Mitchell 1981; RX J0048.4-7332: Kahabka, Pietsch, \\& Hasinger 1994; RX J0058.6-7146: Kahabka et al. 1994; 1E 0056.8-7154: Wang 1991; RX J0513.9-6951: Schaeidt, Hasinger, \\& Tr\u00fcmper 1993; CAL 83: Long, Helfand, \\& Grabelsky 1981; CAL 87: Long et al. 1981; RX J0550.0-7151: Cowley et al. 1993.} Because radiation from SSSs is readily absorbed by the interstellar medium (ISM), we can study large numbers of them only by searching for them in external galaxies. The spatial resolution of \\chandra, and of \\xmm\\ as well, make such studies possible. The great distances to other large galaxies means, however, that the count rates of most of the SSSs we can detect in them are low. Spectral fits are therefore possible for only a small fraction. The primary goal of the work described in this paper is to provide a well-defined and systematic method to select, from a set of X-ray sources dominated by low-count sources, good SSS candidates. In its present form, the procedure works well for both {\\it Chandra} and {\\it XMM} data.  \\subsection{SSSs in Galaxies}   Soft X-ray sources, including the flagship SSS binaries CAL 83 and CAL 87, had been discovered during a survey of the Magellanic Clouds conducted with the {\\it Einstein} X-ray observatory (Long, Helfand, \\& Grabelsky 1981; Seward \\& Mitchell 1981). The large-scale study of SSSs in other galaxies began with {\\it ROSAT} observations of $\\sim 30$ SSSs in the Magellanic Clouds, and in M31. Because M31 is similar in many ways to other large spiral galaxies, it may be a good guide to the sizes of galactic populations of SSSs. {\\it ROSAT} detected $\\sim 15$ SSSs with luminosities greater than $\\sim 10^{37}$ erg s$^{-1}$  in M31\\footnote{Subsequent reanalysis of the {\\it ROSAT} data identified additional SSS candidates, bringing the number to $34$ (Kahabka 1999)}. These observations, combined with simple models for galactic gas distributions, led to an estimate that, at any given time, M31 houses $\\sim 1000$ active SSSs (\\rd\\ \\& Rappaport 1994). A similar analysis of the sources and gas distribution  of the Milky Way found that its  population of SSSs could be comparably large.  The subsequent discovery of the low-luminosity extension ($10^{35}-10^{36}$ ergs s$^{-1}$) of the class (Greiner \\etal 1999, Greiner \\& \\rd 1999, Patterson et al.\\, 1998)  may imply that the Milky Way's population of SSSs is at least an order of magnitude larger. Recent {\\it Chandra} studies (\\rd\\ 2002, 2003) find convincing evidence of a lower-luminosity component of the class of SSSs in M31 as well. If these estimates are correct, and if our Galaxy and M31 are typical, SSSs may form the single largest galactic population of luminous ($L_X > 10^{35}$ erg s$^{-1}$) X-ray binaries. Whatever their fundamental natures, the likely large sizes of SSSs populations indicates that they may be pumping $10^{40}-10^{41}$ ergs s$^{-1}$ of highly ionizing radiation into the ISM of their host galaxies.  \\subsection{From Supersoft Sources to Quasisoft Sources}  In the absence of a single definitive model for SSSs, and with only a small number of known SSSs in the Galaxy and Magellanic Clouds, we did not want to eliminate sources that were somewhat harder than the canonical sources such as CAL 83. For reasons summarized in \\S 2 and in \\S 3.1, we therefore chose a cut-off in $k\\, T$ of $175$ eV for SSSs. We found, however, that the methods we devised to select sources up to any given temperature, inevitably also included sources that either have  (1) somewhat higher temperatures, but generally in the range $175-250$ eV, or (2) SSS-like spectra, but with some harder emission as well.  This suggests two possible results for the method's application to real data. First, if there are few or no soft sources just slightly harder than the canonical SSSs, the procedure will select only SSSs. Second, if there is a supply of somewhat harder sources, we will select them as well. We have now had opportunities to apply the selection criteria to the $4$ galaxies studied in the companion paper (\\rd\\ \\& Kong 2003), to M104 (\\rd\\ et al. 2003a), M31 (\\rd\\ et al. 2003b), and roughly one dozen additional galaxies (\\rd\\ et al. 2003c). We have found that most galaxies have significant populations of both SSSs and sources with somewhat harder spectra (e.g., $k\\, T < 250$ eV). We refer to the latter as quasisoft sources (QSSs).   Our selection procedure distinguishes between SSSs and QSSs according to which step in the algorithm identifies the source as being very soft. In the galaxies we have studied, spectral fits for the brightest SSS and QSS candidates have verified that the algorithmic classification works. To simplify the terminology, we will sometimes use the term ``very soft source\" (VSS) to refer to both SSSs and QSSs.  \\subsubsection{Quasisoft Sources}  The physical significance of this new class is not yet understood, but there are likely to  be several physical models corresponding to QSSs. First, hot SSSs located behind large gas columns will have photons in the medium energy band, $M$ ($1.1-2$ keV), but may have few photons in the soft band, $S$ ($0.1-1.1$ keV). For such sources, the hardness ratios typically used to identify SSSs will have values not normally associated with SSSs, even though their intrinsic characteristics clearly place them in the SSS category. Second if the detector has poorer than anticipated sensitivity to soft photons, soft sources can appear to be harder than they actually are. Thus, some QSSs are likely to have the same physical characteristics as some other sources identified as SSSs. Finally, some QSSs are likely to be genuinely harder than SSSs, so hard that white dwarf models can be ruled out. As we will discuss in \\S 2, intermediate mass black hole models may be appropriate for such systems, but neutron star or stellar mass black hole models should also be considered.  \\subsection{Open Questions}  Below we list some of the questions we hope to answer with studies that compare VSS populations in different galaxies.  \\noindent (1) What are typical galactic populations of SSSs and QSSs? Irrespective of their fundamental natures, the answer to this question will allow us to estimate the influence of soft X-ray sources as ionizers of the ISM.  \\noindent (2) Are any spiral galaxy parameters related to the relative sizes of SSS  and QSS populations? Answering this question can provide insight into the age of the populations that spawn very soft sources, and hence might help to illuminate their nature. \\rd\\ \\& Rappaport (1994) suggested that for spiral galaxies, the size of the SSS population might scale with blue luminosity, but this has not been tested.  \\noindent (3) Do elliptical galaxies house large SSS/QSS populations? Although it has been suggested that the diffuse soft emission in ellipticals may be due to SSSs (see, e.g., Fabbiano, Kim, \\& Trinchieri 1994), we still know very little about SSSs in ellipticals. If accreting WDs form the largest segment of SSS populations, and if a significant fraction of the donor stars have masses small enough to be typical of the stars found in elliptical galaxies, then we may expect SSSs to be important parts of the X-ray source population in ellipticals.  \\noindent (4) Within spiral galaxies, what are the relative populations of SSSs/QSSs in the galaxy bulges and disks?  \\noindent (5) Do galaxies with massive central black holes have more SSSs or QSSs located within 1 kpc of the nucleus than comparable galaxies without massive central black holes? It has been suggested that some SSSs within the central kpc  of galaxies which harbor massive black holes may actually be the stripped cores of stars that have been tidally disrupted (\\rd\\ et al. 2001). Verification of this hypothesis by studying individual SSSs will be difficult, so statistical studies of SSS populations in a large number of galaxies may  provide the best tests.  \\noindent (6) For all galaxies, are the positions of QSSs and SSSs correlated to the positions of other objects, such as HII regions, planetary nebulae, supernova remnants, or globular clusters? The distances to most external galaxies are too large to allow for convincing optical identifications. It is nevertheless useful to identify the types of populations which tend to be associated with VSSs. This can provide clues to their fundamental natures.   \\noindent (7) Are SSSs significant contributors to the rates of Type Ia SNe? The answer to this question can be achieved by combining information about typical total galactic populations with studies of the viability of the accreting WD models.    \\subsection{Previous Work}  Previous studies of SSSs in external galaxies have used a variety of selection criteria. In NGC 4697, e.g., Sarazin, Irwin, \\& Bregman (2001) identified $3$ SSSs by requiring that ${\\tilde {HR1}}=(\\tilde M-\\tilde S)/(\\tilde M+\\tilde S) = -1$ and ${\\tilde {HR2}}=(\\tilde H-\\tilde S)/(\\tilde H+\\tilde S) = -1,$ where $\\tilde S,$ $\\tilde M,$ and $\\tilde H$ represent the numbers of counts in the bands $0.3-1$ keV, $1-2$ keV, and $2-10$ keV, respectively.\\footnote{We use tildes in these expressions because our own definitions of the soft, medium, and hard X-ray bands is slightly different (see \\S 3).} In their studies of the colors of X-ray sources, Prestwich et al.\\, (2002), used the same criteria, which are satisfied by only a handful of the sources they analyzed, drawn from both M101 and M83. Less restrictive criteria were used by Swartz et al.\\, (2002) to identify SSSs in M81. The criteria ${\\tilde {HR1}} < -0.5,$ ${\\tilde {HR2}} < -0.5$ selected $12$ M81 X-ray sources, $2$ of which were eliminated because they are identified with foreground stars, while one is identified with a supernova remnant (SNR). Pence et al.\\, (2002a) identified $10$ SSSs in M101, but did not specify the selection criteria. The possible physical interpretation of the sources seemed to play a role, as one of the galaxy's softest sources was not counted among the SSSs, perhaps because it appears  to be too luminous to be a nuclear-burning WD (Pence et al.\\, 2002b). Kong et al.\\, (2002a) took another approach, requiring $({\\tilde {HR2}} + \\sigma_{{\\tilde {HR2}}} \\leq -1$ {\\it and} ($[{\\tilde {HR1}} < 0,$] {\\it or} [(${\\tilde {HR1}} + \\sigma_{{\\tilde {HR1}}} \\leq -0.8]$). Fourteen sources in the central $17' \\times 17'$ of M31 satisfied these conditions, of which $2$ were apparently identified with SNRs (Kong et al. 2002b) and $3$ with possible foreground stars.  These latter criteria were developed in parallel with a study of the entire SSS population  of M31 as viewed by {\\it Chandra} (Di~Stefano et al. 2002b, 2003). It was clear, however, that not all of the sources identified by these criteria were equally good candidates for the class. We therefore began by developing strict selection conditions (\\rd\\ \\& Kong 2003 a). These criteria, the so-called HR (hardness ratio) criteria, are the first set of conditions applied in the algorithm presented here. The full algorithm we present here allows us to identify SSSs that are, e.g., at the high-T end (usually taken to be $\\sim 100$ eV) of the class, even if they are highly absorbed. It also covers cases in which a small fraction of the emitted radiation is reprocessed and reemitted in the form of photons of higher energy. In addition the full algorithm  allows us to identify sources that may be somewhat hotter than ``classical\" SSSs. This is the first paper to present the full algorithm, and the companion paper tests its efficacy by studying the results of applying it to data from $4$ galaxies (\\rd\\ \\& Kong 2003 b). It has also been applied to {\\it Chandra} data from M31 (\\rd\\ et al. 2003 a), and from M104 (\\rd\\ et al. 2003 b).   In \\S 2 we give an overview of physical models for VSSs, the class of source that motivated this work. In \\S 3 we provide a phenomenological definition of SSSs designed to help select sources with the physical characteristics discussed in \\S 2. In \\S 4 we outline and test a sequential set of selection criteria that can be applied algorithmically to select SSSs as well as sources that are slightly harder (QSSs), in external galaxies. \\S 5 is devoted to summarizing our conclusions. ",
        "conclusions": "The primary focus of the research described in this paper is to develop a systematic procedure to identify SSSs in external galaxies. To accomplish this we began by crafting an operational definition of SSSs. (See \\S 3.2.) It is based entirely on the spectral shape: a blackbody spectral fit with $k\\, T < 175$ eV, or a power law spectral fit with $\\alpha > 3.5$, or any spectral fit where less than $10\\%$ of the energy is carried by photons with energy $> 1.5$ keV. These criteria can be applied directly to sources providing enough counts to allow a good spectral fit. Since most sources in distant galaxies do not provide enough counts, we have designed a set of conditions which can be applied algorithmically to determine if the true spectrum of a low-count source is likely to satisfy the SSS criteria. Because the conditions work for high-count sources as well as low-count rate sources,  we simply apply our algorithm to all sources in each galaxy to identify the SSSs. The sources are then sub-classified according to the specific conditions they satisfy.  The strongest conditions are called the HR and $3\\sigma$ conditions. We have called a source an SSS if it satisfies either the HR or $3\\sigma$ conditions.  We call all sources that satisfy the weaker conditions, QSSs. Some QSSs may be similar to SSSs, but perhaps more highly absorbed, or they may be genuinely harder than the canonical SSSs. The harder sources may have values of $k\\, T$ between roughly $175$ eV and $350$ eV, or may be very soft sources which include a low-luminosity hard spectral component. The QSSs with dominant spectra that are genuinely harder than those of SSSs (e.g., those with $k\\, T > 175$ eV) are almost certainly not WDs, unless the X-ray emission emanates from only a small portion of the photosphere. They may be accreting NSs or, perhaps most likely, accreting intermediate-mass BHs.  Tests with simulated data show that our criteria accomplish the following. (1) They select a majority of SSSs-- $90\\%$ of thermal models. (2) They select only a small fraction of ``hard'' sources ($4\\%$ of models with $k\\, T > 500$ eV). (3) They select a substantial fraction of quasi-soft sources. For blackbody models, quasi-soft sources were those in the range $175$ eV $< k\\, T < 500$ eV; $70\\%$ of these were selected.  The selection of QSSs is difficult to avoid if we are to succeed at identifying possibly absorbed  high temperature SSSs ($k\\, T$ near $100$ eV), some of which could be high-mass accreting WDs on their way to becoming Type Ia supernovae. On the other hand, the presence of QSSs in the data is intriguing, because there appear to be no known X-ray binary analogs with similar luminosities ($> 10^{36}$ erg s$^{-1}$) in the Galaxy or Magellanic Clouds. Some may be SNRs; repeated X-ray observations and optical emission line studies can distinguish SNRs from X-ray binaries. Those that vary in time are likely to be X-ray binaries either of a type or in a state that is not typical of nearby X-ray binaries. The discovery of such sources in real data sets thereby challenges us to understand what they are.  In the companion paper we describe how we tested the selection algorithm by applying it to real data from $4$ galaxies. We verified that the spectra of the sources identified as SSSs tend to be in the range we used to define SSSs. These investigations have also yielded interesting information about extragalactic SSSs, providing some first steps toward answering the $7$ questions posed in \\S 1.3."
    },
    "0311/nucl-th0311083_arXiv.txt": {
        "abstract": "Structure of hot dense matter at subnuclear densities is investigated by quantum molecular dynamics (QMD) simulations. We analyze nucleon distributions and nuclear shapes using two-point correlation functions and Minkowski functionals to determine the phase-separation line and to classify the phase of nuclear matter in terms of the nuclear structure. Obtained phase diagrams show that the density of the phase boundaries between the different nuclear structures decreases with increasing temperature due to the thermal expansion of nuclear matter region. The critical temperature for the phase separation is $\\agt 6$ MeV for the proton fraction $x=0.5$ and $\\agt 5$ MeV for $x=0.3$. Our result suggests the existence of ``spongelike'' phases with negative Euler characteristic in addition to the simple ``pasta'' phases in supernova cores until $T \\alt 3$ MeV. ",
        "introduction": "In the process of the collapse-driven supernova \\cite{bethe}, matter in the core experiences adiabatic compression: the central density increases from $\\sim 10^{9}$ g cm$^{-3}$ at the beginning of the collapse to around the normal nuclear density $\\rho_{0}=0.165$ fm$^{-3}$ just before bounce; the temperature reaches $\\sim O(1)$ MeV at this point.   At subnuclear densities, nuclear matter exhibits the coexistence of a liquid phase with a gas phase due to the internucleon interaction which has an attractive part. In the density region where nuclei are about to melt into uniform matter, it is expected that, at sufficiently low temperatures relevant to neutron star interiors \\cite{review}, the energetically favorable configuration of the mixed phase possesses interesting spatial structures such as rodlike and slablike nuclei and rodlike and spherical bubbles, etc., which are referred to as nuclear ``pasta'' \\cite{rpw,hashimoto}.   This prediction is confirmed by several approaches assuming nuclear shapes such as the liquid drop models \\cite{lorenz,gentaro1,gentaro2} and the Thomas-Fermi calculations \\cite{oyamatsu}, and is also confirmed without assuming nuclear shapes in the framework of the Thomas-Fermi approximation \\cite{williams} and of the quantum molecular dynamics (QMD) \\cite{qmd1,qmd2}. While nuclear pasta at zero temperature is studied by several authors, pasta phases at finite temperatures relevant to supernova inner cores have not been studied yet except for a work by Lassaut {\\it et al.} using the Thomas-Fermi approximation \\cite{lassaut} and brief estimates of thermal fluctuations of the long-wave-length mode \\cite{gentaro1,gentaro2}. It is noted that, at temperatures of several MeV, effects of thermal fluctuations on nucleon distribution would be quite significant at subnuclear densities. However, the mean-field approximation such as the Thomas-Fermi \\cite{tf} and Hartree-Fock \\cite{bonche} approximation is not suitable to incorporate thermal fluctuations.   Finite temperature effects lead to evaporation of nucleons from nuclear liquid region and smoothed nucleon density profiles. At lower temperatures where each nuclei fluctuates a little around an average species, a compressible liquid-drop model \\cite{liquid drop} with incorporating the temperature dependence of its bulk, surface, and Coulomb+lattice components provides a useful way to investigate the pasta phases at finite temperatures. As for the bulk component \\cite{bulk}, binding energy, saturation density, and incompressibility \\cite{incompress}, which are parameters characterizing saturation properties, decrease with increasing temperature while the temperature dependence of the symmetry energy is not significant \\cite{sym}. For surface component, thermal broadening of the nucleon density profile reduces nuclear surface tension \\cite{surf}. Lattice energy is also modified by translational motion of nuclei \\cite{liquid drop}. However, at higher temperatures where the fluctuation of nuclear shape is significant, the above liquid-drop picture no longer holds; we have to call on some {\\it ab-initio} method which does not assume nuclear shape. More interestingly, it might be possible that, at these temperatures, the shape of the nuclear surface fluctuates and nuclei of various sizes and shapes coexist like colloid due to the entropy effect.   These finite temperature effects can be well described by the methods of molecular dynamics (MD) for nucleon many-body systems (see, e.g., Ref.\\ \\cite{feldmeier} for review). QMD \\cite{aichelin}, which is one of them, enables us to treat much larger systems than the other methods of MD do. Furthermore, at temperatures of several MeV, shell effects, which cannot be incorporated by QMD, are less important because they washed away by thermal fluctuations above $\\sim 3$ MeV. Thus QMD is an efficient and trustable method for studying nuclear matter at finite temperatures \\cite{maruyama,peilert}.   Pasta phases in supernova matter (SNM) are expected to affect the neutrino transport and hydrodynamics in supernova cores. Let us first note that the neutrino wavelengths, typically of order 20 fm, are comparable to or even greater than the internuclear spacing, leading to diffractive effects on the neutrino elastic scattering off such a periodic spatial structure of nuclear matter \\cite{rpw}. These effects, induced by the internuclear Coulombic correlations, would reduce the scattering rates and hence the lepton fraction $Y_{L}$. For the bcc lattice of spherical nuclei, such a reduction was examined by Horowitz \\cite{horowitz} by calculating the associated static structure factor. It is also noteworthy that nonspherical nuclei and bubbles are elongated in specific direction. In such direction, the neutrino scattering processes are no longer coherent, in contrast to the case of roughly spherical nuclei whose finiteness in any direction yields constructive interference in the scattering, which leads to the neutrino-trapping \\cite{freedman,sato}. The final point to be mentioned is that the changes in the nuclear shape are accompanied by discontinuities in the adiabatic index, denoting how hard the equation of state of the material is. These discontinuities may influence the core hydrodynamics during the initial phase of the collapse \\cite{lassaut}.   In the present paper, we study the structure of hot dense matter at subnuclear densities within the framework of QMD. Simulations of nuclear matter with proton fraction $x=0.3$ in addition to symmetric nuclear matter ($x=0.5$) are performed because the typical value of the proton fraction for supernova matter is around 0.3 due to the trapping of the electron neutrinos \\cite{freedman,sato}. We draw phase diagrams for $x=0.5$ and 0.3 in the density versus temperature plane, which show the qualitative feature of the finite temperature effects on the structure of nuclear matter. The results of the present study would be helpful to understand the real situation of the interior of the collapsing cores.    The plan of this paper is as follows. In Section \\ref{formulation}, we briefly explain the QMD model used in the present study and then discuss an effective temperature. In this section, a thermostatting method used in the simulations is also explained. In Section \\ref{results}, we show snapshots of some typical nucleon distributions to discuss qualitative features of finite temperature effect. After that, we analyze the structure of matter using two-point correlation functions and Minkowski functionals, and finally, resultant phase diagrams are shown. Summary and conclusion are presented in Section \\ref{conclusion}. ",
        "conclusions": "}   We have investigated the structure of hot nuclear matter at subnuclear densities by QMD simulations. The values of the proton fraction studied in the present work are 0.5 and 0.3; the latter is a typical value in supernova inner cores. The phase diagram obtained for $x=0.3$ (Fig.\\ \\ref{phase diagram x0.3}) tells that the pasta phases and the spongelike intermediate phases survive until $T \\alt 3$ MeV except for the higher density region of $\\rho \\agt 0.45\\rho_{0}$. This result strongly suggests that these phases with exotic nuclear structures can exist in the inner cores in the pre-bounce stage of the collapse-driven supernovae. Thus it would be meaningful to investigate a modification to the neutrino opacity due to the infinitely extending ``pasta'' and spongelike structures in the future \\cite{opacity}.   The present research has provided the general picture of the phase diagram at subnuclear densities in the temperature vs density plane. The qualitative features of the nucleon distribution at finite temperatures, except for higher densities just below which matter becomes uniform even at $T \\sim 0$ MeV, may be summarized as follows. \\begin{itemize} \\item At $T\\simeq$ 1--1.5 MeV, the number of evaporated nucleons starts to be significant. However, the macroscopic structure of matter does not basically change from that of zero-temperature matter except for a small expansion of the nuclear matter regions. \\item At $T\\simeq$ 1.5--2.5 MeV, fluctuations of the nuclear shape start to be significant and, in some cases, the topology of the nuclear structure changes from that in the zero-temperature matter for the same density. \\item At $T\\simeq$ 2.5 -- 3 MeV, single particle excitations are dominant rather than collective excitations, representing nuclear deformations. The density profile of nucleons is smoothed out, and nuclear surface cannot be identified by a threshold density. At higher temperatures, even the density contrast becomes drastically smaller, which leads to the transition to uniform matter. The order of the transition from inhomogeneous matter to uniform matter has yet to be clarified. \\item The critical temperature $T_{\\rm c}$ for the phase separation is $\\agt 6$ MeV for $x=0.5$ and $\\agt 5$ MeV for $x=0.3$. \\end{itemize}   Through the present work and the previous ones \\cite{qmd1,qmd2}, we have depicted the structure of nuclear matter in neutron star inner crusts and supernova inner cores in its equilibrium state. Our studies will be helpful in understanding the realistic situations of the interior of the dense stars. For further work, studies on the dynamical aspects of the nuclear pasta become important. For example, in connection with the possible instability of the long-range order of the planar phase, which cannot be fully incorporated in our simulations, elastic properties \\cite{pp} of the pasta and spongelike phases should be studied further. Structural transitions between the pasta phases induced by compression and decompression is one of the most interesting problems \\cite{formation} which have not yet been studied at all. We believe that our work here, which was performed within a dynamical framework, will open up these new aspects of the study on nuclear pasta."
    },
    "0311/hep-th0311059_arXiv.txt": {
        "abstract": "\\noindent Using the techniques of out-of-equilibrium field theory, we study the influence on the properties of cosmological perturbations generated during inflation on observable scales coming from fluctuations corresponding today to scales much bigger than the present Hubble radius. We write the effective action for the coarse-grained inflaton perturbations integrating out the sub-horizon modes, which manifest themselves as a colored noise and lead to memory effects. Using the simple model of a scalar field with cubic self-interactions evolving in a fixed de Sitter background, we evaluate the two- and three-point correlation function on observable scales. Our basic procedure shows that perturbations do preserve some memory of the super-horizon-scale dynamics, in the form of scale-dependent imprints in the statistical moments. In particular, we find a {\\it blue} tilt of the power-spectrum on large scales, in agreement with the recent results of the {\\it WMAP} collaboration which show a suppression of the lower multipoles in the Cosmic Microwave Background anisotropies, and a substantial enhancement of the intrinsic non-Gaussianity on large scales. ",
        "introduction": "Stochastic inflation provides an efficient approach to study inflationary dynamics and has become a very popular way to describe the growth of density perturbations on scales larger than the Hubble radius. In the first fundamental works \\cite{starobinski,Rey,goncharov,Nakao,Nambu,Nambu2, Linde}, the inflaton field was split into a super-horizon and a sub-horizon part directly in the equation of motion. This splitting is operated in Fourier space through a window function, that separates high from low frequencies. The relevant variable is the long-wavelength part, while the sub-horizon modes are collected in an effective noise term, playing the role of a classical perturbation to the super-horizon dynamics.  The resulting effective equation of motion is then a Langevin-like equation analogous to the one describing {\\it Brownian motion}, where the deterministic evolution is influenced and modified by the stochasticity of the source, whose effects can be taken into account only as a statistical average over time. Indeed, in this formalism there is not any knowledge of the exact form of the noise, but only of its statistical properties.  A more general approach exploits the influence functional method \\cite{morikawa,Hu2}, and operates the frequency splitting at the action level getting rid of the high frequencies via a path-integral over the sub-horizon part of the field. The effective action thus obtained contains some extra terms that can be interpreted as the coupling of the super-horizon field with a classical random noise source, whose configurations are statistically weighted by an appropriate functional probability distribution, becoming the origin of the stochastic character of the Langevin-like equation of motion.  The super-horizon degrees of freedom are then treated as a purely classical field, all the quantum fluctuations being collected in the classical noise term. This feature is claimed not to be a simple computational trick, but an intrinsic characteristic of the system. Stochasticity is thus not only a clue to understand the properties of inflation and the origin of the large-scale structure in the Universe, but also as a way to explain the transition from a quantum to a classical world \\cite{Polarski}. From a formal point of view, the quantum decoherence process in the stochastic inflation framework has been discussed in various works \\cite{Habib,decoherence,argentini1,Matacz,Kiefer}, where it was pointed out that the classicality of the coarse-grained field (implicitely assumed in the first papers) is not necessarily assured, but is subject to some restrictive conditions.  Using standard techniques of stochastic processes \\cite{fokkerplanck,chandra}, the Langevin equation for the field expectation value leads to an evolution equation (the Fokker-Planck equation) for its probability distribution function. In the first works, the noise correlation time is assumed to be infinitesimally short, and the correlation function for different times can therefore be considered as being proportional to Dirac's delta function $\\delta(t-t')$, which sets its {\\it white-noise} properties. This assumption allows to apply a well-known formalism for the derivation of the Fokker-Planck equation and its solution. However, the characteristic of the correlation function strongly depend on the window function, whose choice is not a mere mathematical tool, but has several physical effects \\cite{argentini}. A white noise arises only as a consequence of a sharp momentum cutoff, whereas a smooth weighting avoids highly singular noise correlators and produces a {\\it colored noise}.  The choice of a colored noise is interesting for at least two reasons. The first is that a sharp momentum-space cutoff seems rather unphysical, while a smooth weighting of the modes is much more likely. Actually, the most natural way to integrate out the small-scale fluctuations is to average the field in configuration space, choosing an appropriate finite volume window function. In most cases, this choice results in a smooth weighting in Fourier space (thus producing a colored noise), while the sharp momentum-space cutoff corresponds to a rather complicated infinite volume window function in configuration space. Moreover, it is possible to single out a wide class of window functions for which the shape of the colored noise correlation is asymptotically the same \\cite{Winitzki}. A second reason may be the fact that a colored noise could play, during inflation, an important role in producing intrinsically non-Gaussian density fluctuations as initial conditions for the subsequent evolution of the large-scale structure of the Universe \\cite{Hu}.  In the simplest single-field slow-roll models of inflation, non-Gaussian features in scalar perturbations are produced by either the self-interaction of the inflaton field \\cite{falk}, which are however constrained to be very small by the slow-roll conditions, or by the back-reaction of field fluctuations on the background metric, whose amplitude is also strongly constrained by the slow-roll conditions \\cite{glmm,gupta,acqua,malda}. It has been shown, however, that the most copious source of large-scale non-Gaussianity is stored in the post-inflationary second-order evolution of perturbations, which sets in a universal level of non-Gaussianity for the gauge-invariant gravitational potential, which turns out to be of order unity \\cite{enhance,allmodels}.  In this paper we point out that there is another source of  intrinsic, and generally scale-dependent non-Gaussianity in the fluctuation pattern, which originates from the cross-talk between super and sub-horizon scale perturbation modes. On scales much larger than the Hubble radius, non-Gaussian features generally arise as a consequence of the non-linear multiplicative form of the Langevin equation, when back-reaction effects are accounted for \\cite{salopek,mol,yv1,yv2,yv3}. However, as discussed in Refs. \\cite{salopek,yv1,yv2,yv3,hodges,condiniz,Kofman}, these effects do not directly reflect into the statistical properties of cosmological perturbations on sub-horizon scales. Indeed, in order to deal with fluctuations relative to our patch of the Universe, one cannot simply perform statistical averages over the entire ensemble of possible states, rather one should allow for the observed smoothness of our Universe on large scales. A possible, though approximate, way to take this constraint into account is to replace ensemble averages with averages over the conditional probability density functional that fluctuations on sub-horizon scales assume a certain value, {\\it given that} the inflaton field is spatially homogeneous at $t=\\ts$, {\\it i.e.} about 60 e-folds before the end of inflation (corresponding to a comoving scale slightly larger than the present Hubble radius). This is equivalent to set, for the probability distribution of the fluctuations, the `initial' condition $P(\\df,\\ts)=\\delta(\\df-\\df_{60})$ \\cite{salopek,yv1,yv2,yv3,hodges,condiniz,Kofman,glmm,gupta}. Although this may happen (and in most models it does) well after the beginning of the accelerated expansion, if the noise driving the fluctuations is white their evolution is Markovian, implying that any notion of the previous history is erased. The probability distribution then behaves exactly as if inflation had started at that time and the level of the inflaton non-Gaussianity remains fully negligible. On the contrary, a colored noise has a non-vanishing correlation time: because of this fact the inflaton keeps memory of what happened before the constraint, and its evolution ceases to be a Markov process. In this scenario, the probability distribution evolves in a different way, and also higher moments become important.  Since the solution of the Fokker-Planck equation with colored noise carries several complications and is still a partly unknown matter, in this paper we followed a different approach, trying to perturbatively determine the probability distribution for the inflaton field directly solving the Langevin equation in a statistical way.  The plan of the paper is as follows. In Section \\ref{derivazeqmoto} we briefly describe the derivation of the stochastic equation of motion for the inflaton field averaged over super-horizon scales, using the influence functional method, and evaluate the dependence of the noise on the choice of window function. In Section \\ref{spettri} we then choose a specific Gaussian shaped filter, obtaining the related (colored) noise correlation functions, and the variance and power-spectrum of the coarse-grained fluctuation field. In Section \\ref{interacting}, after introducing a small non-linear (cubic) term in the potential, we evaluate the bispectrum and the third moment of the field. In Section \\ref{distrprob} we investigate the memory effects induced by this colored noise and build up a formalism to quantitatively determine the relevance of the non-Gaussian features of the distribution and their sensitivity to the times before $\\ts$. Finally, in Section \\ref{conclusioni} we draw our conclusions. Some technical aspects of our calculation are reported in five Appendices. ",
        "conclusions": "\\label{conclusioni}  In this paper we used the stochastic inflation approach to address an important theoretical issue, namely that of estimating the possible influence of super-horizon perturbation modes on the statistics of cosmological perturbations on observable sub-horizon scales. To this aim we first modified the standard scheme in order to allow for the cross-talk of perturbations on super and sub-horizon scales. Such an effect is indeed totally absent in the traditional stochastic inflation dynamics; this is an artifact of the sharp k-space filter adopted in the coarse-graining procedure, which is easily removed by adopting a smoother filter function, leading to a colored - rather than white - noise source in the Langevin-like equation which governs the evolution of the coarse-grained inflaton field. This modification implies that the evolution of the coarse-grained inflaton field behaves as a non-Markovian stochastic process, in contrast to the standard case. Perturbations relevant to our smooth local patch of the Universe are then consistently defined by constraining the inflaton field to be homogeneous at a conventional time $t_{60}$ ({\\it i.e.} about 60 e-folds before inflation ends), corresponding to the horizon crossing of a scale slightly larger than the present Hubble radius. As a consequence of its non-Markovianity, the coarse-grained inflaton field preserves some memory of its dynamics prior to the constraint, which means that sub-horizon-scale perturbations have some knowledge of the state of the Universe on super-horizon scales.  Endowed with this {\\it extended} stochastic inflation scheme we are finally able to calculate the conditional second- and third-order moments of inflaton perturbations on observable scales. We perform our calculations in the case of a simple model where a scalar field with a small cubic self-interaction term evolves in a fixed de Sitter background. Our most important findings are: \\begin{itemize} \\item The variance of inflaton fluctuations grows quadratically with time around $t_{60}$, and is therefore smaller than in the standard Markovian case, which is linearly dependent on time; this is equivalent to say that the power-spectrum of density perturbations gets bluer on very large scales, without invoking any ad-hoc new physical input. This generic feature looks very intriguing, since the CMB anisotropy power on the largest angular scales observed by {\\it WMAP} \\cite{spergel} appears to be lower than the one predicted by the standard model of cosmology with almost scale-invariant primordial perturbations arising from a period of inflation.  \\item The skewness of inflaton fluctuations is larger than the standard case at times around $t_{60}$, which is equivalent to an enhancement of the non-Gaussianity level on large scales, for a given value of the self-coupling strength. The possible presence of large-scale non-Gaussianity is an important prediction to be confronted with current observational limits \\cite{komatsu}. Let us also mention that some recent analyses have reported evidence for a positive detection of non-Gaussianity in the {\\it WMAP} 1-year data \\cite{ngdetect1,ngdetect2,ngdetect3,ngdetect4,ngdetect5}. \\end{itemize}  In spite of the simplicity of the considered model, we can conclude that the cross-talk between super- and sub-horizon-scale perturbations is an important effect which would deserve an accurate treatment. The scheme presented here should be considered only as a first step in this direction, which needs to be largely improved and extended in various ways. First of all, one should consistently include metric perturbations in the inflaton dynamics, going beyond our fixed de Sitter background treatment. Second, our choice of a Gaussian filter can be somewhat arbitrary, and it would be important to understand how much the results depend on its choice. The conclusions reached in Ref. \\cite{Winitzki}, according to which there is a wide class of filters whose noise correlators show the same asymptotic behaviour $\\<\\xi(t)\\xi(t')\\>\\propto e^{-2H|t-t'|}$ found in this paper (see Eq. \\eqref{limnoisecorr}), are however quite encouraging in this respect. Actually, since the window function is only a technical device, we would expect our results to be independent of its choice, at least qualitatively. The white-noise exception is not significant, being merely a consequence of the ``bad'' choice of smoothing in configuration space.  The existence of memory effects means that the detailed dynamics of inflation plays a role in the specific form assumed by observable quantities. For instance, in our model one finds a residual dependence on the time when inflation started, $t_{in}$. Nonetheless, it is quite likely that this dependence saturates to a universal value determined by the general asymptotic behaviour, if the overall number of inflation e-folds is very large, as it is usually the case. More difficult is to avoid a dependence on the precise time $t_{60}$ at which we put the homogeneity constraint, and on the specific form of the constraint, which we introduced here through a delta function in  field configuration space (see, however, Ref. \\cite{lesarcs} for a discussion of alternative approaches to the constraint). These are important issues which will certainly deserve further investigation.     \\appendix   \\onecolumngrid  \\vspace{1cm}"
    },
    "0311/astro-ph0311002_arXiv.txt": {
        "abstract": "High energy cosmic ray experiments have identified an excess from the region of the Galactic Plane in a limited energy range around $10^{18}$~eV~(EeV). This is very suggestive of neutrons as candidate primaries, because the directional signal requires relatively-stable neutral primaries, and time-dilated neutrons can reach Earth from typical Galactic distances when the neutron energy exceeds an EeV. We here point out that if the Galactic messengers are neutrons, then those with energies below an EeV will decay in flight, providing a flux of cosmic antineutrinos above a TeV which is {\\it observable} at a kilometer-scale neutrino observatory. The expected event rate per year above 1 TeV in a detector such as IceCube, for example, is 20 antineutrino showers (all flavors) and a $1^\\circ$ directional signal of 4 $\\bar \\nu_\\mu$ events. A measurement of this flux can serve to identify the first extraterrestrial point source of TeV antineutrinos. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311528_arXiv.txt": {
        "abstract": "We report the results of a wide-field narrowband survey for redshift $z\\sim5.7$ Ly$\\alpha$ emitters carried out with the $34' \\times 27'$ field-of-view SuprimeCam mosaic CCD camera on the Subaru 8.3-m telescope. Deep narrowband imaging of the SSA22 field through a 120 \\AA\\ bandpass filter centered at a nominal wavelength of 8150 \\AA\\ was combined with deep multicolor $RIz'$ broadband imaging with SuprimeCam, then supplemented with $BVRZ$ imaging taken with the $42' \\times 28'$ field-of-view CFH12K camera on the Canada-France-Hawaii 3.6-m telescope to select high-redshift galaxy candidates.  Spectroscopic observations were made using the new wide-field multi-object DEIMOS spectrograph on the 10-m Keck\\,II Telescope for 22 of the 26 candidate objects. Eighteen of these objects were identified as $z\\sim5.7$ Lyman alpha emitters, and a further nineteenth object from the candidate list was identified based on an LRIS spectrum.  At the 3.3 \\AA\\ resolution of the DEIMOS observations the asymmetric profile for Ly$\\alpha$ emission with its steep blue fall-off can be clearly seen in the spectra of the identified galaxies. This is by far the largest spectroscopic sample of galaxies at these redshifts, and we use it to describe the distribution of equivalent widths and continuum color break properties for identified Ly$\\alpha$ galaxies compared with the foreground population. The large majority (at least 75\\%) of the lines have rest frame Ly$\\alpha$ equivalent widths substantially less than 240~\\AA\\ and can be understood in terms of young star forming galaxies with a Salpeter initial mass function for the stars.  With the narrowband selection criteria of ($I-N) > 0.7$ and $N<25.05$ (AB mags) we find a surface density of Ly$\\alpha$ emitters of 0.03 per square arcminute in the filter bandpass ($\\Delta z=0.1$) down to a limiting flux of just under $2\\times10^{-17}$ erg cm$^{-2}$ s$^{-1}$. The luminosity function of the Ly$\\alpha$ emitters is similar to that at lower redshifts to the lowest measurable luminosity of $10^{43}$ erg s$^{-1}$ as is the universal star formation rate based on their continuum properties. However, we note that the objects are highly structured in both their spatial and spectral properties on the angular scale of the fields ($\\approx 60$ Mpc), and that multiple fields will have to be averaged to accurately measure their ensemble properties. ",
        "introduction": "\\label{secintro}  The study of high-redshift galaxies in the early universe has seen substantial progress in recent times, with the discovery of objects out to $z=6.5$ (Hu et al. 2002, Kodaira et al. 2003) and with the rapidly increasing number of galaxies discovered at redshifts $z\\gg5$ (\\citealp*{dey98}; Hu, Cowie, \\& McMahon 1998; \\citealp*{wey98}; Hu, McMahon, \\& Cowie 1999; \\citealp*{stern99,ellis,aji02,tan03,rho03,mai03,cuby03,bunk03}). Active galaxies associated with radio galaxies \\citep{wil99} and AGN discovered in deep X-ray exposures with CHANDRA \\citep{cdfn} have been seen out to $z\\sim5.19$, and at optical wavelengths quasars have been found with the Sloan survey out to  $z=6.42$ \\citep{fan01,fan03}. However, because of the relatively sparse distribution of detectable $z > 5$ galaxies it has hitherto been difficult to confirm more than a few at a time, and thus to develop the large samples needed to understand the robustness of candidate selection criteria and the spatial structure of objects on the sky.  Consequently there are large uncertainties in the inferred luminosity function and universal star formation rates.  The availability of mosaic CCDs allows wide-field coverage of the sky on scales of half a degree to a degree at a time \\citep[e.g.,][]{mal02,ouchi03,aji03}. This gives us the opportunity to generate large samples of objects to examine their statistical properties and construct more precise luminosity functions.  This paper is the first in a series which will describe searches for Ly$\\alpha$ emission-line galaxies at $z=5.7$ and $z=6.5$ in a number of such wide fields. It represents a continuation of our narrowband filter survey \\citep{cow98,smitty,la_lum,la6}, which probed for Ly$\\alpha$ emission from $z=3.4\\to6.5$ using smaller area surveys.  We use the wide-field capabilities of SuprimeCam \\citep{supcam} on the Subaru 8.3-m Telescope to obtain samples of narrowband selected Ly$\\alpha$ emitters. Following \\citet{smitty} we use narrowband filters corresponding to $z\\sim 5.7$ and $z\\sim 6.5$ Ly$\\alpha$ emission lines which lie in dark spectral portions of the night sky. We combine these with extremely deep continuum observations to select the  high-$z$ emission-line galaxies in a statistically uniform fashion. We then use the new wide-field DEIMOS spectrograph on the Keck II 10-m telescope to obtain a nearly complete spectroscopic followup.  In the present paper we describe our survey of $z\\sim 5.7$ Lyman alpha emitters in a 700 square arcminute area around the SSA22 field \\citep{cowie_1}.  The existence of extremely deep multi-color $BVRIz'$ continuum images for this field allows a detailed examination of the effectiveness of color selection criteria in identifying Ly$\\alpha$ candidates at $z\\sim5.7$, similar to our studies at $z\\sim 3.4$ \\citet{cow98}.  In particular, we can examine the robustness of color thresholds for the selection of Ly$\\alpha$ emitters as contrasted with \\oiii\\ and \\oii\\ emission-line galaxies.  Discriminating Ly$\\alpha$ emitters from lower $z$ emission-line objects can be difficult at lower redshifts \\citep{stern99}, but at higher redshifts the continuum break signature becomes so extreme that there is little likelihood of a misidentification provided we have sufficiently deep continuum images. Songaila and Cowie \\citeyearpar{gpnot} have used spectra of the bright quasars to measure the average transmission of the Ly$\\alpha$ forest region as a function of redshift over the $z=4\\to6$ range.  Their results translate to a magnitude break across the Ly$\\alpha$ emission line of \\begin{displaymath} \\Delta{m} = 3.8 + 20.3\\ {\\rm log_{10}} \\left ( {1+z}\\over{7} \\right ) \\label{eq:1} \\end{displaymath} For $z=5.7$ systems this corresponds to a 3.4 magnitude break, and color measurements alone can be used to provide a substantial selection of high-$z$ galaxy candidates \\citep[e.g.,][]{iwa03,stan03a,bunk03,leh03,bouw03,dic03}. However, because the continuum magnitudes are extremely faint only the deepest observations currently feasible can find such objects, while the emission-line objects are much easier to find. Furthermore, continuum selected objects can be extremely hard to confirm spectroscopically, while the narrowband selected objects are generally easy to identify. Thus the narrowband filter samples provide an invaluable tool for identifying large homogeneous samples of high-redshift galaxies.  With high-resolution spectroscopic followups of the candidates it is possible to both confirm that the emission is due to red-shifted Ly$\\alpha$ and to estimate the statistics of interlopers. For SSA22 our narrowband and color selections yield an initial sample of 26 candidate $z\\sim 5.7$ emitters in the field.  The powerful new wide-field spectroscopic capabilities of the DEIMOS spectrograph \\citep{deimos03} used on the Keck 10-m telescope, with an approximate field of $16\\farcm7 \\times 5\\farcm0$ and good red sensitivity, allow us to follow up these candidates in an extremely efficient manner. We have now observed 23 of the 26 objects with either this instrument or with the LRIS spectrograph \\citep{lris} on Keck.  This results in the confirmation of 19 new redshift $z\\sim5.7$ galaxies. The 18 candidates observed with DEIMOS  unambiguously show asymmetric profiles with the high resolution observing mode possible with this red-optimized spectrograph.  In Section~\\ref{secdata} we summarize the narrowband and continuum imaging together with the spectroscopic observations. Section~\\ref{secanal} describes the color selection and color-color plots for our selection criteria.  The spectroscopic identifications and colors of candidates are tabulated and labeled in these diagrams. In Section~\\ref{secclass} we discuss the spectroscopic identifications of Ly$\\alpha$ emitters, and the incidence and properties of foreground emitters. In Section~\\ref{secdisc} we discuss the spatial structure of the objects in the fields and show that they are highly correlated. We compare the shapes and equivalent widths of the lines with those of lower redshift samples and discuss their interpretation.  Finally, the Ly$\\alpha$ and UV continuum luminosity functions of the identified $z\\sim5.7$ candidates are constructed and compared with lower redshift samples. We argue that the star formation rates at $z\\sim5.7$ are probably at least comparable to those at lower redshifts, but postpone a more detailed discussion until we can use an ensemble average of a number of wide fields. ",
        "conclusions": "\\label{secconc}  We may summarize the major conclusions of the paper as follows. \\smallskip  1) Combined emission line and color break techniques provide an extremely efficient method of selecting objects at $z=5.7$. Selection criteria of $(N-I) > 0.7$, $(R-z') > 1.8$ and $(N-z') < 0.2$ pick out 24 objects in the field. At least 19 of these are spectroscopically confirmed as Lyman alpha emitters at this redshift. Only one or at most two of the sample selected in this way are lower redshift interlopers. The accuracy is sufficiently good to allow use of color selected samples without the necessity of spectroscopic followup for many purposes.  2) Even on scales as large as the present 26$'$ field the distribution of objects is highly structured in both the spectral and spatial dimensions. Individual fields could easily show factors of several variation in the number density of objects selected in a 100~\\AA\\ bandpass and a substantial number of fields need to be observed to average out these effects.  3) The properties of the Lyman alpha selected galaxies are quite similar to those observed at $z \\approx 3$.  4) Star formation rates at $z=5.7$ are roughly similar to those at $z=3$ and $z=4$ though there are very substantial uncertainties in this estimate."
    },
    "0311/astro-ph0311234_arXiv.txt": {
        "abstract": "A principal candidate for quiescent non-thermal gamma-ray emission from magnetars is resonant inverse Compton scattering in the strong fields of their magnetospheres.  This paper outlines expectations for such emission, formed from non-thermal electrons accelerated in a pulsar-like polar cap potential upscattering thermal X-rays from the hot stellar surface.  The resultant spectra are found to be strikingly flat, with fluxes and strong pulsation that could be detectable by GLAST. ",
        "introduction": "The quiescent, non-thermal X-ray emission of soft gamma repeaters (SGRs) and anomalous X-ray pulsars (AXPs) is observed at levels that are very intense for neutron stars.   It is quite possible that SGRs and AXPs emit in a broad-band, pulsar-like mode, even if the power is not of a rotational origin.  A principal candidate for such emission is inverse Compton scattering, tapping abundant surface X-rays (inferred surface temperatures for AXPs and SGRs are somewhat greater than for canonical pulsars; see Perna et al. 2001) as targets for energetic electrons accelerated in either polar cap or outer gap potentials.  Recently, Cheng \\& Zhang (2001) proposed a model for significant gamma-ray emission from such magnetars, in the context of  the outer gap scenario.  Here, the possibility of magnetar gamma-ray emission from polar cap models is explored.  In the strong fields of their magnetospheres, the scattering is enhanced by the cyclotron resonance, thereby rendering polar cap models even more promising for generating hard gamma-rays that are potentially detectable by the upcoming Gamma-Ray Large Area Space Telescope (GLAST) mission.   The present scenario is similar to previous pulsar (e.g. Sturner, Dermer \\& Michel 1995; Daugherty \\& Harding 1996) and old gamma-ray burst (Dermer 1990; Baring 1994) models. Primary electrons with Lorentz factors up to some maximum \\teq{\\gammax} cool in collisions with the X-ray photons, in the process generating a broad spectrum extending into the gamma-ray band.  This paper outlines expectations for such emission spectra, specifically incorporating relativistic cross sections at supercritical fields that are germane to resonant Compton scattering in the magnetar regime. This investigation will focus on a simple one-zone, uniformly-magnetized region as a preparation for real geometry down the line, where the complexity of pair cascading will be explored.  It is found that a very flat soft gamma-ray continuum results, with spectral cutoffs that depend strongly on the observational perspective to the field.  Spectra differ significantly from those produced in outer gap models.  Estimates of the anticipated flux indicate that GLAST may well detect magnetars in the 100 MeV -- 1 GeV band if they emit efficiently like gamma-ray pulsars. ",
        "conclusions": ""
    },
    "0311/astro-ph0311144_arXiv.txt": {
        "abstract": "We briefly summarize the survey of \\schi 21 cm emission lines to search for shocked atomic gas associated with Galactic supernova remnants (SNRs) in the southern sky. For G347.3$-$0.5, we discuss the distance to the SNR and the implications of the \\schi results. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311372_arXiv.txt": {
        "abstract": "Recent advances in understanding of the basic properties of compressible Magnetohydrodynamic (MHD) turbulence call for revisions of some of the generally accepted concepts. First,  MHD turbulence is not so messy as it is usually believed. In fact, the notion of strong non-linear coupling of compressible and incompressible motions is not tenable. Alfven, slow and fast modes of MHD turbulence follow their own cascades and exhibit degrees of anisotropy consistent with theoretical expectations. Second, the fast decay of turbulence is not related to the compressibility of fluid. Rates of decay of compressible and incompressible motions are very similar. Third, viscosity by neutrals does not suppress MHD turbulence in a partially ionized gas. Instead, MHD turbulence develops magnetic cascade at scales below the scale at which neutrals damp ordinary hydrodynamic motions. The implications of those changes of MHD turbulence paradigm for molecular clouds require further studies. Those studies can benefit from testing of theoretical predictions using new statistical techniques that utilize spectroscopic data. We briefly discuss advances in development of tools using which the statistics of turbulent velocity can be recovered from observations. ",
        "introduction": "It is well known that molecular clouds are magnetized with compressible magnetic turbulence determining most of their properties (see Elmegreen \\& Falgarone 1996, Stutzki 2001). Star formation (see McKee \\& Tan 2002, Elmegreen 2002, Pudritz 2001), cloud chemistry (see Falgarone 1999 and references therein), shattering and coagulation of dust (see Lazarian \\& Yan 2003 and references therein) are examples of processes for which knowledge of turbulence is absolutely essential. There are many excellent reviews that deal the theory of turbulent molecular clouds and numerical simulations (see Falgarone 1999, Vazquez-Semadeni et al. 2000, Mac Low \\& Klessen 2003). This {\\it short} review is focused on the recently uncovered basic properties of MHD turbulence and their implications for understanding of molecular clouds\\footnote{ It is not possible to cite all the important papers in the area of MHD turbulence and turbulent molecular clouds. An incomplete list of the references in a recent review on the statistics of MHD turbulence by Cho, Lazarian \\& Vishniac (2003a; henceforth CLV03a) includes about two hundred entries. The list of references in the molecular cloud review by Mac Low \\& Klessen (2003) is even longer. }. We also briefly deal with recovering of the 3D statistics of turbulent velocity from observations, which is a theoretical problem in itself. A more detailed discussion of MHD turbulence can be found in the companion review by Chandran, while observational aspects of turbulence in molecular clouds are discuss in the review by Falgarone (this volume).   Why do we expect molecular clouds to be turbulent? A fluid of viscosity $\\nu$ becomes turbulent when the rate of viscous dissipation, which is  $\\sim \\nu/L^2$ at the energy injection scale $L$, is much smaller than the energy transfer rate $\\sim V_L/L$, where $V_L$ is the velocity dispersion at the scale $L$. The ratio of the two rates is the Reynolds number $Re=V_LL/\\nu$. In general, when $Re$ is larger than $10-100$ the system becomes turbulent. Chaotic structures develop gradually as $Re$ increases, and those with $Re\\sim10^3$ are appreciably less chaotic than those with $Re\\sim10^8$. Observed features such as star forming clouds are very chaotic for $Re>10^8$. This makes it difficult to simulate realistic turbulence. The currently available 3D simulations containing 512 grid cells along each side can have $Re$ up to $\\sim O(10^3)$ and are limited by their grid sizes. Therefore, it is essential to find ``{\\it scaling laws}\" in order to extrapolate numerical calculations ($Re \\sim O(10^3)$) to real astrophysical fluids ($Re>10^8$). We show below that even with its limited resolution, numerics is a great tool for {\\it testing} scaling laws.  Kolmogorov theory provides a scaling law for {\\it incompressible} {\\it non}-magnetized hydrodynamic turbulence (Kolmogorov 1941). This law provides a statistical relation between the relative velocity $v_l$ of fluid elements and their separation $l$, namely, $v_l\\sim l^{1/3}$.  An equivalent description is to express spectrum $E(k)$ as a function of wave number $k$ ($\\sim 1/l$). The two descriptions are related by $kE(k) \\sim v_l^2$. The famous Kolmogorov spectrum is  $E(k)\\sim k^{-5/3}$. The applications of Kolmogorov theory range from engineering research to meteorology (see Monin \\& Yaglom 1975) but its astrophysical applications are poorly justified and the application of the Kolmogorov theory can lead to erroneous conclusions (see reviews by Lazarian et al. 2003 and Lazarian \\& Yan 2003)  Let us consider {\\it incompressible} MHD turbulence first. There have long been understanding that the MHD turbulence is anisotropic (e.g. Shebalin et al.~1983). Substantial progress has been achieved recently by Goldreich \\& Sridhar (1995; hereafter GS95), who made an ingenious prediction regarding relative motions parallel and perpendicular to magnetic field {\\bf B} for incompressible MHD turbulence. An important observation that leads to understanding of the GS95 scaling is that magnetic field cannot prevent mixing motions of magnetic field lines if the motions are perpendicular to the magnetic field. Those motions will cause, however, waves that will propagate along magnetic field lines. If that is the case, the time scale of the wave-like motions along the field, i.e. $\\sim l_{\\|}/V_A$, ($l_{\\|}$ is the characteristic size of the perturbation along the magnetic field and $V_A=B/\\sqrt{4 \\pi \\rho}$ is the local Alfven speed) will be equal to the hydrodynamic time-scale, $l_{\\perp}/v_l$, where $l_{\\perp}$ is the characteristic size of the perturbation perpendicular to the magnetic field. The mixing motions are hydrodynamic-like\\footnote{ Recent simulations (Cho et al.~2003c) suggest that perpendicular mixing is indeed efficient for mean magnetic fields of up to the equipartition value.}. They obey Kolmogorov scaling, $v_l\\propto l_{\\perp}^{1/3}$, because incompressible turbulence is assumed. Combining the two relations above we can get the GS95 anisotropy, $l_{\\|}\\propto l_{\\perp}^{2/3}$ (or $k_{\\|}\\propto k_{\\perp}^{2/3}$ in terms of wave-numbers). If  we interpret $l_{\\|}$ as the eddy size in the direction of the local magnetic field\\footnote{Incidentally, if we identify the size of molecular cloud with the longer dimension of an eddy, we shall get the Larson relations $v_l\\sim l^{1/2}$ (Larson 1981) for velocity and cloud size.} and $l_{\\perp}$ as that in the perpendicular directions, the relation implies that smaller eddies are more elongated.  GS95 predictions have been confirmed numerically (Cho \\& Vishniac 2000; Maron \\& Goldreich 2001; Cho, Lazarian \\& Vishniac 2002a, hereafter CLV02a; see also CLV03a); they are in good agreement with observed and inferred astrophysical spectra (see CLV03a). However, the GS95 model considered incompressible MHD, but the media in molecular clouds is {\\it highly compressible}. Does any part of GS95 model survives? Literature on the properties of compressible MHD is very rich (see reviews by Pouquet 1999, Cho \\& Lazarian 2003 and references therein). Higdon (1984) theoretically studied density fluctuations in the interstellar MHD turbulence. Matthaeus \\& Brown (1988) studied nearly incompressible MHD at low Mach number and Zank \\& Matthaeus (1993) extended it. In an important paper Matthaeus et al.~(1996) numerically explored anisotropy of compressible MHD turbulence. However, those papers do not provide universal scalings of the GS95 type.  The complexity of the compressible magnetized turbulence with magnetic field made some researchers believe that the phenomenon is too complex to expect any universal scalings for molecular cloud research. High coupling of compressible and incompressible motions is often quoted to justify this point of view.  Below we shall provide arguments that are suggestive that the fundamentals of compressible MHD can be understood and successfully applied to molecular clouds.  In molecular clouds the regular magnetic field is comparable with the fluctuating one. Therefore for most part of our discussion, we shall discuss results obtained for $\\delta V \\sim \\delta B/\\sqrt{4 \\pi \\rho} \\sim B_0/\\sqrt{4 \\pi \\rho}$, where $\\delta B$ is the r.m.s. strength of the random magnetic field. ",
        "conclusions": "~~~~1. Understanding of molecular clouds requires understanding of the basics of MHD turbulence. MHD turbulence is not a mess. Scaling relations for its modes have been established recently.  2. Fast decay of MHD turbulence is not due to strong coupling of compressible and incompressible motions. The transfer of energy from Alfven to compressible modes is small. The Alfven mode develops on its own and decays fast.  3. Doppler shifts imprinted in spectra lines provide an excellent way of testing theoretical expectations. Advances in understanding how this information can be extracted from spectroscopic data allow this.\\\\ {\\bf Acknowledgments}{\\it A.L. acknowledges  the support of NSF Grant AST-0125544}."
    },
    "0311/astro-ph0311622_arXiv.txt": {
        "abstract": "{ We extend our previous analysis of cosmological supernova Type Ia data \\cite{pc03} to include three recent compilation of data sets. Our analysis ignores the possible correlations and systematic effects present in the data and concentrates mostly on some key theoretical issues. Among the three data sets, the first set consists of 194 points obtained from various observations while the second discards some of the points from the first one because of large uncertainties and thus consists of 142 points. The third data set is obtained from the second by adding the latest 14 points observed through HST. A careful comparison of these different data sets help us to draw",
        "introduction": "Current cosmological observations, particularly those of supernova Type Ia, show a strong signature of the existence of a dark energy component with negative pressure \\cite{rfc++98,pag++99,riess00}.  The most obvious candidate for this dark energy is the cosmological constant (with the equation of state $w_X = p/\\rho = -1$), which, however, raises several theoretical difficulties [for reviews, see \\citeN{ss00}, \\citeN{pr03}, \\citeN{padmanabhan03}]. This has led to models for  dark energy component which evolves with time \\cite{rp88,wetterich88,fj98,fhsw95,bm99,bm00b,um00,bcn00,zws99,as00b,btv02}.  Currently, there is a tremendous amount of activity going on in trying to determine the equation of state $w_X(z)$ and other cosmological parameters from observations of high redshift Type Ia supernova \\cite{gjc++98,astier00,srss00,wg01,wa01,gaagp01,wl01,leibundgut01,trentham01,pnr01,cc02,klsw02,mhh02,mbms02,wa02,ge02,rowan-robinson02,lj03,pc03,visser04,cmm04,alcaniz04,wm04,nbs03,kaa++03,zf03,dja04,gcd04,gc04,gd04,gong04,bertolami04,wt04,ccrl04,mcinnes04,sc04,bsss04,la04,np04,ass04,zfh04,ap04}. While there has been a considerable activity in this field, one should keep in mind that there are several theoretical degeneracies in the Friedmann model, which can limit the determination of $w_X(z)$. To understand this, note that the only non-trivial metric function in a Friedmann universe is the Hubble parameter $H(z)$ (besides the curvature of the spatial part of the metric), which is related to the \\emph{total} energy density in the universe. \\emph{Hence, it is not  possible to determine the energy densities of individual components of energy densities in the universe from any geometrical observation.} For example, if we assume a flat universe, and further assume that the only energy densities present are those corresponding to the non-relativistic dust-like matter and dark energy, then we need to know $\\Omega_m$ of the dust-like matter and $H(z)$ to a very high accuracy in order to get a handle on $\\Omega_X$ or $w_X$ of the dark energy. This can be a fairly strong degeneracy for determining $w_X(z)$ from observations.  Recently, we discussed certain questions related to the determination of the nature of dark energy component from observations of high redshift supernova in \\citeN{pc03} [hereafter Paper~I]. In the above work, we reanalyzed the supernova data using very simple statistical tools in order to focus attention on some key issues. The analysis of the data were intentionally kept simple as we subscribe to the point of view that any result which cannot be revealed by a simple analysis of data, but arises through a more complex statistical procedure, is inherently suspect and a conclusion as important as the existence of dark energy with negative pressure should pass such a test. The key results of our previous analysis were:  $\\bullet$ Even if the precise value of $\\Omega_X$ or the equation of state $w_X(z)$ is known from observations, it is {\\em not} possible to determine the nature (or, say, the Lagrangian) of the unknown dark energy source using only kinematical and geometrical measurements. For example, if one assumes that the dark energy arises from a scalar field, then it is possible to come up with scalar field Lagrangians of different forms leading to same $w_X(z)$. As an explicit example, we considered two Lagrangians, one corresponding to quintessence \\cite{pr88,rp88,zws99} and the other corresponding to the tachyonic scalar fields \\cite{padmanabhan02,pc02,fks02,sw02,gibbons02,ft02,mukohyama02,feinstein02,bjp03}. These two fields are quite different in terms of their intrinsic properties; however, it is possible to make both the Lagrangians produce a given $w_X(a)$ by choosing the potential functions in the corresponding Lagrangians [for explicit examples and forms of potential functions, see \\citeN{padmanabhan02}; Paper~I].   $\\bullet$ Although the full data set of supernova observations strongly rule out models without dark energy, the high and low  redshift data sets, individually, admit non-accelerating models with zero dark energy. It is not surprising that the high redshift data is consistent with non-accelerating models as the universe is in its decelerating phase at those redshifts. On the other hand, though the acceleration of the universe is a low redshift phenomenon, the non-accelerating models could not be ruled out using low redshift data {\\it alone} because of large errors. Given the small data set, any possible evolution in the absolute magnitude of the supernovae, if detected, might have allowed the data to be consistent with the non-accelerating models.  $\\bullet$ We introduced two parameters, which can be obtained entirely from theory, to study the sensitivity of the luminosity distance on $w_X$. Using these two parameters, we argued that although one can  determine the present value of $w_X$ accurately from the data, one cannot constrain the evolution of $w_X$. The situation is worse if we add the uncertainties in determining $\\Omega_m$.  All the above conclusions were obtained by analysing only 55 supernova data points from a very simple point of view. In recent times, data points from various sets of observations have been compiled taking into account the calibration errors and other uncertainties. This enables us to repeat our analysis for much larger data sets, and see how robust are the conclusions of Paper~I with respect to the choice of the data points. In this paper, we will compare three such data sets, which differ in their selection criteria for data points and redshift range covered.  The structure of the paper is as follows: In the next section, we describe the three data sets used in this paper, and then analyse them for models with non-relativistic dust-like matter and cosmological constant. Some key points regarding the importance of low and high redshift data are discussed. In Section 3, we briefly discuss the constraints on the dark energy equation of state and its evolution. The results are summarized in Section 4. Finally, the effect of our extinction-based selection criterion on the determination of cosmological parameters is discussed in the Appendix. ",
        "conclusions": "We have reanalyzed the supernova data with the currently available data points and constrained various parameters related to general cosmological models and dark energy. We would like to mention that our analysis ignores the effects of correlation and other systematics present in the data. The main aim of the work has been to focus on some important theoretical issues which are not adequately stressed in the literature. We have used three compiled and available data sets, which are called TONRY (194 points), RIESS(w/o HST) (142 points) and RIESS (156 points). The RIESS(w/o HST) is obtained from the TONRY data set by discarding points with large uncertainties and by reducing calibration errors, while the RIESS data set is obtained by adding the recent points from HST to the RIESS(w/o HST) set. The analysis is an extension to what was performed in Paper~I with a small subset of data points. In particular, we have critically compared the estimated values of cosmological parameters from the three data sets. While the errors on the parameter estimation have come down significantly with all the data sets, we find that there some crucial differences between the data sets. We summarize the key results once more:  $\\bullet$ It has been well known that the supernova data rule out the flat and open matter-dominated models with a high degree of confidence \\cite{rfc++98,pag++99,riess00}. However, for the TONRY and RIESS(w/o HST) data sets, we find that the data favours a model with $\\Omega_{\\rm tot} > 1$ (with probability $\\gtrsim 0.97$) and is in mild disagreement with the ``concordance'' flat models with cosmological constant. This disagreement seem to be less (the probability of obtaining models with $\\Omega_{\\rm tot} > 1$ being $\\approx 0.9$) for the RIESS data set, which includes the new HST points in the redshift range $1 < z < 1.6$,  $\\bullet$ The supernova data on the whole rules out non-accelerating models with very high confidence level. However, it is interesting to note that if we divide the TONRY data set into high and low redshift subsets, neither of the subsets are able to rule out the non-accelerating models. In particular, the low redshift data points are consistent with the non-accelerating models because of large errors on the data. This keeps open the possibility that the evolutionary effects in the absolute luminosities of supernovae might make the entire data set consistent with SCDM model. The situation is quite different for the RIESS(w/o HST) and RIESS data sets, where points with large errors are discarded. The low redshift data alone seem to rule out the SCDM model with high degree of confidence. This means that unless the absolute luminosities of supernovae evolve rapidly with redshift, it might be difficult for the data set to be consistent with the SCDM model. In other words, the RIESS(w/o HST) and RIESS data sets establish the presence of the accelerating phase of the universe regardless of the evolutionary effects.  \\begin{figure*} \\begin{center} \\rotatebox{270}{\\resizebox{0.3\\textwidth}{!}{\\includegraphics{fit_perlmutter_av.ps}}} \\caption{Confidence region ellipses in the $\\Omega_m - \\Omega_{\\Lambda}$ plane for models with non-relativistic matter and a cosmological constant for different selection criteria based on extinction for the TONRY data set. The ellipses corresponding to the 68, 90 and 99 per cent confidence regions are shown. The confidence regions are obtained after marginalizing over ${\\cal M}_1$. The dashed line corresponds to the flat model $(\\Omega_{~m} + \\Omega_{\\Lambda} = 1)$. The left panel shows results when only those points with $A_V < 0.5$ are included, the middle panel considers only points which have $A_V < 0.3$, while the right panel includes all the points irrespective of the value of $A_V$. The values of the best-fit parameters, with 1$\\sigma$ errors are indicated in the respective panels. } \\label{fitav} \\end{center} \\end{figure*}  $\\bullet$ The key issue regarding dark energy is to determine the evolution of its equation of state, $w_X$. We find that although one can constrain the current value of $w_X$ quite well, it is comparatively difficult to determine the evolution of $w_X$. The situation is further worsened when we take the uncertainties in $\\Omega_m$ into account.  $\\bullet$ The supernova data {\\it mildly} favours a dark energy equation of state with its present best-fit value less than -1 which will require more exotic forms of matter (possibly with negative kinetic energy). However, one should keep in mind that the data is still consistent with the standard cosmological constant at 99 per cent confidence level.  $\\bullet$ The analysis of different subsamples of the supernova data set is important in determining the effect of evolution. In this work, we have taken the simple approach of dividing the data roughly around the epoch where the universe might have transited from a decelerating to an accelerating phase, and checked whether the data can be made consistent with the non-accelerating models. In future, it would be interesting to divide the data based on the nature of supernova searches. For example, one can divide the data into three redshift splits: $z < 0.1$, $0.2 < z < 0.8$ and  $z > 0.8$, which roughly correspond to supernovae discovered in shallow searches, ground-based deep searches, and space-based deep searches. It would be interesting to check the cosmological constraints with such a divide."
    },
    "0311/astro-ph0311078_arXiv.txt": {
        "abstract": "{Kinematic and population studies show that bulges are generally rotationally flattened systems similar to low-luminosity ellipticals. However, observations with state-of-the-art integral field spectrographs, such as \\sauron, indicate that the situation is much more complex, and allow us to investigate phenomena such as triaxiality, kinematic decoupling and population substructure,  and to study their connection to current formation and evolution scenarios for bulges of early-type galaxies. We present the examples of two S0 bulges from galaxies in our sample of nearby galaxies: one that shows all the properties expected from classical bulges (NGC\\,5866), and another case that presents kinematic features appropriate for barred disk galaxies (NGC\\,7332). ",
        "introduction": "Lying at the centre of galaxies, bulges are a keystone in our understanding of galaxy formation and evolution. When studying their formation, one is often drawn to examine the order of events during galaxy assembly. In a simplistic approach, either the bulge formed before the disk, or the bulge formed from disk material.  Classical bulges are traditionally understood as spheroids with an $\\mathrm{r}^{1/4}$ surface brightness profile, in most ways identical to elliptical galaxies of the same luminosity. For example, bulges behave like ellipticals in the widely known scaling relations  (Faber \\& Jackson 1976; Terlevich et al. 1981;  Djorgovski \\& Davis 1987; Dressler et al. 1987). Kinematically,  bulges, like low-luminosity ellipticals, are found to be isotropic oblate rotators  (Kormendy \\& Illingworth 1982). Such analogies  lie at the heart of merger-driven bulge formation  models, in which the bulge formed before the disk.  Dynamicists, however, have shown in their numerical simulations that disk instabilities can pump disk material above the plane,  thus generating central structures that also {\\it bulge} over the  thin disk (e.g. Hasan, Pfenniger \\& Norman 1993). These simulations raised the  possibility of an alternative formation scenario in which bulges are formed  via secular evolution processes after the disk  (e.g. Pfenniger 1993).  The presence of bars inside peanut-shaped bulges has been confirmed photometrically (L{\\\" u}tticke, Dettmar \\& Polhen 2000)  and spectroscopically (Kuijken \\& Merrifield 1995; Bureau \\& Freeman 1999; Merrifield \\& Kuijken 1999). Furthermore, some bulges  share strong similarities with disks, on the basis of their low velocity dispersions. The disky nature of bulges is usually discussed in relation to late-type, low-mass bulges (Kormendy 1993, Carollo 1999).  Do bulges come in two flavours, early-types forming in mergers and late-types through disk instabilities? To answer this question, it is vital to increase our knowledge on the structural and dynamical properties that link galaxy bulges to either slowly rotating spheroids akin to elliptical galaxies, or to rapidly-rotating flattened systems more nearly resembling the products of internal  disk transformations.  Advances and discoveries in astrophysics often come together with instrument developments. Photometrically, the detailed pictures HST provided during the last decade, has revealed a large variety  of features (e.g. inner disks or bars) at the centre of galaxies  (Lauer et al. 1995; Carollo 1999; Balcells et al. 2003). Spectroscopically, however, it is the ground-based integral field units (IFUs) that can map the full 2D stellar and gas kinematic signatures of those structures. This allows us to impose tighter constraints on the processes involved in the formation and evolution of galaxies.  The \\sauron\\ survey is a scientific project aiming to understand the formation  and evolution of elliptical and lenticular galaxies and of spiral bulges from 2D observations of a representative sample of E/S0 galaxies and Sa bulges, using a custom-built panoramic integral field spectrograph placed at the 4.2m WHT on La Palma (Spain). Our early results (Bacon et al. 2001; de Zeeuw et al. 2002; Emsellem et al. 2003) reveal a variety of structures much richer than usually recognized in E/S0s galaxies.  We are carrying out a number of collaborative projects to investigate also bulges of later type (e.g. Sbc). These are taken from the samples of Balcells \\& Peletier (1994) and Carollo et al. (1997), which have extensive HST imaging available and in some cases also long-slit spectroscopy (e.g. Carollo et al. 2002; Balcells et al. 2003; Falc\\'on-Barroso et al. 2003a).  Here we present two examples of galaxy bulges observed with \\sauron. One case (NGC\\,5866) shows all the signatures expected on a classical bulge (see Sec. 1), whereas the second one (NGC\\,7332) presents properties of barred disks galaxies. ",
        "conclusions": "Bulges are not just simple scaled-down versions of elliptical galaxies. Instead, they often present properties closely related to bars or disks, that argue for different formation mechanisms. Integral field spectrographs have the ability to map the 2D behaviour of galaxies and allow us to characterize the full dynamical state of the bulges. Features like cylindrical rotation, triaxiality, kinematic decoupling, bars, as well as the amount of rotational support as a function of distance above the plane can be easily studied. Furthermore, the line-strength maps contain essential information on the spatial distribution of the stellar populations, allowing us to determine global population gradients in the radial and vertical directions. All these capabilities, combined with dynamical and stellar population models, will provide unprecedented observational information that can be used to constrain the formation and evolution processes in bulges of galaxies."
    },
    "0311/astro-ph0311552_arXiv.txt": {
        "abstract": "We present deep H$\\alpha +$[N {\\sc ii}] CCD images of selected new Planetary Nebulae (PNe) discovered through an [O {\\sc iii}] 5007$\\AA$~emission line survey in the Galactic Bulge region with $l > 0^{\\rm o}$. In total, we detected more than 200 objects, including 45 new PNe. Deep H$\\alpha +$[N {\\sc ii}] CCD images as well as low resolution spectra were obtained for the new PNe in order to study them in detail. Their first morphological study suggests that most of them are round (60\\%), while 31\\% and 9\\% are elliptical and bipolar, respectively. 63\\% of the round PNe present ring-like structure while 37\\% are compact. ",
        "introduction": "Galactic Planetary Nebulae (PNe) are of great interest because of their important role in the chemical enrichment history of the interstellar medium as well as in the star formation history and evolution of our Galaxy (Beaulieu et al. 2000 and references therein). Many surveys have been made in the past in order to discover new PNe. A list with all PNe catalogues can be found in Boumis et al. (2003a - Paper I). The study of PNe offers the opportunity to determine basic physical parameters (like morphology, kinematics, abundances, distances, masses etc.), which will help to test theoretical models and understand more about the dynamics of the Galaxy. The morphology of PNe provide us the opportunity to have a better picture of how the star evolve. Their morphology is a parameter which is difficult to quantify and several studies have been made in the past to produce a morphological classification scheme for PNe (Balick 1987; Schwarz et al. 1993; Manchado et al. 2000 and references therein). It is accepted that PNe having different morphologies arise from progenitors with different masses. In particular, circular PNe derive from stars with low mass progenitor, bipolar arise from higher mass stars while elliptical derive from all masses of progenitor (Phillips 2003 and references therein).  On a sample of PNe, Stanghellini et al. (1993; 2002) found that the central star distribution is different for bipolar and elliptical. In their latest work, they also confirm that round, elliptical and bipolar PNe have different spatial distributions within the Galaxy and that they might belong to different stellar populations. Here, we present selected H$\\alpha+$[N {\\sc ii}] images of new PNe showing the different morphological types found, and our first results concerning their morphology. ",
        "conclusions": ""
    },
    "0311/astro-ph0311287_arXiv.txt": {
        "abstract": "Recent deep {\\em Chandra} surveys of the Galactic center region have revealed the existence of a faint, hard X-ray source population. While the nature of this population is unknown, it is likely that several types of stellar objects contribute.  For sources involving binary systems, accreting white dwarfs and accreting neutron stars with main sequence companions have been proposed.  Among the accreting neutron star systems, previous studies have focused on stellar wind-fed sources. In this paper, we point out that binary systems in which mass transfer occurs via Roche lobe overflow (RLOF) can also contribute to this X-ray source population.  A binary population synthesis study of the Galactic center region has been carried out, and it is found that evolutionary channels for neutron star formation involving the accretion induced collapse of a massive ONeMg white dwarf, in addition to the core collapse of massive stars, can contribute to this population.  The RLOF systems would appear as transients with quiescent luminosities, above 2 keV, in the range from $10^{31} - 10^{32}$ ergs s$^{-1}$.  The results reveal that RLOF systems primarily contribute to the faint X-ray source population in the Muno et al. (2003) survey and wind-fed systems can contribute to the less sensitive Wang et al. (2002) survey. However, our results suggest that accreting neutron star systems are not likely to be the major contributor to the faint X-ray source population in the Galactic center. ",
        "introduction": "Recent {\\em Chandra} observations of the Galactic center region (e.g., Wang, Gotthelf \\& Lang 2002, hereinafter WGL02; Muno et al. 2003, hereinafter MM03) have led to a renewed interest in the understanding of compact stellar X-ray sources.  Both of these long exposure surveys have revealed a great number of X-ray point sources with luminosities, $L_x \\lesssim 10^{35}$ ergs s$^{-1}$ and spectra which can be fit by an absorbed power law.  Although the nature of these low luminosity sources is unknown, several possible interpretations have been proposed. In particular, Muno et al. (2004) suggest that the intermediate polar subclass of magnetic cataclysmic variables may dominate the faint population characterized by luminosities less than $10^{33}$ ergs s$^{-1}$ and a photon power law index ranging between 0 and 1. Among models involving accreting neutron stars (NS), Pfahl, Rappaport \\& Podsiadlowski (2002, hereinafter PRP02) explored the hypothesis that wind-fed neutron star systems with intermediate and massive main sequence (MS) donors ($M \\gtrsim 3 \\msun$) may account for the majority of sources with $L_x=10^{33}-10^{35}$ ergs s$^{-1}$. Furthermore, the contributions from a population of pre-low mass X-ray binaries in which the neutron stars accrete from the wind of low mass MS companions ($\\lesssim 2 \\msun$) has also  been studied by Bleach (2002) and Willems \\& Kolb (2003). However, it was recently pointed out by Popov (2003) that the neutron stars in this lower mass binary population may not necessarily accrete all the transferred material, implying that only a fraction of these types of systems would be continuous X-ray emitters. Among the non stellar models that have been proposed, Bykov (2003) has suggested that the faint hard sources result from the interaction of fast moving supernova ejecta fragments with the dense interstellar medium of the Galactic center region. These sources, due to their low intrinsic brightness, may only contribute to the deeper MM03 survey.  The aforementioned stellar studies focused on only wind-fed NS systems, however there are a number of RLOF transient systems which, in quiescence, emit at X-ray luminosities very similar to those of the many point sources observed in the Galactic center. Specifically, transients with neutron stars in their quiescent state are observed to have X-ray luminosities in the range of $10^{32}-10^{33}$ ergs s$^{-1}$ (see Campana \\& Stella 2003), while transient systems with black hole (BH) accretors tend to be observed at lower luminosities of $\\sim 2 \\times 10^{31}$ ergs s$^{-1}$ (e.g., Tomsick et al. 2003) in their inactive state.  Although many of the NS transient systems are characterized by soft spectra in their quiescent state and are difficult to detect from the Galactic center region, some fraction have a power law component extending beyond 2 keV with luminosities in the range of $\\sim 10^{31}-10^{32}$ ergs s$^{-1}$ (see Jonker, Wijnands, \\& van der Klis 2004). Such systems may have been detected in the deeper survey of MM03. The systems with BH accretors are also known to emit at energies greater than 2 keV and may also contribute to the faint source population at this luminosity level.  X-ray transient systems with large outburst to quiescent flux ratios are known to have  either a NS or BH accretor and a predominantly low mass companion (classified as a low mass X-ray binary, LMXB).  The outburst phenomenon is generally believed to result from a thermal/viscous instability in an accretion disk surrounding the compact object. Only those systems with mass transfer rates below a critical value (e.g., Dubus et al. 1999; Menou, Perna \\& Hernquist 2002) can exhibit such transient behavior.  In general, the donors of LMXBs are believed to be either low mass MS or giant-like stars.  However, observations of ultra-short period LMXBs (with orbital periods less than about 80 minutes), so called ultracompacts, have revealed that, at least, some of them must host very low mass ($0.01-0.1 \\msun$) white dwarf (WD) donors (e.g., Markwardt et al. 2002, Galloway et al. 2002; see also Deloye \\& Bildsten 2003). In a recent study Belczynski \\& Taam (2004) showed that the population of LMXBs with low mass WD donors (and thus with very low mass transfer rates leading to transient behavior) may be more significant than previously anticipated.  Since the majority of transient sources are expected to be in quiescence due to the very low X-ray duty cycle, we suggest that a fraction of the low luminosity sources observed in the Galactic center region may reflect a population of transient systems in their quiescent state, while the same systems in outburst may contribute to a bright source population.  A binary population synthesis model {\\em StarTrack} (presented in \\S\\,2) is used to determine the number of X-ray binaries (NS or BH accretor, wind-fed or RLOF) in the Galactic center region. The entire mass spectrum of evolved and unevolved donors is considered.  In particular, low, intermediate, and high-mass MS stars are investigated to account for the systems considered in previous studies. The results of our studies are described in \\S\\,3 and their implications are discussed in the final section. ",
        "conclusions": "The population of low luminosity X-ray binaries in the Galactic center region has been investigated within the framework of a population synthesis technique.  Although it has been found that this population may be more heterogeneous than previously predicted, our results suggest that accreting neutron star systems are not likely to be the major contributor to the faint X-ray source population in the Galactic center.  The RLOF systems investigated in this paper provide an additional low luminosity component.  Many of these additional systems are composed of neutron stars with white dwarf companions in ultra-short orbital periods ($\\lesssim 2$\\ hrs) in which mass is transferred at rates ranging from $10^{-12}-10^{-10} \\mpy$ as a result of the action of angular momentum losses associated with gravitational radiation.  Because the mass transfer rates are low, these systems are expected to exhibit transient quiescent/outburst behavior.  For a small duty cycle of $\\sim$\\ 1\\% it is expected that most of these sources would be found in their quiescent state with X-ray luminosities in the range of $\\sim 10^{30}-10^{33}$ erg sec$^{-1}$.  The population of hard X-ray sources ($10^{33}-10^{35}$ erg sec$^{-1}$) discovered in WGL02 survey was interpreted in the context of wind-fed NS-MS binary systems (PRP02). We find that such systems can be a contributor to the X-ray population in this luminosity range provided that $\\eta_{bol} \\sim 1$.  However, our results cannot explain the entire population since the predicted number of such systems ($\\sim 30$) is much lower than the number of observed sources ($\\gtrsim 100$).  The deeper MM03 exposure revealed a large population of faint ($L_x \\sim 10^{30}-10^{33}$ erg sec$^{-1}$) point sources. The majority of these sources are characterized by hard spectra, and are too numerous to be explained by wind-fed NS/BH systems and too hard to be explained by RLOF-fed NS/BH systems.  Muno et al. 2004 have suggested that the intermediate polar class systems can explain both their spectral properties and their observed number.  The nature of the population characterized by softer spectra is unknown, although a number of candidates have been proposed.  We have found that the number of RLOF-fed NS/BH transient systems in quiescence could be as high as 400 in comparison to the 652 observed, but this leads to an overproduction of bright LMXBs in the Galaxy.  The sensitivity of our results to input parameters have been explored, which can lead to a reduced contribution of RLOF sources and to an overall reduction in the number of synthetic soft NS/BH X-ray sources to a level further below that observed in MM03 survey.  The identification of the faint soft X-ray sources will be essential for confirming their nature. Such studies may be fruitfully carried out in the low extinction regions near the Galactic center where the contribution from the low energy component of their spectra (corresponding to less than 1 keV) may be detectable. At these lower energy ranges, the X-ray luminosities of our proposed candidates could be as high as $L_x \\sim 10^{32}-10^{33}$ ergs s$^{-1}$. The Chandra and HST surveys of Baade's window have been carried out (Grindlay et al. 2003), and the forthcoming results may have bearing on the importance of this population."
    },
    "0311/astro-ph0311293_arXiv.txt": {
        "abstract": "The dynamics of a minimally coupled scalar field in the expanding universe is discussed with special reference to phantom cosmology. The evolution of the universe with a phantom field vis-a-vis a quintessence field is compared. Phantom cosmologies are found to have two special features i) occurrence of a singularity where the scale factor, the energy density and Ricci curvature scalar diverge to infinity. This singularity occurs at a finite timexs, depending on the value of  $w$ during cosmic evolution, ii) degeneracy in the determination of $w(z_m)$ for a given  transition redshift $z_m$ which seems to impart similar observational properties to corresponding phantom and quintessence models and makes both of them compatible with the cosmological observations. Although due to the uncertainties in the measurement of the Hubble constant $H_0$, the Hubble dependent observational parameters yield only loose constraints over the range of $w$, the duality in the determination of $w$ with respect to transition redshift may be used to constrain $w$. An observational test, based upon the observations of low redshift galactic clusters, is suggested to  discriminate between the quintessence and phantom dark energy. ",
        "introduction": "The combined analysis of SNe Ia observations~\\cite{1,2}, galaxy cluster measurements~\\cite{3} and the latest CMB data~\\cite{4} provides compelling evidence for the existence of dark energy which dominates the present day universe and accelerates the cosmic expansion. The recent detection of Integrated Sachs-Wolfe effect~\\cite{5} also gives a strong and independent support to dark energy. In principle, any physical field with positive energy density and negative pressure, which violates the strong energy condition, may cause the dark energy effect of repulsive gravitation. Of late, phantom fields~\\cite{6} have emerged as potential candidates for dark energy. Scalar fields with super-negative equation of state $(p=w\\rho  , \\,w< -1)$ are called phantom fields as their energy density increases with the expansion of the universe in contrast to quintessence energy density $(w> -1)$  which scales down with the cosmic expansion. The phantom models violate the dominant energy condition $( p+\\rho) < 0$ as such they may not be physically stable models of dark energy; but, strangely enough, phantom energy is found to be compatible with most of the classical tests of cosmology~\\cite{6} based on current data from SNe Ia observations, CMB anisotropy and mass power spectrum.   The peculiar nature of phantom energy, violation of dominant energy condition and its strange consequences, possible rip-off of the large and small scale structures of matter, occurrence of future singularity and probable decay of phantom e nergy have attracted many cosmologists~\\cite{7}-\\cite{24},~\\cite{52,53,54} and made `phantom cosmologies' a hot topic of research.  In section 2, we have discussed the dynamics of  minimally coupled scalar fields with special  reference to phantom fields. There is extensive literature in cosmology on rolling scalar fields~\\cite{22,23}, quintessence fields and tracker fields~\\cite{24}-\\cite{40},~\\cite{50,51}, the cosmological constant $\\Lambda$~\\cite{25,41,42,49} and other forms of dark energy. The major problem in cosmology is to identify the form of dark energy that dominates the universe today whether it is phantom energy, quintessence, simply $\\Lambda$ or something else. Maor, Brustein and Steinhardt~\\cite{43} have discussed the degeneracy in the measurement of the dark energy parameter $w$ from SNe Ia data, its time variation and pitfalls in taking $w$ to be a constant.  In section 3, we have computed the present age $ t_0 $ of the universe in two steps. First we calculate the expansion age up to the end of matter-dominated era, denoted by $t_m$. In order to supplement it with the expansion age during dark energy dominated era, we express $t_0$ in terms of $t_m$ and thereby we calculate $t_0$. The advantage of this method is that the expansion age in the two segments can be expressed separately in terms of the redshift $z_m$  at the end of matter-dominated era which is again a function of the parameter $w$.  In section 4, assuming $w$ to be constant, we have discussed a kind of degeneracy  in the value of $w(z_m)$ which leads to duality in the behavior of phantom and quintessence models with respect to transition redshift from deceleration to accelerating phase of expansion. In fact two distinct values of parameter $w$, usually one lying in the range of quintessence field and another in the range of phantom field, lead to the same transition redshift $z_m$.  In section 5, we have tried to constrain the range of the dark energy parameter $w$ on the basis of data analysis of Kiselev~\\cite{44}, Freedman and Turner~\\cite{45} Schubnell~\\cite{46} and  the precise observational data  from WMAP~\\cite{47} in combination with SDSS~\\cite{48}. In section 6, we conclude with some remarks on phantom energy. ",
        "conclusions": "The combined analysis of the latest cosmological observations provides a definite clue of the existence of dark energy in the universe but it is difficult to distinguish between the various forms of dark energy at present. So far as phantom energy is concerned, it is found to be compatible with SNe Ia observations and CMB anisotropy measurements but the violation of the `dominant Energy condition' makes phantom models physically unstable; however, phantom models may be considered to be phenomenologically viable provided their age happens to be less than the time scale of the singularity. In case the instability occurs earlier and the dark energy decays into gravitons, as discussed by Carroll et al~\\cite{14}, the universe might escape the ordeal of `Rip-Off'~\\cite{7} and phantom singularity. With the large number of SNe Ia observations expected from SNAP, LOSS and other surveys in the coming years, a clear picture of the dark energy profile is likely to emerge which would reveal whether $\\Lambda$CDM, QCDM or PCDM is the correct cosmology of the universe."
    },
    "0311/astro-ph0311546_arXiv.txt": {
        "abstract": "Submillimeter and near-infrared images of cool dusty debris disks and rings suggest the existence of unseen planets. At dusty but non-imaged stars, semi-major axes of associated planets can be estimated from the dust temperature. For some young stars these semi-major axes are greater than an arc second as seen from Earth. Such stars are excellent targets for sensitive near-infrared imaging searches for warm planets.  To probe the full extent of the dust and hence of potential planetary orbits, SIRTF observations should include measurements with the 160$\\,\\mu$m filter. ",
        "introduction": "At near-infrared wavelengths, adaptive optics on 8~m class ground-based telescopes and the NICMOS camera on the Hubble Space Telescope can probe regions within a few arc seconds of nearby stars. In such regions, for stars younger than $\\sim 100$~Myrs, warm massive planets can be detected. Dozens of such young stars within $\\sim 60$~pc of Earth have been identified \\citep[and references therein]{SZB03,ARAA}.  Extant images at submillimeter and near-infrared wavelengths of cool dusty debris disks at main sequence stars (the so-called Vega phenomenon) usually show substantial spatial structure (e.g., \\citealt{Holland98} \\& \\citeyear{Holland03}; \\citealt{Greaves,Schneider,Krist,Koerner,Wilner}; \\citealt{Weinberger02} \\& \\citeyear{Weinberger03}; \\citealt{Wahhaj,Clampin,Zuckerman01}, and references therein). Specifically, the dusty regions around $\\epsilon $~Eri, Vega, Fomalhaut, $\\beta $~Pic, and HD~141569 all show obvious non-axisymmetric structure. HR~4796 is orbited by a narrow dusty ring. Only the dust at TW Hya has, so far, failed to reveal any structure of note.  Excepting perhaps HD~141569 \\citep[and references therein]{Clampin}, the most likely explanation of the observed structures are gravitational perturbations by planets with semi-major axes comparable to the radius of the dusty rings and disks. For Vega, a 3 Jupiter mass ($M_{J}$) planet is suggested \\citep{Wilner}, while for $\\epsilon $~Eri either a 0.2~$M_{J}$ \\citep{Ozernoy} or a 0.1~$M_{J}$ \\citep{Quillen} planet has been proposed. Indeed, COBE found that Earth is led and trailed in its orbit around the Sun by clumps of dust particles \\citep{COBE}. Additional discussions of planet/disk interactions can be found in \\citet{Holland03}, \\citeauthor{Kenyon} (2002a, 2003), \\citet{Kuchner}, \\citet{WD02}, and \\citet{Wyatt03}.  Alternative mechanisms to generate structure in the dusty disks of other stars have been proposed, but likely account for very few, if any, of the observed structures. \\citet{Kalas} investigated the possibility that a recent, close encounter with a passing star generated the structure observed in the northeast arm of the $\\beta $~Pic disk. While such an encounter might occur on rare occasions, such low-probability events cannot plausibly generate dusty structure in a high percentage of the Vega-like stars \\citep[e.g.,][]{KB02}. \\citet{Takeuchi} proposed that ring structure, such as is observed at HR~4796 could be generated by dust particles migrating in an optically thin gaseous disk.  This mechanism is implausible as an explanation for most, perhaps all, of the observed dusty structures because of (1)\\,their often dramatic non-axisymmetric shape and (2)\\,no evidence of gas at a majority of the above listed stars including even the model prototype HR~4796 \\citep[e.g., ][and references therein]{Zuckerman01}.  As a consequence of the above observations and arguments, one may reasonably assume that most stars with imaged dust emission have at least one planet on a wide orbit (tens of AU). From this assumption it follows that at stars with photometrically detected (but not yet imaged) dust there is a high probability of orbiting planets including at least one with a semi-major axis comparable to that of the dimensions of the dust disk.  Extrapolating from the preceding, one might speculate that young stars that do not have dust detectable by SIRTF do not possess planets on wide orbits (see discussion at end of Section~2). Although dust has not yet been measured directly in the Sun's Kuiper Belt, dust is very likely present and its presence is consistent with the existence of a planet (Neptune) at a roughly comparable distance from the Sun.  Indeed, as has been pointed out by numerous researchers, with but few exceptions \\citep[e.g.,][]{Chen}, the dusty disks detected around main sequence stars (by IRAS and ISO) are likely analogs of the Kuiper Belt.  Twenty years ago, when IRAS first discovered the Vega-phenomenon, astronomers wondered whether a substantial link connected the presence of dust with planets. While stars with imaged dust remain frustratingly few in number, the common and striking structures that are seen now suggest that the link is substantial and that dusty stars represent excellent targets for planet searches.  For main sequence stars with circumstellar dust detected by IRAS and/or ISO, the measured dust temperature is an indication of how far from a star the dust is located. In the following section we outline how the dust can be used to guide imaging planet searches and list specific stars at which to search. ",
        "conclusions": "Recent imaging observations and analysis of structure in the dusty debris disks that surround a handful of nearby stars suggest that they possess planets on wide orbits (tens of AU). Thus, dusty stars, whether imaged or not, are excellent targets for planet searches. Young stars with cool dust make the best targets for adaptive optics and HST imaging programs. Older stars and those with somewhat warmer dust might be better investigated with precision radial velocities. While the earliest searches of the IRAS catalogs revealed mostly A-type Vega-like stars which are not suitable for study by the precision radial velocity technique, a majority of stars in Table 1 are of later spectral types.  Looking toward the future, one anticipates that SIRTF will search for dusty debris at all nearby stars that have been identified as very young (e.g., Zuckerman et al 2001a, Song et al 2003). Since the presence of cool dust would point toward the existence of planets on wide orbits accessible to adaptive optics, it is important that the SIRTF programs include measurements with the 160~$\\mu $m filter."
    },
    "0311/physics0311051_arXiv.txt": {
        "abstract": "We consider stochastic reconnection in a magnetized, partially ionized medium.  Stochastic reconnection is a generic effect, due to field line wandering, in which the speed of reconnection is determined by the ability of ejected plasma to diffuse away from the current sheet along magnetic field lines, rather than by the details of current sheet structure.   As in earlier work, in which we dealt with a fully ionized plasma, we consider the limit of weak stochasticity, so that the mean magnetic field energy density is greater than either the turbulent kinetic energy density or the energy density associated with the fluctuating component of the field.  For specificity, we consider field line stochasticity generated through a turbulent cascade, which leads us to consider the effect of neutral drag on the turbulent cascade of energy.  In a collisionless plasma, neutral particle viscosity and ion-neutral drag will damp mid-scale turbulent motions, but the power spectrum of the magnetic perturbations extends below the viscous cutoff scale. We give a simple physical picture of the magnetic field structure below this cutoff, consistent with numerical experiments.  We provide arguments for the re\\\"emergence of the turbulent cascade well below the viscous cut-off scale and derive estimates for field line diffusion on all scales.  We note that this explains the persistence of a single power law form for the turbulent power spectrum of the interstellar medium, from scales of tens of parsecs down to thousands of kilometers.  We find that under typical conditions in the ISM stochastic reconnection speeds are reduced by the presence of neutrals, but by no more than an order of magnitude.  However, neutral drag implies a steep dependence on the Mach number of the turbulence.  In the dense cores of $H_2$ regions the reconnection speed is probably determined by tearing mode instabilities. ",
        "introduction": "One of the fundamental properties of astrophysical magnetic fields is their ability to change topology via reconnection \\citep[see][]{PF00}.  It is impossible to understand the origin and evolution of large scale magnetic fields without understanding the mobility of magnetic field lines. In a typical astrophysical plasma, resistivity is very small and flux freezing, which follows from assuming zero resistivity, should be an excellent guide to the motion of magnetic fields. In spite of this, fast magnetic dynamo theory \\citep[see][]{P79,M78,KR80} invokes a constantly changing magnetic field topology and motions during which field lines cross each other\\footnote{While standard mean-field dynamo theory has been severely criticized on theoretical and numerical grounds \\citep{VC92,GD94,GD96,CH96,HCK96,B01}, an alternative version can be formulated which evades these criticisms \\citep[][]{VC01} and which assumes fast reconnection only on two dimensional surfaces.}. This assumption is supported by observations of the solar magnetic field \\citep[see][and references contained therein]{D96,IIAW97} which are difficult to explain unless flux freezing is routinely violated on time scales short compared to resistive time scales, at least within thin current sheets.  The Sweet-Parker model of reconnection is the simplest and most robust \\citep{P57, S58}.  In this model, reconnection takes place within a thin current sheet, which separates two large volumes containing uniform, and very different, magnetic fields.  The resulting reconnection speed is less than the Alfv\\'en speed by the square root of the Lundquist number $\\sim {\\cal R}_L^{-1/2}=(\\eta/V_A L)^{1/2}$, where $\\eta$ is the resistivity, $V_A$ is the Alfv\\'en speed, and $L$ is the length of the current sheet, assumed to be determined by the large scale geometry of the problem.  Under typical astrophysical conditions this is very slow (e.g. for the Galaxy as a whole ${\\cal R}_L\\sim 10^{20}$).  This reconnection speed is set by a geometrical constraint. Indeed, plasma tied to the reconnecting magnetic field lines must be ejected from the ends of the narrow current sheet. The disparity of scales, one of which is macroscopic/astrophysical, while the other is microscopic, i.e. determined by ohmic diffusion, makes the reconnection slow.  This evident shortcoming of the Sweet-Parker reconnection scheme has stimulated interest in alternative models that allow fast reconnection. Although the literature on magnetic reconnection is extensive \\citep[e.g.][and references therein]{PF00} it does not successfully address this question. Models that invoke an X-point reconnection geometry \\citep{P64} have been shown to be unstable for sufficiently high $R_L$ \\citep[see][]{B96}, while anomalous resistivity fails to provide rapid reconnection under most astrophysical conditions \\citep[see][]{P79}. A general review of astrophysical magnetic reconnection theory can be found in \\citet{BMW03}.  A notable exception to this discouraging state of affairs is the recent work on fast collisionless reconnection \\citep[see also the discussion by Bhattacharjee, Ma and Wang 2001] {BSD97,SDD98,SD98}. This work indicates that under some circumstances a kind of standing whistler mode can stabilize an X-point reconnection region.  However, these studies have not demonstrated the possibility of fast reconnection for generic field geometries. They assume that there are no bulk forces acting to produce a large scale current sheet, and that the magnetized regions are convex, which minimizes the energy required to spread the field lines.  In addition, this mechanism requires a collisionless environment, where the electron mean free path is less than the current sheet thickness \\citep{TYJKC03, JYHK98}. Unfortunately, in the laboratory the current sheet thickness is comparable to the ion Larmor radius and it is unclear how to generalize this criterion to the interstellar medium, where the Sweet-Parker current sheet thickness is typically much greater than the Larmor radius.  {\\it If} we require that the ion skin depth, the characteristic scale of the standing Whistler mode, be greater than the current sheet thickness, then we have a criterion which is rarely satisfied in the interstellar medium. In this paper we will concentrate on a mechanism that acts in turbulent media and when it works, produces rapid reconnection under a broad range of field geometries, without regard to the particle collision rate.  We will defer all discussion of the relationship between stochastic reconnection and small scale collisionless effects to a later paper.   In a previous paper \\citep[henceforth LV99]{lv99} we discussed `stochastic reconnection', a process which is similar to Sweet-Parker reconnection, except that stochastic wandering of the field lines produces a broad outflow region. The properties of the outflow region are insensitive to the width of the current sheet (and the value of ${\\cal R}_L$), but depend on the level of field line stochasticity.  In a sufficiently noisy environment the reconnection speed becomes a large fraction of the Alfv\\'en speed. In an {\\it extremely} quiet environment the field lines do not enter or leave the current sheet over its entire length and we recover the Sweet-Parker reconnection model. In LV99 we dealt with an inviscid and totally ionized fluid. \\footnote{In this paper we shall show that this approximation is valid for partially ionized collisionless plasma up to a certain percentage of neutrals, and for collisional gases with resistivity larger than viscosity.}  The notion that magnetic field stochasticity might affect current sheet structures is not unprecedented.  In earlier work \\citet{S70} showed that in collisionless plasmas the electron collision time should be replaced with the time a typical electron is retained in the current sheet.  Also \\citet{JM84} proposed that current diffusivity should be modified to include diffusion of electrons across the mean field due to small scale stochasticity.  These effects will usually be small compared to effect of a broad outflow zone containing both plasma and ejected shared magnetic flux.  Moreover, while both of these effects will affect reconnection rates, they are not sufficient to produce reconnection speeds comparable to the Alfv\\'en speed in most astrophysical environments.  It is important to distinguish between stochastic reconnection, as discussed in LV99, and the more conventional notion of turbulent reconnection.  The latter usually involves substituting a turbulent diffusivity for the resistivity, and involves a degree of small-scale mixing which is forbidden on energetic grounds \\citep[see \\S 2 in][and the references contained therein]{P92}.  A common variation of this hypothesis is that instabilities in the current sheet will produce a hugely broadened current sheet and a large effective resistivity within it \\citep[for a review see][] {B00}.  On the contrary, the former is largely a topological effect, using conventional estimates of resistivity, and the only strong mixing associated with it has to do with the polarization of field lines crossing the current sheet.  (That is, one expects sharp gradients, in the current sheet, in the component of the magnetic lines perpendicular to the current sheet.) In this sense, stochastic reconnection belongs to the class of models which try to explain fast reconnection by appealing to a current sheet geometry which is `natural' in some sense, but evades the limits set by the Sweet-Parker model.  (The bulk of the discussion in \\citet{PF00} is centered on laminar three dimensional field configurations which can lead to similarly rapid reconnection speeds.)  In a recent paper Kim \\& Diamond (2001) addressed the problem of stochastic reconnection by calculating the turbulent diffusion rate for magnetic flux inside a current sheet.  They obtained similar turbulent diffusion rates for both two dimensional and reduced three dimensional MHD. In both cases the presence of turbulence had a negligible effect on the flux transport.  The authors pointed out that this would prevent the anomalous transport of magnetic flux within the current sheet and concluded that both 2D and 3D stochastic reconnection proceed at the Sweet-Parker rate even if individual small scale reconnection events happen quickly. If true this would be not only rule out the LV99 reconnection scheme, but also any other fast reconnection scheme. In general astrophysical plasmas are turbulent and if the enhancement of the local reconnection speed, e.g. due to collisionless effects (see Drake et al. 2001), is irrelevant then reconnection must always be slow.  However turbulent diffusion rates within the current sheet are irrelevant for the process of stochastic reconnection or, for broadly similar reasons, fast collisionless reconnection. The basic claim in LV99 is that realistic magnetic field topologies allow multiple connections between the current sheet and the exterior environment, which would persist even if the stochastic magnetic field lines were stationary (\"frozen in time\") before reconnection.  This leads to global outflow constraints which are weak and do not depend on the properties of the current sheet.  In particular, the analysis in LV99 assumed that the current sheet thickness is determined purely by ohmic dissipation and that turbulent diffusion of the magnetic field is negligible inside, and outside, the current sheet.  The major uncertainty in this model is the behavior of the reconnected flux elements, which are nearly perpendicular to the current sheet and must undergo multiple reconnections before being ejected.  We note also that models of collisionless reconnection also evade the objection posed by Kim and Diamond topologically, that is, by stabilizing an X-point reconnection topology, and opening up the rest of the current sheet.  \"Hyper-resistivity\" \\citep{S85,BH86,HB87} is a more subtle attempt to derive fast reconnection from turbulence within the context of mean-field resistive MHD.  The form of the parallel electric field can be derived from magnetic helicity conservation.  Integrating by parts one obtains a term which looks like an effective resistivity proportional to the magnetic helicity current.  There are several assumptions implicit in this derivation, but the most important problem is that by adopting a mean-field approximation one is already assuming some sort of small-scale smearing effect, equivalent to fast reconnection.  \\citet{S88} partially circumvented this problem by examining the effect of tearing mode instabilities within current sheets.  However, the resulting reconnection speed enhancement is roughly what one would expect based simply on the broadening of the current sheets due to internal mixing.  This effect does not allow us to evade the constraints on the global plasma flow that lead to slow reconnection speeds, a point which has been demonstrated numerically \\citep{ML85} and analytically (LV99).  Nevertheless, we show in \\S 4 that this effect may be important in the densest and coldest parts of the ISM.  A partially ionized plasma fills a substantial volume within our galaxy and the earlier stages of star formation take place in a largely neutral medium. This motivates our study of the effect of neutrals on reconnection. The role of ion-neutral collisions is not trivial. On one hand, they may truncate the turbulent cascade, reducing the small scale stochasticity and decreasing the reconnection speed. On the other hand, the ability of neutrals to diffuse perpendicular to magnetic field lines allows for a broader particle outflow and enhances reconnection rates.  Reconnection in partially ionized gases has been already studied by various authors \\citep{NMA92,ZB97}.  In a recent study \\citep[henceforth VL99]{VL99} we studied the diffusion of neutrals away from the reconnection zone assuming anti-parallel magnetic field lines \\citep[see also][]{HZ03a} The ambipolar reconnection rates obtained in VL99, although large compared with the Sweet-Parker model, are insufficient either for fast dynamo models or for the ejection of magnetic flux prior to star formation.  In fact, the increase in the reconnection speed stemmed entirely from the compression of ions in the current sheet, with the consequent enhancement of both recombination and ohmic dissipation.  This effect is small unless the reconnecting magnetic field lines are almost exactly anti-parallel \\citep{VL99, HZ03b}.  Any dynamically significant shared field component will prevent noticeable plasma compression in the current sheet, and lead to speeds practically indistinguishable from the standard Sweet-Parker result.  Since generic reconnection regions will have a shared field component of the same order as the reversing component, the implication is that reconnection and ambipolar diffusion do not change reconnection speed estimates significantly.  None of this work on reconnection in partially ionized plasmas includes the effects of stochasticity. We expect that in the presence of turbulence, reconnection rates will be substantially enhanced, as they are in completely ionized plasmas. To generalize the concept of stochastic reconnection to partially ionized plasmas we need a model for the small scale stochasticity of a turbulent magnetic field in a partially ionized plasma. In this paper we will begin by considering this problem, and then apply our results to the reconnection speed.  Following LV99, we consider reconnection in the presence of a weakly stochastic magnetic field.  Except for the presence of noise, we imagine a reconnection event exactly like a generic Sweet-Parker reconnection event.  Two volumes with average magnetic fields that are of comparable strength, but differing directions are in contact over a surface of length $L$.  Due to the stochastic nature of the fields, field lines come into contact over many small patches (see Fig.~1). For each individual patch the Sweet-Parker reconnection model should be applicable (or at least constitute a minimal reconnection speed). The enhancement of reconnection rates follows from two effects. First, since individual field lines wander out of the narrow current sheet relatively easily, the longitudinal patch size is much smaller than the overall size of the system. This reduces the effective value of ${\\cal R}_L$ and raises the local reconnection speed. Second, whereas in the Sweet-Parker scheme magnetic field lines reconnect sequentially, in the presence of field line wandering field lines many independent patches are brought into direct contact and can reconnect simultaneously. As a consequence, the rate of reconnection of the magnetic flux is increased by a large factor, whose exact value depends on the level of noise in the system.  In LV99 the truncation of the turbulent cascade was assumed to be due to resistivity.  Consequently, the smallest scale of field line wandering decreases as resistivity decreases, and the number of independent patches in contact within the reconnection zone increases. {}From this we concluded that for an idealized inviscid fluid the reconnection rate does not depend on fluid resistivity. It does depend on the level of magnetic field stochasticity.  In the specific case where the field line stochasticity is caused by a turbulent cascade, it depends on the amplitude of the turbulence.  However, there is no necessary connection between turbulent motions and the reconnection speed.  In particular, the rate of turbulent transport of mean magnetic flux is assumed to be negligible in this model, and even the complete absence of turbulent diffusion in the current sheet would not reduce the stochastic reconnection rate \\citep[cf.][]{KD01}.  To quantify stochastic reconnection we have to use a particular description of turbulence. Motions in a magnetized medium can be expanded into incompressible (Alfv\\'en) and compressible (fast and slow) modes. There are theoretical arguments \\citep{gs95,LG01} \\citep[see][for a review]{CLV03a} suggesting that the nonlinear cascade of power for these modes proceeds separately, although not entirely independently. Simulations in \\citet{CL02} support this idea and show that the  \\citet[henceforth GS95]{gs95} scaling is applicable to Alfvenic part of the MHD cascade. For our purposes the scalings of slow and fast modes \\citep{CL02}, \\citep{CL03} are less important since they are subjected to collisionless damping\\footnote{Whether or not those damped modes are important depends on the process studied. For instance, \\citet{YL02} show that for scattering of cosmic rays the residual small amplitude fast modes are much more efficient than the Alfv\\'en modes. This, however, is not true for the field wandering that we consider in this paper.}  Here we assume that the GS95 model describes incompressible turbulence above the ambipolar damping scale\\footnote{We note, however, that our qualitative conclusions for reconnection rates should be valid for other models of MHD turbulence (see LV99) as long as they are in rough agreement with observational constraints and numerical simulations.}. To describe MHD turbulence below the scale of viscous damping we present a new model of magnetic field structure in this regime. This model is in rough agreement with numerical simulations by \\citet[henceforth CLV02b]{CLV02b}.    In \\S 2 of this paper we will consider the effect of a large neutral fraction on a strongly turbulent cascade in a magnetized plasma.  In \\S 3 we apply this to the problem of reconnection in partially ionized plasmas. In \\S 4 we apply this work to various phases in the ISM.  Finally, \\S 5 contains our basic conclusions. ",
        "conclusions": "Our results can be summarized as follows:  \\begin{enumerate} \\item In plasmas characterized by a moderate magnetic Prandtl number (resistivity less than, or of order, viscosity) magnetic reconnection follows the model described in LV99, so that reconnection is fast in the presence of noise.  Collisionless plasmas with small neutral fractions and collisionally dominated plasmas are typically in this class.   In either case the turbulent cascade proceeds to either the resistive scale or the ion Larmor radius. For typical interstellar conditions, it is the ion Larmor radius which is larger.  \\item As the neutral content in plasma increases the turbulent cascade is interrupted at a scale greater than the neutral mean free path. However, the magnetic field perturbations are not suppressed below this scale.  On the contrary, we find that the magnetic field power spectrum flattens out, and exhibits more power on small scales than the Goldreich-Sridhar spectrum. Velocity fluctuations in this regime are driven by magnetic field perturbations, whose energy is much larger than the kinetic energy perturbations on the same scale. The importance of this new regime of magnetic stochasticity goes far beyond understanding the reconnection problem.  \\item On somewhat smaller scales, once the ions and neutrals are decoupled, the turbulent cascade can reappear.  The warm atomic gas in the ISM shows only a slight interruption in the turbulent cascade and the small scale scintillations are close to a continuation of the large scale power.  However, this break will be more conspicuous in colder and denser phases of the ISM.  In this limit the small scale turbulence will be strongly intermittent.  \\item Magnetic stochasticity on scales smaller than the ion-neutral damping cut-off of the Goldreich-Sridhar spectrum promotes fast reconnection. The diffusion of magnetic field lines initially decreases rapidly below this scale, but the decoupling of ions and neutrals leads to a minimum value of the diffusion coefficient on small scales.  The most stringent limit on reconnection speeds comes from considering the largest scales where this minimum value holds.   This is in the range of the turbulent cascade where ions and neutrals are decoupled, but neutral drag is still sufficient to stabilize magnetic perturbations. In most phases of the ISM this gives a reconnection speed which is just slightly smaller than the ionized plasma estimate given in LV99.  \\item Our study shows that the reconnection speed in the ISM is always much faster than the predictions given by the Sweet-Parker rate. Even in very dense and cold regions it is only slightly below the Alfv\\'en speed, and tearing mode instabilities in the current sheet tend to drive it back up to the Alfv\\'en speed. This process needs to be considered in models for magnetic flux removal during star formation, and for studies of the dynamics of magnetized molecular clouds. \\end{enumerate}  Finally, we note that our conclusions are moderately sensitive to the nature of field line diffusivity in the various damped regimes ($k_t<k<\\min[k',k_c]$).  We have argued that the shearing due to motions at $k\\sim k_c$ dominate until we are well into the decoupled damped regime.  Our estimate for a constant field line diffusivity on somewhat smaller scales has not been tested against simulations. If this turns out to be too high (or low) then our reconnection speed estimate will have to be lowered (or raised)."
    },
    "0311/astro-ph0311199_arXiv.txt": {
        "abstract": "{We study dynamical consequences of the kink instability of a twisted coronal flux rope, using the force-free coronal loop model by \\cite{Tit:Dem-99} as the initial condition in ideal-MHD simulations. When a critical value of the twist is exceeded, the long-wavelength ($m\\!=\\!1$) kink mode develops. Analogous to the well-known cylindrical approximation, a helical current sheet is then formed at the interface with the surrounding medium. In contrast to the cylindrical case, upward-kinking loops form a second, vertical current sheet below the loop apex at the position of the hyperbolic flux tube (generalized X line) in the model. The current density is steepened in both sheets and eventually exceeds the current density in the loop (although the kink perturbation starts to saturate in our simulations without leading to a global eruption). The projection of the field lines that pass through the vertical current sheet shows an S shape whose sense agrees with the typical sense of transient sigmoidal (forward or reverse S-shaped) structures that brighten in soft X rays prior to coronal eruptions. The upward-kinked loop has the opposite S shape, leading to the conclusion that such sigmoids do not generally show the erupting loops themselves but indicate the formation of the vertical current sheet below them that is the central element of the standard flare model. ",
        "introduction": "\\label{intro}   Rising filaments, flares, and coronal mass ejections on the Sun often show the phenomenology of a loop-shaped magnetic flux system with fixed footpoints at the coronal base and signatures of magnetic twist. A single twisted magnetic flux rope appears to contain essential elements of the magnetic topology of the unstable, often erupting flux. Since the ratio of kinetic to magnetic pressure, the plasma beta, is very small in the inner corona, $\\beta\\sim10^{-3}\\dots10^{-2}$, the flux system must be nearly force free in the quasi static pre-eruption state, i.e., \\begin{equation} \\nabla\\mbox{\\boldmath$\\times$}\\vec{B}=\\alpha(\\vec{x})\\vec{B}\\,. \\label{eq_forcefree} \\end{equation}  Such major energy release events, which we will refer to collectively as flares in the following, often show a distinctive soft X-ray (SXR) brightening of sigmoidal (S or reverse-S) shape shortly prior or during their early, impulsive stages of development \\cite[e.g.,][] {Rus:Kum-96,Mano:al-96,Pevt:al-96,Aura:al-99,Moor:al-01,Gibs:al-02}. The sigmoidal shape is generally regarded as a signature of coronal currents in a twisted field and shows a hemispheric preference: forward (reverse) S shapes dominate in the southern (northern) hemisphere \\citep{Rus:Kum-96}. Typically the central part of such transient sigmoids is approximately aligned with the neutral line of the normal photospheric field component, often also with a filament. They are bright and more sharply defined in the middle, while their ends fan out more diffusively at opposite sides of the neutral line. Sometimes their middle section is slightly split, giving the structure the appearance of two point-symmetric letters J. Each J may consist of several fibres (Fig.\\,\\ref{fig_TD99f1}). \\begin{figure}    % \\begin{center} \\resizebox{.85\\hsize}{!}                % {\\includegraphics{Fd041_f1.eps}} \\end{center} \\vspace{-3mm} \\caption[] {Transient sigmoid observed in an eruptive flare on 1994 Oct 25 (from \\citeauthor{Mano:al-96}, \\citeyear{Mano:al-96}) in AR 7792, which had primarily left-handed twist \\citep{Aura:al-99,vanDrie:al-00}.} \\label{fig_TD99f1} \\end{figure}  The solar soft X-ray corona often also contains long-lived sigmoidal structures that appear to form a separate class. These typically exist for more than a day, do not show the concentration in the middle, may be composed of many fibres that show the S shape collectively rather than individually, and may form not only within but also between active regions.  The sense of the sigmoidal shape and the handedness of the magnetic twist are correlated: \\cite{Pevt:al-97} found that more than 90{\\%} of all sigmoids in those active regions that have a uniform sign of $\\alpha$ in the photospheric magnetogram show forward (reverse) S shape for $\\alpha>(<)~0$, i.e., for right (left) handed twist. We will refer to this result as $\\sigma$-$\\alpha$ correlation, where $\\sigma>(<)~0$ stands for forward (reverse) S shape. The correlation is consistent with the hemispheric preference, since $\\alpha$ is predominantly $>(<)~0$ in the southern (northern) hemisphere \\citep{Seehafe-90,Pevt:al-95}.  The transient sigmoids that brighten in flares (both confined and eruptive) have so far mostly been interpreted as an illumination of the unstable flux system itself. \\cite{Rus:Kum-96} proposed that they show a sufficiently twisted, kink-unstable flux rope. \\cite{Pevt:al-96}, \\cite{Moor:al-01}, and others suggested further that the unstable flux rope acquired its supercritical twist through the reconnection of two nearly aligned, subcritical flux ropes near their end points, which conforms to the strong brightening in the middle of the sigmoids and may symbolically be written as $J+J\\to{S}$. Transient sigmoids were thus regarded by \\cite{Moor:al-01} as supporting their ``tether cutting'' model of flares (although they noted that another, much fainter structure was moving in sequences of SXR images in some of their events).  The presence of a SXR sigmoid correlates with the likelihood of a flaring active region to cause an eruption \\citep{Canf:al-99}. A ``standard'' model for the main phase of eruptive flares is now widely accepted \\cite[e.g.,][]{Shibata-99}. It includes the formation of a large-scale vertical current sheet below the erupting flux, in which magnetic reconnection leads to the long-lasting energy release. This model is assumed to hold independently of the question whether the event is initiated by an instability of a single flux rope, by reconnection between flux ropes (tether cutting), or by other mechanisms, e.g., ``magnetic breakout'' \\citep{Anti:al-99}.  The canonical instability of a twisted flux rope is the kink instability. In applications to the solar corona, the simplifying assumption of straight, cylindrically symmetric flux ropes has always been used so far \\cite[e.g.,][]{Hood-92}. This is based on the typically large aspect ratio of solar magnetic loops. The instability occurs if the twist, $\\Phi\\!=\\!lB_\\phi(r)/(rB_z(r))$, exceeds a critical value, $\\Phi_\\mathrm{c}$. Here $l$ is the length of the flux rope, $r$ is the (minor) radius, and $B_z$ and $B_\\phi$ are the axial and azimuthal field components, respectively. The critical value depends on the details of the considered equilibrium.  In this paper we address the relationship between the kink instability of an arched twisted flux rope and the formation of current sheets and sigmoidal structures, starting from the analytical magnetic loop equilibrium by \\cite{Tit:Dem-99}, cited as T\\&D in the following (see their Fig.\\,2 for a schematic). This approximate, cylindrically symmetric, force-free equilibrium consists of a toroidal ring current of major radius $R$ and minor radius $a$, whose outward-directed Lorentz self-force is balanced with the help of a field by two fictitious magnetic charges of opposite sign which are placed at the symmetry axis of the torus at distances ${\\pm}L$ to the torus plane. That axis lies below the photospheric plane $\\{z\\!=\\!0\\}$ at a depth $d$. The resulting field outside the torus is current-free and contains a concentric magnetic X line between the torus and its centre. A toroidal field component created by a fictitious line current running along the axis of symmetry is added. This results in a force-free field inside the torus with the twist taking a finite value everywhere in the configuration, whose part in the volume $\\{z\\!>\\!0\\}$ is a model of a coronal loop (or flux rope).  The presence of a toroidal field component turns the neighbourhood of the X line into a hyperbolic flux tube (HFT) which consists of two intersecting quasi-separatrix layers with extremely diverging field lines (for a strict definition of HFTs, see \\citeauthor{Tito:al-02}, \\citeyear{Tito:al-02}). The existence of the HFT is generic to such force-free loop configurations with a nonvanishing net current (\\citeauthor{Tito:al-99}, \\citeyear{Tito:al-99}), which is important for understanding the origin of sigmoidal structures in twisted configurations.  We present two characteristic cases of the kink instability, which show the different possibilities of current sheet and sigmoid formation in the T\\&D equilibrium, leaving a systematic study of the instability for a separate paper \\citep{Toer:al-03}.  T\\&D pointed out that the torus is also unstable with respect to global expansion (growing perturbations ${\\delta}R>0$). This has recently been confirmed by \\cite{Rous:al-03}. We expect that the global expansion instability leads to current sheets and sigmoidal structures in a manner largely analogous to the upward kink instability.      We integrate the compressible ideal MHD equations using the simplifying assumption $\\beta=0$ which has no influence on the qualitative evolution of the kink instability in its linear phase, during which the current sheets and sigmoids are formed: \\begin{eqnarray} \\partial_t\\rho&=& -\\nabla\\cdot(\\rho\\,\\vec{u})\\,,           \\label{eq_rho}\\\\ \\rho\\,\\partial_{t}\\vec{u}&=& -\\rho\\,(\\,\\vec{u}\\cdot\\nabla\\,)\\,\\vec{u} +\\vec{j}\\mbox{\\boldmath$\\times$}\\vec{B} \\,,         \\label{eq_mot}\\\\ \\partial_{t}\\vec{B}&=& \\nabla\\mbox{\\boldmath$\\times$}(\\,\\vec{u}\\mbox{\\boldmath$\\times$} \\vec{B}\\,)\\,,                                         \\label{eq_ind}\\\\ \\vec{j}&=&\\mu_0^{-1}\\nabla\\mbox{\\boldmath$\\times$}\\vec{B}\\,. \\label{eq_cur} \\end{eqnarray}  The equilibrium by T\\&D is used as the initial condition for the field, with the loop chosen to lie in the plane $\\{x\\!=\\!0\\}$. The initial density distribution is specified such that the Alfv\\'en velocity is uniform, $\\rho_0=B_0^2/\\mu_0$. Lengths, velocities, and times are normalized, respectively, by $L$, the initial Alfv\\'en velocity, $v_\\mathrm{a0}=B_0/(\\mu_0\\rho_0)^{1/2}$, and the Alfv\\'en time, $\\tau_\\mathrm{a}=L/v_\\mathrm{a0}$.  A modified Lax-Wendroff scheme is employed on a nonuniform Cartesian grid in a ``half cube'' $[-L_x,L_x]\\times[0,L_y]\\times[0,L_z]$ with closed boundaries except the front boundary, $\\{y\\!=\\!0\\}$, where line symmetry with respect to the $z$ axis is implemented. The latter implies $u_{x,y}(0,0,z)=0$ so that only vertical displacements of the loop apex are permitted. Viscosity is included in Eq.\\,(\\ref{eq_mot}) to improve numerical stability. The numerical details are the same as in \\cite{Toer:al-03}. ",
        "conclusions": "\\label{concl}   The magnetic loop equilibrium by T\\&D is unstable against the long-wavelength ($m\\!=\\!1$) ideal kink mode if the magnetic twist exceeds a critical value. The handedness of the kink perturbation equals the handedness of the field line twist. Consequently, the photospheric projection of rising kink-unstable loops develops a forward (reverse) S shape for left (right) handed twist. This is opposite to the strong observed ($\\sigma$-$\\alpha$) correlation of sigmoid shape with magnetic handedness. Hence, transient SXR sigmoids, which brighten in major energy release events, preferably in eruptive events, cannot in general show the presumed kink-unstable flux itself. We suggest that transient sigmoids outline field lines which pass through the current density enhancement in a vertical current sheet formed below the rising loop. This interpretation leads to agreement with the $\\sigma$-$\\alpha$ correlation (and hemispheric preference) as well as with their alignment with the photospheric neutral line and with their further specific geometric properties. Transient sigmoids thus indicate the formation of the vertical current sheet that is the central element of the standard model for eruptive flares. The question under which conditions a rising kink-unstable loop develops into an eruption requires further study.  Downward-kinking loops take a sigmoidal shape that follows the observed correlation with handedness. They may be related to long-lived sigmoids and to filament formation."
    },
    "0311/astro-ph0311016_arXiv.txt": {
        "abstract": "The relation between the $2-10$ keV, long term, excess variance and AGN black hole mass is considered in this work. A significant anti-correlation is found between these two quantities in the sense that the excess variance decreases with increasing black hole mass. This anti-correlation is consistent with the hypothesis that the $2-10$ keV power spectrum in AGN follows a power law of slope $-2$ at high frequencies. It then flattens to a slope of $-1$ below a break frequency, \\nbh, until a second break frequency, \\nbl, below which it flattens to a slope of zero. The ratio \\nbh/\\nbl\\ is equal to $10-30$, similar to the ratio of the respective frequencies in Cyg X-1. The power spectrum amplitude in the (frequency $\\times$ power) space does not depend on black hole mass.  Instead it is roughly equal to $0.02$ in all objects. The high frequency break decreases with increasing black hole mass according to the relation \\nbh$= 1.5\\times 10^{-6}/$(M/$10^{7}$ M$_{\\odot}$) Hz, in the case of ``classical\"  Seyfert 1 galaxies. The excess variance of NGC~4051, a Narrow Line Seyfert 1 object, is larger than what is expected for its black hole mass and X--ray luminosity. This can be explained if its \\nbh\\ is 20 times larger than the value expected in the case of a ``classical\" Seyfert 1 with the same black hole mass. Finally, the excess variance vs X--ray luminosity correlation is a byproduct of the excess variance vs black hole mass correlation, with AGN accreting at $\\sim 0.1-0.15$ the Eddington limit.  These results are consistent with recent results from the power spectral analysis of AGN. However, as they are based on data from a few objects only, further investigation is necessary to confirm that there is indeed a ``universal\" power spectrum shape in AGN (in the sense that the value of the power spectrum parameters of most AGN will be distributed around the ``canonical\" slope, and amplitude values listed above). One way to achieve this is to determine the excess variance vs black hole relation more accurately, using data from many more objects. This will be possible in the near future, since it is easier to measure the excess variance of archival light curves than to estimate their power spectrum. The excess variance vs black hole relation can therefore play an important role in the study of the X--ray variability scaling with black hole mass in AGN. ",
        "introduction": "Since the beginning of the active galactic nuclei (AGN) X--ray variability studies it was noticed that more luminous sources show ``slower\" variations. Barr \\& Mushotzky (1986) were the first to show that the ``two-folding\" time-scale (i.e.  the time-scale for the emitted flux to change by a factor of two) was faster in lower luminosity objects. Later, Lawrence \\& Papadakis (1993) and Green \\etal\\ (1993), using the results from the power spectral density function (PSD) analysis of the {\\it EXOSAT} ``long looks\", showed that the PSD amplitude at a given frequency decreases with increasing source luminosity. Nandra \\etal\\ (1997) and George \\etal\\ (2000), using the so called ``excess variance\" (\\rms, i.e. the variance of a light curve normalised by its mean squared after correcting for the experimental noise) found an anti-correlation between excess variance and source luminosity.  Similar results were also presented by Leighly (1999) and Turner \\etal\\ (1999), who also found that a particular group of AGN, the so called ``Narrow Line Seyfert 1\" (NLS1) galaxies, show systematically larger variability amplitude when compared to the same luminosity ``classical\" AGN (i.e. Seyferts with predominantly broad permitted lines, BLS1s). Further progress in the study of the longer-term X--ray variability has been afforded by \\xte, which has allowed systematic observations over time-scales longer than before (i.e. months, years). Markowitz \\& Edelson (2001), using 300 day long \\xte\\ light curves of nine Seyfert galaxies, estimated their excess variance and found a significant anti-correlation between \\rms\\ and source luminosity even on these long time-scales.  Most of the previous studies have examined the dependence of the ``variability amplitude\" on the source luminosity.  However, during the last few years, the mass of the central black hole has been estimated for many AGN (see e.g. Woo \\& Urry 2002 for a recent compilation of AGN black hole mass estimates). It is therefore possible to investigate the dependence of the variability amplitude on the black hole mass (\\bhm), a fundamental property of AGN. Recent studies that have followed this approach, have shown that there exists a significant anti-correlation between \\rms\\ and \\bhm, using \\asca\\ light curves of BLS1s and NLS1s (Lu \\& Yu, 2001; Bian \\& Zao, 2003).  The primary goal of the present work is to investigate the relation between excess variance and \\bhm\\ using the \\rms\\ data of Markowitz \\& Edelson (2001), together with the recent \\rms\\ measurement of a radio-quiet quasar, namely \\pg\\ (Papadakis, Reig, \\& Nandra, 2003). Markowitz \\& Edelson (2001) have already shown that \\rms\\ is strongly anti-correlated with the source luminosity. It is important though to see if there exists a correlation between \\rms\\ and \\bhm\\ as well.  In this way, one can relate the observed X--ray variations to the physical system itself, since \\bhm\\ is an important physical parameter of AGN. Based on the fact that the integral of the PSD is equal to the variance of a light curve, the \\rms--\\bhm\\ relation is used in order to examine the intrinsic shape of the PSD in AGN, and how this correlates with \\bhm.  Obviously, a single \\rms\\ value can correspond to different power spectral forms. As a result, the \\rms--\\bhm\\ relation cannot provide us with a unique answer as to what is the intrinsic shape of PSD in AGN, and how this might correlate/change with \\bhm. However, {\\it assuming} a certain PSD shape, the study of the \\rms--\\bhm\\ relation can be a powerful tool in assessing the reality of the assumed PSD shape.  In this work, the \\rms--\\bhm\\ relation is used in order to examine the recent results from the power spectrum analysis of \\xte, \\xmm\\ and \\asca\\ light curves of AGN (e.g. Uttley, McHardy,\\& Papadakis, 2002; Markowitz \\etal, 2003). These results have increased significantly our knowledge of the intrinsic PSD shape of AGN. However, high quality PSD estimation is currently possible for a few objects only, due to the lack of a large number of data sets that meet the necessary requirements (i.e. high signal-to-noise, long, well sampled light curves). The results presented in this work are also based on the use of a small number of objects, which have high quality, long term \\rms\\ measurements. However, although the number of objects with accurate PSD estimation may not increase significantly in the near future, the study of the \\rms--\\bhm\\ relation can be extended soon to incorporate the use of \\rms\\ measurements from a large number of objects. For example, the are numerous AGN observations with \\xte\\ and \\asca, which last typically for $\\sim 1$ day (assuming a $\\sim 30-40$ ksec exposure time). Although the PSD estimation is difficult with these light curves, the measurement of their \\rms\\ is easier. Taking also into account the wealth of the new data that are currently provided by \\xmm, and {\\it CHANDRA}, it will be possible to define the \\rms--\\bhm\\ relation more accurately in the near future. Consequently, the \\rms--\\bhm\\ relation will probably play an important role in the study of the X--ray variability scaling with \\bhm\\ in AGN. ",
        "conclusions": "The main results of this work are the following:  \\noindent (i) There is a significant correlation between the long term excess variance of AGN and their \\bhm, in the sense that \\rms\\ decreases with increasing \\bhm.  \\noindent (ii) The \\rms--\\bhm\\ relation is consistent with the hypothesis of a universal PSD shape in AGN with two frequency breaks, \\nbh\\ and \\nbl. The PSD changes its slope from $-2$ to $-1$ at \\nbh, and from $-1$ to zero at \\nbl.  \\noindent (iii) The high frequency break, \\nbh, decreases with increasing \\bhm\\ as \\nbh$=1.7\\times 10^{-6} {\\rm M}_{7}^{-1}$ Hz. The low frequency break, \\nbl, is $\\sim$ 10 to 30 times smaller than \\nbh.  \\noindent (iv) The PSD peak amplitude in the $\\nu\\times P(\\nu)$ space is $\\sim 0.017$ in all objects.  \\noindent (v) The correlation between \\rms\\ and $L_{X}$ is a byproduct of the \\rms--\\bhm\\ correlation, with AGN radiating at $\\sim 10-15\\%$ of the Eddington luminosity.  \\noindent (vi) The \\rms\\ measurement of NGC~4051, a NLS1 object, is larger than what is expected for its \\bhm\\ and $L_{X}$. This result can be explained if its PSD has only one break (at frequencies higher than $\\sim 4\\times 10^{-8}$ Hz) which is located at a frequency $\\sim 20$ times higher than the frequency expected in the case of a BLS1 with the same \\bhm. An intrinsically higher PSD amplitude (e.g. see McHardy et al. 2003) may also contribute to the larger $\\sigma^{2}_{rms}$ of this source.  These results are based on \\rms\\ measurements of long light curves in the $2-10$ keV band. Since the power spectra of individual AGN appear to be energy dependent (e.g. Papadakis \\& Lawrence 1995, Nandra \\& Papadakis 2001, Vaughan \\etal\\ 2003), if there is indeed a ``universal\" PSD shape, this should also be energy dependent. For example, a slope of $-2$ should characterize the AGN power spectra above the high frequency break in the $2-10$ keV band only.  The results listed above are entirely consistent with the recent results from the power spectral analysis of \\xte, \\xmm, and \\asca\\ light curves. For example, in all cases where a high frequency PSD slope break has been detected, the slope of the PSD below and above this break is roughly equal to $-1$ and $-2$ (see e.g. Uttley \\etal, 2002; Markowitz \\etal, 2003). Furthermore, as discussed in Section 2.3, the \\nbh--\\bhm\\ relation determined from the \\rms--\\bhm\\ correlation is entirely consistent with the existing data.  In fact, the relation of $T_{break}(=1/$\\nbh)=\\bhm$/(10^{6.5}$ M$_{\\odot})$ days, of Markowitz \\etal\\ (2003), is in agreement with the results of the present work. Finally, the observed \\psdamp\\ values are also consistent with the value of 0.017, which is derived from the study of the \\rms--\\bhm\\ relation.  One of the most interesting results presented in this work is that the $2-10$ keV band PSD of BLS1s, at least, has a ``universal\" shape: there two break frequencies, the slope between them is $\\sim -1$, the slope above the high frequency break $\\sim -2$, the high frequency break is mass dependent, but the \\psdamp\\ is not. Instead it has a constant value of $\\sim 0.02$. Since this result is based on the measurement of \\rms\\ and \\bhm\\ of ten objects only (most of which are among the most frequently observed AGN in X--rays) it is then premature to accept that it holds for the Seyfert galaxies as a class. Even if it does, it is not expected that either the high frequency slope slope for example or \\psdamp\\ will be exactly equal to $-2$ and $0.02$ for each individual AGN. Instead, these values should be considered as ``average\" or ``typical\" values around which the respective values of all AGN will be distributed, in the same way that the energy spectral slopes of AGN are distributed around the ``canonical\" value of $\\Gamma\\sim 1.9$. Confirmation of this result can be achieved in two ways. First, by using the results from the PSD analysis of many AGN. This is the direct and most powerful method, as it will allow us to actually determine the distribution of the PSD slope and \\psdamp\\ values around the ``canonical\" values. However, its rather unlikely that the number of AGN with well studied PSDs will increase significantly in the near future.  What the present works shows clearly is that there is a second way, which involves the \\rms\\ estimation of many objects with known \\bhm. Although the \\rms\\ of an individual object cannot constrain its PSD shape in any way, the study of the \\rms--\\bhm\\ relation can be useful in order to investigate whether there is indeed a ``universal\" PSD shape in AGN.  Furthermore, contrary to the former possibility, the use of archival data of many objects for the detailed study of the \\rms--\\bhm\\ relation should be possible in the near future as discussed in the Introduction.  If indeed the PSD of BLS1s AGN has a ``universal\" shape (in the sense that was described above) then this result could constrain physical models proposed to explain the physical process responsible for the X--ray variations.  For example, models should be able to predict the \\nbh--\\bhm\\ relation, and explain the ``universal\" \\psdamp\\ value of $\\sim 0.02$ in AGN (perhaps in GBHs as well). As an example of how this can be achieved, I present below a simple ``exercise\", based on the assumption that the X-ray variations of AGN at frequencies between \\nbl\\ and \\nbh\\ are associated with accretion disc instabilities which cause variations either to the energy release in the innermost regions of the disc or to the soft photons input to the X--ray emitting corona.  As a first step, one should investigate the relevance of physical processes to the origin of X--ray variability by comparing the various physical time-scales in accretion disc to the PSD break time-scales. Let us consider a BLS1 with \\bhm$=10^{7}$ M$_{\\odot}$. According to the \\nbh--\\bhm\\ relation, the break time-scale for this object should be $T_{break}=6.8^{+7.6}_{-3.8}$ days. Let us also consider the case of a standard $\\alpha-$disc (Shakura \\& Sunyaev, 1973) around the central object, and let us assume that $T_{break}$ is associated with one of the characteristic time-scales of the disc, i.e. the orbital ($T_{orb}$), thermal ($T_{th}$), sound-crossing ($T_{sc}$), and viscous time-scale ($T_{visc})$.  Using the relations given by Treves, Maraschi, \\& Abramowicz (1988), one can estimate these time-scales at every radius from the central object. The results show that $T_{break}\\sim T_{orb}$ at $R \\sim 70-80 R_{S}$ (where $R_{S}$ is the Schwarzschild radius). However, this seems rather unlikely as most of the energy in the accretion disc is dissipated at smaller radii. $T_{break}$ could also correspond to $T_{sc}$ at $R\\sim 3-5R_{Sch}$ (assuming that $R\\sim 100 H$, where $H$ is the disc's scale height). However, if $T_{break}$ for a NLS1 with the same \\bhm\\ is $\\sim 20$ times faster, {\\it and} $T_{break}$ corresponds to the {\\it same} physical time-scale in both type of objects, then obviously this cannot be the sound-crossing time-scale. Furthermore, $T_{visc}$ is far too slow to account for $T_{break}$. Assuming that $\\alpha\\le 0.1$ (where $\\alpha$ is the viscosity parameter), and $R\\le H$, then $T_{visc} > 870$ days even at $R=3R_{S}$. Finally, $T_{break}$ could correspond to $T_{th}$ at $R\\sim 30 - 40 R_{Sch}$, for $\\alpha=0.1$. In the case of a NLS1 with the same \\bhm, $T_{break}$ could also correspond to $T_{th}$, but at a smaller radius ($\\sim 4-6 R_{S}$).  Viscosity variations which happen at large radii, and develop on the local $T_{visc}$, can affect the energy release of the accretion disc in its innermost part, and result in a PSD with slope $-1$ between \\nbl\\ and \\nbh\\ (Lyubarskii, 1997).  However, if the PSD breaks are associated with the thermal time-scales, then perhaps as the accretion rate increases at various radii and the disc becomes unstable, thermal instabilities set in. The disc temperature will increase, increasing the local flux, hence the soft photons input to the X--ray corona. These soft photon ``flares\" should then develop on the local thermal time-scale. Assuming an exponential increase of soft photons with time, the superposition of these flares with the different time-scales can lead to a PSD with a flat shape up to \\nbl$=1/T_{th}(R_{out})$, a $-1$ slope up to \\nbl$=1/T_{th}(R_{in})$, and $-2$ slope at higher frequencies (Lehto 1989), where $R_{out}$ and $R_{in}$ are the outer and inner radius of the disc region which is affected by the thermal instabilities. If $R_{out}=200 - 300 R_{S}$, and $R_{in}=30-40 R_{S}$, then the resulting PSD will be in agreement with the present results.  The fact that the $P(\\nu)\\times\\nu$ value between \\nbl\\ and \\nbh\\ is similar in all objects implies that the PSD integral between \\nbl\\ and \\nbh\\ (i.e. the excess variance associated with the components of those frequencies) remains roughly constant, irrespective of \\bhm. In the context of shot noise models where flares occur randomly with a constant rate of say $\\lambda$ flares per unit time, and have the {\\it same} shape and amplitude, then \\rms$\\propto 1/\\lambda$. Considering the case of soft photons flares, which develop on the local thermal time-scale, as \\bhm\\ increases, and $T_{th}$ increases accordingly, it is natural to expect that $\\lambda$ should decrease, as it should take longer time for the instability to develop and the accretion disc to ``relax\" before the next instability starts to build up. In this case we should expect \\psdamp\\ to increase with increasing \\bhm. However, at the same time, as \\bhm\\ increases, the flares should become ``smoother\" (because $T_{th}$ also increases). This fact should result in a decrease in \\rms. Perhaps then, because an increase of the \\bhm\\ causes a decrease of $\\lambda$ and, at the same time, ``smoother\" flares, and because these two effects affect the observed variability in the opposite way, \\rms\\ may turn out to be roughly constant in all objects. However, further work is needed in order to examine if such a model can also predict correctly the \\psdamp\\ value of $\\sim 0.02$, as suggested by this work.  The differences in the variability properties between NLS1s and BLS1s can be explained by the fact that \\nbh\\ is higher in a NLS1 than in a BLS1 with the same \\bhm. The same explanation has also been proposed recently by McHardy \\etal\\ (2003). In the framework of the picture presented above, which involves thermal instabilities at a certain region of the accretion disc, the different \\nbh\\ values should correspond to differences in the location of this region in the two type of objects. In the innermost region of BLS1s (i.e. at radii smaller $30 R_{S}$) the thermal instabilities could be stabilized by a physical mechanism (e.g. the formation of a hot wind). The same mechanism may not be sufficient to stabilize the innermost region of the NLS1s, perhaps because the accretion rate is larger in these objects. Whether NLS1s show a second PSD break at \\nbl\\ $\\sim 10-30$ times smaller than \\nbh\\ is not certain yet. The results of this work show that the long term \\rms\\ measurement of NGC~4051 is not consistent with the presence of a second break, in agreement with the results of McHardy \\etal (2003). As these authors point out, NGC~4051 is the first Seyfert galaxy which shows a PSD similar to the power spectrum of GBHs at high/soft state. However, the \\rms\\ measurement of MCG~-6-30-15 is consistent with the presence of a second break. Ark 564 also shows two breaks in its PSD (Pounds \\etal, 2001; Papadakis \\etal, 2002). Power spectrum analysis of more NLS1s or the study of the \\rms--\\bhm\\ relation of a large number of NLS1s is necessary to clarify this issue."
    },
    "0311/astro-ph0311567_arXiv.txt": {
        "abstract": "We present a kinematic model of the neutral-hydrogen in the spiral galaxy of the lens system PKS1830$-$211, based on a Multi Element Radio-Linked Interferometer (MERLIN) 1.4--GHz radio map and the integrated and redshifted 21--cm hydrogen absorption-line profile as measured with the Westerbork Synthesis Radio Telescope (WSRT). Degeneracies in the models do not allow a unique determination of the kinematic center and forthcoming deeper Hubble Space Telescope observations with the Advanced Camera for Surveys (ACS) are required to break this degeneracy. Even so, we measure the inclination of the hydrogen disk: $i=(17-32)^\\circ$, indicating a close-to face-on spiral galaxy.  The optical depth increases with radius over the extent of the Einstein ring, suggesting HI depletion towards the lens center. The latter could be due to star formation or conversion of HI in to molecular hydrogen because of a higher metalicity/dust content in the galaxy center. The neutral hydrogen optical depth gives $N_{\\rm HI}=2 \\times 10^{21}$~cm$^{-2}$ at $r=5.0 \\,h^{-1}_{70}$~kpc in the disk ($T_{\\rm s}=100$~K), comparable to local spiral galaxies.  Our study shows that planned new radio telescopes (i.e.~ALMA, LOFAR and SKA) -- which will discover large numbers of similar lens systems -- are powerful new tools to probe the internal structure, kinematics and evolution of spiral galaxies in detail to $z\\ga 1$, as well as their neutral-hydrogen (and molecular) content, thereby complementing studies of HI emission from spirals at $z \\la 1$ and damped Ly--$\\alpha$ systems to $z\\gg 1$. ",
        "introduction": "The gravitational lens PKS 1830--211 (e.g. Pramesh Rao \\& Subrahmanyan 1988; Subrahmanyan et al. 1990) is the brightest known radio lens in the sky. The source at $z$=2.51 (Lidman et al. 1999) is extended and lensed into an Einstein ring (e.g. Jauncey et al. 1991) by a high-redshift spiral galaxy at $z$ = 0.89 (Wiklind \\& Combes 1996; Gerin et al. 1997; Mathur \\& Nair 1997; Chengalur, de Bruyn \\& Narasimha 1999). In addition, a time-delay of 26$\\pm$5 days between the two images of the lensed core has been measured (Lovell et al. 1998; see also Wiklind \\& Alloin 2002). New monitoring data in hand is expected to yield an improved measurement of this delay (J.\\,Winn, private communications). The structure of the Einstein ring can be used to place constraints on the mass distribution of the lens (Kochanek \\& Narayan 1992; Nair, Narasimha, \\& Rao 1993; Wiklind \\& Alloin 2002).  Were it not that PKS 1830--211 is seen through the Galactic disk (i.e. $l=+12^\\circ$, $b=-5^\\circ$) -- leading to considerable optical extinction ($A_I=0.8$ mag) and a field crowed by stars (e.g. Winn et al. 2002; Courbin et al. 2002) -- this lens system would be well-suited to determine the Hubble Constant.  Due to the severe extinction it has proven difficult to precisely determine the center of the lens galaxy. A number of mutually inconsistent lens centers have been published over the years (e.g. Kochanek \\& Narayan 1992; Nair et al. 1993; Lehar et al. 2000). Most recently, using the same HST observations, two different interpretations of the lens center were published (Winn et al. 2002; Courbin et al. 2002). The inability to measure the lens center with the currently available observations severely limits the use of PKS\\,1830$-$211 to determine H$_0$. Forthcoming HST--ACS observations, however, will be able to settle this question.  Besides being interesting because of H$_0$, the lens galaxy in PKS\\,1830--211 is the highest-redshift spiral lens galaxy known. The brightness of the lensed source has allowed a unique study of the molecular (e.g. CO and OH) and atomic (i.e. HI) interstellar medium (ISM) content (Wiklind \\& Combes 1996; Gerin et al. 1997; Mathur \\& Nair 1997; Chengalur, de Bruyn \\& Narasimha 1999) in a spiral galaxy as seen $\\approx$8\\,$h_{70}^{-1}$~Gyr ago. Although only a single galaxy, it provides a test case for similar systems anticipated to be discovered with planned instruments such as ALMA, LOFAR and SKA.  The time base-line covered by these galaxies provide enough leverage to study the evolution of the mass-distribution and ISM of spiral galaxies in the redshift-regime where most of their evolution is expected to occur and also connects up to studies of damped Ly--$\\alpha$ systems (e.g. Kauffmann 1996; Mo, Mao \\& White 1998).  In this letter, we make a first attempt to model the kinematics of HI in this high-redshift spiral galaxy, improving upon a previous more simple analysis by Chengalur et al. (1999). One of the initial aims, i.e. the determination of an independent kinematic lens center, has proven too difficult given the current data, but constraints on the radial distribution of HI, its column density and kinematics are obtained.  \\begin{inlinefigure} \\vspace{0.3cm} \\begin{center} \\resizebox{0.95\\textwidth}{!}{\\includegraphics{fig1.eps}} \\end{center} \\figcaption{The MERLIN 1.4--GHz continuum maps of PKS\\,1830--211 (North is up and East is left). Contours start at 25 mJy/beam and increase by factors of two. Overplotted is the best-fit kinematic model for the lens center of Nair et al. (1993). The two crosses indate the lensed cores A (NE) and B (SW). The thin elliptical contours indicate constant optical depth increasing by factors of 2 outward. The optical depth contour passing through image B has $\\tau$=0.064.} \\end{inlinefigure} ",
        "conclusions": "We have modeled the neutral hydrogen (HI) distribution and kinematics in the spiral galaxy of the lens system PKS 1830$-$211 at $z$=0.89, using a MERLIN 1.4--GHz continuum map of the Einstein ring and an WSRT HI absorption line. The main numerical results are summarized in Sect.~4.  In general, we find that the HI disk is seen close-to face-on, as previously suggested (e.g. Winn et al. 2002). We find evidence for HI depletion toward the galaxy center and a column density of $N_{\\rm HI}\\approx 2\\times 10^{21}$~cm$^{-2}$ at 5\\,$h^{-1}_{70}$~kpc, comparable to those of low redshift spiral galaxies. Our data are insufficient to discriminate between the different lens centers as published by different authors and we have to await new HST--ACS observations.  We emphasize some cautionary notes, listed in Sect.4, regarding our results.  Even so, we believe that our analysis is a good first attempt to constrain the kinematics of neutral hydrogen in a spiral galaxy at relatively high redshift, within the current observational limits, and should be regarded as such.  Our analysis shows that the combination of HI observations with gravitational lensing -- in particular radio Einstein rings -- potentially provide a strong tool to study spiral galaxies to high redshifts.  A similar program to combine lensing and stellar kinematics of E/S0 galaxies out to $z\\sim1$ is currently underway (i.e. the Lenses Structure \\& Dynamics (LSD) Survey; e.g. Koopmans \\& Treu 2002, 2003; Treu \\& Koopmans 2002).  We anticipate that {\\sl spatially resolved} HI absorption line profiles and radio continuum images at the redshifted HI-line frequency can be obtained of many similar systems in the future with e.g. ALMA, LOFAR and SKA. This would resolve many of the issues outlined in the paper and provide us with an unprecedented probe into the internal structure of spiral galaxies at high redshifts. By probing spiral lenses over a range of redshifts, their evolution can be studied (e.g. Koopmans \\& de Bruyn 1999), providing a complementary tool to study spiral galaxies between those at low redshifts using HI emission and damped Ly-$\\alpha$ systems at very high redshifts.  We note, however, that none of the other studies can provide {\\sl two independent} constraints, i.e. from lensing and kinematics, on the same mass distribution of the spiral galaxy. Their combination can break degeneracies inherent to each technique separately. Our results, admittedly based on a single system, show that very interesting constraints can be obtained for spiral galaxies at high redshifts."
    },
    "0311/astro-ph0311421_arXiv.txt": {
        "abstract": "Models of multi-field inflation exhibiting primordial non-Gaussianity have recently been introduced. This is the case in particular if the fluctuations of a light field scalar field, transverse to the inflaton direction, with quartic coupling can be transferred to the metric fluctuations. So far in those calculations only the ensemble statistical properties have been considered. We explore here how finite volume effects could affect those properties. We show that the expected non-Gaussian properties survive at a similar level when the finite volume effects are taken into account and also find that they can skew the metric distribution even though the ensemble distribution is symmetric. ",
        "introduction": "The observations of Cosmic Microwave Background anisotropies and of the large-scale structure of the universe offer a window to the physics of the inflaton. For instance the detailed measurement of the shape of the power spectrum can give constraints on the shape of the inflaton potential. On the other hand the detection of Non-Gaussian metric fluctuations could signal the existence of effective couplings between the inflaton and other fields. There have been a series of observational progresses in view of constraining the deviation from Gaussianity using CMB data on large scales ($\\theta\\sim10^o$)~\\cite{large}, intermediate ($\\theta\\sim1^o$)~\\cite{intermediate} and small scales ($\\theta\\sim10'$)~\\cite{small}, as well as large scales structures~\\cite{lss}. The recent CMB data by WMAP were also used by Komatsu {\\em et al.}~\\cite{komatsu} that concluded that the CMB spectrum was compatible with Gaussianity on the basis of an analysis of the bispectrum. Similar analysis using Minkowski functionals~\\cite{colley} and 3-point correlation function~\\cite{gaztanaga} reached the same conclusions. More recently, an analysis of the WMAP data using measurements of the genus and its statistics~\\cite{park} conluded that the Gaussianity of the CMB field was ruled out at a 99\\% level. As stressed after, non-Gaussianity of primordial origin may still have escaped detection and further investigations are needed.  From a theoretical point of view, Gaussianity is a generic prediction of slow-roll single field inflation~\\cite{maldacena,bu1}. A series of models including feature in the inflationary potential~\\cite{features}, existence of seeds such as $\\chi^2$~\\cite{chi2}, axion~\\cite{axion} or topological defects~\\cite{topdef}, curvaton scenario~\\cite{curvaton}, varying inflaton decay rate~\\cite{decay} have been shown to be able to generate some primordial non-Gaussianity.  In Ref.~\\cite{bu1}, we proposed a general mechanism to produce non-Gaussianity in the adiabatic mode and explicit realizations of such models were presented in Ref.~\\cite{bu2}. The mechanism is based on the generation of non-Gaussian isocurvature fluctuations which are then transferred to the adiabatic modes through a bend in the classical inflaton trajectory. Natural realizations were shown to involve quartic self-interaction terms. The statistical properties of the resulting metric fluctuations were then shown to be the superposition of a Gaussian and a non-Gaussian contribution of the same variance, the relative weight of the two contributions being related to the total bending of the trajectory in field space. The non-Gaussian probability distribution function (PDF) was also computed and shown to be described by a single new parameter so that generically only two new parameters suffice in describing this class of models.  In order to infer constraints from e.g. CMB data on these kind of models predicting non-Gaussianity one needs at least to go through 2 steps. \\begin{enumerate} \\item One first needs a precise prediction of the statistical properties of the curvature in order to construct estimator adapted to the detection of this kind of non-Gaussianity. For instance the analysis Komatsu {\\em et al}.~\\cite{komatsu} uses the bispectrum and Minkowski functional and constrains a $\\chi^2$ deviation from Gaussianity. The gravitational potential was parametrized as $\\Phi=\\Phi_{\\rm L} + f_{\\rm NL}(\\Phi^2_{\\rm L} - \\left<\\Phi^2_{\\rm L}\\right>)$ where $\\Phi_{\\rm L}$ is the Gaussian linear perturbation of zero mean and it was concluded that $-58<f_{\\rm NL}<134$ at 95\\% CL. This constrains only a very peculiar type of non-Gaussianity and does not apply e.g. to the models of Refs.~\\cite{bu1,bu2}. \\item Even if the form of the PDF is predicted theoretically, one needs to investigate what is measured and how the measurements are related to the theoretical predictions. \\end{enumerate}  The goal of this article is to investigate the observational implications of the theoretical predictions of Ref.~\\cite{bu1}. It requires further investigations in particular on the effect of finite volumes of the survey while the detailed comparison of our model to existing dataset, beyond the scope of this work, is still left for further studies. In particular, we want to show that these finite volume effects do not suppress the predicted non-Gaussianity.  Section~\\ref{sec_pb} formulates the problem, focusing on the quantities that can actually be observed. In Section~\\ref{sec_mod}, we recall basics of the models we consider in this paper and give simple consequences. We then develop in Section~\\ref{sec_langevin} some arguments to understand the effect of the mean value of the non-Gaussian component on the size of the survey. Section~\\ref{sec_finitevol} is devoted to the computation of the lowest order measured cumulants. In particular, we will show that there appears a non vanishing skewness that is directly induced by the finite volume effects as stressed in our concluding remarks, Section~\\ref{sec_concl}. ",
        "conclusions": "In this article we have focused on the phenomenology of the non-Gaussianity generated in models developed in Refs.~\\cite{bu1,bu2}. Interestingly, whereas the metric perturbation statistics involve only two microscopic parameters respectively related to the weight of the non-Gaussian component and to its PDF, the finite volume effects imply that the statistical properties of any observational quantity will involve a third parameter. This new parameter arises from the fact that the mean value, $\\chib$, of the field on the size of the observable universe does not vanish a priori. Obviously $\\chib$ cannot be determined on the basis of any observation. As described in \\S~\\ref{sec_langevin}, it implies that the originally symmetric PDF can be skewed and that this skewness is directly proportional to $\\chib$ that needs to be considered as a new parameter of the PDF while dealing with observations.  These results open the way for more detailed phenomenological studies. To see how those properties translate for the temperature and local density fields is not easy. Such investigation will probably require numerical tools such that developed in Ref.~\\cite{nonGaussnum}.\\\\    {\\bf Acknowledgements}: We thank Nabila Aghanim and Tristan Brunier for discussions as well as Alain Riazuelo for early discussions on the project."
    },
    "0311/astro-ph0311617_arXiv.txt": {
        "abstract": "{The 4.62\\,$\\mu$m absorption band, observed along the line-of-sight towards various young stellar objects, is generally used as a qualitative indicator for energetic processing of interstellar ice mantles. This interpretation is based on the excellent fit with OCN$^-$, which is readily formed by ultraviolet (UV) or ion-irradiation of ices containing H$_{2}$O, CO and NH$_{3}$. However, the assignment requires both qualitative and quantitative agreement in terms of the efficiency of formation as well as the formation of additional products. Here, we present the first quantitative results on the efficiency of laboratory formation of OCN$^{-}$ from ices composed of different combinations of H$_{2}$O, CO, CH$_{3}$OH, HNCO and NH$_{3}$ by UV- and thermally-mediated solid state chemistry. Our results show large implications for the use of the 4.62\\,$\\mu$m feature as a diagnostic for energetic ice-processing. UV-mediated formation of OCN$^{-}$ from H$_{2}$O/CO/NH$_{3}$ ice matrices falls short in reproducing the highest observed interstellar abundances. In this case, at most 2.7\\% OCN$^{-}$ is formed with respect to H$_{2}$O under conditions that no longer apply to a molecular cloud environment. On the other hand, photoprocessing and in particular thermal processing of solid HNCO in the presence of NH$_{3}$ are very efficient OCN$^-$ formation mechanisms, converting 60\\%--85\\% and $\\sim$100\\%, respectively of the original HNCO. We propose that OCN$^-$ is most likely formed thermally from HNCO given the ease and efficiency of this mechanism. Upper limits on solid HNCO and the inferred interstellar ice temperatures are in agreement with this scenario. ",
        "introduction": "\\indent The chemical composition of interstellar ices acts as a tracer of the chemical and physical history of molecular clouds. Based on spectral absorption features in the 2--16\\,$\\mu$m region, the ices are found to consist of simple neutral molecules. In addition, observational evidence now exists for the presence of ions such as OCN$^-$, HCOO$^-$ and NH$_4^+$ (Grim et al. 1989; Allamandola et al. 1999; Novozamsky et al. 2001; Schutte \\& Khanna 2003). The formation mechanisms underlying their presence, however, are still uncertain for most species. Understanding these processes is of key importance to improve our knowledge of the ice history and the evolution of molecular clouds.\\\\ \\indent The 4.62\\,$\\mu$m (2165 cm$^{-1}$) laboratory feature of OCN$^-$ matches well an interstellar absorption band, observed towards various young stellar objects (YSO's). Ever since its first detection in the infrared spectrum of W\\,33\\,A by Soifer et al. (1979), its assignment has remained controversial. First laboratory studies by Moore et al. (1983) reproduced the interstellar feature by proton irradiation of interstellar ice analogs composed of H$_{2}$O, CO and NH$_{3}$. Subsequently, Lacy et al. (1984) produced a similar feature by Vacuum ultraviolet (VUV) photoprocessing of CO/NH$_{3}$ ice. This `XCN' feature was assigned to OCN$^{-}$ by Grim \\& Greenberg (1987), and Schutte \\& Greenberg (1997) strengthened the hypothesis by looking at the band shift using different isotopes. In addition, Demyk et al. (1998) showed that OCN$^{-}$ could also form from HNCO by simple acid-base chemistry with NH$_3$, serving as a proton acceptor. This process has recently been studied in more detail by Raunier et al. (2003a, b). Although HNCO has not been identified directly in interstellar ices, it is observed in `hot cores' surrounding massive protostars and is a likely precursor of OCN$^-$. In spite of the strong case for assigning the interstellar 4.62\\,$\\mu$m feature to OCN$^{-}$, many other carriers were proposed. A summary is given by Pendleton et al. (1999). Some carriers, like silanes, thiocyanates, isothiocyanates and ketenes, cannot be excluded but to date none are serious candidates.\\\\ \\indent The ease with which the interstellar 4.62 $\\mu$m feature can be reproduced in the laboratory by the energetic processing of simple ices led to the direct association of this band in interstellar spectra with the presence of energetic processes. This is, however, a tentative assumption. A careful interpretation of interstellar data requires a thorough quantitative study of the efficiency of the OCN$^{-}$ formation together with an exhaustive investigation of the formation of additional products. Moreover, renewed interest has stemmed from the possibility that CN-bearing species may get included in protoplanetary discs to become an important component of prebiotic matter. This adds to the importance in the quest to understand the underlying formation mechanism of one of its simplest compounds, OCN$^{-}$ (Tegler et al. 1995; Pendleton et al. 1999; Whittet et al. 2001).  \\\\ \\indent This study will look in detail at the formation efficiency of OCN$^{-}$ by UV- and thermal processing of interstellar ice analogs and the evolution of additional products under laboratory conditions. The outline of the paper is as follows. Sect. 2 describes the experimental procedure, leading to the results in Sect. 3. The initial precursor abundance and the influence of photon-energy distribution and ice-thickness on the efficiency of OCN$^-$ formation are addressed. Also alternative C- or N-sources to form OCN$^-$ other than CO or NH$_3$ are considered, namely, H$_2$CO, CH$_4$ and N$_2$. Special attention is paid to the formation of OCN$^{-}$ from HNCO or CH$_3$OH. Sect. 4 discusses the astrophysical implications of the results in relation to the 4.62\\,$\\mu$m `XCN' feature observed in ices towards various YSO's and aims to provide a more solid base for the use of this feature as a potential diagnostic of energetic processing of ices in space. Finally in Sect. 5 our conclusions are summarised. ",
        "conclusions": "We have presented a laboratory study on the efficiency of thermal versus UV-mediated OCN$^-$ formation to qualitatively investigate the interstellar 4.62\\,$\\mu$m feature. OCN$^{-}$ is easily formed by UV-photolysis of a classic ice matrix of H$_{2}$O/CO/NH$_{3}$ and also when CO is replaced by CH$_{3}$OH or HNCO. Nevertheless, it remains questionable whether sufficient UV-photons are present inside dense clouds. UV photolysis also requires high initial abundances of NH$_{3}$ to reproduce the optical depth of the 4.62\\,$\\mu$m feature, particularly towards YSO's where observed OCN$^{-}$ abundances exceed 1.2\\% with respect to H$_{2}$O. Moreover, when starting from a CO ice mixture, the amount of NH$_{3}$ remaining in the ice after photoprocessing is too high to be consistent with the conservative upper limit of 5\\% determined toward W\\,33\\,A. In a CH$_{3}$OH ice mixture a lower NH$_{3}$ content is needed but the UV fluence required is approximately 3$\\times$ higher with respect to a CO-containing ice. The most favorable OCN$^-$ formation route emerging from this study is the thermal processing of HNCO-containing ices. Provided that HNCO is present in interstellar ices, thermal processing is $\\sim$100\\% efficient and is consistent with the data in particular toward W\\,33\\,A, NGC 7538 IRS9 and GL 2136. Also the inferred interstellar ice temperatures along the line-of-sight towards other OCN$^-$-containing YSO's are in agreement with this route of formation.\\\\ In the case of sources for which OCN$^-$ has not been detected, little insight can be gained about its formation mechanism. They could be either too young (too little UV exposure) or colder than 10 K. In these cases, HNCO can be present but only in H$_2$O-rich ices at temperatures below 10 K. At higher temperatures, the presence of sufficient amounts of proton acceptors like NH$_3$ rapidly induces OCN$^-$ formation. A non-detection for OCN$^-$ would in that case indicate lack of HNCO or NH$_3$ formation. Future searches for weak OCN$^-$, NH$_3$ and HNCO features are needed to further constrain these chemical routes."
    },
    "0311/astro-ph0311171_arXiv.txt": {
        "abstract": "{We present the results of a three-dimensional Monte Carlo radiative transfer code for starless molecular cloud cores heated by an external isotropic or non-isotropic interstellar radiation field. The code computes the dust temperature distribution inside model clouds with specified but arbitrary density profiles. In particular we examine in detail spherical (Bonnor-Ebert) clouds, axisymmetric and non-axisymmetric toroids, and clouds heated by an external stellar source in addition to the general interstellar field.  For these configurations, the code also computes maps of the emergent intensity at different wavelengths and arbitrary viewing angle, that can be compared directly with continuum maps of prestellar cores.  In the approximation where the dust temperature is independent of interactions with the gas and where the gas is heated both by collisions with dust grains and ionization by cosmic rays, the temperature distribution of the gas is also  calculated. For cloud models with parameters typical of dense cores, the results show that the dust temperature decreases monotonically from a maximum value near the cloud's edge (14--15~K) to a minimum value at the cloud's center (6--7~K).  Conversely, the gas temperature varies in a similar range, but, due to efficient dust-gas coupling in the inner regions and inefficient cosmic-ray heating in the outer regions, the gradient is non-monotonic and the gas temperature reaches a maximum value at intermediate radii.  The emission computed for these models (at 350~$\\mu$m and 1.3~mm) shows that deviations from spherical symmetry in the density and/or temperature distributions are generally reduced in the simulated intensity maps (even without beam convolution), especially at the longer wavelengths. \\keywords Radiative transfer; ISM: clouds, dust, extinction} ",
        "introduction": "Understanding the structure of pre--protostellar cores is an essential step towards an understanding of protostellar evolution. The density distribution immediately prior to the onset of gravitational collapse defines the initial conditions for collapse and hence one has a strong motivation to attempt to derive this density distribution from the observations.  Since these objects are most commonly observed by means of their millimeter dust emission, one has a strong interest in understanding the temperature distribution of the dust in the pre--protostellar core and its influence on the emergent intensity distribution at millimeter-submillimeter wavelengths.  In this article, we present a radiative transport code which has the interpretation of mm-submm maps of pre--protostellar cores as its main goal.  The approach which we have adopted is influenced by the fact that, as a rule, the density structure of these objects shows evidence for large deviations from spherical symmetry.  More precisely, the maps of the millimeter dust emission show clear departures from circular symmetry (see e.g.  Andr\\'e et al.~2000; Caselli et al.~2002a,b; Tafalla et al. 2002). In some cases there is also evidence for polarization and hence for a magnetic field (Ward-Thompson et al.~2000).  It seems likely to us that such behavior is caused by magnetic fields of energy densities sufficiently large to influence the core structure (see also Shu et al. 1987) and cause flattening along field lines (Basu 2000, Jones \\& Basu~2002).  We can attempt to infer the density distribution on the basis of mm-submm continuum maps if we can derive the temperature distribution in regions where departures from spherical symmetry are important. A first attempt in this direction was made by Zucconi et al.~(2001, hereafter ZWG, see also Evans et al.~2001 for the spherically symmetric case) but these authors neglected heating of dust grains due to re-absorption of photons emitted by the cloud itself.  Here we present a Monte--Carlo code without this limitation but which is capable of handling a variety of geometries.  We expect that this type of analysis will be especially useful as a tool for the interpretation of the high class data we expect to come from future instruments like ALMA and HERSCHEL.  Some high quality data are already available and are consistent with a picture in which the dust temperature of pre-protostellar cores decreases towards the center (Ward-Thompson et al.~2002, Bianchi et al.~2003).  ISO observations in the mid-IR (Bacmann et al.~2000) and in the far-IR (Ward-Thompson et al.~2002) suggest dust temperature gradients consistent with heating from the external interstellar radiation field (ISRF). The available models predict a factor of $\\sim 2$ increase in temperature from center to edge (ZWG, Evans et al.~2001, Andr\\'e et al.~2003), with the gradient dependent on the cloud structure.  The predicted emission of several cores is in  reasonable accord  with the data, and implies that the assumed equilibrium configurations are either unstable or maintained by a magnetic field.  Additional observational constraints on the thermal structure of prestellar cores are given by spatially resolved measurements of the gas temperature from NH$_3$ observations. Low-mass cores show fairly uniform gas temperature (e.g. Tafalla et al.~2002 for L1517B and L1498), whereas massive quiescent cores in Orion show significant temperature drops from edge to center (Li et al.~2003). Thus another check on our understanding of core structure is possible if radiative transfer models are also able to predict the gas temperature distribution resulting from the balance of the relevant heating and cooling mechanisms.  We therefore have also developed the capability to predict gas temperature distributions for our model cores.  While techniques have been developed for the study of radiative transfer in cold cores in one dimension (e.g. Leung~1976, Rowan-Robinson~1980, Ivezic et al.~1997), their extension to three-dimensional geometries is not straightforward. A Monte Carlo technique has the advantage that it deals easily with any general geometry, independently of the density distribution and the anisotropy of the incident radiation field. The increasingly faster workstations available render this technique very useful, and fully 3-D Monte Carlo radiative transfer codes have been developed recently (e.g. Wolf et al.~1999, Niccolini et al.~2003). However, in these latter studies the radiative sources are internal, and this is a computationally distinct problem from a core heated from the outside for two main reasons. First, as we are interested in cold cores, it is reasonable to expect that they will be optically thin to their own emission, and therefore, for methods that use iteration, the temperature distribution quickly relaxes to the final value; alternatively, for methods that follow individual packets until they exit the domain, each individual packet will interact with the dust often only once. Second, as packets are launched from the outside, a wide range of spatial scales demands a prohibitively large photon statistics, as the probability of a packet being launched in the direction of the innermost cells becomes increasingly small.  In general, our paper confirms and supports the results recently obtained for embedded prestellar cores by Stamatellos \\& Whitworth~(2003) with an independent but similar Monte Carlo radiative transfer code. While Stamatellos \\& Whitworth~(2003) focus on the important effect of the ambient medium in which (spherical) cores are embedded, our paper mostly emphasizes the consequences of deviations from spherical symmetry in the density distribution and anisotropies of the ISRF. The two papers represent therefore complementary attempts to model radiative transfer in molecular cloud cores in less idealized situations than those considered so far.  The structure of this paper is as follows. In Sect. 2, we set the general problem of radiative transfer in a dusty cloud. In Sect. 3, we describe the Monte Carlo method we have used. In Sect. 4, we present our choice of opacities and interstellar radiation field and briefly discuss their effect in the results. In esections 5 and 6 we present our results for two- and three-dimensional models, respectively. Finally, in Sect. 7 we summarize our conclusions. ",
        "conclusions": "In this paper we have presented a Monte Carlo model for radiative transfer applicable to the study of the physical conditions in starless molecular cloud cores. The method described is able to compute accurately and efficiently the dust and gas temperature distribution in clouds of specified but arbitrary density profiles heated by the ISRF, an external stellar source, or a combination thereof.  In addition, the method allows the computation of synthetic maps of the emission at near-infrared and submillimeter wavelengths for arbitrary viewing angles, and to predict the emergent spectral energy distribution.  We have successfully tested our code by detailed comparison with one-dimensional calculations in spherical symmetry (ZWG, Galli et al.  2002) and have shown that the assumption of effective optical thinness is a good one for prestellar cores.  As preliminary applications of our code, we have considered the radiative transfer in model clouds of increasing degrees of asymmetry and in a spherical cloud core heated by a neighboring star.  With our method, we confirm the results of ZWG, Evans et al.~(2001), and Stamatellos \\& Whitworth~(2003) that the dust temperature in cloud models with physical characteristics typical of dense cores varies monotonically between a minimum value at the center (6--7~K) and a maximum value near the cloud's edge (14--15~K) for the standard ISRF.  In general, maps at submillimeter wavelengths closely reproduce the column density distribution, even for the most non-isothermal configurations. In the Rayleigh-Jeans regime, in fact, temperature differences of a factor $\\sim 2$ along the line-of-sight are washed out by the much larger variation of the density.  However, the decrease of the dust temperature at the cloud center has the consequence that submillimeter maps sample preferentially the low-density envelope of prestellar cores, and are therefore unable to distinguish centrally peaked (pivotal) configurations from more flattened density profiles. We also note  that our three-dimensional model based on the work of Galli et al.~(2001) produces in natural fashion the cometary structure observed towards many prestellar cores.  For a spherical core heated by an external stellar source close to the cloud's surface, the temperature distribution is non-spherical, but the submillimeter emission generally reproduces the core's column density profile.  Deviations from circular symmetry become extremely apparent on the other hand shortward of 100~$\\mu$m where one samples mainly the heated surface layer."
    },
    "0311/astro-ph0311492_arXiv.txt": {
        "abstract": "{In the present paper we present $JHK$ photometric results of the young embedded cluster NGC\\,2316. We construct the cluster radial profile from which we determine a radius of 0.63\\,pc. We find 189\\,$\\pm$\\,29 cluster members in an extinction limited sub-sample of the survey, 22\\,$\\pm$\\,19 of which are possibly substellar. An average extinction of 4.5 visual magnitudes is derived using $(H - K)$ colours of control fields. This extinction is due to the presence of residual parental molecular cloud. NGC\\,2316 presents 16\\% source fraction of excess emission which is consistent with other results from clusters with an age of 2 -- 3\\,Myr. This age is consistent with the distribution of sources in the colour-magnitude diagram when compared to theoretical isochrones, and the overall shape of the cluster KLF. The substellar population of the cluster is similar or smaller than that observed for other embedded clusters and the stellar objects dominate the cluster membership. ",
        "introduction": "Embedded clusters may be the fundamental units of star for\\-ma\\-tion (Lada \\& Lada, 2003). They contain statistically si\\-gni\\-fi\\-cant samples of young stars of similar age and composition spanning a wide range of mass, providing excellent la\\-bo\\-ra\\-to\\-ries for investigating a number of important issues in star formation such as the form and universality of the initial mass function (IMF) and the frequency and lifetimes of protoplanetary disks. Surveys for near infrared (NIR) excess in embedded clusters can provide knowledge of the disk frequency and evolution because a protoplanetary disk is easier to detect than a planetary system with a similar mass of solid material (Beckwith \\& Sargent, 1996). In addition, embedded clusters provide the smallest spatial scale for investigating the nature of the IMF, its form and universality; their mass func\\-tion is in fact an initial mass function since young clusters have not lost significant numbers of members to either dynamical or stellar evolution.\\\\NGC\\,2316 is a young partially embedded cluster with coordinates $(\\alpha,\\delta)(J2000)=(6^\\mathrm{h}59^\\mathrm{m}40^\\mathrm{s},-7\\degr46\\arcmin36\\arcsec)$ at a distance of 1.1\\,kpc according to Felli et al. (1992) and Hodapp (1994). There is an HII region centered on a B3 ZAMS star (Noguchi et al., 1993; Fukui et al., 1993),  associated with the \\object{IRAS 06572-0742} source. The UV radiation is apparently creating a dense photodissociated region (Ryder et al., 1998), showing some physical signatures of ongoing star formation, such as $H_2O$ masers (Felli et al., 1992) and $CO$ outflow with spectral index characteristic of optically thin free-free emission (Beltr\\'an et al., 2001). NGC\\,2316 was first uncovered by Parsamian (1965), followed by other studies (e.g. Hodapp, 1994), however certain parameters are still not well established. To better characterize this cluster we obtained deep multi wavelength NIR images of the region to better constrain the cluster's age and spatial distribution, as well as provide estimates of membership number, fraction of NIR excess sources and brown dwarf po\\-pu\\-la\\-tion. ",
        "conclusions": "Circumstellar disks emit at NIR wavelengths. Studying the NIR excess emission of young cluster sources provides thus information on the frequency of circumstellar disks. Our estimate of the NIR excess emission for NGC\\,2316 according to the colour-colour diagram is 16\\% $\\pm$3\\%.\\\\ We conclude that NGC\\,2316's members span a wide range of masses, from B--stars down to the substellar regime (40\\,M$_\\mathrm{Jup}$), and is significantly embedded ($\\langle A_\\mathrm{V}\\rangle\\,=\\,4.5$) in its parental molecular cloud. We estimate a size of 0.63\\,pc for NGC\\,2316 from the radial profile assuming a distance of 1.1\\,kpc. Using an extinction limited sub-sample of the survey, we derive a membership for the entire cluster of 189\\,$\\pm$\\,29 objects, of which 22\\,$\\pm$\\,19 are beyond the HBL.\\\\ Regarding the crude estimate of the substellar population we note that studies of clusters at distances around and beyond 1\\,kpc will be hampered by an overwhelming background contamination at the faintest magnitudes (that cannot be reduced by deeper observations). Nevertheless, the estimated fraction of substellar objects in NGC\\,2316 (2-22\\,\\%) is comparable, within the errors, to that derived for the Trapezium cluster (20-25\\,\\%\\, Muench at al., 2002) and IC\\,348 (14-20\\,\\%\\, Muench et al., 2003) from a similar KLF analysis. This is additional evidence that brown dwarfs do not make a significant contribution to the overall cluster, either by numbers or mass. In other words, the population of these clusters is dominated by stellar rather than substellar members.\\\\ The NIR excess emission obtained for NGC\\,2316 from $JHK$ photometry is likely  an underestimate. Significant emission of a circumstellar disk is in the $L$--band (3.4\\,$\\mu$m), arising from warm circumstellar material within a few stellar radii, so $JHKL$ photometry would provide a more complete census of excesses, (\\cite{Liu}). IC\\,348, a cluster of a\\-ppro\\-xi\\-ma\\-tely 3\\,Myr, presents a fraction of NIR excess sources of 21\\% using $JHK$ photometry and 65\\% using $JHKL$ pho\\-to\\-me\\-try, (\\cite{haisch,lada3}). On the other hand, NGC\\,2316 has a bright UV source, a central B3 star, which might accelerate the dissipation of circumstellar disks by photoevaporation and consequently decrease the number of observed sources with NIR excesses.\\\\ The comparison of the NGC\\,2316, IC\\,348 and Trapezium KLFs in section 4.4 shows good agreement between their shapes; in particular, the first two are very similar (apart from the number of sources), as they roughly show the same broad peak at the same absolute magnitudes. Assuming an universal IMF, (e.g. \\cite{Lada2}), we infer that NGC\\,2316 has approximately the same age as IC\\,348, about 2--3\\,Myr (Lada \\& Lada, 1995; Muench et al., 2003), but is more evolved and older than the Trapezium."
    },
    "0311/astro-ph0311201_arXiv.txt": {
        "abstract": "A new approach to simulate fluorescence photons produced in extensive air showers is described. A Monte Carlo program based on CORSIKA produces the fluorescence photons for each charged particle in the development of the shower. This method results in a full three-dimensional simulation of the particles in a shower and of the fluorescence light generated in the atmosphere. The photons produced by this program are tracked down to a telescope and the simulation of a particular detector is applied. The differences between this method and one-dimensional approaches are quantitatively determined as a function of the distance between the telescope and the shower core. As a particular application, the maximum distance at which the Auger fluorescence telescope might be able to measure the lateral distribution of a shower is determined as a function of energy. ",
        "introduction": "Introduction}  There are several questions yet to be solved about cosmic rays with energy above $10^{18} \\ eV$. The most exciting topics concern the existence of a cut off in the energy spectrum, the possible anisotropy of the sources, the chemical composition and the propagation mechanisms.  Particles with these energies arrive on Earth with a very low flux and therefore can only be studied by the detection of the extensive air showers (EAS) they produce in the atmosphere. The present HiRes \\cite{bib:hires} and the AGASA \\cite{bib:agasa} experiments have contributed to our understanding of the nature of high energy cosmic rays and future detectors (Auger \\cite{bib:auger-nim}, EUSO \\cite{bib:euso-artigo}, OWL \\cite{bib:owl-artigo} and  Telescope Array \\cite{bib:telescope:array}) are very promising due to the increase in collection area and the use of new technologies.  The fluorescence technique has been used successfully by the HiRes experiment and is also going to be used by the new generation of detectors: Auger, EUSO, Telescope Array and OWL. Fluorescence detectors measure the EAS by counting, as a function of depth, the photons produced by the excitation of nitrogen molecules due to the passage of the charged particles in the shower. The determination of the longitudinal profile and the time the photons have hit the detector allow the reconstruction of the energy and direction of the primary particle.  Monte Carlo simulations have been widely used by cosmic ray physicists to optimize their detectors and to determine reconstruction methods. One of the most popular and tested programs known in the cosmic ray community is CORSIKA \\cite{bib:corsika}. CORSIKA simulates the development of the particles in a shower in a high level of sophistication including three-dimensional particle transport, electromagnetic interactions, decay of unstable particles, multiple scattering, deflection in the Earth's magnetic field and several high energy hadronic interaction models. Comparisons of CORSIKA simulations and experimental data can be found in \\cite{bib:corsika-teste}.  A shower with energy $10^{19} \\ eV$ has more than $10^{10}$ particles and a Monte Carlo simulation would take a long time to follow all of them. Hillas \\cite{bib:thin} has overcome this problem by using a thinning algorithm in Monte Carlo simulations. This procedure follows only a fraction of the particles with energy below a certain limit and conserve the total energy in the shower. The use of this method has reduced the computational time to practical limits without decreasing significantly the quality of the simulations.  In contrast to the simulation of particles in a shower, the simulation of the fluorescence light produced by the passage of these particles in the air has been done with almost no detail. Calculations made so far were done based on one-dimensional simulations of the fluorescence photons following an average function given by a Gaisser-Hillas \\cite{bib:gaisser-hillas} profile. Apparently, the crude treatment given to the simulation of the fluorescence component was justified by the optics and electronics of the present generation of telescopes. However, for the future detectors this justification is certainly not going to be valid.  In reference \\cite{bib:sommers_size}, Sommers has made analytic calculations to show the possibility of detecting the lateral distribution of a shower with a fluorescence detector. His ideas have been exploited by G\\'ora et al. \\cite{bib:gora_size} using simulations based on a Gaisser-Hillas longitudinal profile and a NKG \\cite{bib:nkg,bib:nkg_crsk} lateral distribution function.  More recently, other groups have announced new fast simulations in development \\cite{bib:seneca,bib:casc1d} which may be used to produce fluorescence photons. These programs are based on a combination of numerical and Monte Carlo approaches and special attention must be taken in order not to reduce the intrinsic shower fluctuations.  The work presented here explains how we generate fluorescence photons in the CORSIKA framework. Our approach allows the simulation of this component of the shower with great detail. For each charged particle produced in the shower development, the corresponding fluorescence light is produced from the ionization energy loss and the are tracked to the detectors. Atmospheric attenuation may be applied and the particles continue to be followed by CORSIKA down to the observational level resulting in a full hybrid (fluorescence + particles) simulation. In particular, this simulation can be useful for the Auger Collaboration due to its hybrid design.  The atmospheric attenuation is done by a simple implementation that simulates the extinction of photons according to particular atmospheric properties. The user is allowed to implement different atmospheric extinction coefficients. Excluding the showers simulated in section \\ref{sec:certification} where no attenuations has been applied, all simulations in this paper were done with a midlatitude US Standard Atmosphere and scattering of light by aerosols and molecules was taken into account. The aerosol horizontal attenuation length and vertical scale height were set to $12 \\ km$ and $1.2 \\ km$, respectively.  The mean free path for Rayleigh scattering was set to $2974 \\ g/cm^2$.  Scattering of the fluorescence photons is not simulated in the present version which could be an important effect for some specific geometries of detection and therefore should be accounted for in future implementations.  Historically, the main advantage of using Monte Carlo simulations comparing to analytic models has been considered to be the possibility to reproduce the event by event fluctuations. Nevertheless, the development of the softwares has extended the differences beyond that, among which, we stress the possibility to inject almost all kind of primary particles, the detailed description of the atmosphere and the treatment of each shower component (hadronic and electromagnetic) in an independent way with parameters updated from the accelerator experiments.  Besides that, the simulation we has developed have some other improvements if compared to analytic or one-dimensional models regarding the fluorescence photons production. In a deeper level of analysis, the approach we present here is able to simulate the particle by particle fluctuation in the production of the fluorescence photons, since each charged particles is treated independently.  Moreover, the time information is available for each single photon taking into account its exact position in the atmosphere. This feature would allow for the first time a detailed study of the snapshots proposed in \\cite{bib:sommers_size} with a good raytracing resolution.  Recently proposed experiments \\cite{bib:fy:campinas,bib:fy:airfly} are planning to measure the fluorescence yield in  more details than done so far by Kakimoto et al. \\cite{bib:kakimoto}. The fluorescence yield varies with the particle energy and may be completely different for each particle type. The simulation presented here can easily take into account the dependencies of the yield as a function of energy and particle type for each particle in the shower after they have been measured by the proposed experiments. This great level of detail would be practically impossible to be accounted for in a analytical model or in an one-dimensional Gaisser-Hillas approach.  In this paper, we describe the simulation, test it against old and simple approximations and illustrate the use of the program by calculating the lateral size of the shower as a function of the distance to a ground based telescope. This application shows that the lateral distribution of showers with energies $10^{18}$, $10^{19}$ and $10^{20} \\ eV$ measured by detectors similar to the Auger telescopes configuration can not be neglected when distances from the detector to the shower axis are shorter than $4 \\ km$, $8 \\ km$ and $14 \\ km$, respectively. ",
        "conclusions": "We have presented here a new approach to simulate fluorescence photons in an atmospheric air shower. A three-dimensional and fully hybrid (particles+photons) simulation tool was developed under the CORSIKA framework. In our routines, fluorescence photons were produced for all charged particles in a shower and for electrons and positrons with energy below the CORSIKA cut-off.  This program is going to be useful for the new generation of detectors (Auger, EUSO, Telescope Array and OWL) due to the improvement in the technology used by them.  As shown here, for the new observatories the shower can not be considered one-dimensional in all cases and therefore a three-dimensional approach should be used to optimize the construction and in the elaboration of reconstruction methods.  In this paper, the number of photons produced along the longitudinal development of the shower and the arrival time of the photons in the detector determined by our program have been compared to analytical calculations based on the energy deposited profile given by CORSIKA.  A good agreement can be seen in figures \\ref{fig:perfil} and \\ref{fig:time}.  The famous $320 \\ EeV$ shower detected by the Fly's Eye experiment was also simulated according to the geometry given in \\cite{bib:fly51joule}. The consistency of our simulation can be seen as a test of the approach proposed here from one hand and as confirmation of the profile reconstructed by them from the other hand.  Our code was approved in several comparisons of simulated and real data. The arrival time, the longitudinal profile and the energy determined by our simulation show a very good agreement to what has been published so far.  The fluctuation due to thinning algorithm introduced in section \\ref{sec:simulation} has been tested in respect to its influence on the lateral distribution of particles. No effect of the thinning was found in a range from $10^{-4}$ to $10^{-7}$ of the thinning factor when detectors similar to the Auger telescopes configuration are used.  Applications of the code were performed. We determined the sensitivity of a fluorescence telescope of the kind used by the Auger collaboration to the lateral distribution of the particles in a shower. Such a telescope could distinguish a real shower from a one dimensional approximation for distances of the core closer than $4 \\ km$ for $10^{18} \\ eV$ showers, $8 \\ km$ for $10^{19} \\ eV$ showers and $14 \\ km$ for $10^{20} \\ eV$ showers.  These results represent a limit on the possibility to measure the lateral distribution of particles with a fluorescence telescope based on time information as suggested by \\cite{bib:sommers_size}. However, a great number of showers with energy beyond $10^{19} \\ eV$ with core closer $8 \\ km$ can be detected and reconstructed by the present and future fluorescence telescopes. For this subset of events we suggest a detailed analysis of the data in order to possibly measure the lateral distribution of the particles."
    },
    "0311/astro-ph0311215_arXiv.txt": {
        "abstract": "{Recent UV observations of the globular cluster NGC~2808 (Brown et al. \\cite{brsw01}) show a significant population of hot stars fainter than the zero-age horizontal branch (``blue hook'' stars), which cannot be explained by canonical stellar evolution. Their results suggest that stars which experience unusually large mass loss on the red giant branch and which subsequently undergo the helium core flash while descending the white dwarf cooling curve could populate this region.  Theory predicts that these ``late hot flashers'' should show higher temperatures than the hottest canonical horizontal branch stars and should have helium- and carbon-rich atmospheres. As a test of this late hot flasher scenario, we have obtained and analysed medium resolution spectra of a sample of blue hook stars in NGC~2808 to derive their atmospheric parameters.  Using the same procedures, we have also re-analyzed our earlier spectra of the blue hook stars in $\\omega$ Cen (Moehler et al. \\cite{mosw02}) for comparison with the present results for NGC~2808.  The blue hook stars in these two clusters are both hotter (\\teff $\\ge$35,000~K) and more helium-rich than canonical extreme horizontal branch stars in agreement with the late hot flasher scenario. Moreover, we find indications for carbon enhancement in the three most helium-enriched stars in NGC~2808.  However, the blue hook stars still show some hydrogen in their atmospheres, perhaps indicating that some residual hydrogen survives a late hot flash and then later diffuses to the surface during the horizontal branch phase. We note that the presence of blue hook stars apparently depends mostly on the total mass of the globular cluster and not so much on its horizontal branch morphology. ",
        "introduction": "Low-mass stars burning helium in a core of about 0.5~\\Msolar\\ and hydrogen in a shell populate a roughly horizontal region in the colour-magnitude diagrams of globular clusters, which has earned them the name ``horizontal branch\" (HB) stars.  The Galactic globular clusters show a great variety in horizontal branch morphology, i.e., in the temperature distribution of their HB stars.  The temperature of an HB star depends -- at a given metallicity -- on the mass of its hydrogen envelope, with the hottest or extreme HB (EHB) stars (\\teff\\ $>$ 20,000~K) having extremely thin  $\\lesssim$0.01\\Msolar) envelopes. The increase in the bolometric correction with increasing temperature turns the blue HB into a vertical blue tail in optical colour-magnitude diagrams with the faintest blue tail stars being the hottest and least massive (see Fig.~\\ref{cmd} for a prominent example of a blue tail). It is still unclear how some stars manage to lose nearly all of their envelope mass and still undergo the helium core flash, as indicated by the long blue tails in a number of globular clusters (e.g., NGC~6752, Moehler et al. \\cite{mosw00}). The globular cluster NGC~2808 represents a good template cluster, since it has an extremely long blue tail extending to very high temperatures and correspondingly low envelope masses.  NGC 2808 has been the subject of four recent deep ground-based and Hubble Space Telescope (HST) photometric studies (Sosin et al.\\ \\cite{sosi97}, Walker \\cite{walk99}, Bedin et al.\\ \\cite{bepi00}, Brown et al.\\ \\cite{brsw01}).  The ground-based studies have shown that the blue tail of NGC~2808 extends to extremely faint magnitudes ($M_V \\approx$ \\magpt{5}{5}), with gaps at $M_V \\approx$ \\magn{3} and \\magpt{4}{5}. These gaps can be seen at V $\\approx \\magpt{18}{5}$ and $\\magpt{20}{0}$ in Fig.~\\ref{cmd}. The brighter gap is thus at a similar position as the underpopulated region in the colour-magnitude diagram of NGC~6752, which separates the EHB stars from the classical blue HB stars (Moehler et al.  \\cite{mohe97b}, \\cite{mosw00}; Momany et al. \\cite{mopi02}), while the fainter gap coincides with the hot end of the blue tail in NGC~6752.  The blue tail in NGC~2808 thus extends to hotter temperatures than that of NGC~6752.  That in itself would not be a problem, but the spectroscopic analyses of Moehler et al. (\\cite{mohe97b}, \\cite{mosw00}) show that the blue tail stars in NGC~6752 already populate the EHB to the hot end predicted by canonical HB models.  \\begin{figure}[h] \\vspace*{9.cm} \\special{psfile=0505_1.ps hscale=45 vscale=45 hoffset=-20 voffset=-65 angle=0} \\caption{ Colour-magnitude diagram of NGC~2808 (Walker \\cite{walk99}) with our spectroscopic targets marked.  Our targets were chosen to span the gap at V = \\magpt{20}{0} which, according to the hot flasher scenario, separates the canonical EHB stars from the fainter blue hook stars. \\label{cmd}}  \\end{figure}  Observations of NGC~2808 in the far- and near-UV (Brown et al. \\cite{brsw01}, their Fig.\\ 3) show that the stars below the faint gap form a hook-like feature that extends up to \\magpt{0}{7} below the hot end of the zero-age HB (ZAHB). Such ``blue hook'' stars are so far known in only three other globular clusters: $\\omega$ Cen (D'Cruz et al. \\cite{dcoc00}, Moehler et al. \\cite{mosw02}), NGC~6388 and possibly NGC~6441 (Busso et al. \\cite{bupi03}), all of which are very massive.  Within the framework of canonical HB theory there is no way to populate this region of the UV colour-magnitude diagram without requiring an implausibly large decrease in the helium-core mass.  Thus canonical theory fails to explain both the faint UV luminosities and expected high temperatures of the blue hook stars.  Brown et al.\\ (\\cite{brsw01}) have proposed a scenario to explain the blue hook stars.  According to this scenario the stars brighter than the gap at $M_V \\approx$ \\magpt{4}{5} are canonical EHB stars which lost all but $\\sim$0.01\\Msolar\\ of their envelope mass prior to undergoing the helium core flash.  As discovered by Castellani \\& Castellani (\\cite{caca93}), stars which lose more mass than this will leave the red giant branch and evolve to high effective temperatures before undergoing the helium core flash, producing the so-called ``hot flashers\".  Sweigart (\\cite{swei97}) subsequently showed that hot helium flashes can occur either as a star evolves from the tip of the red giant branch to the top of the white dwarf cooling curve (``early\" hot flasher), or later as a star descends the white dwarf cooling curve (``late\" hot flasher).  For even higher mass loss, a star will die as a helium white dwarf without ever igniting helium.  These different evolutionary paths are illustrated in Moehler et al. (\\cite{mosw02}) and Sweigart et al. (\\cite{swbr02}).  D'Cruz et al. (\\cite{dcdo96}, \\cite{dcoc00}) proposed that the blue hook stars could be the progeny of such hot flashers, but unfortunately the D'Cruz et al. models were, at most, only $\\approx$\\magpt{0}{1} fainter than the canonical ZAHB, much less than required by the observations.  More recently, Brown et al. (\\cite{brsw01}) have explored the evolution of both the early and late hot flashers through the helium core flash to the EHB in more detail.  Their models show that a late hot helium flash on the white dwarf cooling curve will induce substantial mixing between the hydrogen envelope and helium core, leading to helium-rich EHB stars that are much hotter than canonical ones, as found previously by Sweigart (\\cite{swei97}).  This result has been confirmed by the calculations of Cassisi et al.  (\\cite{casc03}), who were able to follow the evolution of a late hot flasher completely through the helium core flash.  Brown et al.\\ (\\cite{brsw01}) suggest that such flash mixing may be the key for understanding the evolutionary status of the blue hook stars. Such mixing may also be responsible for producing the helium-rich, high gravity field sdO stars (Lemke et al. \\cite{lehe97}), whose origin is otherwise obscure.  The flash mixing scenario predicts a gap of about 6000~K between the canonical EHB stars (i.e., stars without flash-mixing, including early hot flashers) and the late hot flashers as well as a helium dominated atmospheric composition for the late hot flashers. The hydrogen deficiency results in a flux redistribution where more flux is emitted shortward of the Lyman edge at 912 \\AA\\, and less flux is emitted at ultraviolet (1500 \\AA) and longer wavelengths, which makes flash-mixed stars fainter than canonical EHB stars at these wavelengths.  Spectroscopic observations of blue hook stars in $\\omega$ Cen by Moehler et al. (\\cite{mosw02}) showed that these stars reach effective temperatures of more than 35,000~K, i.e.,  well beyond the hot end of the canonical EHB.  In addition, most of them show at least solar helium abundances with the helium abundance increasing with effective temperature, in contrast to canonical EHB stars such as those studied in NGC~6752 by Moehler et al. (\\cite{mosw00}). Contrary to the predictions of Brown et al. (\\cite{brsw01}), however, only one star showed a helium abundance of \\loghe\\ $>$ 0. These results may indicate that flash mixing is less efficient than assumed, or that some residual hydrogen survives flash mixing and then diffuses outward to the surface. It is unclear whether the lower than expected helium abundances can still reproduce the UV properties of the stars.  UV data have been presented for the blue hook stars in the central region of $\\omega$ Cen by D'Cruz et al. (\\cite{dcoc00}),  but due to crowding these stars are not accessible to ground-based spectroscopy.   UV data are available for the blue hook stars in the outer region of $\\omega$ Cen from the Ultraviolet Imaging Telescope (UIT, Stecher et al. \\cite{stec97}), but  unfortunately, these data are too noisy to permit a quantitative test.  Considering these somewhat ambiguous results and the fact that $\\omega$ Cen is not a typical globular cluster, we decided to test the late hot flasher scenario by observing  the more typical globular cluster NGC~2808. We discuss our observational data in Sect.~\\ref{sec-observ} and then derive the parameters of the blue hook stars (temperatures, gravities and helium abundances) in Sect.~\\ref{sec-analysis}.  In Sect.~\\ref{sec-discuss} we compare our results with the predictions of the flash-mixing scenario. ",
        "conclusions": "The high temperatures and high helium abundances reported here for the blue hook stars in NGC~2808 and $\\omega$ Cen provide general support for the flash-mixing scenario of Brown et al. (\\cite{brsw01}). The three most helium-rich stars in NGC~2808 also show evidence for carbon enrichment.  However, all of our targets show some hydrogen in their atmospheres, and only one target in each cluster has \\loghe $> 0$.  This is partially a selection effect, as our targets are preferentially the brighter blue hook stars, and the most helium-rich star in each cluster is either the faintest (W20334 in NGC~2808) or second-faintest (D10763 in $\\omega$ Cen).  However, it is also likely that some hydrogen survives a late hot flash, and subsequently diffuses to the surface during the HB phase.  Moehler et al. (\\cite{mohe97a}) reported an isolated helium-rich sdB star in M15 with the expected properties (\\teff\\ = 36,000~K, M$_V$ = \\magpt{+4}{7}, \\loghe = +0.82) of a flash-mixed star.  However, significant numbers of blue hook stars have been found only in the most massive Galactic globular clusters: $\\omega$ Cen (M$_V$ = \\magpt{-10}{29}), NGC~2808 (M$_V$ = \\magpt{-9}{39}), NGC~6388 (M$_V$ = \\magpt{-9}{42}), and possibly NGC~6441 (M$_V$ = \\magpt{-9}{64}), all data for M$_V$ from Harris (\\cite{harr96}). Other globular cluster in the same mass regime are 47~Tuc (M$_V$ = \\magpt{-9}{42}), NGC~2419 (M$_V$ = \\magpt{-9}{58}), and NGC~6715 (M$_V$ = \\magpt{-10}{01}). The optical colour-magnitude diagrams of NGC~2419 (Harris et al. \\cite{habe97}) and NGC~6715 (Momany, priv. comm., Rosenberg et al. \\cite{rore03}) show stars faint enough to be blue hook stars. The situation for 47~Tuc is unclear due to contamination by the Small Magellanic Cloud.  Apparently the number of blue hook stars depends on the total mass of the globular cluster and not on the number of other hot HB stars (see Rosenberg et al. \\cite{rore03} for an independent discussion of this effect). For instance, NGC~6752, which is a factor of 100 less luminous than NGC~2808 and has a very long and well populated blue tail, contains 59 EHB stars with \\magpt{3}{0} $\\le M_V \\le$ \\magpt{4}{5} (photometry of Buonanno et al. \\cite{buca86}), compared to 34 such stars observed in NGC~2808 (photometry from Walker \\cite{walk99}).  Even more EHB stars can be seen in the data of Momany et al. (\\cite{mopi02}) for NGC~6752.  However, while NGC~2808 has 28 blue hook stars with \\magpt{4}{5} $\\le M_V \\le$ \\magpt{5}{5}, {\\em no} such stars are found in NGC~6752.  Likewise, the UV colour-magnitude diagram of Brown et al. (\\cite{brsw01}) shows that the population of blue hook stars dominates over the population of EHB stars in NGC~2808 (46 vs. 29 stars, respectively).  As the mass loss ranges populating the hot end of the canonical extreme HB and the blue hook region are not too different, it is hard to understand why globular clusters with large numbers of canonical EHB stars do not show at least some blue hook stars as well."
    },
    "0311/gr-qc0311064_arXiv.txt": {
        "abstract": "It is known that  the imaginary parts of the quasi normal mode (QNM) frequencies for the Schwarzschild black hole are evenly spaced with a spacing that depends only on the surface gravity. On the other hand, for massless minimally coupled scalar fields, there exist no QNMs in the pure DeSitter spacetime. It is not clear what the structure of the QNMs would be for the Schwarzschild-DeSitter (SDS) spacetime, which is characterized by two different surface gravities. We provide a simple derivation of the imaginary parts of the QNM frequencies for the SDS spacetime by calculating the scattering amplitude in the first Born approximation and  determining its poles. We find that, for the usual set of boundary conditions in which the incident wave is scattered off the black hole horizon, the imaginary parts of the QNM frequencies have a equally spaced structure with the level spacing depending on the surface gravity of the black hole. Several conceptual issues related to the QNM are discussed in the light of this result and comparison with previous work is presented. ",
        "introduction": "It is well known that the quasi normal modes (QNMs, hereafter) are crucial in studying the gravitational and electromagnetic perturbations around black hole spacetimes (for comprehensive reviews and exhaustive list of references see \\cite{nollert99,ks99}). The QNMs also seem to have an observational significance as the gravitational waves produced by the perturbations can, in principle, be used for unambiguous detection of black holes. The early studies of the QNMs concentrated on numerical computations (see, for example, \\cite{vishveshwara70,press71}) as analytical solutions of the perturbation equations were difficult to obtain. Essentially, one solved the perturbation equations numerically for different initial conditions in a given black hole spacetime. It was found that the results were, in general, independent of the initial conditions and depended mostly on the parameters characterizing the black hole horizon. For example, for the Schwarzschild black hole, the spectrum of the QNMs was found to have the structure \\cite{nollert93,andersson93,hod98,ga03,cly03,xsws03} \\be k_n =   i\\kappa\\left(n+\\frac{1}{2}\\right)+\\frac{\\ln 3}{2\\pi} \\, \\kappa  + \\mathcal{O}[n^{-1/2}] \\label{eq:kn_num} \\e which depends only on the surface gravity of the black hole $\\kappa$. Similarly, for charged (Reissner-Nordstr\\\"{o}m) and rotating (Kerr) black holes, the spectrum was found to depend only on the parameters characterizing the horizon, i.e., the mass, charge and angular momentum of the black hole \\cite{detweiler80,mashhoon85,leaver85,leaver90,kokkotas91,aas94,omoi96,klpa97,onozawa97,bcko03,bk03,hod03,hod03b}.  The numerical studies were followed by a series of analytical calculations which were mainly based upon some approximation scheme, like for example, (i) approximating the scattering potential with some simple functions so that the problem becomes exactly solvable \\cite{fm84,bm84,molina03}, (ii) using WKB-like techniques \\cite{sw85,iyer87,iw87,ks88,gwks90,si90,gm92,sw92,leaver92b,al92,ah03,konoplya03}, (iii) Born approximation \\cite{mmv03,padmanabhan03b} and (iv) others \\cite{leaver86,pm87,mp89,ffah92,ns92,aas93,aas93b,ans93,lm96,clsy95,clsy95b,pz99,solodukhin03,birmingham03}. Nearly all of these schemes came with their own limitations and rarely reproduced the exact form of the QNM spectrum as given in equation (\\ref{eq:kn_num}) -- in particular, there was hardly any explanation for the $\\ln 3$ term. Analytical proofs of equation (\\ref{eq:kn_num}) were provided only recently, using the methods of continued fractions \\cite{leaver85,leaver92,nollert93,motl02} and also by studying the monodromy of the perturbation continued to the complex plane \\cite{mn03,neitzke03}.  While most of the analytical work concentrated on the real part of QNM (due to the interest in the $\\ln 3$ factor),  the structure of the imaginary part is equally intriguing and requires physical understanding. In particular, we need to understand  why  the imaginary part of $k_n$ in equation (\\ref{eq:kn_num}) has a simple, equally spaced structure, with $\\kappa$ determining the level spacing. Physical quantities with a quantized spectrum are always of interest when they have constant spacing. In the case of horizon area, for example, one can attempt to relate an equally spaced spectrum (obtained in several investigations, e.g., see \\cite{pp03} and references therein) to the intrinsic limitations in  measuring length scales smaller than Planck length \\cite{padmanabhan87}. But the uniform spacing of QNM frequencies is a purely {\\it classical} result and hence is harder to understand physically.  It turns out, however,  that one can make {\\it some} progress in this direction. There are some simple derivations, based on the Born approximation, which reproduce this structure for the imaginary part of the QNM frequencies for a large class of spherically symmetric spacetimes like the Schwarzschild black hole \\cite{padmanabhan03b, mmv03}. The relation obtained, valid for large $k$ (i.e., for large $n$), has the form \\be k_n = i n \\kappa ~~~ (\\mbox{for } n \\gg 1) \\e (In the imaginary part of $k_n$, one obtains a factor $n$ instead of $n + (1/2)$ since  the Born approximation is expected to be valid only for $n \\gg 1$). The important point to note is that ---  though it is not possible to obtain the real part of the QNM frequencies using the Born approximation --- this simple derivation does give the correct level spacing for imaginary values of $k_n$. Hence the possibility of extending the formalism to spacetimes with more complicated structure is worth examining. Following this approach, we use the Born approximation to examine the level spacing for spacetimes with two horizons.  The simplest spacetime with two horizons is that of a black hole in a spacetime with a cosmological constant, described by the Schwarzschild-DeSitter (SDS) metric. The metric is characterized by the presence of a black hole event horizon and a cosmological horizon. In recent times, studying such a spacetime has acquired further significance because of the cosmological observations suggesting the existence of a non-zero positive cosmological constant (\\cite{pag++99}; for reviews, see \\cite{padmanabhan03,ss00}). While the observations can be explained by a wide class of models (see, e.g \\cite{pc03}), including those in which the cosmic equation of state can depend on spatial scale \\cite{pc02,bjp03}, virtually all these models approach the DeSitter (DS) spacetime at late times and at large scales. Like in the other cases, the QNMs for the SDS spacetime have been studied both numerically (see, for example, \\cite{bckl97,bclp99,yf03,cl03,brink03,zhidenko03}) and analytically (see, for example, \\cite{mm90,mn02,suneeta03}). However, since the SDS spacetime is characterized by two different surface gravities (corresponding to the two horizons), the dependence of the level spacing of the imaginary part of the QNM frequencies on the surface gravities is not obvious. The numerical studies fail to give a ``clean'' result like equation (\\ref{eq:kn_num}) as they seem to vary depending on the relative values of the two surface gravities. A calculation based on the monodromy of the perturbation continued to the complex plane \\cite{ca03}, similar to what is done for the Schwarzschild case, gives a result of the form $k_n = i \\kappa_- \\left(n + \\f{1}{2}\\right) + \\mbox{real part}$, where $\\kappa_-$ is the surface gravity of the black hole horizon. The real part has a term analogous to $\\ln 3$ whose actual form depends on the values of the surface gravities. The above result is found to be in qualitative agreement with numerical results for the near-extremal case (i.e., when the surface gravities are very close to each other), although the exact numerical coefficients seem to differ \\cite{yf03}. It is not clear whether the problem is with the numerics, or the analytical derivation misses out on some issues. The above calculation indicates that the QNM spectrum is independent of the surface gravity of the cosmological horizon $\\kappa_+$. On the other hand the first version of \\cite{mmv03}, based on Born approximation, claimed that the spectrum of the imaginary part of the QNMs should behave as $k_n = i (n_- \\kappa_- + n_+ \\kappa_+)$ , where $n_{\\pm}$ are two integers. [This arXiv submission has been revised subsequently, after the appearance of the current work, removing this claim]. Analytical calculation based on approximating the scattering potential by some simple form gives QNMs proportional to both the surface gravities, depending on the time scale one is interested in \\cite{suneeta03}. Currently, we do not seem to have any general consensus on this particular issue!  There are a few more conceptual issues because of which such a study (using a prototype metric with two horizons) is important. First is the possible  connection between the QNMs and thermodynamics of horizons. It was pointed out in \\cite{padmanabhan03b} that the Born approximation result arises from integrals which are very similar to those which arise in the case of horizon thermodynamics (see e.g., equation (23) of \\cite{padmanabhan03c}). In the case of spacetimes with multiple horizons (like SDS), there is no unique temperature except when the ratio of surface gravities is a rational number \\cite{cp03}. If the QNMs are related to the horizon thermodynamics (in some manner which is not yet clearly understood), then it would be interesting to see whether the level spacing in the SDS spacetime can give any idea about the temperature of spacetimes with multiple horizons. Second, it is well known that thermodynamics of gravitating systems depend crucially on the ensemble which is used (see e.g., \\cite{padmanabhan90}) which translates into the boundary conditions on the horizon (see e.g., \\cite{gt03}). If QNMs are related to horizon thermodynamics, we will expect some similar kind of dependence on the boundary condition for the wave modes in case of SDS spacetime.  We shall see that this expectation is indeed borne out. Finally, there were also some discussion in the literature as to whether the QNMs depend on the region beyond the horizons. We will see that we can make some comments regarding this issue.  The structure of the paper is as follows. In Sec. II, we briefly review the results for the Schwarzschild metric. The main problem of interest, the SDS metric, is taken up in Sec. III. We derive the explicit form of the scattering amplitude using Born approximation and discuss the structure of the QNM spectrum. The main conclusions are summarized and compared with other results in Sec. IV. ",
        "conclusions": "We have used the first Born approximation to obtain the QNM spectrum for the SDS spacetime. The approximation gives the correct level spacing for the imaginary values of the QNM frequencies for the Schwarzschild black hole, and the spacing is related to the temperature corresponding to the horizon. On the other hand, there exist no QNMs for a massless  minimally coupled scalar field in a pure DS spacetime. It turns out that the situation is more complicated in the SDS spacetime and depends on the type of scattering one is interested in, i.e., on the type of boundary conditions one imposes.  One can start with the usual set of conditions where an incident wave is propagating towards the black hole and calculate the scattering amplitude. The poles of the amplitude will then represent boundary conditions appropriate for QNMs. It turns out that for this case the QNM level spacing depends only on the surface gravity of the black hole $\\kappa_-$, as expected. Thus the introduction of a black hole in the DS spacetime brings along the appropriate QNMs.  However, there exists another set of boundary conditions in which one starts with an incident wave propagating towards the cosmological horizon. As shown in Appendix C, it is possible to choose the boundary conditions and the definition of the scattering amplitude such that the QNM level spacing in this case depends only on the surface gravity of the cosmological horizon $\\kappa_+$. We do not believe these boundary conditions are physically relevant.  It was found earlier, based on the monodromy of the perturbation continued to the complex plane \\cite{ca03} that the imaginary part of the QNM frequencies have an equally spaced structure, with the spacing dependent only on $\\kappa_-$. It is not clear whether there exists any extensions of the above procedure for obtaining the other set of QNMs which are dependent on $\\kappa_+$. Analytical calculations, based on approximating the potential by a Poschl-Teller form \\cite{suneeta03}, gave a QNM spectrum which depends on {\\it both} the surface gravities $\\kappa_-$ and $\\kappa_+$ depending on the time-scale one is interested in. Since we are studying a time-independent situation, it is difficult to comment on time-dependence of the QNM spectrum -- however, our analysis indicates that if one starts with an incident wave packet which is composed of monochromatic waves propagating in both directions (i.e., terms of the form $e^{i k r_*}$ and $e^{-i k r_*}$), then one might obtain a QNM spectrum which depends on both the surface gravities. There is one more crucial difference between our results and those obtained by the Poschl-Teller potential \\cite{suneeta03}-- the level spacing in the later case is $2 \\kappa_{\\pm}$ rather than $\\kappa_{\\pm}$. Such a difference was noted in the case of QNMs obtained by Born approximation for the Schwarzschild black hole \\cite{padmanabhan03b} while comparing with results obtained by approximating the potential \\cite{solodukhin03}; the difference is probably related to the incorrect use of $q=k_i$ rather than $q = 2 k_i$ for momentum transfer in the Born approximation.  Most of the numerical computations regarding the SDS spacetime concentrate on the near extremal case (where $\\kappa_+ \\approx \\kappa_-$), and it is found that the imaginary part of the QNM frequencies have an equally spaced structure with the spacing given by either of the surface gravities (which are anyway equal to the lowest order) \\cite{cl03,brink03,yf03}. This means that, to the lowest order, we do not find any disagreement between our results and numerical computations. In other numerical computations, where the values of the two surface gravities are taken to be widely different \\cite{bckl97,bclp99}, one obtains two sets of QNM spectra proportional to the two surface gravities, each valid at different time-scales. At this stage, we are unable to make any direct comparison with such results since the issue of time-scales is difficult to settle in our approach.  In future studies for the SDS spacetime, it would be interesting to calculate the scattering amplitude using some more rigorous technique (like, say, what is done for the pure Schwarzschild case \\cite{neitzke03}) and see how the real parts of the QNM frequencies depend on the different surface gravities and on the type of scattering.   \\appendix"
    },
    "0311/astro-ph0311023_arXiv.txt": {
        "abstract": "We investigate the relation between the central velocity dispersion, \\sigc , and the circular velocity, \\vcirc , in galaxies. In addition to previously obtained data, we consider an observationally homogeneus sample of 52 high surface brightness and 11 low surface brightness spiral galaxies. We performed a straight line regression analysis in a linear scale, finding a good fit, also for low \\sigc\\ galaxies, always rejected in the previous studies. Low surface brightness galaxies seem to behave differently, showing either higher values of \\vcirc\\ or lower values of \\sigc\\ with respect to their high surface brightness counterparts. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311486_arXiv.txt": {
        "abstract": "We study the suppression of the small-scale power spectrum due to the decay of charged matter to dark matter prior to recombination. Prior to decay, the charged particles couple to the photon-baryon fluid and participate in its acoustic oscillations.  However, after decaying to neutral dark matter the photon-baryon fluid is coupled only gravitationally to the newly-created dark matter.  This generically leads to suppression of power on length scales that enter the horizon prior to decay.  For decay times of $\\sim$$3.5$ years this leads to suppression of power on subgalactic scales, bringing the observed number of Galactic substructures in line with observation.  Decay times of a few years are possible if the dark matter is purely gravitationally interacting, such as the gravitino in supersymmetric models or a massive Kaluza-Klein graviton in models with universal extra dimensions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/hep-ph0311336_arXiv.txt": {
        "abstract": "} \\nc{\\eab}{ Although often chosen because of simplicity, a single scalar field does not provide a general parametrization of an MSSM flat direction. We derive a formalism for a class of gauge invariant polynomials which result in a multifield description of the flat directions. In contrast to the single field case, the vanishing of the gauge currents yields an important dynamical constraint in the multifield framework. We consider in detail the example of the $H_uL$ flat direction and study the dynamical evolution during and after inflation. We highlight the differences between the single and the multifield flat directions. We show that in the multifield case the field space has an intrinsic curvature and hence unsuppressed non-minimal kinetic terms for the flat direction scalars arise. Also the phases of the individual components evolve non-trivially right after inflation, charging the components of the condensate and producing an enhanced entropy after the decay of the condensate, which is due to cross-coupling of different lepton flavours in the F term. However, the qualitative features of the single field Affleck-Dine baryogenesis, such as the  produced total charge, remain largely unchanged. ",
        "introduction": "Supersymmetric gauge theories come with gauge invariant polynomials along which the scalar potential vanishes at a classical level \\cite{Buccella:1981ib}. These correspond to the flat directions in the moduli space of the scalar fields.  Within the Minimal Supersymmetric Standard Model (MSSM) the classical potential is given by the sum of the $F$ and $D$ term contributions, which vanish individually along these flat directions.  Flatness is however lifted by soft supersymmetry breaking, but a non-renormalization theorem guarantees that the flat direction does not obtain any renormalizable perturbative superpotential correction \\cite{gwr79}. However, non-perturbative (supergravity-induced) corrections will also lift the flatness.  MSSM flat directions can give rise to a host of cosmologically interesting dynamics (for a review, see \\cite{Enqvist:2003gh}). These include Affleck-Dine baryogenesis \\cite{Affleck:1984fy,Dine:1995uk,Dine:1995kz}, the cosmological formation and fragmentation of the MSSM flat direction condensate and subsequent $Q$-ball formation \\cite{Kusenko:1997zq,Enqvist:1997si,Jokinen:2002xw,Kasuya:1999wu}, reheating the Universe with $Q$-ball evaporation~\\cite{Enqvist:2002rj}, generation of baryon isocurvature density perturbations \\cite{Enqvist:1998pf}, as well as curvaton scenarios where MSSM flat directions reheat the Universe and generate adiabatic density perturbations \\cite{Enqvist:2002rf}. Adiabatic density perturbations induced by fluctuating inflaton-MSSM flat direction coupling  has also been discussed~\\cite{Enqvist:2003uk}.   In the MSSM all the flat directions have been parameterized by gauge invariant monomials \\cite{Dine:1995kz,Gherghetta:1995dv}. In such a case the MSSM flat direction is some linear combination of the MSSM scalars and can be thought of as a trajectory in the moduli space described by a single scalar degree of freedom. A simple example is the $H_uL$ flat direction $\\phi$, defined as $(H_u)_1 = L_2 =\\phi/\\sqrt{2}$ (indices here label SU(2)). However, a flat direction can also be constructed as a gauge invariant polynomial, where the previously mentioned monomials generate the polynomial. This means that the flat direction is a non-linear combination of the MSSM fields. It no longer can be described by a single scalar field. Such a gauge invariant polynomial, rather than a monomial, is needed e.g. when  for $H_uL$ instead of a single leptonic generation one considers simultaneously all three.  The purpose of the present paper is to illustrate the dynamics of Affleck-Dine condensate described by a gauge invariant polynomial. Previously Senami and Yamamoto \\cite{Senami:2001qn} have considered such a case in MSSM and its extensions but they used an extrinsic point of view without gauge invariant polynomials. We will discuss the dynamics from an intrinsic point of view {\\it i.e.} using coordinates that describe only the behaviour of the flat direction. Given that the flat direction is the minimum of energy the two approaches are equivalent.  We will also discuss how a non-renormalizable superpotential correction can lift such a multi-field flat direction. We will highlight the formation and charge conservation of the condensate along with the gauge conditions required to be satisfied. In the present paper we will mostly restrict ourselves to a simple example of a flat direction involving $L$ and $H_{u}$ (\"the $H_{u}L$ flat direction\"), which is a linear combination of superfields, and study its dynamical evolution through the cosmic history of the Universe.  We begin our discussion with the basic definitions of the flat directions in $N=1$ supersymmetry in Section 2. In Section 3 we first construct $H_{u}L$ flat direction parameterized by a gauge invariant monomial. We then discuss gauge invariant polynomials which are $F$ and $D$ flat at the level of unbroken supersymmetry and write down the equation of motion for the $H_{u}L$ scalars with the help of two projection operators, defined in the Appendix. Vanishing of the gauge currents is carefully accounted for. to In Section 4 we write down the multifield scalar potential, accounting for supersymmetry breaking and non-renormalizable superpotential corrections. In Section 5 we discuss qualitatively the dynamical behavior of the $H_{u}L$ condensate consisting, comparing it with single field case. Section 6 presents the results numerical studies for the multifield trajectories, charge densities and charge asymmetries. In Section 7 we make concluding remarks and discuss the physical implications of the results. ",
        "conclusions": "The numerical studies indicate the produced baryon-to-entropy ratio is not significantly changed when including all the lepton families to the $H_uL$ flat direction. It thus seems that the simple one-field case captures the essential features of the Affleck-Dine mechanism, at least for $H_uL$.  However, there are some interesting details that do change. The first is that the flat direction has non-minimal kinetic terms, which is due to the fact that the sub-manifold spanned by the flat direction is curved. In the one-field case non-minimality is absent because all the one-dimensional manifolds are flat. Hence there is always a choice of coordinates where the field metric is given by just the Kronecker delta. Another difference is that there is phase motion of the fields already right after inflation. This is due to the F-terms containing cross terms that break individual lepton symmetries. However, the F-terms conserve the total lepton charge which is produced only through the A-term. The individual charges may in general have opposite signs and can give substantial contribution to the entropy, which is proportional to the sum of absolute values of the individual charge densities. The total charge is the sum of individual charge densities and thus cancellations can occur.  Although the behaviour of a homogenous $H_uL$ background with multiple scalar fields is not qualitatively different from the corresponding one-field case, this may not be true if inhomogeneities are taken into account. For small cosmological perturbations and multiple fields the gauge current condition of \\eq{grad323} becomes highly non-trivial. This can be seen as follows. $H_uL$ is described by three complex scalar fields or six real degrees of freedom. The current condition \\eq{grad323} defines a two-dimensional submanifold in the six-dimensional space, so that there are four real fluctuation degrees of freedom that would not obey \\eq{grad323}. Hence there would exist matter currents. This implies that the pure gauge field no longer provides a solution for the equations of motion. As a consequence, there would arise electric and magnetic gauge fields which could contribute to the CMB perturbations. These might be of paramount importance for condensate fragmentation, since the fragments (the ground states of which are Q-balls) result from the non-linearities of the perturbations \\cite{Kusenko:1997zq,Enqvist:1997si}. The Q-ball properties \\cite{Coleman:1985ki} have been considered only for condensates with minimal kinetic terms.  In this paper we focused on the $H_uL$ direction. For other directions more complications may arise. For example, the $d=6$ flat direction $(LL)(\\bar{d}\\bar{d}\\bar{d})(\\bar{u}\\bar{d}\\bar{d})$ has six different superpotential terms that contribute to the lifting of the flat direction: $d=4$ term $\\bar{u}\\bar{u}\\bar{d}\\bar{e}$, $d=5$ terms $H_uL\\bar{u}\\bar{d}\\bar{d}$ and $H_dL\\bar{d}\\bar{d}\\bar{d}$ and $d=6$ terms $\\bar{u}\\bar{d}\\bar{d}QL\\bar{d}$, $(\\bar{u}\\bar{d}\\bar{d})(\\bar{u}\\bar{d}\\bar{d})$ and $LL\\bar{e}\\bar{u}\\bar{d}\\bar{d}$. The $d=5,~6$ terms have $B-L=-2$ whereas the $d=4$ term  has $B-L=0$. These will give rise to cross terms in the F term which has a non-zero $B-L$ and therefore contributes to the total charge. Such a situation occurs whenever the various superpotential terms with different $B-L$ have at least one field in common.    \\subsection*{Acknowledgements} A.J. is supported by Magnus Ehrnrooth foundation. K.E. is supported partly by the Academy of Finland grant no. 75065, and A.M. is a CITA-National fellow. We thank Karsten Jedamzik for helpful discussion.  \\appendix"
    },
    "0311/hep-th0311144_arXiv.txt": {
        "abstract": "We consider the effects of a two-form field on the late-time dynamics of brane gas cosmologies. Assuming thermal equilibrium of winding states, we find that the presence of a form field allows a late stage of expansion of the Universe even when the winding degrees of freedom decay into a pressureless gas of string loops. Finally, we suggest to understand the dimensionality of the Universe not as a result of the thermal equilibrium condition but rather as a consequence of the symmetries of the geometry. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311603_arXiv.txt": {
        "abstract": "We are carrying out a search for Ly$\\alpha$ emitting galaxies at $z = 6.5$ employing slitless spectroscopy at the VLT. In our 43 arcmin$^2$ field we find three isolated single emission lines with SEDs consistent with line emitting galaxies at $z = 6.5$. ",
        "introduction": "Since the discovery of $z\\sim3$ Lyman Break Galaxies (LBGs~\\cite{ste96}), many efforts are being made to select the whole population of star-forming galaxies at the highest possible redshifts (see~\\cite{ste99} for a review), and to constrain the star formation history of the universe (e.g.~\\cite{mad96}). The detection of $z > 6$ galaxies allows to study the modes of early galaxy formation and the interplay between the first galaxies and the intergalactic medium (IGM).  The redshift range $6 < z < 7$ is a very intriguing time during cosmic evolution, when hydrogen reionization is believed to be basically complete and the IGM starts to be polluted with metals \\cite{gne00,cia03}.  Because a significant part of the bolometric luminosity of primeval star-forming objects escapes as Ly$\\alpha$ emission, especially if dust is not yet ubiquitous, it seems profitable to search for the redshifted Ly$\\alpha$ line of high redshift galaxies. The advent of ten meter class telescopes made this approach successful and efficient, and both narrow-band imaging and optical \\emph{serendipitous} spectroscopy found Ly$\\alpha$ emitters at $3 < z < 6$ \\cite{ste99,hu00,mal01}. Thanks to the recently increased sensitivity of CCDs at $\\lambda > 0.8$\\,\\umu m, the discovery of $z > 6$ galaxies has also become possible. Using narrow band filters sensitive to wavelengths in the range $9100 < \\lambda < 9250$\\,\\AA, i.e.\\ Ly$\\alpha$ emission at $z \\sim 6.5$ and subsequent spectroscopy, two groups have detected the first three galaxies known at $z > 6$. The first of these was found due to the strong lensing amplification (4.5) by the cluster Abell 370~\\cite{hu02}. The other two galaxies were found in a survey made with Subaru Suprime-Cam instrument which provides a huge field of view (814 arcmin$^2$~\\cite{kod03}). ",
        "conclusions": "The number of possible Ly$\\alpha$ emitters at $z = 6.5$ detected in our field is consistent with the number densities of $z > 5.5$ galaxies reported by other authors~\\cite{ste99,hu00}.  In particular, the recent Subaru survey by~\\cite{kod03} resulted in the detection of 73 Ly$\\alpha$ candidates at $z = 6.5$ (with $i' - z' > 1.3$) in an 814 arcmin$^2$ field.  Based on this result, we expect to find in our available 43 arcmin$^2$ field three candidate Ly$\\alpha$ emitters at $z = 6.5$.  Slitmask spectroscopic observations are needed to confirm the identity of the candidate Ly$\\alpha$ emitters and the LBGs.  These observations will provide higher signal to noise and higher spectral resolution to measure the characteristic asymmetric line profile of Ly$\\alpha$ at high redshift.  The larger wavelength coverage will allow the detection of emission lines from the presumed LBGs outside the $z_{sp}$ filter.  Time has been allocated to carry out these observations with FORS2."
    },
    "0311/astro-ph0311329_arXiv.txt": {
        "abstract": "We show that large, high-redshift ($z>10$) galaxies with virial temperature in excess of $10^4$K may be mostly comprised of cold atomic clouds which were formerly minihalos. % These clouds move at a speed of $\\sim 15-30\\kms$ and collide with one another on a time scale of $10^7$yr. The supersonic collisions result in spatially distributed star formation with high efficiency. We then show that most, subsequent star formation in cold atomic clouds may be triggered by shockwaves launched from the first stars formed in collisions. Those shockwave-compressed clouds are even more widespread spatially because of large imparted velocities and some can possibly escape into the intergalactic medium. More importantly, with respect to cosmological reionization, widespread star formation would allow a much higher ionizing photon escape fraction. These favorable conditions may form the physical basis to enable the standard cosmological model to produce a reasonably high Thomson optical depth $\\tau_e=0.10-0.14$. A chain reaction of star formation in minihalos in the intergalactic space may be triggered by explosions in the intergalactic medium, if minihalos are strongly clustered. In this case, a still higher $\\tau_e$ would be achievable. ",
        "introduction": "The Wilkinson Microwave Anisotropy Probe (WMAP) polarization observations measured Thomson optical depth $\\tau_e=0.17\\pm 0.04$ ($68\\%$) (Kogut \\etal 2003), requiring a full reionization to have occured first at $z=20^{+10}_{-9}$, such as in the double reionization picture (Cen 2003a). This high Thomson optical depth is further confirmed by a variety of combined data sets (Tegmark \\etal 2003). To ionize the universe at such a high redshift demands both a high formation efficient ($\\epsilon$) for Population III stars and a large escape fraction ($f_{esc}$) for the produced ionizing photons. Quantitatively, one would require the product of the two to be greater than $0.01$ in order to produce $\\tau_e > 0.11$ (e.g., Cen 2003b; Haiman \\& Holder 2003; Wyithe \\& Loeb 2003; Ciardi, Ferrara \\& White 2003; Ricotti \\& Ostriker 2003) to be consistent with the WMAP observations at $\\sim 1.5\\sigma$ level. Conventional wisdom may view these assumptions somewhat too optimistic at first sight.  In this {\\it Letter} we suggest that star formation at high redshift ($z>10$) may occur in massive starbursts, triggered by supersonic collisions between constituent cold atomic clouds and blastwaves produced by explosions of first stars. The consequences are (1) high star formation efficiency and (2) high ionizing photon escape fraction. ",
        "conclusions": "In the cold dark matter model large, high-redshift ($z>10$) galaxies with virial temperature $T_v>10^4$K are shown to be mostly comprised of cold atomic clouds, which were originally minihalos and accreted into larger halos. It is shown that supersonic collisions of cold atomic gas clouds on time scale of $10^7$yr will shock heat and compress the gas, which then cools and form \\h2 molecules to further cool to $\\sim 100$K. Massive star formation ensues. The resulting star formation is spatially distributed.  It is then shown that most of the star formation in minihalos may be triggered by shockwaves produced by the first stars formed in collisions. Those shockwave compressed minihalos will be more widespread spatially due to large imparted velocities and some can possibly escape into the intergalactic medium. If minihalos in the intergalactic space are strongly clustered, a chain reaction of star formation in minihalos may be triggered initiated by explosions of these stray stars. In this case, far more ionizing photons can be produced than large halos alone can. Separately, whether cooling shells of swept-up mean density gas in the intergalactic medium can form stars and possibly trigger an explosion-induced chain reaction even in the absence of strong minihalo clustering is unclear.  More importantly with respect to cosmological reionization, widespread star formation would allow an ionizing photon escape fraction much higher than in the conventional picture where star formation occurs in the central regions of galaxies and whose the escape fraction is estimated to be below $1\\%$ for galaxies more massive than $10^7\\msun$.  Such star formation may be the norm at $z>10$ and may create the physical basis to reconcile the standard cosmological model with the WMAP polarization observations. An optical depth $\\tau_e=0.10-0.14$ can be achieved realistically even in the absence of chain reactions forementioned. A still higher optical depth may indicate that some further enhancement in star formation has occurred, such as through possible chain reactions. Detailed simulations of star formation triggered by collisions and blastwaves at high redshift galaxies would be required to firm up this picture and quantify key parameters, including star formation efficiency and ionizing photon escape fraction.  Finally, we note that these widespread star formation outlined above may give rise to metal enrichment of the intergalactic medium far more uniform than any other processes that have been proposed (e.g., Gnedin 1998; Aguirre \\etal 2001; Cen \\etal 2003).  This research is supported in part by grants AST-0206299 and NAG5-13381."
    },
    "0311/astro-ph0311439.txt": {
        "abstract": "We present near infrared photometry of the Type Ia supernova 1999ee;  also, optical and infrared photometry of the Type Ia SNe 2000bh, 2000ca, and 2001ba.  For SNe 1999ee and 2000bh we present the first-ever SN photometry at 1.035 $\\mu$m (the $Y$-band).  We present K-corrections which transform the infrared photometry in the observer's frame to the supernova rest frame.  Using our infrared K-corrections and stretch factors derived from optical photometry, we construct $JHK$ templates which can be used to determine the apparent magnitudes at maximum if one has some data in the window $-$12 to +10 d with respect to T($B_{max}$). Following up previous work on the uniformity of $V \\; minus$ IR loci of Type Ia supernovae of mid-range decline rates, we present unreddened loci for slow decliners.  We also discuss evidence for a {\\em continuous} change of color at a given epoch as a function of decline rate. ",
        "introduction": "\\parindent = 9 mm  Type Ia supernovae (SNe) are well known to the {\\em cognoscenti} as standardizable candles for determining extragalactic distances.  They are considered to be the explosions of CO white dwarfs in binary systems.  Using optical light curves, Phillips (1993), Hamuy et al. (1996a), and Phillips et al. (1999)  established the relationship between the absolute magnitudes at maximum and the decline rates of these SNe.  The characteristic parameter is $\\Delta$m$_{15}$($B$), the number of magnitudes that the SN declines in the first 15 days after $B$-band maximum.\\footnote[15]{The method actually uses the $B$-, $V$- and $I$-band light curves to obtain a characteristic value of $\\Delta$m$_{15}$($B$).} Riess, Press, \\& Kirshner (1996), Riess et al. (1998), and Jha, Riess, \\& Kirshner (2004) have elaborated three versions of the Multicolor Light Curve Shape (MLCS) method. This method now correlates the light curve shapes in $UBVRI$ with the absolute magnitudes.  In the MLCS method the characteristic parameter is $\\Delta$, the number of $V$-band magnitudes that a Type Ia SN at maximum is brighter than ($\\Delta < 0$), or fainter than ($\\Delta > 0$) a fiducial object of M$_V$(max) = $-$19.47 on a scale where H$_0$ = 63 km s$^{-1}$ Mpc$^{-1}$ . Finally, the stretch method of Perlmutter et al. (1997) is used to stretch or shrink $B$- and $V$-band light curves in the time axis to match a fiducial.\\footnote[16]{The stretch method does not work with $R$-and $I$-band light curves, except around T($B_{max}$), because of the shoulder in the $R$-band light curves and the secondary maximum in the $I$-band ones. See Fig. \\ref{i_stretch} below.  For roughly 90 percent of Type Ia SNe, the strength of the secondary hump is correlated with $\\Delta$m$_{15}$($B$) (Krisciunas et al. 2001), hence with the absolute magnitudes at maximum.} Since brighter SNe are more slowly declining, the stretch factor is less than 1.0 for the intrinsically brighter SNe, and greater than 1.0 for the intrinsically fainter ones.  Typically, $\\Delta$m$_{15}$($B$) = 1.1 corresponds to $\\Delta$ = 0 and stretch factor $s$ = 1.0.  While the optical light curves of most Type Ia SNe fit into some kind of an ordering scheme, with the marked increase of SN discoveries in the past few years, more and more objects are being found which do not fit the patterns. SN~1999ac was a slow riser, fast decliner (Labbe et al. 2001, Phillips et al. 2003).  SN~2000cx was a fast riser, slow decliner and had unusually blue $B-V$ colors after T($B_{max}$) + 30 d (Li et al. 2001, Candia et al. 2003).  Recently, Li et al. (2003) describe the even stranger SN~2002cx, which had premaximum spectra like SN 1991T (the classical hot Type Ia SN), a luminosity like SN 1991bg (the classic sub-luminous object), a slow late-time decline, and unidentified spectral lines. These ``pathological'' examples should motivate breakthroughs in the modeling of Type Ia SNe, allowing us to understand better the more normal objects.  Following the pioneering papers of Elias et al. (1981, 1985) and the photometry of SN~1986G (Frogel et al. 1987), essentially no infrared light curves of Type Ia SNe were obtained until the appearance of SN~1998bu. %\\footnote[17]{If we are allowed an historical analogy, %we might call this the ``Brezhnev era'' $-$ not a whole lot bad %happened, but very little progress occurred.  If there are any %Czech astronomers reading this, we apologize for our historical %analogy.} (See Meikle 2000 for a compilation of all the IR photometry available three years ago.)  In 1999 two of us (MMP and NBS) started regular observing campaigns at Las Campanas Observatory and Cerro Tololo Inter-American Observatory to obtain optical and infrared light curves of supernovae, primarily Type Ia's.  For this we have used mainly the LCO 1-m Swope telescope and the CTIO 1-m Yale-AURA-Lisbon-Ohio (YALO) telescope.  In this paper we report optical and infrared photometry of the Type Ia SNe 2000bh, 2000ca, and 2001ba.  We present infrared photometry of SN~1999ee. Previous extensive $UBVRIz$ optical photometry of SN~1999ee has been published by Stritzinger et al. (2002).  A significant fraction of the data reported in this paper was obtained for the Supernova Optical and Infrared Survey (SOIRS), a project organized and carried out in 1999 and 2000 (Hamuy principal investigator).  See Hamuy et al. (2001) for further details.  SOIRS also included a photographic search for supernovae carried out at Cerro El Roble (Maza 1979). Amongst SOIRS discoveries are three of the four SNe reported here.  We note that the optical and IR photometry and spectroscopy of SN~1999ee (one of the prime SOIRS targets) is the largest such dataset ever obtained for a Type Ia SN. ",
        "conclusions": "In this paper we presented well sampled light curves for four Type Ia SNe. This includes the first-ever SN data published at 1.03 $\\mu$m, though the $Y$-band data lack a firm basis in calibration.  Given that the maxima in the infrared light curves occur typically 3 to 4 days before the time of $B$-band maximum, it is a considerable challenge to obtain light curves that cover the IR maxima.  We achieved this in the cases of SNe~1999ee, 2000ca, and 2001ba.  We have now observed a sufficient number of Type Ia SNe near the time of maximum that we can investigate the question of uniform templates for the infrared.  If we stretch the light curves in the time domain according to stretch factors based on the $B$- and $V$-band relationships of Jha (2002, Fig. 3.8), we obtain reasonably uniform templates from $-$12 to +10 days after maximum (in ``stretched days'').  The {\\sc rms} uncertainties of the fits are $\\pm$~0.062, 0.080, and 0.075 mag in $J$, $H$, and $K$, respectively.  This allows us to determine the maximum magnitudes of other Type Ia SNe as long as we have some data in the $-$12 to +10~d window, and the light curves are not obviously ``abnormal'' (e.g. like SN~2000cx).  Primarily on the basis of observations carried out at CTIO and Las Campanas since the beginning of 1999, we can now draw the $V \\; minus$ IR color curves for Type Ia SNe which are mid-range decliners and slow decliners.  We can utilize the present observational data, backed up by modeling for the mid-range decliners (Krisciunas et al. 2003, \\S 3.4), to determine the extinction suffered by these objects.  This is true even for highly reddened objects that occurred in galaxies with different dust properties than we measure for the dust in our Galaxy.  The motivation behind using IR data is that extinction is considerably lower than in the optical, and if we know the intrinsic $V \\; minus$ IR colors of Type Ia SNe, then we can determine A$_V$ by slightly scaling the $V \\; minus$ IR color excesses.  This should give us more accurate extinction corrections compared to deriving $B-V$ color excesses and scaling them by an assumed value of R$_V$.  However, while we can enumerate unreddened loci for Type Ia SNe of mid-range and slow decline rates, for the slow decliners the ``uniformity'' of $V-J$ and $V-H$ color is limited to within nine days of the time of $B$-band maximum, just about what we found for the construction of our $JHK$ templates.  The available data can also be used as evidence that there is a {\\em continuous} change of $V \\; minus \\; $ IR color at $t$ = +6 d vs. the decline rate.  The $V_{max}-J_{max}$ and $V_{max}-H_{max}$ colors are further evidence that Type Ia SNe are progressively redder as a function of decline rate.  Given that we already have evidence for a spectroscopic sequence as function of temperature (Nugent et al. 1995), the further calibration of the optical-IR color relationships will be relevant to SN modeling.  As large scale surveys of the sky are planned and implemented, we can and should make use of the light curve and color templates presented here to derive the maximum magnitudes of Type Ia SNe and correct the data for the effects of dust along the line of sight.  Since Type Ia SNe are the most important distance indicator at redshifts $z \\; >$ 0.01, we may further solidify the foundation of the cosmological distance ladder with these useful cosmic beacons.  %We now have the means to eliminate a potentially serious source of %systematic error in the distances to Type Ia SNe.  Until recently the %scatter in the Hubble diagram of Type Ia SNe was a combination of %photometric errors, uncertainties in extinction, and the intrinsic %scatter of the absolute magnitudes as a function of decline rate.  With %CCD imaging and image subtraction techniques we can eliminate problems %with the photometry, even for SNe buried in the light of their hosts. %With the $V \\; minus$ IR loci described here and in our previous papers %we have the tools to eliminate large uncertainties in the extinction %corrections even if the SNe are highly reddened and their host galaxies %have unusual dust properties in the optical.  Thus, the scatter in the %Hubble diagram of Type Ia SNe is now largely a function of the scatter %of intrinsic luminosities.  %qqq"
    },
    "0311/astro-ph0311090_arXiv.txt": {
        "abstract": "We present a possible orbit for the Southern Stream of stars in M31, which connects it to the Northern Spur.  Support for this model comes from the dynamics of planetary nebulae (PNe) in the disk of M31: analysis of a new sample of 2611 PNe obtained using the Planetary Nebula Spectrograph reveals $\\sim 20$ objects whose kinematics are inconsistent with the normal components of the galaxy, but which lie at the right positions and velocities to connect the two photometric features via this orbit.  The satellite galaxy M32 is coincident with the stream both in position and velocity, adding weight to the hypothesis that the stream comprises its tidal debris. ",
        "introduction": "\\label{sec:intro} It is now generally accepted that mergers play a key role in the formation of galaxies \\citep{wr78}.  Since galaxy evolution is an ongoing process, we might therefore expect to catch a number of systems in the nearby Universe mid-merger.  Indeed, dramatic major mergers with complex tidal tails have been documented for quite some time \\citep[see, for example, ][]{s86}.  Perhaps of greater importance to the more passive evolution of galaxies, and as a possible cause of phenomena such as thick disks \\citep{qhf93}, evidence for the more common minor mergers is now coming to light.  In these cases, the detritus of the events has a high enough surface brightness be visible as a stellar stream.  The Sagittarius Dwarf Galaxy provides a fine example in the Milky Way \\citep{maj03}, and one need look no further than the closest good-sized galaxy, M31, to find another dramatic stream of stars \\citep{ib01}, which is presumably the remnant of a similar minor merger.  The proximity of these streams indicates quite how common this phenomenon must be, but it also offers prime laboratories for studying the merger process in detail.  The Sagittarius Dwarf is inconveniently placed behind the centre of the Milky Way, and also presents the usual problems of geometry when trying to study an object inside our own galaxy, so the M31 stream is probably the best candidate for analysis.  However, even this is not without its problems: the large extent of M31 necessitates a very extensive data set for any complete survey, and the low surface brightness of the feature renders the stream hard to detect particularly against any of the brighter parts of M31.  The faintness of the stream and its constituent stars also makes it very challenging to obtain the kinematic observations that would tie down the full dynamical structure of the stream, and thus unequivocally demonstrate its nature.  In this paper, we seek to overcome these difficulties by supplementing the existing photometric data with a new sample of 2611 planetary nebulae in the disk of M31.  The kinematics of these discrete stellar tracers can be used to pick out the star stream right through the brighter parts of M31's disk as well as helping determine its dynamics.  The remainder of this paper is laid out as follows.  In Section~\\ref{sec:photo} we review the existing photometric data on the M31 stream features and the hypothesis as to how they may be connected.  Section~\\ref{sec:kin} presents the kinematic data, which strengthens the case for such a link.  Section~\\ref{sec:orbit} shows an orbit model that quantitatively marries these photometric and kinematic features, and presents the evidence that the stream may arise from M32's tidal debris.  Section~\\ref{sec:conc} concludes.  \\begin{figure*} \\includegraphics[width=177mm]{figure1.eps} \\caption{Composite illustration of the data on the star stream in M31, presented in a coordinate system in which $X$ lies along the major axis of M31, increasing toward the SW, $Y$ increases along the minor axis toward the NW, and $Z$ increases along the line of sight.  The greyscale shows the star-count data, as published in Ferguson et al.\\ (2002), with the Southern Stream and Northern Spur annotated.  The black ellipse indicates the extent and inclination of M31's disk, and the locations of NGC205 and M32 are also marked.  The points show the PNe detected in the PN.S survey, with the size of the symbols encoding their line-of-sight velocities relative to M31 (crosses receding, boxes approaching).  The circled points are ``friendless'' PNe that lie more than $4\\sigma$ in velocity from their 30 nearest neighbours.} \\label{fig:data} \\end{figure*} ",
        "conclusions": "\\label{sec:conc}  We have demonstrated that the star stream to the south of M31 and the spur to the north of the system can be modeled as parts of a single coherent merger stream stretching over some hundreds of kiloparsecs. As well as explaining where the Southern Stream goes after disappearing into M31, this model also eliminates the awkward need to invoke the most extreme stellar warp known as an explanation for the Northern Spur.  This is not to say that the stream is unrelated to the gaseous warp, however, as such a merger event could well play a role in exciting the warp \\citep{qhf93}.  The model also makes a prediction as to where one might be able to detect the stream kinematically, and observations of the PNe in the disk of M31 show that there are stars at exactly the velocities that one would expect.  These data trace the stream in the bright disk region where there is no chance of detecting it photometrically, and provide a direct link between the two photometric features at larger radii.  It is at present unclear whether the satellite whose debris makes up the Stream is still present or not. Both M32 and And~VIII are possible parents, though they imply different orbit solutions. Direct measurements of the radial velocity of the Southern Stream will be required to distinguish between the possibilities."
    },
    "0311/astro-ph0311573_arXiv.txt": {
        "abstract": "We present new dynamical models of weakly self-gravitating, finite dispersion eccentric stellar disks around central black holes for the double nucleus of M31.  The disk is fixed in a frame rotating at constant precession speed, and is populated by stars on quasi-periodic orbits whose parents are numerically integrated periodic orbits in the total potential. A distribution of quasi-periodic orbits about a given parent is approximated by a distribution of Kepler orbits dispersed in eccentricity and orientation, using an approximate phase-space distribution function written in terms of the integrals of motion in the Kepler problem.  We use these models, along with an optimization routine, to fit available published kinematics and photometry in the inner $2\\arcsec$ of the nucleus.  A grid of 24 best-fit models is computed to accurately constrain the mass of the central black hole and nuclear disk parameters.  We find that the supermassive black hole in M31 has mass $M_{BH} = 5.62 \\pm 0.66 \\times 10^7 \\msun$, which is consistent with the observed correlation between the central black hole mass and the velocity dispersion of its host spheroid.  Our models precess rapidly, at $\\Omega = 36.5 \\pm 4.2 \\kmspc$, and possess a characteristic radial eccentricity distribution, which gives rise to multi-modal line of sight velocity distributions along lines of sight near the black hole.  These features can be used as sensitive discriminants of disk structure. ",
        "introduction": "\\label{sec:CHAP1}  It is widely believed that most, if not all, galaxies have supermassive black holes (BHs) in their centers.  Estimates of the total mass density in quasar remnants (Soltan 1982, Chokshi \\& Turner 1992), models for the evolution of the quasar luminosity function in hierarchical structure formation scenarios  (Haehnelt \\& Rees 1993), and the large numbers of nearby galaxies with low-luminosity nuclear activity (Ho, Filippenko \\& Sargent 1997) are consistent with this belief, assuming that active galactic nuclei (AGNs) are powered by accretion of matter onto a BH (Lynden-Bell 1969, Rees 1984).  The discovery that the BH mass ($M_{BH}$) correlates with certain host galaxy properties has made obtaining accurate masses for these objects a high priority in extragalactic studies.  Kormendy \\& Richstone (1995) and Magorrian et al. (1998) find that BH mass is proportional to the mass or luminosity of the host spheroidal component, though with significant scatter.  Ferrarese \\& Merritt (2000) and Gebhardt et al. (2000a) find that the BH mass correlates with the velocity dispersion ($\\sigma$) of the stellar component, with much less scatter than the previous correlation; from a sample of 31 galaxies with secure BH mass estimates, Tremaine et al. (2002) find that the correlation can be written as $\\log(M_{BH}/\\msun) = \\alpha + \\beta \\log(\\sigma/\\sigma_0)$, where $\\alpha = 8.13 \\pm 0.06$ and $\\beta = 4.02 \\pm 0.32$ for a reference dispersion of $\\sigma_0 = 200 \\kms$.  Since $\\sigma$ is measured outside the radius of influence of the BH, defined to be $r_{h} = G M_{BH}/\\sigma^2$, the $M_{BH}$ - $\\sigma$ correlation demonstrates a fundamental relationship between the BH and its host spheroid.  Such a correlation has important implications for theories of BH and galaxy formation and evolution.  It is thus important to confirm and strengthen the correlation by providing highly accurate BH masses for a large number of galaxies.  BH masses can be found using a variety of techniques, including: measuring the kinematics of individually resolved stars (Eckart \\& Genzel 1996, Ghez et al. 1998, Sch\\\"{o}del et al. 2002, Ghez et al. 2003), dynamical modeling of spatially resolved stellar absorption-line kinematics near the BH (see Kormendy \\& Richstone 1995, Verolme et al. 2002, Gebhardt et al. 2003), measuring rotation curves from optical (Harms et al. 1994, Macchetto et al.  1997, Bower et al. 1998, van der Marel \\& van den Bosch 1998, Marconi et al. 2003) or maser (Miyoshi et al. 1995, Ishihara et al. 2001) emission lines from orbiting gas, reverberation mapping in active galaxies (Peterson \\& Wandel 1999, Peterson \\& Wandel 2000, Gebhardt et al. 2000b), and modeling of line profile widths (Vestergaard 2002).  The kinematics of resolved stellar motions and small maser disks provide the most reliable mass estimates.  However, the motions of individual stars can only be resolved in the Milky Way (see Sch\\\"{o}del et al. 2002), and regular maser emission is only found in a few galaxies (Hagiwara et al. 2003).  Of the other techniques, stellar-dynamical modeling provides the most secure BH measurements;  gas near the BH can be subject to non-gravitational forces, unlike the stellar component.  The kinematic data must be resolved inside the region where Keplerian motion dominates, however, to ensure that those stars fully contribute to the line of sight velocity distribution (LOSVD), and not just in its tails.  Simple calculations suggest that Keplerian motion should dominate within a region of radius $r_k = k r_h$, where $k\\approx 0.1$ - $0.3$, depending on the BH mass and stellar radial density profile.  Using $k = 0.3$, it is easy to show that $r_k \\leq 1.3 \\times 10^{-3} M_{BH} / \\sigma^2$ (pc), with $M_{BH}$ in solar masses and $\\sigma$ in km/s.  Using values for $M_{BH}$, $\\sigma$, and distance given in Tremaine et al. (2002), or from the $M_{BH}$ - $\\sigma$ relation, $r_k$ subtends an angle of $11 \\farcs 60$, $0\\farcs 74$, and $0\\farcs 15$ for the Milky Way, M31, and M32, respectively.  Other than the Milky Way, M31 is the only nearby galaxy in the Local Group with a resolved $r_k$ at the resolution of the Hubble Space Telescope (HST) Space Telescope Imaging Spectrograph (STIS; $\\sim 0\\farcs 1$); M33 is consistent with having a maximum BH mass of $\\sim 3000 \\msun$ (Merritt et al. 2001, Gebhardt et al. 2001), so its $r_k$ subtends an angle $< 0\\farcs005$.  M31 offers a unique opportunity to obtain a secure BH mass from spatially resolved stellar kinematics inside $r_k$.  M31 is also the nearest galaxy with a normal bulge (Kormendy 1993), and it has a nucleus; that is, a small-scale stellar component which is photometrically and dynamically distinct from the bulge and the large-scale galactic disk (Kormendy \\& Richstone 1995, and references therein).  Galaxy nuclei are poorly understood, as is the dynamical connection between the nuclear stars and the central BH.  M31's nucleus is $\\sim 2\\arcsec$ in radius (Light, Danielson \\& Schwarzschild 1974), which is fully within the sphere of influence of the BH;  $r_h \\simeq 2\\farcs5$, if M31 is located at a distance of $770$ kpc and has a BH mass of $5.5 \\times 10^7 \\msun$, as implied by the $M_{BH}$ - $\\sigma$ relation.  Thus, M31's nucleus allows for a more detailed dynamical study than is possible for any other galaxy. Even more enticing is the fact that the nucleus is shown to be double; the photometric profile shows two brightness peaks, one of which is off-center with respect to the outer bulge isophotes.  Kinematic profiles also show strong asymmetries.  Thus, standard axisymmetric dynamical modeling techniques (e.g. in Gebhardt et al. 2003) are inappropriate.  New modeling methods are needed to obtain an accurate measure the BH mass in M31.  M31 was first shown to have a photometrically asymmetric nucleus by Light et al. (1974), using the Stratoscope II balloon-borne telescope.  Nieto et al. (1986) confirmed those observations with groundbased data, and found that the brightest point in the nucleus was offset from the center of the bulge by $\\sim 0\\farcs 4$.  The HST Wide-field and Planetary Camera (WFPC1) later resolved the nucleus into two brightness peaks (Lauer et al. 1993, hereafter L93), as did more recent HST images taken with WFPC2 (Lauer et al.  1998, hereafter L98).  The optically brighter peak, P1, is offset $0\\farcs 49$ from the bulge photometric center, which coincides with the fainter peak, P2.  P1 and P2 have central $V$ band surface brightnesses of $13.4 \\magarc$ and $13.7 \\magarc$, respectively, when averaged over a $0\\farcs22$ wide slit (L93).  P1 is compact, with a major-axis core radius of $\\sim 0\\farcs4$; P2 has a weak stellar cusp, unlike P1 (L93).  Near-IR (Mould et al. 1989, Davidge et al. 1997, Corbin et al. 2001), optical (L93, L98), and far-UV (King et al. 1995, hereafter K95) images all show that the asymmetric or double-peaked structure of the nucleus is not caused by dust absorption.  Along with absorption-line strengths from long-slit spectra (Kormendy \\& Bender 1999, hereafter KB99), they also demonstrate that P1 has a similar stellar content as the rest of the nucleus, which is unlike any globular cluster or dwarf elliptical.  Thus, P1 is an intrinsic part of the nucleus, and not an interloping star cluster (K95).  The $V-I$ color of the nucleus is not the same as that of the bulge, implying a difference in stellar populations (L98, Bacon et al. 2001, hereafter B01); line strengths in KB99 also show the same.  The color difference is not agreed upon, however; L98 find that the nucleus is redder than the bulge, whereas B01 find the opposite.  Sil'chenko et al. (1998) argue that the nucleus is more metal rich than the bulge, and, using $H_\\beta$ lines to disentangle metallicity and age, find that the nucleus is a factor of three younger than the bulge.  P2 is brighter than P1 in the UV, as a result of an embedded UV-bright source (hereafter the UV peak; K95, Brown et al. 1998, L98) whose center is located $0\\farcs 076$ toward P1 from the center of P2 in the $I$-band (B01).  The UV peak is resolved, with a half-power radius of $\\sim 0\\farcs2$ (Brown et al.  1998, L98, B01).  Brown et al. (1998) show that the UV peak is consistent with being comprised of extreme horizontal branch stars with masses between $0.47$ and $0.53$ \\msun, but not consistent with a majority contribution from main sequence stars, blue stragglers, or post-asymptotic giant branch stars more massive than $0.56 \\msun$.  Unpublished spectra from STIS also suggest that the UV peak is dominated by starlight (E. Emsellem, private communication), rather than a low-level AGN (K95).  The UV peak is thought to be the location of the photometric center of the bulge, and the supermassive BH (K95, KB99, B01; Peng 2002, hereafter P02).  Hereafter in this paper, ``the nucleus'' refers to P1 and P2 together, but does not include the UV peak, which is a separate nuclear star cluster.  The locations of P1, P2, and the UV peak with respect to the nucleus as a whole are shown in Figure \\ref{f.nucleus}.  \\begin{figure} \\plotone{f1.eps} \\caption{\\footnotesize WFPC2/HST photometry in the $I$ (F814W) band from L98.  The image has been boxcar smoothed for clarity.  North points toward the top of the page, while East points to the left.  Arrows show the locations of P1, P2, and the UV peak; the center of the UV peak is denoted by an asterisk at the origin.  The solid line shows the P1-P2 line ({\\it PA}$_d=42\\arcdeg$), or the major axis of the nucleus.  The dotted line shows the kinematic major axis, which is the line joining the velocity extrema in the two-dimensional kinematic map ({\\it PA}$_K=56.4\\arcdeg$). } \\label{f.nucleus} \\end{figure}  Groundbased observations at $\\sim 1\\arcsec$ (FWHM) resolution by Dressler (1984), Kormendy (1988), Dressler \\& Richstone (1988), and van der Marel et al. (1994) were the first to show that the stellar component in M31's nucleus rotates rapidly, and that there is a significant velocity dispersion peak (hereafter the dispersion spike) in the central few parsecs, both possibly indicating the presence of a central BH of mass $\\sim 10^7 \\msun$.  The data show the dispersion spike to be centered $\\sim 0\\farcs 6$ away from the peak in brightness (P1) in the Stratoscope II photometry, and the nucleus to be colder than the bulge on both sides of the dynamical center.  Two-dimensional kinematic maps obtained by Bacon et al. (1994) at similar resolution ($0\\farcs 87$ FWHM), using the TIGER integral field spectrograph on the Canada-France-Hawaii Telescope (CFHT), are consistent with most of the earlier observations;  however, the dispersion spike in that data set is located $\\sim 0\\farcs 7$ from P2 on the anti-P1 side of the nucleus, which places it $\\sim 0\\farcs 6$ farther away from P1 than found previously.  The deconvolved TIGER rotation curve is asymmetric about the rotation center, which is near P2; the maximum amplitude on the anti-P1 side is roughly $60 \\kms$ greater than that on the P1 side.  Observations at better spatial resolution ($0\\farcs 64$ FWHM) and higher signal-to-noise (S/N), taken with the Subarcsecond Imaging Spectrograph (SIS) on CFHT (KB99), show the dispersion spike offset from the UV peak by $\\sim 0\\farcs 2$, roughly $0\\farcs 4$ less than the Bacon et al. (1994) offset;  the spike's amplitude is $248 \\pm 5 \\kms$ before bulge subtraction and $287 \\pm 9 \\kms$ after.  The nucleus is cold on both sides of the UV peak, as in Kormendy (1988);  for example, the dispersion at $r=0\\farcs92$ from the UV peak on the P1 side is $123 \\pm 2 \\kms$ with the bulge, and $\\sim 100 \\kms$ without.  KB99 find that the bulge-subtracted maximum rotation velocity is $-236 \\pm 4 \\kms$ on the anti-P1 side, but only $179 \\pm 2 \\kms$ on the P1 side, confirming the $\\sim 60 \\kms$ rotation amplitude asymmetry of Bacon et al. (1994).  When the bulge is added, the asymmetry is only $\\sim 7 \\kms$, with a maximum velocity on the P1 side of $152 \\pm 3 \\kms$.  The zero velocity crossing is displaced from the UV peak toward P1 by $0\\farcs 051 \\pm 0\\farcs 014$.  Slit-averaged velocity profiles from the OASIS integral field spectrograph on CFHT (B01), which has about twice the spatial resolution ($\\sim 0\\farcs 4$ - $0\\farcs 5$ FWHM) as TIGER, are consistent with the SIS observations.  B01 measure the kinematic major axis, or the line joining the velocity extrema in the two-dimensional map, to be at position angle {\\it PA}$_K = 56.4\\arcdeg \\pm 0.2\\arcdeg$, which is not on the P1-P2 line ({\\it PA}$_d = 42\\arcdeg$; see Figure \\ref{f.nucleus}).  The kinematic observations with the best resolution to date ($\\sim 0\\farcs 1$ FWHM) come from the f/48 long-slit spectrograph of the HST Faint Object Camera (FOC; Statler et al. 1999).  The rotation curve is resolved through the rotation center with a projected velocity gradient of $\\sim 300 \\kmspc$;  the zero velocity crossing is offset by $0\\farcs 16 \\pm 0\\farcs 05$ from P2 towards P1, or $\\sim 0\\farcs 135$ from the UV peak (using the spatial registration suggested by B01).\\footnote{B01 determined that spatial shifts must be applied to the FOC and SIS data to register them to the center of the UV peak in the F300W band, which is the reference center for the OASIS and STIS data.  The origin defined in Statler et al. (1999) must be shifted $0\\farcs 25$ toward P1, whereas the origin defined in KB99 must be shifted $0\\farcs 031$ away from P1.  They also found that a positive shift of $30 \\kms$ must be applied to the FOC velocity profile for consistency with the STIS profile;  this amounts to adding $30 \\kms$ to the systemic velocity.} The rotation curve is asymmetric, as in the SIS data, with an amplitude asymmetry of at least $60 \\kms$ and a P1-side maximum of $\\sim 240 \\kms$.  The dispersion spike has amplitude $440 \\pm 70 \\kms$, but is only offset from P2 by $0\\farcs 06$, in contrast to the $\\sim 0\\farcs 2$ offset found with SIS and OASIS.  High-resolution kinematic data from STIS/HST (B01) also show similar asymmetries.  The rotation curve has an amplitude asymmetry possibly as high as $\\sim 90 \\kms$, with a maximum rotation amplitude on the P1-side of $201 \\pm 5 \\kms$.  The velocity gradient is $\\sim 220 \\kmspc$ through the zero velocity crossing, which occurs $0\\farcs 09$ from the UV peak;  both of these values are lower than for FOC.  The dispersion spike has amplitude $321 \\pm 33 \\kms$, and is located $0\\farcs 235$ from the UV peak on the anti-P1 side.  B01's STIS dispersion spike is substantially more offset than that in the FOC data.  We refer the reader forward to figures in Section \\ref{s.bestfits} to see FOC, STIS, SIS and OASIS kinematic profiles.  Two hypotheses have been explored to account for the photometric and kinematic asymmetries observed in M31's nucleus:  first, that P1 represents a captured star cluster orbiting around a stellar disk and central BH (Emsellem \\& Combes 1997); second, that P1 is produced by orbit crowding at apoapsis in an eccentric disk of stars on apse-aligned Kepler orbits about a BH at P2 (Tremaine 1995, hereafter T95).  Of the two hypotheses, the evidence strongly points toward the second as correct.  The eccentric disk picture naturally explains the nearly uniform colors of the nucleus, since P1 and P2 are the same stellar population. It is difficult to explain the colors with the orbiting cluster picture, since the colors of P1 are unlike any globular cluster or dwarf elliptical.  A further strike against the cluster picture is demonstrated in self-consistent N-body simulations by Emsellem \\& Combes (1997);  they find that the timescale for disruption is only $\\sim 10^5$ years, so it is unlikely that such a configuration would be observed.  T95's original model for an eccentric nuclear disk in M31 consists of three nested and aligned Keplerian ringlets with outwardly decreasing eccentricities.  Random velocities are roughly accounted for by convolution with a Gaussian point spread function in the plane of the sky. The model fits the photometry of L93 and is broadly consistent with the ground-based kinematics of Kormendy (1988) and Bacon et al. (1994). Though simple, the model predicts many of the asymmetries seen in the more recent kinematic profiles from SIS, FOC, STIS, and OASIS, including the displaced rotation center, the asymmetric rotation amplitudes, the low velocity dispersion at $r \\sim 1\\arcsec$ on the P1 side of the nucleus, and the presence of a dispersion spike near P2.  In its original form, the T95 model is too limited to be used to constrain the mass of the central BH in M31.  The model ignores self-gravity, which is necessary to maintain apse-alignment in the disk against differential precession (T95; Statler 1999, hereafter S99).  Also, the model does not include a realistic treatment of velocity disperison, which is needed for an accurate prediction of the dispersion profile. Both of these ingredients need to be included self-consistently.  Hints at how such a model can be constructed were first given by Sridhar \\& Touma (1999).  They compute orbits in nearly-Keplerian potentials with lopsided perturbations and find a family of periodic loop orbits elongated in the same sense as the perturbation.  They suggest that the nearly elliptical periodic parents of such orbits can be used as the backbone around which an eccentric disk with self-gravity and finite dispersion can be built.  S99 computes periodic loop orbits for a continuous, uniformly precessing T95-like disk model, and shows that the requirement of uniform precession has important consequences for the disk structure.  He finds that the periodic orbits follow a non-monotonic radial eccentricity distribution, in which a steep negative eccentricity gradient though the densest part of the disk is followed by a reversal of the arrangement of pericenter and apocenter with respect to the BH.  S99 suggests that approximate self-consistent equilibria can be constructed around such a sequence of numerically integrated closed periodic orbits, by approximating a distribution of quasi-periodic orbits about a given periodic parent with a distribution of Kepler orbits dispersed in eccentricity and orientation.  Salow \\& Statler (2001) use this approximation to construct radially truncated models that reproduce many of the features seen in FOC kinematics and one-dimensional HST photometry within $0\\farcs 5$ of the UV peak; the models are built by iteration with a phase space distribution function (DF) written in terms of the integrals of motion in the Kepler problem (S99).  They find that the backbone orbits follow an eccentricity distribution similar to that in S99, which gives rise to distinctive multi-peaked LOSVDs near the UV peak.  Several authors have constructed self-consistent eccentric disk models by other methods.  Jalali \\& Rafiee (2001) construct integrable models whose potentials are of the St\\\"{a}ckel form in elliptic coordinates.  They show that models with double nuclei are sustained by four general types of regular orbits (butterflies, nucleophilic bananas, horseshoes, and aligned loops).  Their models, however, require that both P1 and P2 have density cusps, which is not seen in the data.  B01 and Jacobs \\& Sellwood (2001) perform N-body simulations of lopsided ($m=1$) modes in a cold disk orbiting a central BH, and are able to find models that reproduce some of the observed features of the nucleus.  More importantly, they demonstrate that lopsided stellar disks can be long-lived, giving further support to the eccentric disk picture.  Sambhus \\& Sridhar (2002, hereafter SS02) construct models using a Schwarzschild-type method (Schwarzschild 1979) with an orbit library composed of both prograde and retrograde orbits. They find that the latter are needed to better fit the kinematics and photometry near P2;  Touma (2002) argues that a small percentage of retrograde orbits is all that is needed for a Keplerian disk to grow an unstable lopsided mode.  Both SS02 and B01 find an eccentricity distribution different than that found by S99 and Salow \\& Statler (2001). Orbits follow a steep negative eccentricity gradient through the dense part of the disk, but do not switch their apoapses to the anti-P1 side of the disk afterward.  Peiris \\& Tremaine (2003, hereafter PT03) have recently shown how a T95-like model can be extended to three-dimensions.  They construct models comprised of non-interacting Kepler orbits in the gravitational field of the BH.  They draw orbital elements from a Monte-Carlo scheme, and populate the disk with a parametric DF;  orbits are dispersed in eccentricity, orientation, and inclination, rather than just the first two, as in Salow \\& Statler (2001).  Their models are able to reproduce most of the important features in HST photometry and SIS and unpublished STIS (Bender et al.  2003) kinematics within $\\sim 1\\arcsec$ of the UV peak.  However, their models are missing self-gravity and gravity-induced precession in the disk.  In this paper we extend the self-gravitating, finite dispersion models of Salow \\& Statler (2001) to include a greater radial extent, in order to rigorously model the double nucleus of M31.  Along with an optimization routine, these models are used to fit FOC, STIS, and SIS one-dimensional kinematics, OASIS two-dimensional kinematics, and one and two-dimensional WFPC2/HST photometry.  Best-fit disk parameters and BH masses are found for a grid of 24 models by minimizing a chi-square merit function which assesses agreement between model and data.  The primary result of this paper is an accurate mass for the BH in M31.  Secondarily, we present the properties of the disk that best fits the nucleus.  The plan of this paper is as follows.  In Section \\ref{sec:CHAP2}, we give the details of model construction.  We then provide a description of the necessary assumptions and instrument specifications needed to find models that best fit data from the nucleus of M31 in Section \\ref{sec:CHAP3}.  In Section \\ref{sec:CHAP4} we present results from a grid of 24 best-fit models for M31's nucleus, including the BH mass in M31 and disk parameters and properties.  Section \\ref{sec:CHAP5} discusses the connection with other work.  Finally, Section \\ref{sec:CHAP6} presents some brief concluding remarks. ",
        "conclusions": "\\label{sec:CHAP6}  Our models of eccentric stellar disks around central black holes incorporate self-gravity, finite velocity dispersion, and gravity-induced precession, in a self-consistent way.  We have used these models to perform the first detailed fit to the nucleus of M31 which includes both one and two-dimensional kinematics and photometry;  the data set includes FOC, STIS, and SIS one-dimensional kinematics, OASIS two-dimensional kinematics, and one and two-dimensional WFPC2 photometry.  The primary result of this modeling effort is an accurate measurement of the mass of the central black hole in M31.  We find that $M_{BH} = 5.62 \\pm 0.66 \\times 10^7 \\msun$.  This value is consistent with the $M_{BH} - \\sigma$ correlation (Tremaine et al. 2002), which gives a value of $5.5 \\pm 1.5 \\times 10^7 \\msun$.  We find eccentric disks with large precession rates ($\\Omega = 36.5 \\pm 4.2 \\kmspc$) and small maximum eccentricities ($e_{max} = 0.21 \\pm 0.05$) for the backbone orbit sequence, $e_0(a)$.  The backbone orbits possess a characteristic non-monotonic distribution with a steep negative eccentricity gradient ($de/da < 0$) through the densest part of the disk, which gives rise to distinctive multi-modal LOSVDs for lines of sight near the central BH.  Such features may be used to further constrain model parameters when LOSVDs from upcoming high S/N STIS observations of M31's nucleus become available.  Although our models provide an accurate estimate for the BH mass, there is room for improvement.  We have assumed that the disk in M31 is thin. The disk may have non-negligible vertical structure, however, which could slightly alter our BH estimate.  Vertical velocity dispersion can be added to our models by including dispersions in inclination, $i$, and in the longitude of the ascending node, $\\Omega_n$, in our prescription for populating quasi-periodic orbits about the backbone orbit sequence.  We would then have a DF which includes all five integrals of motion in the three-dimensional Kepler problem; that is, $f(a,e,\\omega,i,\\Omega_n)$; PT03 demonstrate how the third dimension can be included in this way. Further improvement may require that new bulge models be considered, to better fit the behavior of the models within the central $0 \\farcs 4$ of the nucleus.  More flexible versions of $F(a)$ and a population of retrograde orbits can also help in this regard.  Retrograde orbits should be included self-consistently with their own DF and parameters, to determine if the location of the dispersion spike can be adjusted by their presence.  The eccentric disk picture for the nucleus of M31 is clearly the correct one, given the success that our models, along with those of other investigators, have at reproducing most of the asymmetric features in the kinematic and photometric data.  A complete description of the available data should be possible with models like ours, which already include most of the key physical ingredients, once they are extended as mentioned above.  Such models should be flexible enough to probe the connection between the BH and nuclear stars in greater detail, and may yield new clues about the formation of eccentric disks around BHs.  With the knowledge gained from detailed study of these models, other systems exhibiting features like those seen in M31, which do not have resolved Keplerian-dominated regions ($r_k$), can be investigated with confidence, as long as the sphere of influence of the BH ($r_h$) is resolved.  Galaxies with central properties similar to those in M31 are already known to exist.  NGC 4486B, a low-luminosity E1 companion of M87, is already known to possess two brightness peaks (Lauer et al. 1996).  This galaxy has no distinct nucleus, however, and its P1-P2 separation is $\\sim 10$ - $13$ pc ($\\sim 0\\farcs 15$ at $D=16$ Mpc), which is about six times that in M31;  also, its P1 and P2 have nearly the same brightess, unlike in M31.  Models like ours may be applicable, with some modification to the density structure to account for lack of a distinct nucleus.  Dynamical modeling of spectroscopic data from SIS/CFHT suggests that NGC 4486B harbors a BH of mass $\\sim 6 \\times 10^8 \\msun$ (Kormendy et al. 1997). Kormendy et al. find $\\sigma \\simeq 130 \\kms$, which, along with the aforementioned BH mass, implies that $r_h \\simeq 2\\arcsec$ (probably as an upper limit), which can be resolved by STIS/HST.  Further examples are shown by Lauer et al. (2002), who recently discovered six early-type galaxies with surface brightness profiles that decrease inward near their centers, reminiscent of the central dip in surface brightness found in M31.  These galaxies harbor torus-like brightness distributions, rather than double nuclei. Such structures can possibly be fit by a thin disk with low $\\epsilon / \\Omega$, which is only slightly asymmetric.  The sharpest structure observed in the sample is in the S0 galaxy NGC 3706, which has a bright stellar torus of radius $r \\sim 20$ pc ($0\\farcs 12$ at $D = 35$ Mpc).  Using the velocity dispersion observations from Carollo \\& Danziger (1994), the $M_{BH}$ - $\\sigma$ correlation gives $M_{BH} \\simeq 6 \\times 10^8 \\msun$ for this galaxy.  With this mass, $r_h \\simeq 0\\farcs 15$.  Modeling this galaxy will require observations with higher resolution than is currently available, but which may be available in the near future."
    },
    "0311/astro-ph0311059_arXiv.txt": {
        "abstract": "We discuss upcoming opportunities with cosmic microwave background (CMB) observations during the post-WMAP era. The curl-modes of CMB polarization probe inflationary gravitational waves (IGWs). While a significant source of confusion is expected from cosmic shear conversion of polarization related to density perturbations, higher resolution observations of CMB anisotropies can be used for a lensing reconstruction and to separate gravitational-wave polarization signature from that of lensing. Separations based on current lensing reconstruction techniques allow the possibility to probe inflationary energy scales below 10$^{15}$ GeV in a range that includes grand unified theories. The observational detection of primordial curl-modes is aided by rescattering at late times during the reionized epoch with optical depth to electron scattering at the level of 0.1 and above. An improved measurement of this optical depth is useful to optimize experimental parameters of a post-WMAP mission attempting to target the IGW background. ",
        "introduction": "The cosmic microwave background (CMB) is now a well known probe of the early universe. The temperature fluctuations in the CMB, especially the so-called acoustic peaks in the angular power spectrum of CMB anisotropies, capture the physics of primordial photon-baryon fluid undergoing oscillations in the potential wells of dark matter (Hu et al. 1997). The associated physics --- involving the evolution of a single photon-baryon fluid under Compton scattering and gravity --- are both simple and linear, and many aspects of it have been discussed in the literature since the early 1970s (Peebles \\& Yu 1970; Sunyaev \\& Zel'dovich 1970). The gravitational redshift contribution at large angular scales (e.g., Sachs \\& Wolfe 1968) and the photon-diffusion damping at small angular scales (e.g., Silk 1968) complete this description.  By now, the structure of the first few acoustic peaks is well studied with NASA's Wilkinson Microwave Anisotropy Probe (WMAP) mission (e.g., Spergel et al. 2003), while in the long term, ESA's Planck surveyor\\footnote{http://astro.estec.esa.nl/Planck/}, will extend this to a multipole of $\\sim$ 2000 with better frequency coverage. Beyond temperature fluctuations, detections of the polarization anisotropy, related to density or scalar perturbations, have now been made by the DASI experiment (Kovac et al. 2003), followed up with an improved study of the temperature-polarization power spectrum with WMAP (Kogut et al. 2003).  An interesting result from WMAP data is related to the very large angular scale polarization signal, which probes the local universe and associated astrophysics instead of physics at the last scattering surface probed with polarization at tens of arcminute scales. This local universe contribution arises when the universe reionizes again and the temperature quadrupole begins to rescatter to produce a new contribution to the polarization anisotropy at angular scales corresponding to the horizon at the new scattering surface (Zaldarriaga 1997). Such a signal was measured in the WMAP temperature-polarization cross-correlation power spectrum and was interpreted as rescattering with an  optical depth to electron scattering, inferred from the amplitude of the large angular scale CMB polarization, of $0.17 \\pm 0.04$ (Kogut et al. 2003).  Such an optical depth suggests early reionization; for example, if the universe reionized completely and instantaneously at some redshift, the measured optical depth suggests a reionization redshift of $\\sim$ 17 ($\\pm 5$). Such a high optical depth, in the presence of a $>$ 1\\% neutral Hydrogen fraction from z $\\sim$ 6 Sloan quasars (Fan et al. 2002), suggests a complex reionization history. A complex, and patchy, reionization, however, is expected if the reionization is dominated by the UV light from first luminous objects, though to explain the high level of the WMAP's optical depth requires high star formation efficiency at redshifts of order 25 with a possibility for a population of metal-free massive Pop III stars (e.g., Cen 2003).  In general, one can only extract limited information related to the reionization history from  polarization peak at tens of degrees or more angular scales (e.g., Hu \\& Holder 2003). One reason for this is the large cosmic variance associated with measurements at multipoles of $\\sim$ 10. To extract detailed information of the reionization process, one can move to temperature fluctuations generated during the reionization era at angular scales corresponding to few arcminutes and below. At these small angular scales, potentially interesting contributions from high redshifts include a scattering contribution related to moving electrons in ionized patches and fluctuations in the moving electron population (e.g., Santos et al. 2003), and a thermal Sunyaev-Zel'dovich effect, just as  in massive galaxy clusters containing hot electrons (Sunyaev \\& Zel'dovich 1980), related to the first supernovae bubbles (Oh et al. 2003). The same redshift ranges are expected to be illuminated in the near-infrared band between 1 to 3 $\\mu$m in the form of spatial fluctuations to the cosmic IR background (Cooray et al. 2003). Potentially interesting studies to extract more information related to the reionization process include cross-correlation studies between these small scale effects in CMB as well as other wavelengths. Upcoming experiments such as the South Pole Telescope and the Atacama Cosmology Telescope will provide some of the first possibilties in this direction.  \\begin{figure}[!h] \\centerline{\\psfig{file=cooray1.eps,height=20pc,angle=-90}} \\caption{CMB polarization anisotropies in the gradient (E) mode due to scalars and curl (B) modes due to the inflationary gravitational wave background (IGB, with the number in parenthesis showing the normalization in terms of the tensor-to-scalar ratio, which in terms of the energy scale of inflation is $2 \\times 10^{16}$ GeV (0.1), top curve, and $2 \\times 10^{15}$ GeV (10$^{-5}$), bottom curve). Note the effect due to reionization where rescattering produces new large angular scale anisotropies. For optical depths at the level of 0.1, the peak of the power related to IGB curl modes is at tens of degree angular scales and not at the degree scale ($l \\sim 100$) corresponding to the bump at recombination. Note that this enhancement, amounting to over two magnitudes in power at $l \\sim$ few is significant and will dominate the potential detectability of the IGB via polarization. The IGB detection is confused with cosmic shear conversion of  a fractional E-mode $\\rightarrow$ B-mode by the intervening large scale structure. The residual noise curve related to cosmic shear is after doing lens-cleaning in polarization maps with an arcminute scale CMB experiment with instrumental polarization sensitivity at the level of 1 $\\mu$K $\\sqrt{\\rm sec}$. For reference, we show an estimated level of contamination from dust and synchrotron emission (at 150 GHz);  this noise level can be improved with foreground cleaning in multifrequency data with sensitivity better than the foreground noise contributions. The curve labeled Planck is the noise curve related to upcoming Planck (HFI) polarization observations as a function of each multipole. } \\label{fig:cl} \\end{figure} ",
        "conclusions": ""
    },
    "0311/astro-ph0311053_arXiv.txt": {
        "abstract": "The possibility of a gaseous halo stream which was stripped from the Sagittarius dwarf galaxy is presented. The total mass of the neutral hydrogen along the orbit of the Sgr dwarf in the direction of the Galactic Anti-Center is $4 - 10 \\times 10^6$ \\Msun\\ (at 36 kpc, the distance to the stellar debris in this region). Both the stellar and gaseous components have negative velocities in this part of the sky, but the gaseous component extends to higher negative velocities. We suggest this gaseous stream was stripped from the main body of the dwarf 0.2 - 0.3 Gyr ago during its current orbit after a passage through a diffuse edge of the Galactic disk with a density $> 10^{-4}$ cm$^{-3}$.  The gas would then represent the dwarf's last source of star formation fuel and explains how the galaxy was forming stars 0.5-2 Gyr ago. ",
        "introduction": "Our Galaxy is in the process of consuming its small neighbors. Though this may sound rather grim, it is its only source of fuel and has allowed our Galaxy to grow into the large spiral it is today. One of the closest examples of a recognizable accreting satellite is the Sagittarius dwarf galaxy (28 kpc; hereafter Sgr dwarf; Ibata et al. 1994). The evidence continues to accumulate that contiguous streams of leading and trailing stellar debris are being pulled from the Sgr dwarf as it spirals into the Milky Way (e.g., Newberg et al. 2003), with the recent 2MASS data tracing this trail around the majority of the Galaxy (Majewski et al. 2003; hereafter M03).  The stars associated with the Sgr dwarf span a wide range of ages, with the youngest population between 0.5 - 3 Gyr old (Layden \\& Sarajedini 2000; Dolphin 2002; M03).  This indicates that within the past Gyr, and during its current orbit about the Galaxy, the Sgr dwarf was forming stars and had a source of star formation fuel.  Neutral hydrogen is a principal source of star formation fuel for a galaxy.  Galaxies which contain HI are commonly currently forming stars (e.g., Lee et al. 2002) and those without detectable HI tend to have primarily an older stellar population and thus appear to have exhausted their star formation fuel (e.g., Gavazzi et al. 2002).  This is evident in the dwarf galaxies of the Local Group.  The stars in the Local Group dwarfs vary from being almost entirely ancient ($>$ 10 Gyr; e.g., Ursa Minor) to a number of systems which are actively forming stars (e.g., WLM, Phoenix, LMC). The majority of the Local Group galaxies with HI have formed stars within the past 2 Gyr and those with no evidence for recent star formation do not contain detectable HI (e.g., Mateo 1998; Dolphin 2002; Bouchard et al. 2003).   HI observations of the central region of the Sgr dwarf ($\\alpha$, $\\delta$ = 19$^{h}$ 00$^{m}$, -30\\deg\\ 25\\arcmin\\ (J2000) indicate that our closest satellite galaxy does not currently contain a significant amount of star formation fuel (\\Mhi\\ $< 1.5 \\times 10^{4}$ \\Msun\\ (3$\\sigma$); Koribalski, Johnston \\& Otrupcek 1994).  The search for HI associated with the Sgr dwarf was continued by Burton \\& Lockman (1999); but they also found no associated gas over 18 deg$^2$ around the core of the Sgr dwarf with limits of \\Mhi\\ $< 7 \\times 10^{3}$ \\Msun\\ (3$\\sigma$).  These results are surprising considering the Sgr dwarf was forming stars within the last Gyr.  Since the orbit of the Sgr dwarf is approximately 0.7 Gyr (Ibata \\& Lewis 1998), one might expect to find this fuel stripped along the dwarfs orbit, possibly at a similar location to the stellar trail.  The trailing stellar tidal tail of the Sgr dwarf has recently been found to extend for over 150\\deg\\ across the South Galactic Hemisphere with a mean distance between 20 to 40 kpc from the Sun (M03).  Here we present HI data from HIPASS (HI Parkes All-Sky Survey\\footnote{The Parkes Telescope is part of the Australia Telescope which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO.}) along the entire Sgr dwarf galaxy orbit to investigate the possibility that a gaseous Sgr trail is also present. The gas detected represents a potential method of tracing the history, make-up, and classification of the Sgr dwarf galaxy, as well as the construction of our Galaxy. ",
        "conclusions": "The present orbit of the Sgr dwarf is estimated to be 0.7 Gyr (Ibata \\& Lewis 1998).   Using this orbit, the core of the Sgr dwarf was at the position of the HI complex noted here approximately 0.2 to 0.3 Gyr ago.  It would make sense if the gas was part of the Sgr dwarf within that timescale considering the age of the Sgr dwarf stellar population (0.5 - 3 Gyr; Layden \\& Sarajedini 2000; Dolphin 2002; M03). These gas clouds are along a sightline relatively close to the plane of our Galaxy, and at this position the Sgr dwarf stellar debris is approximately 45 kpc from the Galactic Center (M03). This distance is significantly beyond the typical radius quoted for our Galaxy ($\\sim$26 kpc), however it is possible that an extended ionized disk exists (e.g., Savage et al. 2003), as found in other systems (Bland-Hawthorn, Freeman \\& Quinn 1997). The passage through an extended disk of our Galaxy, in addition to the tidal forces already obviously at work as evident from the stellar tidal stream, might have been enough to disrupt the HI in the core of the galaxy and cause the dwarf to lose all remaining star formation fuel.  The gas will be stripped from a galaxy if $\\rho_{IGM} v^2 > \\sigma^2 \\rho_{gas} / 3$ (Mori \\& Burkert 2000; Gunn \\& Gott 1972). We can use this equation to estimate the density needed at the edge of the disk to strip the gas from the core of the Sgr dwarf.  We use a tangential velocity of 280 \\kms\\ for the Sgr dwarf and a velocity dispersion of 11.4 \\kms\\ (Ibata et al. 2003; Ibata \\& Lewis 1998).  If the column densities and size of the HI distribution in the core of the Sgr dwarf were on the order of $5 \\times 10^{20}$ cm$^{-2}$ (averaged over the core) and 1 kpc, the typical $\\rho_{gas}$ is 0.16 cm$^{-3}$. An extended disk density greater than $3 \\times 10^{-4}$ cm$^{-3}$ is then needed to strip the gas via ram pressure stripping.  Based on previous estimates of the density of the Galactic halo, this density would be easily achieved in the plane of Galaxy, 20 kpc from the currently observed edge of the disk (e.g., Sembach et al. 2003; Maloney 2003). Moore \\& Davis (1994) proposed a similar model to create the Magellanic Stream.    The gas that was most likely once part of the Sgr dwarf is the stream of clouds orientated along the orbit of the Sgr dwarf between $\\alpha = 2-4.5$h and v$_{\\rm LSR} =$ -380 to -180 \\kms~ (Figure 1). At 36 kpc from the sun this HI gas has a mass of $4 \\times 10^6$ \\Msun. This gas is at higher negative velocities than the carbon stars with velocity determinations in this region, but this is not unexpected with the addition of ram pressure forces to the tidal. If the filament extending perpendicular to the orbit is included as once being part of the Sgr dwarf, the total HI mass goes up to $10^7$ \\Msun. A typical total HI mass for a dwarf galaxy is on the order of $10^7$ to $10^8$ \\Msun\\ (Grebel et al. 2003), so this gas may represent anywhere from the majority to 10\\% of the total amount of gas once associated with the Sgr dwarf.  Using a total Sgr dwarf mass of $5 \\times 10^8$ \\Msun\\ (M03), this gas represents 1-2\\% of its total mass. The Sgr dwarf most likely originally had more than $10^6 - 10^7$ \\Msun\\ of neutral hydrogen associated with it, as the outer gaseous component would have been the first thing stripped from our closest satellite (e.g., Yoshizawa \\& Noguchi 2003). This outer gas would have had column densities between $10^{18} - 10^{19}$ cm$^{-2}$ and has most likely either already dispersed or been ionized, although remnants of this gas may be present as small HVCs along the Sgr orbit.   Since there is currently no HI associated with the core of the Sgr dwarf (Koribalski et al. 1994), and the dwarf has stars which are 0.5 - 3 Gyr old, the HI gas presented here most likely represents the high column density gas from the core of the Sgr dwarf which was finally stripped when the dwarf passed through the medium in the extended Galactic disk.  This is supported by the relatively high peak column densities of the HI gas currently ($\\sim 10^{20}$ cm$^{-2}$). For the Sgr HI gas to survive for 0.2-0.3 Gyr (the time since the last passage of the core Sgr dwarf), not to mention the Magellanic Stream which is thought to be $>$ 0.5 Gyr old, the gas should either be confined by an existing halo medium or associated with significant amounts of dark matter.  The association of HI gas with the Sgr stellar stream suggests that the dwarf spheroidal classification for the Sagittarius galaxy may need to be reconsidered. It also suggests that slightly offset HI and stellar streams may be a common feature of disrupted satellites in the Galactic halo.  With the accretion of this Sgr HI stream, the Magellanic Stream (Putman et al. 2003), Complex C (Wakker et al. 1999), and possibly other HVCs, there is ample fuel for our Galaxy's continuing star formation and understanding the distribution of stellar metallicities (e.g., the G-dwarf problem is not a problem). Determining the metallicity, distance, and ionization properties of the HI gas presented here will aid in confirming if this HI gas was indeed stripped from the Sgr dwarf during its current orbit.  Since we propose this is the last gas from the core of the Sgr dwarf, we will also look for molecular gas and dust associated with these clouds and complete simulations to determine the detailed parameters of gaseous stripping from the Sgr dwarf.  \\begin{figure} \\plotfiddle{putman.fig1.eps}{0pt}{0}{55}{55}{-190}{-180} \\vskip 2.4in \\caption{The negative high-velocity HI ($\\sim -85$ to $-400 $\\kms; LSR) along the orbit of the Sgr dwarf galaxy in Celestial coordinates.  The current position of the Sgr dwarf is shown by the solid point, and the orbit of the Sgr dwarf (Ibata \\& Lewis 1998) is plotted as the solid line through this point. The negative velocity gas attributed to the Magellanic Stream and Galactic Center is labeled, and the Galactic Plane is indicated by the solid line with the two dotted lines on each side representing $b = +5$\\deg\\ and $-5$\\deg. The negative velocity carbon stars extend from approximately $\\alpha, \\delta = 0^{h}, -20$\\deg\\ to $12^{h}, 20$\\deg\\ along the Sgr orbit (Ibata et al. 2001).  Contours represent column density levels of 0.5, 1.0, 5.0, and 10.0 $\\times 10^{19}$ cm$^{-2}$. }  \\end{figure}"
    },
    "0311/astro-ph0311579_arXiv.txt": {
        "abstract": "We present new velocity dispersion measurements for a set of 12 spiral galaxies and use them to derive a more accurate $v_c-\\sigma$ relation which holds for a wide morphological range of galaxies. Combined with the $M_{BH}-\\sigma$ relation, this relation can be used as a tool to estimate supermassive black hole (SMBH) masses by means of the asymptotic circular velocity. Together with the Tully-Fisher relation, it serves as a constraint for galaxy formation and evolution models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311265_arXiv.txt": {
        "abstract": "We present a study of the nearby post-merger giant elliptical galaxy, \\cena\\ (Centaurus A), in which we use the properties of its globular cluster (GC) and planetary nebula (PN) systems to constrain its evolution.  Using photometric and spectroscopic data for 215 GCs presented in Paper I, we study trends in age, metallicity, and kinematics for the GC system.  We confirm that the GC metallicity distribution is bimodal, and show that these two sub-populations have different properties.  Using spectral line index measurements of the brightest clusters, the metal-poor GCs have old ages like the Milky Way globular clusters, while the metal-rich GCs have \\Hb\\ line-strengths that could be interpreted as a mean age of $\\sim5^{+3}_{-2}$~Gyr.  Both populations appear to have [Mg/Fe] ratios consistent with that of the Galactic GC system, although this quantity is not very well-constrained.  The kinematics of the metal-rich GCs are similar to those of the planetary nebulae, exhibiting significant rotation about a misaligned axis, while the metal-poor GCs have a higher velocity dispersion and show a weaker kinematic correlation with the field stars.  The total gravitating mass of \\cena\\ derived from the GCs is in excellent agreement with the value derived from stellar (PN) kinematics.  We suggest that these and other data support a picture in which the main body of \\cena\\ was formed 3--8~Gyr ago by the dissipational merger of two unequal-mass disk galaxies supplemented by the continual accretion of both gas-rich and gas-poor satellites. ",
        "introduction": "The fossil record of old star clusters that is present in nearly every galaxy provides a window onto the vigorous epochs of star formation that mark its evolutionary history.  In nearby galaxies ($D\\lesssim20$~Mpc), where detailed investigations are possible, globular cluster systems can provide important leverage on the chemical and dynamical history of their hosts and their environment.  In spheroidal stellar populations such as elliptical galaxies and bulges, the observed color distributions of globular clusters (GCs) alone require either a significant amount of merging, multiple epochs of star formation (Forbes, Brodie, \\& Grillmair 1997), or both (Ashman \\& Zepf 1992).  The hierarchical merging that has been invoked to explain these systems can be purely dissipationless (C{\\^o}t{\\'e}, Marzke, \\& West 1998), or also involve gaseous bursts of star formation (Beasley \\etal 2002)  While much has been learned from the metallicity distributions of these GC systems, relatively little work has been done on their kinematics and their ages.  This is mainly because such studies are difficult for a large fraction of nearby galaxies, even with the new generation of 6--10 meter telescopes. These properties, however, offer to extend the power of GC systems as galaxy diagnostics by probing the dynamical states of galactic sub-components, and the time scales on which they formed.  Recently, there is evidence that while metal-poor GCs are universally old, the metal-rich GCs in ellipticals can have a large spread in age (Puzia 2002; Larsen \\etal 2003).  The story told by the few GC systems with substantial kinematic data is less consistent.  In general, the metal-poor GC sub-populations have a higher velocity dispersion, but rotation in either population varies on a galaxy-by-galaxy basis (Zepf \\etal 2000, C\\^{o}t\\'{e} \\etal 2001, Perrett \\etal 2002, C\\^{o}t\\'{e} \\etal 2003).  Globular clusters are in many ways ideal test particles for probing the dynamics of early-type galaxies.  While integrated long slit spectroscopy of the galaxy light is limited to $r<2r_e$ due to rapidly declining surface brightness, GCs are essentially unresolved and are easily observed at large distances from the galaxy center. As single-age, single-metallicity entities, GCs also offer the opportunity to correlate age and chemical enrichment with dynamics, providing a powerful probe of galaxy evolution.  Moreover, luminous ellipticals can host thousands of GCs, providing ample statistics. Although GCs are often considered ideal tracers of ancient star formation because they are simple stellar populations, they still only make up less than 1\\% of the light in a typical galaxy. Comparisons to the field star population are essential for any complete picture of galaxy formation.  One of the best galaxies for conducting an in-depth study is the nearby post-merger elliptical \\cena\\ (Centaurus~A).  At a distance of 3.5~Mpc (Hui \\etal 1993), \\cena\\ is the nearest giant elliptical galaxy, and provides a unique opportunity to investigate both the nature of spheroids and the effects of the merging process. In ellipticals, only for \\cena\\ has the metallicity distribution of the {\\it field} population been determined (W.Harris \\& G.Harris 2002).  Likewise, only a handful of galaxies have both globular cluster and planetary nebula (PN) kinematic data (Romanowsky \\etal 2002). In this paper, we present results based on the GC survey described in Paper I (Peng, Ford, \\& Freeman 2004b).  We combine these results with those presented in our survey of planetary nebulae (Peng, Ford, \\& Freeman 2004a) to try and elucidate the formation history of this interesting and important galaxy. ",
        "conclusions": "While one must always be cautious about trying to infer too much from a sample of one galaxy, detailed studies of \\cena\\ can teach us much about the formation of giant elliptical galaxies.  As the nearest gE galaxy, radio galaxy, and recent large merger, \\cena\\ is an accessible laboratory where we can study the processes that likely played important roles forming most elliptical galaxies at higher redshift.  The stellar populations and kinematics of \\cena\\ show its history to be consistent with a hierarchical scenario in which most metal-poor GCs formed early and merged with the galaxy or its progenitors over time.  The major event that shaped the galaxy as we see it today was likely a disk-disk merger, perhaps with unequal masses, that produced a disky, rapidly rotating remnant with a metal-rich GC population.  This occurred at intermediate times ($\\sim5$~Gyr ago, or $z\\sim0.5$), and may also be responsible for the gas currently in the galaxy, although a subsequent merger of another gas-rich system with the newly-formed elliptical is not ruled out.  The gas that resulted from this merger currently feeds the active nucleus, which in turn triggers star formation in the halo.  Future work that increases the number of known GCs, and obtains better ages and metallicities for all GCs will be important for corroborating or extending this formation picture.  As well, scenarios and simulations that provide chemodynamical predictions will be essential for making full use of these kinds of data sets. Ellipticals are likely to have diverse formation histories that are a function of their mass and environment.  It will be desirable to extend detailed studies of GC and PN systems to a larger sample of galaxies in the nearby universe."
    },
    "0311/astro-ph0311115_arXiv.txt": {
        "abstract": "Microlensing and pixel-lensing surveys play a fundamental role in the searches for galactic dark matter and in the study of the galactic structure. Recent observations suggest the presence of a population of old white dwarfs with high proper motion, probably in the galactic halo, with local mass density in the range $1.3\\times 10^{-4}-4.4\\times 10^{-3}~M_{\\odot}$ pc$^{-3}$, in addition to the standard galactic stellar disk and dark halo components. Investigation of the signatures on microlensing results towards the LMC of these different lens populations, with particular emphasis to white dwarfs, is the main purpose of the present paper. This is done by evaluating optical depth and microlensing rate of the various lens populations and then calculating through a Montecarlo program, the probability that a lens which has caused a microlensing event of duration $t_{\\rm E}$ belongs to a certain galactic population. Data obtained by the MACHO Collaboration allow us to set an upper bound of $1.6\\times 10^{-3}~M_{\\odot}$ pc$^{-3}$ to the local mass density of white dwarfs distributed in spheroidal models, while for white dwarfs in disk models all values for the local mass density are in agreement with observational results. ",
        "introduction": "MACHOs (Massive Astrophysical Compact Halo Objects) have been detected since 1993 in microlensing experiments towards the Large and Small Magellanic Clouds (LMC and SMC) as well as towards the galactic center. Although the events detected towards the SMC seem to be a self-lensing phenomenon \\cite{sahu,ss,st,gyuk}, a similar interpretation of all the events discovered towards the LMC looks unlikely \\cite{alcock2000}. Indeed, no purely LMC self-lensing models may produce optical depths as large as that reported by MACHO and EROS Collaborations. However, due to large uncertainties and the few events at disposal, it is not possible to draw at present sharp conclusions about the lens nature and localization. Indeed, the most plausible solution is that the LMC events detected so far are not due to a single lens population but to lenses belonging to different intervening populations: low mass stars in the LMC or in the galactic disk and MACHOs in the galactic halo. The average MACHO mass from observations results to be $\\simeq 0.5~M_{\\odot}$, \\cite{jms}, so that the white dwarf hypothesis looks to several authors as the best explanation for MACHOs. However, the resulting excessive metallicity of the halo makes this option untenable, unless the white dwarf contribution to halo dark matter is not substantial. \\cite{gm,binney} So, some variations on the theme of white dwarfs have been explored.  An option is that the galactic halo resembles more closely a minimal halo (i.e. the Milky Way is well described by a maximal disk model) rather than an isothermal sphere, in which case the average MACHO mass gets decreased so that most of them may still be brown dwarfs. In this connection, two points should be stressed. First, a large fraction of microlensing events (constituting up to $50\\%$ in mass) can be binary systems - much like ordinary stars - thereby counting as twice more massive objects. \\cite{depaolismnras,distefano} Second, MACHOs in the galactic halo could be brown dwarfs - with mass substantially larger than $\\simeq 0.1~M_{\\odot}$ \\cite{HANSEN} - since a slow accretion mechanism from cloud gas is likely to occur. \\cite{lcs} An alternative possibility has been pointed out: since the stellar initial mass function may change with the galactocentric distance, \\cite{taylor} it can well happen that brown dwarfs substantially contribute to the halo mass density without however dominating the microlensing optical depth. \\cite{ke}  On the other hand, it has been suggested that, in addition to main sequence stars, brown dwarfs and white dwarfs, also a population of black holes may exist in the Galaxy. These black holes may have been already observed in microlensing surveys. \\cite{bennett,quinn} Indeed, at least 6 extremely long events exhibiting very strong microlensing parallax signals have been detected by MACHO, GMAN and MPS Collaborations, indicating that a substantial fraction of the galactic lenses can be massive stellar remnants with masses up to $\\simeq 10~M_{\\odot}$.  Quite recently, faint blue objects discovered by the Hubble Space Telescope have been understood as old halo white dwarfs lying closer than $\\sim 2$ kpc from the Sun: \\cite{hansen,hansen01,ibata} they look as a good candidate for MACHOs. Moreover, very recently it has been found that the sample of local white dwarfs is largely complete out to 13 pc and that the local number density of white dwarf stars is $n_{\\rm wd}(R_0)=(5.0 \\pm 0.7) \\times 10^{-3}$ pc$^{-3}$. \\cite{holberg}  If these white dwarfs make a relatively important fraction of the galactic dark matter, as it seems, then they should also show up in the microlensing searches. In this context, we consider either a spheroid or a disk-like distribution of white dwarfs in addition to the two (thin and thick) stellar disks, the bulge and the standard halo components. \\cite{bss} Investigation of the signatures on microlensing results towards the LMC of the different lens populations, with particular emphasis to white dwarfs, is the main aim of the present paper.   To this end, we shall consider all available informations from gravitational microlensing experiments towards the LMC: optical depth, microlensing rate and event duration to try solving the problem of distinguishing among different lens populations. Note that we do not take into account the informations from the two observed parallax events towards the SMC which imply that these events are most likely self-lensing. From this fact one cannot infer that also all LMC events are self-lensing events since it is known that the SMC is extended along the line of sight (so that microlensing is dominated by self-lensing)  while there is little evidence that the LMC is similarly extended. Furthermore, since we are interested in the matter distribution in our galaxy and are not modelling the LMC mass distribution, we do not calculate the expected number of self-lensing events but use the estimation by \\cite{jms} to subtract these events from the total number of microlensing events given by the MACHO Collaboration. \\cite{alcock2000}  Regarding the problem of distinguishing between different lens populations, we also note that since white dwarfs lie relatively nearby the stellar disk, their transverse velocity should be different with respect to the typical transverse velocity of the halo MACHOs. This should give, in principle, a ``signature'' of the events due to white dwarfs in microlensing searches.  In Section 2 we review the current status of high velocity white dwarf observations from which the local white dwarf mass density value is estimated. In Section 3 we present the adopted mass distribution for the stellar, white dwarf and halo components while in Section 4 we briefly review the standard equations for the microlensing optical depth, event rate and event duration. In Section 5 we present our model results which, compared with microlensing observations towards the LMC, allow to constrain the white dwarf local mass density. Finally, in Section 6 we draw our main conclusions. ",
        "conclusions": "In the present paper we have considered, in addition to the standard stellar disk, bulge and halo lens populations, gravitational microlensing towards the LMC also from a population of white dwarfs distributed according to two different models: thick disk and spheroidal shape with various flatness values.  By comparing optical depth theoretical results and observational data towards the LMC we find for the local mass density of white dwarfs in the spheroidal distribution an upper limit of about $\\rho_{\\rm wd}(R_0)\\simeq 1.6\\times 10^{-3}~M_{\\odot}$ pc$^{-3}$. For white dwarfs distributed according to disk models, instead, available observations are unable to set an upper limit on the white dwarf local mass density value.  We have then calculated the expected microlensing rate from each lens population model, the average event duration $<t_{\\rm E}>_{\\rm i}$ and the expected number of events $N_{\\rm ev, i}$ for each galactic lens population by taking into account the detection efficiency given by the MACHO Collaboration. \\cite{alcock2001}  In the Table our main results for selected lens models are presented. As one can see, the expected number of events for white dwarfs in the spheroidal model results in the range $0.4-4.4$, for the lower and upper local mass density values, respectively. For white dwarfs in the disk model the expected number of events is $0.1-1.2$, for the lower and upper local mass density values. Clearly, if one considers that the total number of observed events is $13-17$ \\cite{alcock2000} and that $3-4$ events can be accounted for by self-lensing, \\cite{jms} the obtained results may be considered as consistent with observations. The event durations for the various adopted models are given in the Table. From this Table and Fig. \\ref{figeventi} it is clear how it is difficult to distinguish between the various lens populations, since the obtained distributions look similar. Therefore, microlensing observations alone cannot allow at present to understand the nature of the lens.  Indeed, in order to distinguish among different lens populations, one should eliminate the degeneracy in microlensing observations (which give as a result the event duration and the source amplification) since a direct measurement of the MACHO distance, mass or transverse velocity is generally not possible. The best way to solve the issue is to perform parallax observations of microlensing events since it allows a direct measurement of the MACHO distance to the observer. \\cite{gg,hk,soszy} On the other hand, this requires telescopes in orbit around the Sun, which is probably rather far in the future.  In any case, present and future microlensing and pixel lensing observations towards different directions - e.g. M31 \\cite{calchi} or other nearby galaxies - should allow to increase the available statistics of microlensing events towards several directions getting, through an accurate data analysis, more information about the possible presence of different lens populations in the Galaxy. Moreover, it will be possible to find an increasing number of microlensing events for which parallax measurements (or a direct imaging of the lens through the next generation of space-based telescopes) may allow to solve the parameter degeneracy allowing a direct estimate of the lens mass and distance."
    },
    "0311/astro-ph0311323_arXiv.txt": {
        "abstract": "We have discovered three globular clusters beyond the Holmberg radius in Hubble Space Telescope Advanced Camera for Surveys images of the gas-rich dark matter dominated blue compact dwarf galaxy NGC\\,2915.  The clusters, all of which start to resolve into stars, have $M_{V606} = -8.9$ to --9.8 mag, significantly brighter than the peak of the luminosity function of Milky Way globular clusters. Their colors suggest a metallicity $[{\\rm Fe/H}] \\approx -1.9$ dex, typical of metal-poor Galactic globular clusters.  The specific frequency of clusters is at a minimum normal, compared to spiral galaxies.  However, since only a small portion of the system has been surveyed it is more likely that the luminosity and mass normalized cluster content is higher, like that seen in elliptical galaxies and galaxy clusters.  This suggests that NGC\\,2915 resembles a key phase in the early hierarchical assembly of galaxies - the epoch when much of the old stellar population has formed, but little of the stellar disk.  Depending on the subsequent interaction history, such systems could go on to build-up larger elliptical galaxies, evolve into normal spirals, or in rare circumstances remain suspended in their development to become systems like NGC\\,2915. ",
        "introduction": "\\label{s:intro}  All galaxies with massive old stellar populations are thought to contain globular clusters (GCs).  They are particularly noticeable where the old stellar population is dominant such as in elliptical (E), dwarf elliptical (dE), and central Dominant (cD), as well as spiral galaxies with prominent bulges.  Dwarf spheroidal galaxies generally do not contain GCs, probably because they have insufficient mass to make their formation likely.  The exceptions are the most massive dwarf spheroidals Fornax \\citep{h61} and Sagitarius \\citep{igi94} which each have at least 4 GCs.  Disk galaxies dominated by population I stars contain fewer GCs per unit luminosity, presumably because of star formation in the disk after the formation of the population II component.  Galaxies with a high \\MHILB\\ ratio have yet to form much of their baryonic mass into stars.  They typically are blue and not considered likely hosts for populous GC systems.  NGC\\,2915 is an extreme gas rich galaxy having $\\MHILB = 1.7\\, \\Msun/L_{B,\\odot}$ \\citep[][hereafter MCBF96]{mcbf96}.  Its regularly rotating \\HI\\ disk extends to over 5 times beyond the readily detectable optical emission providing an excellent dynamical tracer for the mass distribution; not coincidentally NGC\\,2915 has one of the highest known mass-to-light ratios in a single galaxy (MCBF96). Furthermore, while its optical morphology is that of a blue compact dwarf \\citep[BCD;][hereafter MMC94]{mmc94}, its \\HI\\ disk clearly shows spiral arms which are not apparent in the optical.  In this Letter, we report the discovery of three luminous GCs found in Hubble Space Telescope Advanced Camera for Surveys \\citep[ACS;][]{ford_acs02} images of NGC\\,2915 that were obtained in order to look for a stellar heating source for the \\HI\\ disk.  That issue will be discussed in a separate article (Meurer \\etal\\ 2003; in preparation, hereafter Meu03). ",
        "conclusions": "\\label{s:disc}  \\subsection{Cluster formation efficiency}  Because of the blue core and gas rich nature of NGC\\,2915 we had not expected to find GCs in our images.  However, in retrospect this discovery should not have been surprising. GCs have previously been discovered around morphologically similar systems; \\citet{obr98} find a considerable population of at least 65 old GCs about the Blue Compact (albeit not dwarf) galaxy ESO~338-IG04.  In addition, NGC\\,2915 is not a dwarf system in terms of mass, and optical imaging shows that it to be dominated by an older stellar population for $R > 0.5$ Kpc (MMC94). These facts suggest that old GCs might have been expected.  Are the number of clusters found anomalous?  To address this we estimate the total number of clusters in the system.  Unfortunately, our images sample only a small fraction of the galaxy.  If the GCs are distributed spherically out to $R_{\\rm HI} = 10'$ (where the \\HI\\ distribution ends; MCBF96), then we have only surveyed 3.5\\%\\ of the available area.  Here we consider two cases for the possible distribution of old clusters. Case (1) assumes that we are lucky and have managed to observe all the clusters in NGC\\,2915.  While this is unlikely, it provides a strict lower limit to the cluster formation efficiency.  The more likely case (2) is that the globular cluster system is like that of more luminous systems - having a spherically symmetric power-law radial distribution in number per unit area - $N(R) \\propto\\ R^\\alpha$ \\citep{h91}.  We assumes this extends between $R = 0.5$ Kpc and $R_{\\rm HI}$, where the minimum $R$ insures that the predictions are finite and was chosen to correspond to the size of the central star forming population (MMC94). For $-2.5 < \\alpha < -0.5$ we estimate that the total number of clusters in the system is 7 to 12 times higher than found on our images.  For case (2) we adopt a correction factor of 9.  Hence there are at least 3 GCs in the system (case 1), with a more likely number being $\\sim 27$ (case 2).  In principle, we should correct the total GC estimate for the finite luminosity sampling of the images.  However, SExtractor \\citep{ba96} catalogs of our images show that we can detect extended objects down to $\\vacs \\approx 27$ which corresponds to an $M_{V606} \\sim -2$ for sources in NGC2915.  This is well below the peak of the GC luminosity function $M_{V} \\sim -7.5$ (Secker 1992; revised to agree with the {\\em Hipparcos\\/} RR Lyrae zero point, e.g. Carretta \\etal\\ 2000)\\nocite{s92,cgcfp00}.  Hence we apply no luminosity sampling correction. The fact that we only have found GCs much brighter than this peak is somewhat puzzling.  We will consider three measures of cluster formation efficiency.  First, the specific frequency, $\\SN = N_T\\, 10^{0.4(M_V + 15)}$, is the $V$ luminosity normalized cluster content \\citep{hvdb81}.  \\SN\\ is typically around 1 for spiral galaxies and increases towards earlier galaxy types, with E galaxies having $S_N \\sim 4$.  In the center of galaxy clusters \\SN\\ ranges from 2.5 to 12.5 \\citep{btm97,jpb99}. We consider two measures of the mass normalized contribution: \\etagc\\ the number of clusters per $10^9$ \\Msun\\ in dynamical mass (Blakeslee \\etal\\ 1997), and \\epscl, the fractional baryonic (gas and stars) mass in GCs \\citep{McL99}.  As done by \\citet{McL99} we assume each cluster has an average mass $\\langle {\\cal M}_{\\rm cl}\\rangle = 2.4 \\times 10^5\\, \\Msun$ when calculating \\epscl.  Blakeslee \\etal\\ (1997) find an average $\\etagc = 0.7$ with a scatter of 30\\%, while \\citet{McL99} computes an average $\\epscl = 0.0026$ with a 20\\%\\ uncertainty.  Table~\\ref{t:gcs} tabulates \\SN, \\etagc\\ and \\epscl\\ for NGC\\,2915 under the above two cases.  Relevant quantities used for these calculations are given in Table~\\ref{t:prop}.  Table~\\ref{t:gcs} also lists literature values and uncertainties of the efficiencies for ``normal'' systems.  The uncertainty of our estimates are large: for a Poissonian distribution yielding a count of 3, the 95\\%\\ confidence limits on the cluster formation efficiencies are 0.3 and 2.3 times the estimated value.  We find that at a minimum (case 1) \\SN\\ is close to normal for a spiral galaxy while while \\etagc\\ and \\epscl\\ are low compared to normal globular cluster systems.  In the more likely case (2), \\SN\\ is very high compared to normal gas rich systems.  \\etagc\\ and \\epscl\\ are also somewhat high compared to literature values, but in reasonable agreement considering the Poissonian uncertainty of our measurements.  While we can not rule out the possibility that NGC\\,2915 is like a normal spiral galaxy in terms of \\SN, it is more likely that it has a high luminosity normalized cluster content, whereas the mass normalized content is closer to normal.  \\subsection{A missing link in galaxy evolution?}  It is interesting to interpret these results within the framework of hierarchical evolution through comparison with other systems. \\citet{jpb99} has discussed how the \\SN\\ of the young Milky Way must have been fairly high, similar to the values for cluster E galaxies, after the formation of the Galactic halo but before stellar disk formation.  The increase in luminosity from later star formation in the disk would cause a decrease in the Galactic \\SN\\ to its present low value, while leaving the number of GCs per unit mass unchanged.  In NGC\\,2915 we have an example of a present-day galaxy with an old stellar component, including GCs, but a spiral disk that is still mainly in the form of gas.  In this sense, NGC\\,2915, like the centers of rich galaxy clusters, has an elevated value of \\SN\\ compared to typical spirals because of a lower efficiency for converting gas into stars.  Clusters typically have 2--5 times as much mass in gas as in stars \\citep{arbvv92}, similar to what we see in NGC\\,2915.  The difference, however, is that the gas in clusters now resides in the hot intracluster medium, while the gas in NGC\\,2915 still retains the potential to be converted into stars.  To a large extent, the fate of the gaseous disk (and the future evolution of \\SN\\ in NGC\\,2915) must depend on the surrounding environment.  Simulations show that when a disk galaxy enters the environment of a rich galaxy cluster, much of the gas in its disk and halo will be removed by the combination of tidal and ram pressure stripping \\citep{amb99,bcs02,bcdg01,g03}.  For example, if a system similar to the Galaxy were to fall into a rich cluster it would have its gas disk truncated down to a radius of a few kpc within a few tens of Myr by ram pressure stripping alone (Abadi \\etal\\ 1999).  Indeed, models of present day evolution of galaxy clusters invoke mildly truncated star formation in field galaxies accreting onto the cluster as the cause of the Butcher-Oemler effect \\citep{bnm00,kb01}.  It is likely that in their {\\em early evolution\\/} clusters were assembled from building blocks similar to NGC\\,2915 which had their ISM stripped from them by these mechanisms.  The stripped gas was then virialized to become the cluster X-ray halo.  In this scenario galaxy clusters have high \\SN\\ because they formed out of subclumps which had already efficiently formed star clusters but which never had the chance to form disks.  In a less hostile environment, a building block similar to NGC\\,2915 could go on to form a normal spiral galaxy.  If the \\HI\\ disk of NGC\\,2915 were to be perturbed enough to efficiently form stars in a fairly quiescent fashion this would result in an additional $7 \\times 10^8 \\Msun$ of stars forming with no additional GCs.  Assuming a $M/L_V \\sim 1\\, \\Msun/L_{V,\\odot}$ for the additional stars then the system would evolve towards $S_N \\sim 2$, a fairly normal value for spiral galaxies, while \\etagc\\ and \\epscl\\ remain fixed at their normal values. The \\SN\\ in NGC\\,2915 is then anomalously high because the formation of its disk has not proceeded, presumably due to a lack of external perturbations (MCBF96). If so, then we may expect other galaxies with high \\MHILB\\ ratios and extended \\HI\\ disks to also have a significant GC population, especially if they have a strong old population.  While this scenario seems compelling, we caution that it is possible that NGC\\,2915 only superficially resembles the building block we describe.  We have not proven that the number of clusters is anomalously high. Furthermore, we can not yet rule out the possibility that the clusters are of higher metallicity and younger than Galactic GCs.  This is shown in figure~\\ref{f:2cd} where we overplot \\citet{bc03} population models on the two color diagram illustrating the strong age-metallicity degeneracy for the filters used here.  If the clusters are not old, they may represent the remnants of a starburst occurring as recently as a few Gyr ago.  We are undertaking additional observations (imaging and spectroscopic) to get a more accurate census of star clusters in the NGC\\,2915 system and determine their nature.  The results of these studies should determine whether NGC\\,2915 is just a gas rich galaxy with a peculiar star formation history, or whether it truly presents us with a rare local view of how galaxies looked in the epoch of cluster assembly."
    },
    "0311/astro-ph0311609_arXiv.txt": {
        "abstract": "In this paper we investigate the properties of low X-ray--to--optical flux ratio sources detected in a wide area ($2.5 \\rm deg^2$) shallow [$f_X(\\rm 0.5 - 8 \\,keV) \\approx 10^{-14} \\,erg \\,s^{-1} \\,cm^{-2}$] XMM-{\\it Newton} survey. We find a total of  26 sources ($5\\%$ of the total X-ray selected population) with  $\\log f_X/f_{opt} < -0.9$ to the above flux limit.  Optical spectroscopy is available for 20 of these low X-ray--to--optical flux ratio objects. Most of them are found to be associated with Galactic stars (total of 8) and broad line AGNs (total of 8).  We also find two sources with optical spectra showing absorption and/or narrow emission lines and X-ray/optical properties suggesting AGN activity. Another two sources are found to be associated with low redshift galaxies with narrow emission line optical spectra, X-ray luminosities $L_X(\\rm 0.5-8\\,keV)\\rm \\approx 10^{41}\\,erg\\,s^{-1}$ and $\\log f_X/f_{opt} \\approx -2$ suggesting `normal' star-forming galaxies. Despite the small number statistics the sky density of `normal' X-ray selected star-forming galaxies at the flux limit of the present sample is low consistent with previous ROSAT HRI deep surveys. Also, the number density estimated here is in good agreement with both the $\\log N -\\log S$ of `normal' galaxies in the Chandra Deep Field North (extrapolated to bright fluxes) and model predictions based on the X-ray luminosity function of local star-forming galaxies. ",
        "introduction": "X-ray surveys with the {\\it ROSAT} and more recently with the {\\it Chandra} and  the XMM-{\\it Newton} observatories have demonstrated that the X-ray source  population is a heterogeneous mix of objects comprising (i) powerful AGNs/QSOs, (ii) galaxy groups/clusters, (iii) low luminosity AGNs, (iv) `normal' galaxies with X-ray emission dominated by stellar processes and (v) a small number of Galactic stars (Lehmann et al. 2001; Barger et al. 2002; Hornschemeier et al. 2003).  Powerful AGNs are identified in X-ray surveys by their enhanced X-ray--to--optical flux ratios   $\\log f_X/f_{opt} \\ga -1$. These sources are responsible for the bulk of the diffuse X-Ray Background (XRB) and are therefore important for constraining models on the XRB origin.  In addition to powerful AGNs, deep surveys with the {\\it Chandra} observatory reaching fluxes well below  $f(\\rm 0.5 - 2 \\,keV) \\la 10^{-15} \\,erg \\,s^{-1} \\,cm^{-2}$ have revealed large numbers of low X-ray--to--optical flux ratio sources $\\log  f_X/f_{opt} \\la -1$. These objects although too faint to significantly contribute to the XRB are detected in increasing numbers with decreasing flux and are likely to outnumber  powerful AGNs below  $f(\\rm 0.5 - 2 \\,keV) \\approx 10^{-17} \\, erg\\, s^{-1}\\, cm^{-2}$ (Hornschemeier et al. 2003). The low X-ray--to--optical flux ratio regime is thought to be populated by Low Luminosity AGNs (LLAGNs) and `normal' galaxies (Hornschemeier et al. 2003). Indeed, sources with $-2\\la \\log f_X/f_{opt}\\la   -1$ exhibit enhanced activity attributed to either starbursts or LLAGNs (Alexander et al. 2002; Bauer et al. 2002;  Hornschemeier et al. 2003; Georgakakis et al. 2003b). Even lower X-ray--to--optical flux ratios,  $\\log  f_X/f_{opt}  \\la -2$, are  believed to be quiescent Milky  Way type galaxies (Hornschemeier et al. 2002; Georgakakis et al. 2003a; Hornschemeier et al. 2003) although the presence of heavily  obscured AGNs cannot be excluded.  Deep Chandra surveys have provided valuable information on the nature of low $f_X/f_{opt}$ sources at fluxes $f(\\rm 0.5 - 2 \\,keV) \\la 10^{-17} - 10^{-16} \\,erg \\,s^{-1} \\,cm^{-2}$. These surveys are not however suitable for the study of low $f_X/f_{opt}$ sources at brighter fluxes  because of the {\\it Chandra} small field-of-view ($\\rm \\approx0.07\\,deg^{2}$). Indeed, due to the low surface density of these sources at bright fluxes wide area surveys are required to compile large statistically complete samples to elucidate their nature.  The {\\it ROSAT}  satellite with a field-of-view of $\\rm \\approx0.30\\,deg^{2}$ per pointing has probed the bright flux regime [$f(\\rm 0.5 - 2 \\,keV)  \\approx  10^{-15}\\,erg \\,s^{-1} \\,cm^{-2}$] but has identified only a small number of low X-ray--to--optical flux ratio sources (Griffiths et  al. 1995; Griffiths et al. 1996; Georgantopoulos et al. 1996; Lehmann et al. 2001). This can be attributed to (i) the poor positional accuracy of ROSAT rendering the optical  identification of X-ray sources difficult and (ii) its low sensitivity making wide area surveys to the limit $f(\\rm 0.5 - 2 \\,keV) \\approx 10^{-15}\\,erg  \\,s^{-1} \\,cm^{-2}$ time  consuming. The latter is particularly true due to the low surface density of $f_X/f_{opt}\\la -1$ sources to the flux limit above requiring large surveyed regions to compile statistically complete samples. Moreover, the insensitivity of ROSAT to hard X-rays ($>2.5\\rm\\,keV$) is a major drawback since heavily obscured AGNs (believed to populate the $\\log f_X/f_{opt}\\la   -1$ regime) are expected to emit most of their X-ray energy output in this spectral band.  In this paper we employ a shallow  (2-10\\,ks) wide area ($\\rm \\approx 2.5\\,deg^{2}$) XMM-{\\it Newton} survey to explore the nature of the $\\log f_X / f_{opt} \\la -1$  sources to the limit $f(\\rm 0.5 - 8 \\,keV) \\approx 10^{-14}\\,erg \\,s^{-1} \\,cm^{-2}$ similar to that probed by previous {\\it ROSAT} surveys. Compared to {\\it ROSAT}, XMM-{\\it Newton} has the advantage of significantly higher sensitivity over a wide energy range (0.2-10\\,keV) and improved positional accuracy ($\\la3.5$\\,arcsec; Hasinger et al. 2001; McHardy et al. 2003) facilitating the optical identification of X-ray sources. Our XMM-{\\it Newton} observations (hereafter referred to as the XMM/2dF survey) cover an area of $\\rm \\approx 2.5\\,deg^{2}$ much  larger than any previous {\\it ROSAT} surveys at the same flux  limit. This is essential to compile a large sample of low X-ray--to--optical flux ratio sources to explore their nature and to constrain the relative fraction of the different X-ray  populations. An advantage of the data presented here is that unlike deep Chandra samples (e.g. Brandt et al. 2001) limited by the small field--of--view,  the XMM/2dF survey has sufficient areal coverage to probe the low-$z$ ($z\\la0.1$)  universe. This is essential to explore the nature of the relatively nearby X-ray population and to provide a link between local ($\\rm <100\\,Mpc$) and more distant samples.  Moreover, our XMM/2dF survey has the advantage of high quality follow-up optical spectroscopic and photometric observations. Apart from our own spectroscopic campaign the  XMM/2dF overlaps with large scale spectroscopic programs: the Sloan Digital Sky Survey (SDSS), the 2dF Galaxy and QSO  Redshift Surveys (2dFGRS and 2QZ respectively). These unprecedented spectroscopic databases are complemented by a homogeneous set of multi-waveband ($ugriz$) photometric data from  the SDSS available for part of our XMM/2dF survey.  Section \\ref{sec_survey} presents the XMM/2dF survey, section \\ref{sec_xray} describes the reduction of the X-ray data while sections \\ref{sec_phot} and \\ref{sec_spec} detail the optical photometric and spectroscopic observations respectively. In section \\ref{sec_id} we outline the optical identification method while section \\ref{sample} presents the sample employed in this study. The results are discussed in section \\ref{discussion}.  Finally section \\ref{conclusions} summarises  our conclusions.  Throughout this paper we adopt $\\rm H_{o}=65\\,km\\,s^{-1}\\,Mpc^{-1}$, $\\rm \\Omega_{M}=0.3$ and $\\rm \\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "In this paper we employ a wide field ($\\rm 2.5\\,deg^{2}$) shallow ($f_X\\rm (0.5 - 8\\,keV)\\approx 10^{-14} \\, erg \\, s^{-1}$) XMM-{\\it Newton} survey to investigate the nature of low X-ray--to--optical flux ratio sources, $\\log f_X / f_{opt}<-0.9$. Our sample comprises 26 objects of which 20 have optical spectroscopic information. The majority of the sources without spectroscopic information have X-ray/optical properties that strongly suggest Galactic stars. Most of the spectroscopically identified systems are Galactic  stars (total of 8) and broad line AGNs (total of 8).  We also find four sources with  narrow emission and/or absorption line optical spectra. Two of them have X-ray/optical properties suggesting LLAGN activity. The remaining two sources have narrow emission line optical spectra, X-ray luminosities $L_X\\rm \\approx 10^{41}\\,erg\\,s^{-1}$ and X-ray--to--optical flux ratios $\\approx-2$ suggesting `normal' galaxies powered by star-formation activity.  The small number of `normal' galaxies found in our wide field  XMM/2dF survey is in agreement with the results from the 1\\,Ms ROSAT HRI survey of the Lockman Hole. Despite  the poor statistics the estimated number density of `normal' galaxies at the flux limit probed here is consistent with  the $\\log N / \\log S$ of deeper {\\it Chandra} surveys  extrapolated to bright fluxes. Using the X-ray luminosity function of local star-forming galaxies we find that the predicted $\\log N - \\log S$ is in fair agreement with both our shallow and the deeper {\\it Chandra} samples. The pure luminosity evolution model is in better agreement with the observations compared to the no-evolution prediction providing some evidence for X-ray evolution of star-forming spirals. However, the difference between the two models even at the flux limits of the 2\\,Ms Chandra Deep Field North survey is small and does not allow firm conclusions to be drawn. This is because of the low mean redshift ($z\\approx0.3$) of the `normal' galaxies probed by this ultra deep survey."
    },
    "0311/astro-ph0311489.txt": {
        "abstract": "We estimate energy spectra and fluxes at the Earth's surface of the cosmic and Galactic neutrino backgrounds produced by thermonuclear reactions in stars. The extra-galactic component is obtained by combining the most recent estimates of the cosmic star formation history and the stellar initial mass function with accurate theoretical predictions of the neutrino yields all over the thermonuclear lifetime of stars of different masses. Models of the structure and evolution of the Milky Way are used to derive maps of the expected flux generated by Galactic sources as a function of sky direction. The predicted neutrino backgrounds depend only slightly on model parameters. In the relevant 50 keV-10 MeV window, the total flux of cosmic neutrinos ranges between 20 and 65 ${\\rm cm}^{-2}\\,{\\rm s}^{-1}$. Neutrinos reaching the Earth today have been typically emitted at redshift $z\\sim 2$. Their energy spectrum peaks at $E\\sim 0.1-0.3$ MeV. The energy and entropy densities of the cosmic background are negligible with respect to the thermal contribution of relic neutrinos originated in the early universe. % In every sky direction, the cosmic background is outnumbered by the Galactic one, whose integrated flux amounts to 300-1000 ${\\rm cm}^{-2}\\,{\\rm s}^{-1}$. The emission from stars in the Galactic disk contributes more than 95 per cent of the signal. ",
        "introduction": "Almost all our current knowledge of the universe derives from the detection of electro-magnetic quanta of different energies. However, the hot dense regions lying at the centre of most astrophysical sources are completely opaque to photons, and their physical properties can only be inferred indirectly. % Additional information on astrophysical objects is gathered by studying the properties of cosmic rays. Yet, these are charged particles and the presence of galactic magnetic fields makes nearly impossible to trace them back to their sources. % Another limitation which applies to observations of both high-energy photons and charged particles is given by the so called Greisen-Zatsepin-Kuz'min effect, namely the creation of electron-positron pairs on the cosmic infrared radiation background (for photons) and of pions on the cosmic microwave background (for protons and heavier nuclei). In other words, the universe is optically thick to photons above $\\sim 10$ TeV and to protons above $\\sim 10 $ EeV on a scale of few tens of Mpc. % Neutrino astronomy, although still in its infancy, might provide a way to circumvent all these limitations. In fact, neutrinos are stable (so that they can cross long distances), weakly interacting (so that they are able to penetrate regions which are opaque to photons), and electrically neutral (so that their trajectories are not affected by magnetic fields). It is then reasonable to expect that neutrino telescopes would allow us to get spectacular new views of the universe.  Regrettably, the same reasons which make neutrinos interesting astrophysical probes make their detection extremely difficult. At present, only the Sun and Supernova 1987A have been detected through their neutrino emission. This is not surprising, since, as like as in the electro-magnetic visible band, the Sun is expected to largely outshine any other astronomical source.  Anyway, the detection of solar neutrinos has already led to important advances. By cross-checking the experimental results against the predictions of standard solar models it was realized that only a fraction (ranging between a third and two thirds depending on the particle energy) of the expected solar neutrino flux had been detected (see e.g. \\cite{BAH02} and references therein). A possible interpretation is that part of the solar electron neutrinos transform into different particles. This indication, demonstrated by the Sudbury Neutrino Observatory (SNO) detection of non-electron neutrinos from the sun \\cite{AHM01} and confirmed by the disappearance of reactor antineutrinos in the Kamioka Liquid Scintillator Anti-Neutrino Detector (KamLAND) \\cite{EAL03}, requires modifying the minimal standard model of electro-weak interactions so as to allow for neutrino flavor mixing and oscillations.  An important problem for the development of neutrino astronomy is the determination of sky backgrounds as a function of energy. Many different sources might contribute, ranging from distant galaxies to Galactic supernovae. % In this paper we estimate fluxes and energy spectra at the Earth's surface of the cosmic and Galactic neutrino backgrounds generated by thermonuclear activity in stars.  The nuclear fusion reactions which turn hydrogen into helium produce two electron neutrinos and release about 27 MeV of energy (including the neutrino budget). The generation of a solar luminosity by hydrogen-core burning then implies the emission of about $2\\times 10^{38}$ neutrinos per second. In consequence, we expect the existence of a diffuse and nearly isotropic neutrino background originating from all stars in the universe that are and/or have been thermonuclearly active. At the same time, the fact that we live within the disk of a spiral galaxy should imprint characteristic features on the angular distribution of the neutrino flux.  Despite the large effort aimed at estimating the amplitude and spectrum of the cosmic background generated by Type II supernovae \\cite{HW97,KSW00,AST03} and Gamma-ray bursts \\cite{NAL03}, the current literature still lacks of accurate predictions concerning the cosmic and Galactic stellar neutrino backgrounds. A first attempt in this direction has been performed by Hartmann et al. \\cite{HAL95} who, however, only considered main-sequence stars and used semi-analytic stellar models to compute the neutrino emission rates.  Accurate numerical determinations of the neutrino yields over the whole H and He burning phases have been presented by Brocato et al. \\cite{BAL98}. These results have been used by the same authors to estimate the present neutrino flux at the Earth's surface due to Galactic sources. Their Monte Carlo models give fluxes ranging from 3 to 24 ${\\rm cm}^{-2}\\,{\\rm s}^{-1}\\,{\\rm deg}^{-1}$ depending on Galactic longitude. A rough estimate of the cosmic background has been obtained by assuming that the present-day Milky Way is representative of the whole galaxy population at all times.  With this simple approach, Brocato et al. \\cite{BAL98} found a background number density of $n\\sim 10^{-9}$~cm$^{-3}$ and a flux $c\\cdot n\\sim$ 40~cm$^{-2}$ s$^{-1}$ (with $c$ the speed of light in vacuum). As a matter of fact, the solar neutrino flux at Earth ($\\sim 10^{11}$~cm$^{-2}$ s$^{-1}$) is orders of magnitude larger than both the Galactic and the cosmic neutrino backgrounds.  From the experimental point of view, it is presently impossible to discriminate the non-solar contribution from backgrounds in the detectors. For this reason, the cosmic and Galactic backgrounds will remain undetectable until technological advancements will make possible the construction of highly-efficient directional neutrino detectors at energies ranging between 0.1 and 15 MeV.  In this paper, the cosmic neutrino background is obtained by combining the most recent estimates of the cosmic star formation history and the stellar initial mass function with accurate theoretical predictions of the neutrino yields all over the thermonuclear lifetime of stars of different masses. We show that, even though this background is hardly of any cosmological relevance, its estimated amplitude, combined with observations of the $\\gamma$-ray background, sets an upper limit to the radiative lifetime of neutrinos. % By adopting state-of-the-art models of the structure and evolution of the Galaxy, we also derive maps of the expected flux generated by Galactic sources as a function of sky direction. % Beyond contributing to the ``noise\" level for the detection of single neutrino sources, these backgrounds are themselves interesting since they carry information about the past history and structure of the Galaxy and our Hubble volume.  Even though their detection is still beyond present-day technical capabilities, detailed predictions of their properties may motivate and inspire future experimental efforts. %  The plan of the paper is as follows. In Section 2 we discuss the cosmic stellar neutrino background. After introducing the notation and summarizing the basic properties of thermonuclear neutrino sources, we present our results and discuss their cosmological relevance and the related uncertainties. In Section 3 we estimate the amplitude of the Galactic neutrino background and derive maps of its variation on the sky. Our conclusions are listed in Section 4. ",
        "conclusions": "We have estimated two neutrino backgrounds on the sky, in the 50 keV-10 MeV window: the Galactic component from stars presently burning Hydrogen and Helium in the Milky Way and the cosmic contribution from thermonuclear reactions in other galaxies. On the grounds of the recent findings of Super-Kamiokande, SNO, and KamLAND, these neutrinos, born as $\\nu_e$, propagate as a superposition of (at least) two mass eigenstates. When (and if) interacting with matter, they will show up as neutrinos of all flavors, with comparable probabilities.  Depending only slightly on some model details (namely, the assumed functional form of the SFR and the IMF), the background intensity of the cosmic component peaks at $E\\simeq 0.1-0.3$ MeV, where $\\bar{I}^{(0)}_\\nu=3-7\\, {\\rm cm}^{-2}\\,{\\rm s}^{-1}\\,{\\rm MeV}^{-1}\\,{\\rm sr}^{-1}$, and sharply drops by 2 orders of magnitude above 1-2 MeV. The corresponding flux at the Earth's surface (integrated over all energies) ranges between 20 and 65 ${\\rm cm}^{-2}\\,{\\rm s}^{-1}$, while the associated energy density is $\\simeq 5\\times 10^{-4} \\,{\\rm eV}\\,{\\rm cm}^{-3}$. Assuming that the sum of the masses of the three neutrino flavors is at least 1/30 eV (as suggested by the oscillation patterns of solar and atmospheric neutrinos) implies that this energy density is negligibly small with respect to the contribution from the thermal ($T_{\\rm rel}=1.95$ K)  background of relic neutrinos originated in the early universe  ($u_{\\rm rel} > 3.7\\, {\\rm eV}\\,{\\rm cm}^{-3}$). Similarly, the number and entropy densities of cosmic thermonuclear neutrinos ($n\\simeq 10^{-9}\\,{\\rm cm}^{-3}$, $s\\simeq 10^{-7}\\,k_{\\rm b}\\,{\\rm cm}^{-3}$) are overwhelmed by the corresponding quantities for the primordial background ($n_{\\rm rel} \\simeq 10^2\\,{\\rm cm}^{-3}$, $s_{\\rm rel}\\simeq 10^{3}\\,k_{\\rm b}\\,{\\rm cm}^{-3}$).  Thermonuclear neutrinos reaching the Earth today have been typically produced by stars around $z\\sim 2$. Even though they have been around for a significant fraction of the life of the universe, they hardly interacted with anything. As far as we know, the thermonuclear neutrino background is absolutely transparent to any diffuse radiation impinging onto it since the mean free path $\\lambda =(n\\, \\sigma)^{-1}$ exceeds the Hubble length for any reasonable value for the interaction cross section $\\sigma$.  Attributing a fraction $f$ of the observed $\\gamma$-ray background in the sub-MeV region to neutrino radiative decay, one can derive a lower bound for the radiative lifetime, valid for any neutrino type and mass $m$, $\\tau_\\gamma/m \\geq (2-6)\\,f^{-1} \\times 10^{12}$ s ${\\rm eV}^{-1}$. This can be easily translated into a limit for the electromagnetic form factors, which, after accounting for the contribution of active galactic nuclei (i.e. by assuming $f\\sim 10^{-2}$), is competitive with the results obtained from the analysis of the SN1987A neutrino burst.    Independently of the direction of observation on the sky, the Galactic thermonuclear background dominates its cosmic counterpart. Even at high latitudes, Galactic neutrinos outnumber the cosmological background by a factor which ranges between a few and 10 depending on the details of the Galaxy model (see Figures \\ref{maps} and \\ref{phi}). % The integrated flux of Galactic neutrinos over the whole sky ranges between 300 and 1000 ${\\rm cm}^{-2}\\,{\\rm s}^{-1}$. %, compared to the cosmic contribution of %20-65 ${\\rm cm}^{-2}\\,{\\rm s}^{-1}$. % The emission from stars in the Galactic disk contributes more than 95 per cent of the signal.  Is there any prospect for detecting the Galactic component and, in the long term, to derive an observational neutrino map of the Galaxy? This clearly requires directional detectors, since one has first to distinguish the Galactic neutrino flux from the solar component. % In the last ten years, large Cherenkov detectors have been built and successfully operated (Kamiokande, Super-Kamiokande and SNO), recording the directional signal of target electrons scattered from the impinging neutrinos. Analyzing the angular distribution of the events with respect to the Sun direction, the peak due to solar neutrinos is clearly discernible over a flat background \\cite{FAL01}. So far, this method has been applied to the detection of the most energetic solar neutrinos originating from Boron decay (corresponding to a measured flux of $2.4\\times 10^{6}\\,{\\rm cm}^{-2}\\, {\\rm s}^{-1}$). In this case, the decay of $^{222}$Rn in water, natural radioactivity, reactor antineutrinos, and radioactive decay of muon induced spallation products account for most of the instrumental background. After applying all possible cuts, background counts above 5 MeV are about ten times the number of events from solar neutrinos. Even though this is not the optimal energy window for the detection of Galactic neutrinos, it is interesting to note that, for $E_\\nu>5$ MeV, their flux amounts to $\\sim 1$~cm$^{-2}$ s$^{-1}$, which is nearly 7 orders of magnitude smaller than the instrumental background. This shows how far we are from their detection.  Dedicated experimental designs and techniques should then be devised to detect cosmic and Galactic neutrinos. % Let us consider, for instance, the cosmic (anti)neutrino background from Type II supernovae. Kaplinghat et al. \\cite{KSW00} found an upper bound to the flux of 54~cm$^{-2}$ s$^{-1}$ while, more recently, Ando, Sato, \\& Totani \\cite{AST03} estimate $c\\cdot n\\simeq 11$~cm$^{-2}$ s$^{-1}$. Even though there is no energy region in which the contribution from supernova neutrinos dominates over the backgrounds, their presence might be detectable with Cherenkov telescopes as a distortion in the spectrum of electron and positrons produced in the decay of invisible muons. Ten years of integration at SuperKamiokande (or one at the proposed HyperKamiokande) should already guarantee a $1\\, \\sigma$ detection of the supernova neutrino background at $\\sim 20$ MeV \\cite{AST03}.  Over the next decade, the focus of solar neutrino experiments will shift gradually from the relatively high energy $^{\\mathrm 8}$B neutrinos to the low energy, $^{\\mathrm 7}$Be, p-p, and CNO neutrinos. This is an encouraging perspective for the detection of a Galactic signal since, for $E<15$ MeV, the number density of thermonuclear neutrinos in the cosmic and Galactic backgrounds is expected to exceed that generated by Type II supernovae. % Regrettably, instrumental backgrounds tend to increase in the MeV region. Therefore, in order to use the solar neutrino experiments to measure a Galactic signal, it is mandatory to suppress the noise in the detectors by several orders of magnitude by using more efficient screening procedures and accurate selection of radiopure materials. Note that extremely long times of integration might be anyway necessary to detect Galactic neutrinos with high statistical significance.  In summary, thermonuclear reactions in the Galaxy produce an anisotropic neutrino background in the MeV region. At the Earth's surface, this corresponds to an integrated flux of $\\sim1000 \\ {\\rm cm}^{-2}\\,{\\rm s}^{-1}$. Even though current detectors are unable to distinguish this signal from other backgrounds, it is not inconceivable that Galactic neutrinos might be revealed in the relatively near future with dedicated experimental setups.  \\begin{ack} %"
    },
    "0311/astro-ph0311445_arXiv.txt": {
        "abstract": "A combination of spectral and moments analysis of the continuous X-ray flux data is used to show consistency of statistical properties of the powerful solar flares with 2D BTW prototype model of self-organized criticality. ",
        "introduction": "The solar flares are the manifestation of a sudden, intense and spatially concentrated release of energy in the solar corona. According to the Parker's idea \\cite{parker} stochastic photospheric fluid motions shuffle the footpoints of magnetic coronal loops. The high electrical conductivity of the coronal gas implies that the magnetic field is frozen-in, so that the subsequent dynamical relaxation within the loop results in a complex, tangled magnetic field, essentially force-free everywhere except in numerous small electrical sheets which form spontaneously in highly-stressed regions. As the current within these sheets is driven beyond some threshold, reconnection sets in and releases magnetic energy, leading to localized heating. In this picture there exists a separation of timescales between energy input to the system (minutes to hours, for photospheric motions), and energy release (seconds to minutes, for reconnection and thermalization under coronal conditions). It means that the coronal loop is slowly driven by footpoint motions. Thus, the coronal field does meet the requirements for the appearance of self-organized criticality (SOC): self-stabilizing local threshold instability; open boundaries; and separation of time scales between driving and avalanching. Following to this picture the authors of Ref. \\cite{lh} suggested that SOC could describe the main features of the hard X-ray flares with avalanches of the reconnection events as the flares (see for further development \\cite{char}-\\cite{hug} and references therein). Power law probability density functions (PDF) observed for peaks, waiting and duration times of the X-ray flares present the main experimental support of this hypothesis. This experimental information, however, turned out to be non-sufficient to distinguish between the SOC and fluid (MHD) turbulence processes mixed in the solar corona \\cite{boff}. There exist substantial differences between SOC and turbulence mechanisms for explanation of the power law PDF's \\cite{boff}. Such differences can be ascribed to the lack of long time correlations for the avalanches in the monoscale SOC models. The distinction between the two mechanisms is however lost in the 2D BTW (Bak-Tang-Wiesenfeld) model \\cite{btw},\\cite{dhar} which obeys a peculiar form of multiscaling for the PDF's of several avalanche measures \\cite{mst}-\\cite{ms}. The 2D BTW model has long time correlated bursts and other turbulence-like intermittent properties if studied at the wave time scales. It is suggested in \\cite{ms} that such SOC models would be candidates to describe the solar flares. It should be also noted, that recent investigations \\cite{dp},\\cite{snc} show that the lack of exponential distribution for the waiting times between the solar flares (which is mentioned in Ref. \\cite{boff}) cannot be considered as a reason for discarding the SOC models as underlying dynamics. The long-range correlations of different nature can result in the non-Poissonian behavior of the waiting times for the SOC systems. The long-range correlations for the solar flares dynamics can have both internal and external reasons (a correlated drive related, for instance, to planetary motions affecting the Sun dynamics \\cite{snc}). The problem now is: How can one distinguish between fluid (MHD) turbulence and the SOC models with long-time correlations in the observed flares? We consider {\\it continuous} X-ray solar flux measured during a remarkable two-week period from 9 July to 23 July 2000 year. This period can be called as a flares festival: 98 C-class, 39 M-class and 3 X-class flares. For example, in analogous period of 1997 year the only one low-power B-class flare was observed. In the minute time scale we have more than 21000 data points presumably dominated by the powerful X-ray flares. We will use a combined spectral-moments analysis of the signal to extract information about underlying physical mechanisms. ",
        "conclusions": "The BTW prototype model of self-organized criticality is defined on a square lattice ($L\\times L$) box \\cite{dhar}.  On site $i$ $z_i=0,1,2,\\ldots$ is the number of ``grains''. If $z_i<4$, $\\forall i$, the configuration is stable. Grain addition to a stable configuration is made by selecting at random a site $k$ where $z_k \\to z_k+1$. If then $z_k\\ge 4$, toppling occurs, i.e. $z_k \\to z_k-4$, while each nearest neighbor, $l$, of site $k$ gets one grain ($z_l \\to z_l +1$). If $k$ is at the border grains are dissipated. The toppling of site $k$ may cause instabilities in the neighbors, leading to further topplings at the next microscopic time step, and so on. Thus, an avalanche made by a total number $s\\geq 0$ of topplings occurs before a new stable configuration is reached and a new grain is added. After many additions the system reaches a stationary critical state in which avalanche properties are sampled.  A key notion in the approach to the BTW is that of wave decomposition of avalanches \\cite{ikp}. The first wave is obtained as the set of all topplings which can take place as long as the site of addition is prevented from a possible second toppling. The second wave is constituted by the topplings occurring after the second toppling of the addition site takes place and before a third one is allowed, and so on.  The total number of topplings in an avalanche is the sum of those of all its waves.  The wave size has a pdf satisfying finite size scaling $P_w(s,L)\\sim s^{-\\tau_w} f_w(s/{L^{D_w}})$, where $\\tau_w=1$, $D_w=2$, and $f_w$ is a suitable scaling function.  The multiscaling of the dissipation field $$ \\epsilon_t=\\sum_{k=1}^t (s_{k+1}-s_k)^2 /t $$ was observed in the Ref. \\cite{ms} for the 2D BTW model and the multiscaling exponents for the BTW dissipation calculated in \\cite{ms} are given in figure 4 as triangles. One can see good correspondence between the exponents calculated in \\cite{ms} and the observed exponents.  The dissipation threshold $\\epsilon \\simeq 10^{-12}~ W^2/cm^4s^2$, for which the solar X-ray data agree with the 2D BTW model both for the flux spectrum and for the dissipation multiscaling, corresponds to transition from the low intensity M-class flares to the high intensity M-class flares. At present time we do not know why just at this level of the dissipation rate the SOC (presumably BTW) mechanism overcomes all other mechanisms and appears so clear in the X-ray flux data. In any way, we can see that the multiscale SOC takes place there and even appears as a dominating mechanism at a certain stage.\\\\  The authors are grateful to the NOAA for providing the data.         \\newpage"
    },
    "0311/astro-ph0311359_arXiv.txt": {
        "abstract": "Some attempts to predict the very distant future of Radio Astronomy are given. It is not easy to predict a list of the first priority problems which may appear, but the facilities potential is more predictable. It is suggested, that in addition to the \"dedicated for Radio Astronomy\", facilities may be extended greatly by integration with the next generation living standards facilities, connected with People-to-People communications through the global networks and by incorporating of  the \"Natural facilities\", such as grav. lensing,  maser amplification in the ISM etc. As an examples of the extreme cases of the  $10^9 m^2$ class of the new generation Radio Telescopes, utilization of the personal dipole size communication facilities by SKA type instrument, and array from the asteroids first \"Frehnel zones\" will be mentioned. Radio Astronomy from the secondary to optical facilities tool will be the only tool in the exploration of the $z>10$ Universe. The reality of all predictions depend mostly on the way, the Civilization will prefer : \"Ahead, to HOMO SAPIENCE\" or \"BACK TO PRIMATES\". ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311390_arXiv.txt": {
        "abstract": "Numerical simulations of the interiors of radiation dominated accretion disks show that significant density inhomogeneities can be generated in the gas. Here, we present the first results of our study on X-ray reflection spectra from such heterogeneous density structures. We consider two cases: first, we produce a number of toy models where a sharp increase or decrease in density of variable width is placed at different depths in a uniform slab. Comparing the resulting reflection spectra to those from an unaltered slab shows that the inhomogeneity can affect the emission features, in particular the \\fe\\ and \\oeight\\ Ly$\\alpha$ lines. The magnitude of any differences depends on both the parameters of the density change and the ionizing power of the illuminating radiation, but the inhomogeneity is required to be within $\\sim2$ Thomson depths of the surface to cause an effect. However, only relatively small variations in density (on the order of a few) are necessary for significant changes in the reflection features to be possible. Our second test was to compute reflection spectra from the density structure predicted by a simulation of the non-linear outcome of the photon bubble instability. The resulting spectra also exhibited differences from the constant density models, caused primarily by a strong 6.7~\\kev\\ iron line. Nevertheless, constant density models can provide a good fit to simulated spectra, albeit with a low reflection fraction, between 2 and 10~\\kev. Below 2~keV, differences in the predicted soft X-ray line emission result in very poor fits with a constant density ionized disk model. The results indicate that density inhomogeneities may further complicate the relationship between the \\fe\\ equivalent width and the X-ray continuum. Calculations are still needed to verify that density variations of sufficient magnitude will occur within a few Thomson depths of the disk photosphere. ",
        "introduction": "\\label{sect:intro} The reprocessing of hard X-rays in relatively cold, optically thick material is a major component of the current X-ray phenomenology of active galactic nuclei (AGN) and Galactic black hole candidates (GBHCs). The evidence for this is based on the shape of the observed spectral continuum. Starting with \\textit{Ginga} observations in the late 1980s and continuing on with data collected from \\textit{ASCA}, \\textit{RXTE} and \\textit{Beppo}SAX, it was found that the continuum of almost all type 1 Seyfert galaxies exhibit a spectral hardening beyond $\\sim$10~\\kev\\ \\citep{pou90,np94} before rolling over at energies $\\ga 100$~\\kev\\ \\citep{per02,ris02}. This fact combined with the near ubiquitous presence of a \\fe\\ line \\citep{np94,rey97} is completely consistent with the predictions of Compton reflection of X-rays in dense, optically thick material \\citep*{gf91,mpp91}. Moreover, the equivalent width (EW) of the \\fe\\ lines are usually large enough (EW $\\sim$ 50--300~eV) that the reflector must cover about half the sky as seen from the X-ray source \\citep{gf91}, but remain out of the line of sight (in order to not produce significant absorption). In the complex environment of the central engine of an AGN, there are many potential locations for the reprocessing gas, such as broad-line region clouds, and the obscuring material of the unification schemes. Indeed, reflection from one or both of these sites may be a common occurrence in Seyfert galaxies \\citep*[e.g.][]{yaq01,kas01,py03}. However, in at least some objects (most famously, \\mcg; e.g., \\citealt{fab02}) it was found that the \\fe\\ line exhibits a broad red-wing, consistent with originating from the inner regions of the accretion disk \\citep*{fab89,fab00,rn03}. Thus, in these sources, a detailed comparison of the X-ray spectrum with models of Compton reflection may yield information on the structure, metallicity, and ionization state of the accretion disk at small Schwarzschild radii ($R_{\\mathrm{S}} = 2 GM/c^2$, where $M$ is the black hole mass). Evidence for broad \\fe\\ lines in GBHCs is becoming more common \\citep*[e.g.,][]{mar02,mil02a,mil02b,mil02c,mil03}, so the possibility of performing similar analyses also exist for the Galactic accreting black hole systems.  Calculations of X-ray reflection from optically-thick material have increased in sophistication over the last decade. The initial Monte-Carlo calculations by \\citet{gf91} and \\citet{mpp91} assumed a neutral, constant density slab, and were able to quantify the observed strength of the \\fe\\ fluorescence line as a function of viewing angle and abundance. Soon after, \\citet{ros93} and \\citet{zyc94} allowed the gas to be ionized by the incoming X-rays and produced reflection spectra that included recombination lines and edges as well as fluorescence lines. \\citet{mfr93,mfr96} studied the effects of ionization on the \\fe\\ line, while \\citet{ros99} clearly showed that the entire reflection continuum changed shape as a function of the ionization parameter ($\\xi = 4\\pi F_{\\mathrm{X}} / n_{\\mathrm{H}}$, where $F_{\\mathrm{X}}$ is the incident X-ray flux, and $n_{\\mathrm{H}}$ is the density of the reprocessor). More recently, the assumption of constant density slabs has been replaced with more complicated density distributions, such as hydrostatic balance \\citep*{nkk00,nk01,brf01,roz02} or constant pressure \\citep{dum02}. The major difference between these variable density models and the constant density ones is that under certain conditions (when the Compton temperature of the radiation field is very high) the illuminated atmosphere can be subject to a thermal instability which can split the gas into an outer, low density, completely ionized zone and a deeper, high density, completely recombined zone \\citep*[cf.,][]{kmt81}. When this instability occurs the reflection spectrum exhibits only neutral features formed in the lower layer that are slightly broadened due to Compton scattering by the hot surface layer. Otherwise, the reflection spectra looks much like a diluted version of an ionized constant density model \\citep{brf01,dn01}. The features are weaker due to both Compton broadening by the hot surface gas, and the lower emissivity (proportional to density) of the emitting material.  Over the last few years it has been possible to fit models which self-consistently predict both the reflection continuum and line emission to real X-ray data. The constant density models of \\citet{ros93} have been applied to both narrow-line and broad-line Seyfert~1 galaxies \\citep*{bif01,orr01,der02,lon03} as well as GBHCs \\citep{mil03}. While these are not as sophisticated as the hydrostatic models, they are relatively quick to compute and provide a good parameterization of the data which allows easy comparison between objects. On the whole, when combined with a relativistic blurring function, the models do provide a reasonable description of the hard X-ray spectrum for many accreting black holes. The exceptions to this, such as the narrow-line Seyfert~1 1H~0707-495 \\citep{bol02} and \\mcg\\ \\citep*{bal03} may indicate that more complicated geometries are required than a simple flat accretion disk.  As interesting as these fits are, a lingering question remains: are the constant density or hydrostatic assumptions representative of real accretion disks? Where radiation pressure exceeds gas pressure, as likely occurs in the inner regions of AGN disks accreting at more than 0.3\\% of the Eddington rate \\citep{ss73}, models with stress assumed proportional to total pressure are viscously \\citep{le74}, thermally \\citep{ss76}, and convectively \\citep{bkb77} unstable.  These instabilities may lead to space or time variation in disk structure. Furthermore, the stresses driving accretion in ionized, differentially-rotating disks are now thought to be due to magnetic forces in turbulence resulting from the magneto-rotational instability or MRI \\citep{bh91}. In MHD simulations, gas motions in the turbulence vary over an orbital period \\citep{bh98}, which is similar to the timescale needed for establishing vertical hydrostatic balance \\citep*{fkr02}, Therefore, hydrostatic reflection models may not accurately represent the structure of turbulent accretion disks. Simulations of the turbulence which include effects of radiation diffusion indicate that density fluctuations may exceed an order of magnitude \\citep{tsks03}. Radiation-supported disks with a magnetic field are subject to another dynamical instability, the photon bubble instability or PBI \\citep{ar92,gam98}. Photon bubble modes having wavelengths shorter than the gas pressure scale height can grow faster than the orbital frequency \\citep{bs01,bs03}.  The non-linear development of the instability may lead to trains of shocks propagating through the disk surface layers \\citep{beg01}, leading to a way of producing super-Eddington luminosities \\citep{beg02}. Density contrasts between shocked and inter-shock regions can exceed one hundred (N.J. Turner \\etal\\ 2004, in preparation).  Will density fluctuations caused by the MRI and PBI affect the reflection spectrum? \\citet{fab202} and \\citet*{rfb02} suggested that they may strengthen the emission features and thus produce an apparent large reflection fraction. But several unknowns remain: since the reflection features are formed within 10 Thomson depths ($\\tau_{\\mathrm{T}}$) of the surface, what is the magnitude of any density change and where must it lie in the illuminated gas in order to alter the observed reflection spectrum? Do the accretion disk simulations produce the necessary fluctuations? The answers to these questions are crucial to the interpretation of X-ray spectra of accreting black holes.  In this work, we investigate the properties of X-ray reflection spectra from heterogeneous atmospheres. This paper (Paper 1) first considers a toy model where a density cut of variable depth and width has been sliced out of a constant density slab (\\S~\\ref{sect:toy}). This will allow a systematic investigation of the effects of a simple density inhomogeneity. We will then consider a more complicated structure produced by calculations of the photon-bubble instability in a radiation-dominated atmosphere with no turbulence (\\S~\\ref{sect:bubbles}). Reflection spectra calculated from all these structures will then be compared with the standard constant density ones. In \\S~\\ref{sect:discuss} we discuss the results from this paper and form our conclusions in \\S~\\ref{sect:concl}. In a companion paper (Paper 2), we apply the knowledge gained here to reflection by simulated radiation dominated accretion disks. ",
        "conclusions": "\\label{sect:concl} Our conclusions from this first paper in our study on reflection from heterogeneous accretion disks can be summarized as follows:  \\begin{enumerate} \\item The reflection spectrum from accretion disks can be altered by clumps or voids in the gas just beneath the photosphere. The effects range from negligible to large and depend on the structure of the inhomogeneity and the illuminating continuum.  \\item The greatest effects on the reflected emission for most ionization parameters occurs when the density change is within $\\sim 2$ Thomson depth of the surface.  \\item The change in density does not have to be large, even an increase or a decrease by a factor of a few can significantly alter the spectrum from that of a constant density model.  \\item Density inhomogeneities beneath the disk surface are a possible explanation for the apparent disconnectedness between the \\fe\\ line and the X-ray continuum.  \\item The soft X-ray emission lines (below 1~\\kev) are more sensitive to the presence of an inhomogeneity than the harder emission.  \\item Fitting the spectra produced from models with inhomogeneities with constant density models shows that a mixture of ionization parameters are present, which often results in reflection fractions smaller than unity.  \\item Constant density models may still be a good means of parameterizing reflection spectra, but only for energies $\\ga 2$~\\kev. They will have much more difficulty with broadband spectra that cover the soft X-ray lines emitted from the disk. However, in practice this may only be a problem for sources with weak or non-existent warm absorbers.  \\end{enumerate}  In Paper 2, we continue our investigation by computing reflection spectra from the predictions of 3-D radiation dominated accretion disk simulations. Future work will also include a consideration of the dynamical effects of the incident X-rays and multi-dimensional reflection."
    },
    "0311/astro-ph0311029_arXiv.txt": {
        "abstract": "We present followup observations of five hard X-ray sources from the ASCA 5-10 keV SHEEP survey, which has a limiting flux of $\\sim 10^{-13}$ erg cm$^{-2}$ s$^{-1}$. Chandra data have been obtained to improve the X-ray positions from a few arcmin to $<1''$, which allows unambiguous optical identification. While the objects almost certainly house AGN based on their X-ray luminosity, optical spectroscopy reveals a variety of properties. The identifications indicate that the SHEEP survey samples the same populations as deeper surveys which probe the origin of the X-ray background, but because the SHEEP sources are far brighter, they are more amenable to detailed followup work. We find a variety of classifications and properties, including a type II QSO, a galaxy undergoing star formation, and a broad-line AGN  which has a very hard X-ray spectrum, indicating substantial absorption in the X-ray but none in the optical. Two objects have X-ray/optical flux ratios which, were they at an X-ray flux level typical of objects in Chandra deep surveys, would place them in the ``optically faint'' category. They are both identified with broad line QSOs at z$\\sim 1$. Clearly this survey - which is relatively unbiased against obscured objects - is revealing a set of remarkable objects quite different to the familiar classes of AGN found in previous optical and soft X-ray surveys. ",
        "introduction": "\\label{Sec:Introduction}  The majority of the X-ray background in the 2-8 keV band has been resolved into discrete sources, thanks mainly to Chandra deep surveys (Mushotzky et al. 2000; Brandt et al. 2001; Giacconi et al. 2001; Cowie et al. 2002). The majority of the objects are almost certainly AGN, and if so they are more numerous than active galaxies found by other surveys. These newly discovered objects, many of which have hard X-ray spectra, are therefore the dominant AGN population and our understanding of black hole accretion, and its evolution through cosmic time, depends on their detailed study. Mutiwavelength followup of these objects (e.g. Mushotzky et al 2000; Barger et al. 2001, 2002; Rosati et al. 2002) has shown that they would have been very difficult to identify as AGN without the X-ray observations. The reasons are twofold: 1) Spectroscopic followup in the optical shows many objects have little or no evidence for the high excitation and/or broad emission lines hitherto considered the defining characteristic of AGN and 2) Many are too faint optically for spectroscopic followup at all. Nonetheless, the X-ray luminosity of these objects, typically $L_{\\rm X} \\sim 10^{43}$~erg s$^{-1}$ (Cowie et al. 2003), is strongly indicative of AGN activity.  The precise nature of the X-ray populations are therefore in doubt, but the best bet is that they are obscured AGN at moderate redshift (Alexander et al. 2001). Definitive answers will not be obtained unless they are observed spectroscopically with larger telescopes, or we can select brighter examples for study with the current generation of ground-based facilities. The most promising way currently is to investigate the bright end of the new X-ray populations is by using large area surveys such as the BeppoSAX HELLAS survey (Fiore et al. 1999; 2001; Comastri et al. 2001) and our own equivalent  with ASCA, SHEEP (the Search for the High Energy Extragalactic Population; Nandra et al. 2003, hereafter paper I). The hard detection bandpass (5-10 keV) is relatively unbiased to absorption, and the wide area means rarer, bright objects are found.  The big disadvantage of the HELLAS and SHEEP survey data is the relatively poor angular resolution of the BeppoSAX and ASCA telescopes, which make it very difficult to identify the X-ray sources unambiguously. To improve on this, we are obtaining Chandra data for a subset of the SHEEP objects, which should give arcsec position and therefore secure optical counterparts for optical spectroscopic and other followup. The first results from this program are the subject of this Letter. ",
        "conclusions": "X-ray surveys are the most efficient way of finding AGN, and hard X-rays in particular have great importance as they are not strongly biased against obscured objects unless the column densities are huge. For example, in the SHEEP detection band of 5-10 keV a column density of $10^{23}$~cm$^{-2}$ (equivalent to an extinction of $A_{\\rm V} \\sim 50$) only depresses the flux by $\\sim 15\\%$ at $z=0$, and even less at higher redshifts. Thus the content of the survey should more closely represent the true AGN population than, say, optical or soft X-ray surveys. Chandra and XMM do the same, but for fainter sources. Our preliminary results show that the SHEEP objects bear little resemblence to standard AGN classes.  We find S1 to be very hard in the X-ray but shows broad optical lines. Two similar objects were found in Paper 1, FBQS J125829.6+35284 and CRSS J1429.7+4240  (see also Comastri et al. 2001; Georgantopoulos et al. 2003), hinting that this class of object - optical type 1s with heavily absorbed X-ray spectra - may be very common. The implication is that the obscuring material is dust free and/or very close to the nucleus. This is also the situation in some bright local AGN (e.g. NGC 4151 - see paper I). Our survey is showing that this phenomenon could be quite common in more typical AGN making up the X-ray background, and which may evolve differently to unobscured, broad line QSOs. This suggests that the physical processes very close to the central black hole in galactic nuclei evolve with cosmic time. This fascinating possibility clearly deserves further investigation.  S2, also a broad line QSO, must also be obscured, based both on its Chandra colour and the fact it was observed but not detected by ROSAT. What is perhaps most interesting, however, is its relative optical faintness. The detection of a broad line in the rest-frame UV shows that the object cannot be heavily obscured in the optical, and the object is therefore faint simply due to a high \\fxfopt\\ ratio. Indeed, a QSO with the same \\fxfopt\\ ratio detected at the flux limit of the Chandra 1Ms surveys would have R$>26$, making it very difficult to identify spectroscopically. This shows the power of the SHEEP survey in finding brighter objects which can be studied in more detail.  It also hints that many of the \"optically faint\" sources found in the Chandra deep surveys (Mushotzky et al. 2000; Alexander et al. 2001) may not have particularly exotic properties.  This conclusion is strengthened when we consider S3. The optical spectrum is of poor quality, but if our identification is correct it is also a broad-line QSO.  Again one might classify this type of object as \"optically faint\" if it were detected in a deep survey.  The \\fxfopt\\ ratio based on the Chandra flux implies R=25 at the CDF-S flux limit. If the higher ASCA flux is more representative we would expect the object to appear at R$>28$ in a Chandra 1Ms survey. One must also ask why the ASCA and Chandra fluxes are so discrepant. Watanabe et al. (2002) concluded that large amplitude variability was at play when they observed similar behaviour in Chandra followup of two hard ASCA sources. To conclude that in our case would be premature, but we will be able to much more about this phenomenon when we have an analysis of the whole SHEEP sample. For example, if some SHEEP objects were spuriously detected (e.g. due to Eddington bias) we should observe the discrepancy only in the least significant sources, like S3. Otherwise the prospect of widespread transience or outburst in hard X-ray selected objects is most interesting, and would impact to a degree on the conclusions drawn from all flux-limited X-ray surveys.  According to its optical spectrum, S4 is intermediate between an AGN and an H{\\sc ii} galaxy. A  naive conversion from flux to luminosity would then make it the most X-ray luminous starburst known (c.f. Zezas, Georgantopolous \\& Ward 1998; Moran et al. 1999). The very hard X-ray colours argue for substantial obscuration too, making the intrinsic luminosity even higher. It is perhaps more likely that the object hosts an obscured accreting black hole. Substantial obscuration can occur in the star-forming regions, blocking soft X-rays optical light from the nucleus.  The \"intermediate\" classification may also be in large part due to contamination of the AGN by the (extended) galaxy. Indeed, our photometry indicates that the whole galaxy is almost 2 magnitudes brighter than a 2\" region centered on the nucleus (roughly where the spectrum was obtained). A similar object observed at higher redshift would presumably therefore be completely dominated by stellar light. This dilution may go some way toward accounting for the fact that deeper surveys reveal many strong X-ray sources without AGN signatures (e.g. Mushotzky et al. 2000).  Our final object, S5, is a candidate member of a long-sought after population: type II QSOs. These have been known to exist for quite some time (e.g. Hines \\& Wills 1993) but excitement, and controversy, over their numbers have surrounded them due to their possible contribution to the X-ray background. They have been found in Chandra surveys (Stern et al. 2002; Norman et al. 2002), but they do not appear to be particularly common. On the other hand, Steidel et al. (2002) have shown that narrow-line objects are the dominant population of AGN in faint UV-selected galaxy samples at high redshift. S5 is also interesting because its X-ray colors from ASCA show relatively little obscuration: certainly less than the broad-line QSO S1 (see also Pappa et al. 2001). This emphasises an emerging phenomenon: that in hard X-ray surveys the X-ray and optical spectroscopic properties of AGN show almost no correspondence. This suggests it is time to revisit AGN classification schemes based on low redshift Seyferts: they are likely to tell us very little about black hole accretion at higher redshift.  The SHEEP survey was designed to pick out bright examples of objects which make up the X-ray background. The survey seems to be fulfilling that promise, but the first indications are that these objects are quite unlike classical AGN. Given the relative lack of bias in hard X-ray selected samples, and their efficiency in finding accreting black holes, this indicates that much of what we think we know about what makes an AGN needs to be revised. Robust statistical conclusions await the results of the full survey, which should reveal much about the underlying AGN population."
    },
    "0311/astro-ph0311503_arXiv.txt": {
        "abstract": "We study the properties of cosmological density perturbations in a multi-component system consisting of a scalar field and a perfect fluid. We discuss the number of degrees of freedom completely describing the system, introduce a full set of dynamical gauge-invariant equations in terms of the curvature and entropy perturbations, and display an efficient formulation of these equations as a first-order system linked by a fairly sparse matrix. Our formalism includes spatial gradients, extending previous formulations restricted to the large-scale limit, and fully accounts for the evolution of an isocurvature mode intrinsic to the scalar field. We then address the issue of the adiabatic condition, in particular demonstrating its preservation on large scales. Finally, we apply our formalism to the quintessence scenario and clearly underline the importance of initial conditions when considering late-time perturbations. In particular, we show that entropy perturbations can still be present when the quintessence field energy density becomes non-negligible. ",
        "introduction": "The material content of the universe is commonly assumed to be a mixture of fluids, such as radiation or non-relativistic matter, and scalar fields, either driving a period of early universe inflation \\cite{infbook} or playing the role of dark energy (quintessence) in the present universe \\cite{quint,RP,CDS}. The latter possibility has motivated a number of works devoted to the study of cosmological perturbations in a multi-component system consisting of fluids and a scalar field, for instance Refs.~\\cite{RP,CDS,quintpert,nonminquipert,PB,BMR,AF,SA,ML2,DMSW}. Nevertheless, the literature contains some contradictory statements concerning the properties of such perturbations.  In this paper we aim to resolve these discrepancies and will provide a comprehensive analysis of the problem. We will study the role of intrinsic entropy perturbation in the scalar field, and whether the notion of adiabaticity is preserved by the dynamics of the multi-component system when the evolution of such an intrinsic entropy perturbation is explicitly accounted for. In the process of doing so, we will discuss the number of degrees of freedom which completely describe the system and we will find a highly-efficient formulation of the perturbation equations including the effects of spatial gradients.  Finally, using our formalism we will specifically discuss the quintessence scenario. We will correct some common misconceptions, discuss the evolution of entropy perturbations, and clearly show the importance of initial conditions when considering late-time perturbations. In particular, we will show that entropy perturbations can be enhanced by the evolution of the field and may still be present when its density is no longer negligible. This is an important result which has generally been overlooked when studying dark energy models. ",
        "conclusions": "We have explored the nature of scalar perturbations for a universe filled with both a scalar field and a perfect fluid. We have introduced a useful set of variables and have provided a full analysis including spatial gradients. In particular, we have focused on the isocurvature perturbation modes and on the degrees of freedom which completely characterize the system. While for the case where only the scalar field is present its intrinsic entropy perturbation is forced to vanish at linear order for superhorizon scales, the presence of a fluid --- even if sub-dominant --- allows the possibility for such an intrinsic contribution to be present on large scales. However, in the case of a very sub-dominant fluid the intrinsic entropy of the scalar field decays, dynamically recovering the single scalar field situation.  We have recast the basic evolution equations in a rather simple matrix formalism in terms of the gauge-invariant variables for the adiabatic and isocurvature perturbations, taking into account the dynamics of the perturbations when a given wavelength re-enters the horizon. In particular, we have obtained an equation for the intrinsic entropy perturbation which shows that, on large scales, an initial adiabatic condition is indeed preserved, regardless of the evolution. Only when the perturbations approach the horizon are the adiabatic and entropy perturbations fully coupled together.  Finally, we have applied our formalism to the quintessence scenario. In this case we have analyzed the large-scale evolution of the adiabatic and entropy perturbations in the different regimes which the quintessence scalar field dynamics may feature. As is well known, entropy perturbations are suppressed during the tracking regime, but, as already shown in Ref.~\\cite{ML2}, during the kinetic regime one entropy mode undergoes an exponential growth. We have shown that it is then exactly compensated by an exponential decay during the first potential regime and then remains constant during the second potential regime. Therefore, after the three regimes preceding the tracker regime, entropy modes are neither enhanced nor suppressed. However, we discussed the remaining possibility of an enhancement during the short transition phase preceding the kinetic regime. We have studied two special cases numerically and have confirmed our analytical analysis. Moreover, we have observed that, in one of the cases, at the beginning of the tracker regime, entropy perturbations are larger than initially, and therefore we have concluded that entropy mode enhancement is possible.  To summarize, we have shown that in general it is incorrect to assume that the observational imprint of quintessence perturbations will be independent of the initial conditions, because entropy perturbations can still be present when the quintessence energy density is no longer negligible. Note that this can happen as soon as the tracker starts, long before quintessence domination. In this case, entropy perturbations would feed curvature perturbations, but then would slowly decay to become negligible today."
    },
    "0311/hep-ph0311022_arXiv.txt": {
        "abstract": "In this paper we consider the possibility of neutrino mass varying during the evolution of the Universe and study its implications on leptogenesis. Specifically, we take the minimal seesaw model of neutrino masses and introduce a coupling between the right-handed neutrinos and the dark energy scalar field, the Quintessence. In our model, the right-handed neutrino masses change as the Quintessence scalar evolves. We then examine in detail the parameter space of this model allowed by the observed baryon number asymmetry. Our results show that it is possible to lower the reheating temperature in this scenario in comparison with the case that the neutrino masses are unchanged, which helps solve the gravitino problem. Furthermore, a degenerate neutrino mass patten with $m_i$ larger than the upper limit given in the minimal leptogenesis scenario is permitted. ",
        "introduction": "\\subsection{The model}  We extend the minimal seesaw model by introducing a new coupling between Quintessence scalar field, $Q$, and the RH neutrinos. The Lagrangian relevant to leptogenesis is given by \\begin{equation} \\mathcal{L}=\\mathcal{L}_{lep}+\\mathcal{L}_Q\\, , \\end{equation} with \\begin{equation} -\\mathcal{L}_{lep}= Y_{ij}\\bar{L}_i\\tilde{H} N_j + \\frac{1}{2}M_i(Q)N^T_iCN_i + h.c.\\, , \\end{equation} and \\begin{equation} \\mathcal{L}_Q=\\frac{1}{2}\\partial_\\mu Q\\partial^\\mu Q-V(Q)\\, , \\end{equation} where $Y_{ij}$ is the Yukawa coupling of the RH neutrinos, $M_i(Q)$ is the Majorana masses of the RH neutrinos which is now a function of the value of $Q$, and $V(Q)$ is the potential of the Quintessence. We have taken the basis that the RH Majorana mass matrix is diagonal. The form of $V(Q)$ for a specific model and an explicit form of $M(Q)$ will be given in Sec. IV.  The Quintessence dependent masses of the right-handed neutrinos give rise to different Majorana masses at the early epoch from that at the present epoch, which is the key for the difference of our model from the usual leptogenesis scenario.  \\subsection{The CP asymmetry}  Generally, the Yukawa coupling $Y$ contains CP violating phases and lead to different branching ratios for $N$ decays into lepton and antilepton. The asymmetry for $N_1$ decays is\\cite{early,vari} \\begin{eqnarray} \\epsilon_1 &=& \\frac{\\Gamma(N_1\\to LH) - \\Gamma(N_1\\to \\bar{L}\\bar{H})}{\\Gamma(N_1\\to LH) + \\Gamma(N_1\\to \\bar{L}\\bar{H})} \\nonumber \\\\ & \\approx & -\\frac{3}{16\\pi} \\frac{1}{(Y^\\dagger Y)_{11}} \\sum_{i=2,3} \\text{Im}\\left[ (Y^\\dagger Y)_{1i} \\right]^2\\frac{M_1}{M_i}\\ , \\;\\; \\text{for} \\;\\; M_1\\ll M_2, M_3 \\; . \\end{eqnarray}  Since the light neutrino masses are generated by the seesaw mechanism\\cite{seesaw} and constrained by the neutrino oscillation experiments, the CP asymmetry is also constrained by the neutrino data. In the case of hierarchical neutrinos, i.e., $m_3\\approx \\sqrt{\\Delta m_{atm}^2} \\approx 0.05 eV \\gg \\sqrt{\\Delta m_{sol}^2} \\sim 0.008 eV \\gg m_1$, an approximate upper bound on the CP asymmetry\\cite{ibarra} is \\begin{equation} |\\epsilon_1| \\lesssim \\frac{3}{16\\pi} \\frac{M_1\\sqrt{\\Delta m_{atm}^2}} {v^2} \\ , \\end{equation} with $v=174 GeV$ being the vacuum expectation value of the neutral Higgs boson. For the degenerate case, i.e., $m_1\\approx m_2\\approx m_3 \\approx \\overline{m}/\\sqrt{3} \\gg \\sqrt{\\Delta m_{atm}^2}$, the upper bound is given by \\begin{equation} |\\epsilon_1| \\lesssim \\frac{3\\sqrt{3}}{16\\pi} \\frac{M_1\\Delta m_{atm}^2} {v^2\\overline{m}} \\ , \\end{equation} where $\\overline{m}=\\sqrt{m_1^2+m_2^2+m_3^2}$ is the quadratic mean of the neutrino masses. One can see that the maximal $\\epsilon_1$ is proportional to $M_1$ while inversely proportional to $\\overline{m}$. In the following we always take the upper bound of $\\epsilon_1$ to study the maximal value of $\\eta_B$ on the parameter space.  The final lepton asymmetry is determined in general by \\begin{equation} \\label{yl} Y_L\\equiv \\frac{n_L-n_{\\bar{L}}}{s} = \\kappa \\frac{\\epsilon_1}{g_*}\\ , \\end{equation} where $s$ is the entropy density, $g_*$ represents the number of the relativistic degrees of freedom at the time when $N_1$ decays and $\\kappa < 1$ represents the washout effects for the lepton number asymmetry in the thermal bath. The lepton asymmetry in Eq. (\\ref{yl}) will be converted partly into baryon asymmetry by the electroweak sphaleron process taking into account the gauge, Yukawa interaction and the QCD sphaleron effects\\cite{qcdspha}, $Y_B = a Y_{B-L}\\approx -a Y_L$, with $a=28/79$\\cite{therm,chem}.  \\subsection{The washout effect}  The factor $\\kappa$ in Eq. (\\ref{yl}) can be calculated by solving the Boltzmann equations\\cite{buch,therm}. At the leading order of $\\epsilon_1$ one has the Boltzmann equations for $Y_{N_1}=n_{N_1}/s$ and $Y_{B-L}=n_{B-L}/s$, \\begin{eqnarray} \\label{boln1} \\frac{sH}{z}\\frac{dY_{N_1}}{dz} &=& -\\left( \\frac{Y_{N_1}}{Y_{N_1}^{eq}} -1\\right)\\left( \\gamma_D+2\\gamma_{H,s}+4\\gamma_{H,t} \\right)\\, ,\\\\ \\label{bolbl} \\frac{sH}{z}\\frac{dY_{B-L}}{dz} &=& -\\epsilon_1 \\gamma_D \\left( \\frac{Y_{N_1}}{Y_{N_1}^{eq}} -1\\right) - \\frac{Y_{B-L}}{Y_L^{eq}} \\gamma_W\\, , \\end{eqnarray} with \\begin{equation} \\gamma_W=\\frac{1}{2}\\gamma_D +2\\gamma_N+2\\gamma_{N,t}+2\\gamma_{H,t}+ \\frac{Y_{N_1}}{Y_{N_1}^{eq}}\\gamma_{H,s}\\, , \\end{equation} where $\\gamma$ is the space time density of the scatterings in thermal equilibrium. In the equations above, $\\gamma_D$ refers to the reaction density for the process of $N_1$ decays (and inverse decays), $N_1\\leftrightarrow LH (\\bar{L}\\bar{H})$, $\\gamma_{H,s}$ and $\\gamma_{H,t}$ for the processes of $\\Delta L=1 $ scatterings via exchanging Higgs boson, $N_1L(\\bar{L})\\leftrightarrow \\bar{t}(t)Q(\\bar{Q})$ in $s$ channel and $N_1t(\\bar{t}) \\leftrightarrow \\bar{L}(L)Q(\\bar{Q})$ in $t$ channel, $\\gamma_N$ and $\\gamma_{N,t}$ for the processes of $\\Delta L=2 $ scatterings via exchanging the RH neutrinos, $LH\\leftrightarrow \\bar{L}\\bar{H}$ in both $s$ and $t$ channels and $LL\\leftrightarrow \\bar{H}\\bar{H}$, $\\bar{L}\\bar{L}\\leftrightarrow HH$ in $t$ channel.  In Eq. (\\ref{boln1}), the minus sign on the right-hand side means that all the reactions drive the $Y_{N_1}$ to its equilibrium value $Y_{N_1}^{eq}$. In Eq. (\\ref{bolbl}), the first term on the right-hand side is the source term, which gives rise to non-zero $Y_{B-L}$, while the second term represents washout effects that lead $Y_{B-L}$ to zero.  The $\\gamma_N$ consists two parts, the resonant part, $\\gamma_{N,res}$ and the non-resonant part, $\\gamma_{N,non}$, i.e., $\\gamma_N=\\gamma_{N,res}+\\gamma_{N,non}$. The $\\gamma_D$, $\\gamma_{H,s}$, $\\gamma_{H,t}$ and $\\gamma_{N,res}$ are all proportional to the effective neutrino mass\\cite{buch} \\begin{equation} \\widetilde{m}=\\frac{(Y^\\dagger Y)_{11} v^2}{M_1}\\ , \\end{equation} while $\\gamma_{N,non}$ and $\\gamma_{N,t}$ are proportional to $M_1\\overline{m}^2$. These properties are essential in understanding the behavior of the solutions of the Boltzmann equations. When $\\delta \\gamma_W =\\gamma_{N,non}+\\gamma_{N,t}\\propto M_1\\overline{m}^2$ is negligible, i.e., at $M_1(\\overline{m}/0.1 eV)^2 < 10^{13} GeV$\\cite{buch}, the washout factor $\\kappa$ is mainly determined by $\\widetilde{m}$ and nearly independent of $M_1$ and $\\overline{m}$. The effect of $\\widetilde{m}$ has two different aspects. When it is very small ($\\lesssim 10^{-3} eV$), the RH neutrinos will not be able to reach thermal equilibrium via the scattering with the thermal bath, so $\\kappa$ will be small. Conversely, if $\\widetilde{m}$ is too large the RH neutrinos can be brought into thermal equilibrium rapidly. However, the washout effect will also be large, consequently we have a small $\\kappa$. Around $\\widetilde{m}\\sim 10^{-3} eV$, $\\kappa$ reaches its maximal value. When $\\delta \\gamma_W$ is sizable $\\kappa$ will be suppressed exponentially, since the washout effect is given by \\begin{equation} \\kappa (z) = \\int_{z_{in}}^z dz' \\gamma_D(z') \\Delta (z') e^{-\\int_{z'}^z dz'' z''\\gamma_W(z'')/(sHY_L^{eq})}\\, , \\end{equation} with $\\Delta(z)= (Y_{N_1}/Y_{N_1}^{eq}-1) z/(sH)$. In this case, $\\kappa$ becomes nearly independent of $\\widetilde{m}$.  In short, $\\epsilon_1$ (as well as $Y_L$) increases with $M_1$ while decreases with $\\overline{m}$ linearly when $\\delta \\gamma_W$ is negligible, $M_1(\\overline{m}/0.1 eV)^2 < 10^{13} GeV$\\cite{buch}. However, when $\\delta \\gamma_W\\propto M_1\\overline{m}^2$ is sizable, $M_1(\\overline{m}/0.1 eV)^2 > 10^{13} GeV$, the final asymmetry $Y_L$ will be suppressed exponentially\\cite{buch}. The $\\kappa$ reaches maximal at $\\widetilde{m}\\sim 10^{-3} eV$ when the effect of $\\delta \\gamma_W$ is negligible. It is shown in Ref. \\cite{buch} that to generate the required value of the baryon number asymmetry, $\\eta_B$, one needs $M_1\\gtrsim 10^{10} GeV$ and $\\overline{m}\\lesssim 0.2 eV$ in general.  \\subsection{The effect of Quintessence}  With the interaction between the Quintessence and the RH neutrinos, $M_1(Q)$ is now a function of the Quintessence field. $M_1$ at the epoch when it decays will differ from that at the present epoch. We introduce a parameter $K$ \\begin{equation} K(Q^D)=\\frac{M_1^0}{M_1^D}, \\end{equation} where $M_1^0=M_1(Q^0)$ and $M_1^D=M_1(Q^D)$ denote the Majorana masses of the RH neutrinos at present and at the decay time. For simplicity we take $K$, as well as $M_1$, as a constant when solving the Boltzmann equations. In section IV, we will show that $Q$ is quite flat during the period of leptogenesis for the model we consider  and this is a reasonable assumption. We then have \\begin{equation} m_\\nu^D=-Y\\frac{1}{M^D}Y^Tv^2=m_\\nu^0\\cdot K\\ ,\\;\\; \\overline{m}^D=\\overline{m}^0\\cdot K\\ ,\\;\\; \\widetilde{m}^D=\\widetilde{m}^0\\cdot K\\ , \\end{equation} where the masses with the superscripts ``D\" and ``0\" correspond to the values evaluated at the leptogenesis and at the present time respectively. ",
        "conclusions": "The explanation of the baryon number asymmetry of the Universe remains a challenge for cosmological and particle physics. Among various theoretical models, leptogenesis is one of the most attractive scenarios. Especially, the minimal thermal leptogenesis only needs the minimal extension of the standard model, while the extension seems necessary in order to explain the neutrino oscillation experiments.  The minimal thermal leptogenesis is closely connected with the properties of the light neutrinos. Detailed studies show that, to generate the observed baryon asymmetry, the constraints $M_1 \\gtrsim 10^{10} GeV$ and $m_i \\lesssim 0.1 eV$ should be satisfied. These results disfavor the degenerate neutrino mass spectrum strongly and require a high reheating temperature.  We investigate in this paper the implications on the minimal thermal leptogenesis of a new scenario that the RH neutrinos couple with the Quintessence, whose evolution causes the masses of the RH neutrinos vary. We study the interplay of this coupling between the evolution of Quintessence and the leptogenesis. By solving the Boltzmann equations numerically we show that $M_1$, at the decay time, can be as low as about $10^8 GeV$ and $m_i$, at the present epoch, can be as large as $m_i\\sim 0.2 eV$."
    },
    "0311/astro-ph0311517_arXiv.txt": {
        "abstract": "We present results from cryogenic tests of a Volume-Phase Holographic (VPH) grating at 200 $K$ measured at near-infrared wavelengths. The aims of these tests were to see whether the diffraction efficiency and angular dispersion of a VPH grating are significantly different at a low temperature from those at a room temperature, and to see how many cooling and heating cycles the grating can withstand. We have completed 5 cycles between room temperature and 200 $K$, and find that the performance is nearly independent of temperature, at least over the temperature range which we are investigating. In future, we will not only try more cycles between these temperatures but also perform measurements at a much lower temperature (e.g., $\\sim 80~K$). ",
        "introduction": "Volume-Phase Holographic (VPH) gratings potentially have many advantages over classical surface-relief gratings (Barden, Arns, \\& Colburn 1998; see also Barden et al. 2000), and are planned to be used in a number of forthcoming instruments (e.g., AA$\\Omega$; Bridges et al. 2002). While applications to optical spectrographs only are currently being considered, VPH gratings will also be useful to near-infrared spectrographs if the performance at low temperatures is satisfactory. In particular, its diffraction efficiency and angular dispersion should be confirmed. Contraction of dichromated gelatin with decreasing temperature could cause variations in the line density and profile of diffraction efficiency (the thickness of the gelatin layer is one of the parameters defining diffraction efficiency). Since cooling and heating cycles might cause some deterioration of a VPH grating and reduce its life time, we also need to see whether these characteristics vary with the successive cycles.  In this paper, results from measurements of a sample grating at 200 $K$ and at room temperature are presented. A picture of the grating investigated is shown in Figure \\ref{vph}. This grating was manufactured by Ralcon Development Lab, and its diameter is about 25 cm. The line density is 385 lines/mm and thus the peak of diffraction efficiency is around 1.3 $\\mu$m at the Bragg condition when the incident angle of an input beam to the normal of the grating surface is 15$^{\\circ}$. The measurements are performed at wavelengths from 0.9 $\\mu$m to 1.6 $\\mu$m. The target temperature, size and line density of the grating, and wavelengths investigated are nearly the same as those adopted for the Fibre-Multi Object Spectrograph (FMOS; e.g., Kimura et al. 2003), which is one of the next generation instruments for the 8.2m Subaru telescope with commissioning expected in 2004: this instrument will be observing at wavelengths from 0.9 $\\mu$m to 1.8 $\\mu$m, and a VPH grating will be used as an anti-dispersing element in the near-infrared spectrograph which is operated at $\\sim$ 200 $K$ to reduce thermal noise. ",
        "conclusions": "In this paper, results from the cryogenic tests of a VPH grating at 200 $K$ are presented. The aims of these tests were to see whether diffraction efficiency and angular dispersion of a VPH grating are significantly different at a low temperature from those at a room temperature, and to see how many cooling and heating cycles the grating can withstand. We have completed 5 cycles between room temperature and 200 $K$, and find that diffraction efficiency and angular dispersion are nearly independent of temperature. This result indicates that VPH gratings can be used in spectrographs cooled down to 200 $K$ such as FMOS without any significant deterioration of the performance.  In future, we will be trying more cycles between 200 $K$ and 280 $K$ to mimic more realistic situations of astronomical use. Measurements at a much lower temperature (e.g., $\\sim$ 80 $K$) will also be necessary to see whether VPH gratings are applicable to spectrographs for use in the $K$-band. We will report on these issues in a forthcoming paper."
    },
    "0311/astro-ph0311498_arXiv.txt": {
        "abstract": "We discuss the discovery of `notch-like' features in the mean pulse profile of the nearby, bright pulsar B0950+08.  We compare these low-level features with those previously seen in the pulse profiles of pulsars J0437$-$4715 and B1929+10. While J0437$-$4715 is a binary millisecond pulsar and B0950+08 and B1929+10 are isolated, slow pulsars, all three pulsars are nearby and very bright. Furthermore, all three have detectable emission over an unusually wide range of pulse phase.  We describe the similar properties of the notch features seen in all three pulsars and discuss possible interpretations. ",
        "introduction": "Detailed studies of the structure of radio pulsar pulse profiles are essential for understanding the pulsar emission mechanism and the forms of the pulsar emission beam. Pulse profiles are remarkably stable over timescales ranging from a few pulses to many years, indicating that these profiles are determined by permanent physical characteristics of the neutron star and its environment. Many profile properties can be explained very well by the ``hollow-cone'' model of pulsar emission developed by Radhakrishnan \\& Cooke (1969) and Komesaroff (1970), shortly after the discovery of the first pulsar. This ``magnetic-pole'' model, further developed by Sturrock (1971) and Ruderman \\& Sutherland (1975), involves charges streaming from the polar cap along the open field lines, radiating in the direction of their motion and forming an emission cone aligned with the magnetic axis of the star.  As more pulsars were discovered, many were found to have two, three or even more pulse components, prompting various efforts to classify and interpret the various forms (\\eg. Backer 1973; Rankin 1983, 1993; Lyne \\& Manchester 1988).  In this paper we draw attention to the strange ``notches'' in the profiles of B0950+08 and two other pulsars, J0437--4715 and B1929+10.  All three stars exhibit very similar double ``notch\" features which appear to have virtually nothing to do with the component structures of their very different profiles.  B0950+08 and B1929+10 have rotational periods of 253 and 227 ms, whereas J0437--4715 is a binary millisecond pulsar with a period of only 5.75 ms.  Significantly, these stars are some of the very closest among the pulsar population---all with dispersion measures (DMs) of about 3 ~pc cm$^{-3}$ and parallax-determined distances of some 262$\\pm$5, 331$\\pm$10 and 170$\\pm$3 pc (Brisken \\etal\\ 2000, 2002)---and, in spite of the fact that they all have rather ordinary spindown values, all have been detected in the optical and/or soft x-ray region (Mignani \\etal\\ 2002; De Luca \\etal\\ 2003; Wang \\etal\\ 1997; Zavlin \\etal\\ 2002). B0950+08 and B1929+10 are among the brightest pulsars in the radio sky\\cite{ppcat}, and J0437--4715 is well known as the brightest millisecond pulsar.   In \\S\\ref{sec:obsres} we discuss our observations of B0950+08, present its average pulse profiles and describe the unusual features.  In \\S\\ref{sec:discuss} we compare these features with similar ones seen in the profiles of J0437$-$4715 and B1929+10 and discuss possible physical interpretations. ",
        "conclusions": "\\label{sec:discuss}  Unusual features similar to those presented in \\S\\ref{sec:obsres} have been observed in two other pulsars. Rankin \\& Rathnasree (1997) reported a double notch-like feature in the average profile of B1929+10 at 430~MHz, and a display similar to theirs using the same observations is given in Fig.~\\ref{fig:1929}.  As can be seen there, the ``notches'' follow the main pulse peak by roughly $100^\\circ$ and have a total width of approximately $10^\\circ$.  More accurate scaling reveals that they have widths of some 2.3$\\deg$ and that their centers are separated by just twice this interval.  Note also that B1929+10's profile is generally fully linearly polarized, so that the features are most clearly seen in both the total power and total linear polarization\\footnote{The possible origin of this strange ``superpolarization'', that was also seen by Phillips (1990), is discussed in Rankin \\& Rathnasree (1997).  Phillips, however, seems to have plotted his Figs.~1 \\& 2 with some unreported smoothing, so that the ``notches'' are not visible, though there is a broad dip in the power just were they should be.}, where they represent about a 40\\% diminuation of the power at their centers.  The ``notches'' do appear to be just visible in a sensitive 430-MHz observation of Blaskiewicz \\etal\\ (1991: see their Fig.~22), though they are not seen at 1.4 GHz either in this paper or in Rankin \\& Rathnasree.  This may be due to the overall weakening of the broad low-level emission feature that follows the MP as well as the poorer signal-to-noise of the 21-cm observations.  Neither, unfortunately, were the recent observations of Weisberg \\etal\\ (2003) sufficiently sensitive to detect the features.  Similar features have also been observed by Navarro \\etal\\ (1997) in the average profile of the millisecond pulsar J0437$-$4715.  These ``notches'' fall on the trailing edge of this star's very broad profile, some 70$\\deg$ after the bright central component.  Here, the pair is separated by about 3.3$\\deg$, the half-power widths are about a degree, and the fractional intensity decrease at their centers is about 50\\%.  These characteristics are clearly seen in the 1512-MHz polarization profile of Navarro \\etal\\ (their Fig.~4) which we reproduce here as Fig.~\\ref{fig:0437a}. However, this paper reports additional detections at frequencies of 438 and 660~MHz, and the three total-power profiles appear together in our Fig.~\\ref{fig:0437b}, aligned (according to their Fig.~5 caption) on the basis of the anti-symmetric circularly polarized signature associated with the central component.  One can see that the ``notches'' at the three frequencies do not quite align, falling slightly earlier at 660 and 1512 MHz. However, were the profiles aligned more nearly according the the centroids of the central component, the ``notches'' might well fall at just the same profile phase.  It would be interesting to see the results of a timing alignment as carried out for B0950+08 above.  Finally, the features do appear to evolve somewhat with frequency, becoming wider and more separated at lower frequencies.  It is noteworthy that, for all three pulsars, these unprecedented features have roughly the same angular width, implying that their overall scale is some nearly fixed fraction of the rotational cycle of the star [as, apparently, is also the case for pulsar microstructure (\\eg, Popov \\etal\\ 2002)]. Furthermore, for the two pulsars in which the features have been detected at more than one frequency, they seem to appear at exactly the same longitude.  Again, this is a well-known feature of pulsar microstructure, implying that the two phenomena derive from the same physical regions in the pulsar magnetosphere. Pulsar microstructure generally shows the same width regardless of frequency, however, there is evidence for the width of the notch features decreasing with frequency.  \\begin{figure} \\begin{center} \\epsfig{file=B1929notches.ps,height=8.5cm,angle=-90.} \\caption{Full-period average polarization profile of B1929+10 at 430 MHz after Rankin \\& Rathnasree (1997).  The total power (Stokes $I$) is plotted first at full scale and then at a $\\times$250 expanded scale, so that only features below 0.4 percent of the MP amplitude are now visible.  The linear ($L$) and circular ($V$) polarization are then plotted at the expanded scale. The PA is plotted in the lower panel.  Note the double ``notch'' feature which follows the MP by about 100$\\deg$.  This region exhibits complete linear polarization (see text), so the ``notches'' are clearly visible in both $I$ and $L$, representing an intensity decrease at their centers of some 40\\%. \\label{fig:1929}} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\epsfig{file=navarro.f4.ps, height=6.0cm} \\caption{Polarized pulse profile for J0437$-$4715 at 1512~MHz from Navarro \\etal\\ (1997); their Fig.~4. \\label{fig:0437a}} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\epsfig{file=navarro.f5.ps, height=5.0cm} \\caption{Total-intensity profiles for J0437$-$4715 at 438, 660 and 1512~MHz, which have been aligned according to the zero-crossing point of the anti-symmetric circular polarization under their central components from Navarro \\etal\\ (1997); their Fig.~5.  Note that the ``notches'' almost align---and that they could even do so more closely if only a slightly different alignment had been used. \\label{fig:0437b}} \\end{center} \\end{figure}  \\medskip  B0950+08 and B1929+10 have remarkably similar properties. Both are very bright, have quarter-second rotational periods, and are very nearby. Both pulsars have detectable emission over an unusually broad fraction of their rotation cycle and both exhibit interpulses.  With a spin period of 5.75~ms, J0437$-$4715 has very different spin-down properties; however, it is also nearby, extremely bright, and, while it has no interpulse, it has detectable emission over more than $180^\\circ$ of longitude.  The identification of these features in the three brightest pulsars with broad pulse profiles suggests that such features may be common among pulsars with broad profiles. As they are apparently seen over a substantial portion of the radio band and evolve little with frequency, it would seem that they cannot be attributed to absorption in any usual sense.  Thus it would seem that they are either due to some very puzzling aspect of the emission process or represent an equally puzzling obscuration of the emission along our sight line. While we cannot more than speculate about the causes of these features, a very novel mechanism is being suggested by Wright (2003) in a companion paper.  Clearly, few pulsars show emission so far from the main pulse or interpulse, so observing this feature in more pulsars will be very difficult until the advent of more sensitive pulsar instruments such as the SKA. However, future more sensitive multi-frequency observations of the pulsars discussed above, in addition to polarimetric observations, may yield more clues as to the origin of these unusual features.  \\bigskip"
    },
    "0311/astro-ph0311101_arXiv.txt": {
        "abstract": "The Galactic Anticenter Stellar Structure (GASS) has been identified with excess surface densities of field stars in several large area sky surveys, and with an unusual, string-like grouping of five globular clusters.  At least two of these are diffuse, young ``transitional'' clusters between open and globular types.  Here we call attention to the fact that four {\\it younger} open or transitional clusters extend the previously identified, string-like cluster grouping, with at least one having a radial velocity consistent with the previously found GASS velocity-longitude trend.  All nine clusters lie close to a plane tipped 17$\\arcdeg$ to the Galactic plane.  This planar orientation is used to forage for additional potential cluster members in the inner Galaxy, and a number are found along the same plane and string-like sequence, {\\it including almost all fifteen known outer, old open clusters}. Tidal accretion of a dwarf satellite galaxy on a low inclination orbit --- perhaps the GASS system --- appears to be a plausible explanation for the origin of the outer, old open and transitional clusters of the Milky Way.  We use these clusters to explore the age-metallicity relation of the putative accreted GASS progenitor.  Finally, we provide the first radial velocity of a star in the cluster BH 176 and discuss its implications. ",
        "introduction": "Excesses of stars beyond the apparent limit of the Galactic disk have been used to argue for the presence of a distinct, extended stellar structure wrapping around the disk at low latitudes (Newberg et al. 2002; Ibata et al. 2003, hereafter I03; Majewski et al., hereafter M03; Yanny et al. 2003, hereafter Y03; Rocha-Pinto et al., hereafter R03). Collectively, these surveys suggest that the structure spans $|b|<30\\arcdeg$ and at least $122\\arcdeg<l<225\\arcdeg$ at a mean $R_{GC}\\sim16$ kpc (I03, R03) and radial thickness $\\lesssim4$ kpc (Y03, I03).  However, because of unfortunate placement behind considerable extinction, it has been difficult to get information on the system's true shape, orientation, breadth, etc.  Even the location of the structure's center (presumably corresponding to a ``nucleus'') remains uncertain; thus we refer to the entire system here as the Galactic anticenter stellar structure (GASS).  The origin of GASS --- originally described as a ``ring'' around the Galaxy (I03, Y03) --- is also not definitively established, with a number of potential scenarios outlined, e.g., by I03: a tidally disrupted satellite galaxy, an outer spiral arm, or (their preference) a resonance induced by an asymmetric Galactic component.  From among possibilities involving accretion, \\citet{helmi} explore the extremes of dynamically young and old tidal debris, and \\citet[ hereafter C03]{crane} argue, as did Y03 earlier, that the most straightforward interpretation is that GASS is a disrupted satellite galaxy, resembling in many ways the Sagittarius (Sgr) dwarf galaxy (e.g., M03) system. As evidence, C03 point to: (1) a velocity-longitude trend indicating a slightly non-circular orbit, (2) a velocity dispersion smaller than even that of thin disk stars, (3) a wide metallicity spread from [Fe/H] $=-1.6 \\pm 0.3$ (Y03) to at least $=-0.4 \\pm 0.3$ dex, and (4) at least four star clusters apparently associated with the stream based on position {\\it and} radial velocity (RV).  These clusters (Pal 1, NGC 2808, NGC 5286, NGC 2298), plus a fifth having no RV measurement (BH 176), lie in an unusual, arc-like configuration not seen elsewhere among low latitude, outer globular clusters (GCs), but one resembling configurations expected for tidal debris systems \\citep[e.g.,][]{bellazzini}.  Several unusual GCs are identified by C03 as potential GASS members. Pal 1 is both very small ($M_v = -2.5$) and relatively metal-rich ([Fe/H] $= -0.6$) for an $R_{GC} > 8$ kpc GC.  \\citet{rosenberg} derive a Pal 1 age of $8\\pm2$ Gyr and suggest it is either the youngest GC or one of the oldest open clusters (OCs) in the Galaxy.  \\citet[ hereafter PS03]{phelps_schick} find BH 176 to be young ($7.0 \\pm 1.5$ Gyr) and metal-rich ($-0.20 \\le$ [Fe/H] $\\le +0.20$) and suggest it is ``transitional'' between a young, metal-rich GC and a massive, metal-rich old OC.  Prompted by this hint that GASS may contain younger, more metal-rich star clusters, and by our new distance for OC Saurer A (Sau A) \\citep{fp02} placing it near the R03 and C03 tracings of GASS, we search here for other {\\it open clusters} coincident with GASS, and find an interesting potential connection with the Milky Way (MW) old OC system. ",
        "conclusions": "How the relatively high $Z_{GC}$-distributed, old OC system was formed has remained a challenging problem.  Among the two most plausible models, \\citet{friel} concludes that old OC creation during evolution of the MW disk requires ``fine tuning'' of formation and destruction processes, whereas in accretion ``one finds a natural mechanism for open cluster formation'', particularly the high $|Z_{GC}|$ OCs.  In this context we find several interesting coincidences regarding the 15 known OC/TC objects with $d_{\\sun}>7$ kpc and $R_{GC}>7$ kpc: (1) Remarkably, 13 of these 15 clusters are confined to the third and fourth Galactic quadrants, something that would occur by chance only 4.3\\% of the time.  Such a lopsided distribution is easily accommodated by an accretion origin, but not by a disk formation model. (2) Eleven of these 15 clusters are among the GASS candidate clusters in Table 1, while one more, Berkeley 22, lies right in the M giant GASS tracing by R03 and would have been included in our sample had we used a planar distance limit only 20 pc larger.  Although clearly old (0.7 to 7 Gyr) for OCs, the Table 1 TC/OC objects are generally poorly studied and have unknown RVs; their proposed association with one another and with the GASS field star overdensities must therefore be considered tentative. Nevertheless, their arcing spatial sequence and planar alignment (\\S4) is tantalizingly suggestive of an origin relating to the interaction of a satellite galaxy with the MW.   \\begin{figure} \\epsscale{1.00} \\plotone{f3.eps} \\caption{AMR of Table 1 clusters plus Be 22 (see \\S5).  Symbols as in Fig.\\ 2.} \\end{figure}  An alternative view might hold that the arcing spatial sequence of star clusters and the correlated overdensities of field stars defining GASS merely represent an outer MW spiral arm. However, if a spiral arm, it is strangely inclined by 17${\\arcdeg}$ to the GP. Additionally, spiral arms characteristically have {\\it young} star clusters, yet no Table 1 cluster is younger than 0.67 Gyr. There is also no GASS correlation with the Galactic warp (C03).  On the other hand, star clusters in the Fornax (For) and Sgr satellite galaxies have sizes and luminosities \\citep[e.g.,][]{mg03} that span those of the typical old OC \\citep{friel} and Table 1 objects; that For and Sgr clusters are called ``globular'' seems mainly to reflect a difference in age, but {\\it young} clusters of similar $M_V$ are prevalent \\citep{hunter} in the Magellanic Clouds (MCs).  Under the premise that the Table 1 objects represent the cluster system of a dwarf galaxy, one can use them to explore the age-metallicity relation (``AMR'') of that putative system.  Fig.\\ 3 shows an AMR typical of that expected for an independently evolving, ``closed-box'' system with protracted star formation \\citep[compare Fig.\\ 3 to the similar AMR of Sgr field stars and clusters in Fig.\\ 18 of][]{layden}. Apart from BH 176, the Fig.\\ 3, outer TC/OC clusters show a relatively tight AMR (even including Be 22), especially compared to that for {\\it all} old MW OCs \\citep[e.g., Fig.\\ 8 of][]{friel} and in comparison to the LMC cluster AMR \\citep[e.g., Fig.\\ 2a of][]{bica}.  A large metallicity spread among these clusters mimics the spread among non-cluster GASS stars discussed by C03.  Together, the various ensemble properties of the GASS clusters lend further circumstantial support to the ``tidal debris'' explanation for GASS.  However, while sharing a similar overall AMR to Sgr, GASS appears to have star clusters much younger than in the MW-accreting cluster system of Sgr, for which the youngest and most metal-rich known cluster is Terzan 7 (age $8.3\\pm1.8$ Gyr, [Fe/H] $=-0.82\\pm0.15$; Layden \\& Sarajedini 2000). Clearly such differences in the selective production and/or destruction of clusters within MW satellites is not a problem since differences in the distribution of cluster ages are already observed between the systems in For, Sgr and the MCs, all galaxies with continuing star formation up to the near present.  But the lower inclination and apparently smaller orbit of GASS compared to these other MW satellites suggests consideration of an additional mechanism for the constant production of new star clusters through the continuous interaction of a ``GASSeous'' dwarf galaxy with the gaseous disk of the MW.        We acknowledge funding by NSF grant AST-0307851, NASA/JPL contract 1228235, the David and Lucile Packard Foundation, and the Celerity Foundation."
    },
    "0311/astro-ph0311271_arXiv.txt": {
        "abstract": "{Results of optical and NIR spectral and photometric observations of a sample of candidate double-barred galaxies are presented. Velocity fields and velocity dispersion maps of stars and ionized gas, continuum and emission-line images were constructed from integral-field spectroscopy observations carried out at the 6m telescope (BTA) of SAO RAS, with the MPFS spectrograph and the scanning Fabry-Perot Interferometer. \\object{NGC 2681} was also observed with the long-slit spectrograph of the  BTA. Optical and NIR images were obtained at the BTA and at the 2.1m telescope (OAN, M\\'exico). High-resolution images were retrieved from the  HST data archive. Morphological and kinematic features of all 13 sample objects  are described in detail. Attention is focused on the interpretation of observed non-circular motions of gas and stars in circumnuclear (one kiloparsec-scale) regions. We have shown first of all that these motions are caused by the gravitational potential of a large-scale bar. \\object{NGC 3368} and \\object{NGC 3786} have nuclear bars only, their isophotal twist at larger radii being connected with the bright spiral arms. Three cases of inner polar disks in our sample (\\object{NGC 2681}, \\object{NGC 3368} and \\object{NGC 5850}) are considered. We found ionized-gas counter-rotation in the central kiloparsec of the lenticular galaxy \\object{NGC 3945}. Seven galaxies (\\object{NGC 470}, \\object{NGC 2273}, \\object{NGC 2681}, \\object{NGC 3945}, \\object{NGC 5566}, \\object{NGC 5905}, and \\object{NGC 6951}) have inner mini-disks nested in large-scale bars. Minispiral structures occur often in these nuclear disks. It is interesting that the majority of the observed, morphological and kinematical, features in the sample galaxies can be explained without the secondary bar hypothesis. Thus we suggest that a dynamically independent secondary bar is a rarer  phenomenon than  follows from isophotal analysis of the images only. ",
        "introduction": "\\label{intro}  More than a quarter of a century ago de Vaucouleurs (\\cite{vauc}) found that an inner bar-like structure was nested in the large-scale bar in the optical image of \\object{NGC 1291}. The first systematic observational study of this phenomenon was made by Buta \\& Crocker (\\cite{butacr}), who published a list of 13 galaxies, where the secondary (inner) bar was arbitrarily oriented with respect to the primary (outer) one. In the following years the isophotal analysis technique made it possible  to detect secondary bars in optical (Wozniak et al., \\cite{woz}; Erwin \\& Sparke, \\cite{erw99}; Erwin \\& Sparke, \\cite{erw03}) and NIR images of barred galaxies (Friedli et al., \\cite{fri96}; Jungwiert et al., \\cite{jung}; Laine et al., \\cite{laine02}). Up to now, we have found 71 candidate  double-barred galaxies in the literature (see references in Moiseev, \\cite{mois01b}). A new revised catalog of 67 barred galaxies including 50 double-barred ones is presented by Erwin (\\cite{erwcat}). Although various theoretical studies exist, the secondary bar dynamics is still far from being well understood. Shlosman et al. (\\cite{shlosman89}) have shown that an additional circumnuclear bar may be formed as a result of instabilities of the gaseous ring in the Inner Lindblad Resonance (ILR) within a large-scale bar. Heller et al. (\\cite{heller01}) present a detailed simulation of this process, which, unfortunately, results in the appearance of an elliptical gaseous ring instead of the ``real'' stellar-gaseous bar. Maciejewski \\& Sparke (\\cite{mac00}) found some families of closed orbital loops, able to support both bars. Similar independently rotating structures sometimes appear in stellar-gaseous simulations of galactic disks (Pfenninger \\& Norman, \\cite{pfenninger}; Friedli \\& Martinet, \\cite{frimar}). New hydrodynamical simulations of the gas behavior in double bars, with an analysis of the problems arising during modeling and interpretation of the double-barred structures are presented by Maciejewski et al. (\\cite{mac02}) and Shlosman \\& Heller (\\cite{shlosman02}).  Various observational data indicate that we are probably seeing a new structural feature of barred galaxies in the case of double bars. However, the majority of observational results has been based on photometric measurements only, when an elongated elliptical structure appears inside the large-scale bar. Moreover, using the isophotal analysis formalism (Wozniak et al., \\cite{woz}) has allowed  even ``triple bars'' (Friedli et al., \\cite{fri96}; Jungwiert et al., \\cite{jung}; Erwin \\& Sparke, \\cite{erw99}) without any arguments about the dynamical behavior of such complex stellar systems. At the same time, the structures seen in direct images can be explained in a less exotic manner, without using the hypothesis of double or triple bars. Many structural features would be able to create an illusion of a secondary bar (Friedli et al., \\cite{fri96}; Moiseev, \\cite{mois01b}; Petitpas \\& Wilson, \\cite{petit}), such as for example, an oblate bulge, a complex distributions of dust and star formation regions in the circumnuclear region, or elliptical rings at the ILRs of a primary bar (Shaw et al., \\cite{shaw}). Therefore new observational data are required  to test our understanding of the secondary bar phenomenon.   Kinematic data (the stellar and gas  line-of-sight velocities and velocity dispersion) over the circumnuclear  regions of barred galaxies are of crucial importance for this kind of studies. Recently, Emsellem et al. (\\cite{emsellem01}) have studied the kinematics of double-barred galaxy candidates and reach a conclusion on dynamical decoupling of the circumnuclear regions in \\object{NGC 1097}, \\object{NGC  1808} and \\object{NGC 5728}. This conclusion is based on the fact that peaks of the relative line-of-sight velocities and stellar velocity dispersion drops  are seen in the circumnuclear regions from long-slit cross-sections. Wozniak (\\cite{woz99}) has supposed that ``counter-rotation'' of stars, observed in the nuclear region of \\object{NGC 5728}, may be associated with secondary bar influence on stellar kinematics. All these features observed in the one-dimensional data may be explained without using the hypothesis of secondary bars. This  may be related to the specific mass distributions in the inner kiloparsec and to non-circular motions of stars and gas within the primary bar only ($x_2$-families of stellar orbits lead to  ``counter-rotation'' properties). Also Wozniak et al. (\\cite{woz03}) in their self-consistent N-body simulations have shown that the velocity dispersion drops observed by Emsellem et al. (\\cite{emsellem01}) can be independent on the presence of a secondary bar and are reproduced in the model of a single-barred galaxy.   Since the motions of stars and gas inside the bar region are strongly non-circular, and such objects are non-axisymmetric by definition,  panoramic (also named 2D, 3D or integral-field) spectroscopy should give us  more detailed information on circumnuclear kinematics in compared to  ``classical'' (long-slit) spectroscopy. Integral-field spectroscopy makes it possible to map two-dimensional fields of line-of-sight velocities and velocity dispersion. Therefore, during 2000-2002 we have carried out a series of observational programs on 2D-kinematics and morphology of double-barred galaxies candidates. The main goal was to find an answer to the question -- \\textit{Are the secondary bars dynamically decoupled systems?}  This problem is complicated and ambiguous because of the lack of theoretical studies, where the ``observed'' velocity fields in the double-barred galaxies had been modeled. It is important to collect a full sample of candidate double-barred galaxies with observable 2D-kinematics of stars and gas. These observations must be used to search for common dynamical (kinematic) features in these kinds of galaxies. Recently, Shlosman, \\& Heller (\\cite{shlosman02}) have written: ``..clearly, the most promising method in detecting the nuclear bars is two-dimensional spectroscopy of the central kiloparsec, which can reveal the underlying kinematics''.  Our observational program was finished in 2002. Independently similar observations  were started at the 3.6m CFHT  using the integral-field spectrographs OASIS and TIGER. Recently, the velocity fields of stars and ionized gas in some double-barred candidates have been presented by Emsellem \\& Friedli (\\cite{eric00}), Emsellem et al. (\\cite{emsellem01b}) and Emsellem (\\cite{eric470}). The velocity fields of the molecular gas were also investigated in double-barred candidates \\object{NGC 2273} and \\object{NGC 5728} (Petitpas \\& Wilson, \\cite{petit}) and \\object{NGC 4303} (Schinnerer et al., \\cite{eva}).  In this work, we present the observational data (integral-field spectroscopy and NIR/optical surface photometry) for 13 galaxies and discuss the problem of central region dynamical decoupling. The paper is structured as follows. In Sect.~\\ref{sec_obs} we describe the observations and data reduction technique. The methods of analysis of the velocity fields and surface brightness distributions are considered in Sect.~\\ref{method}. In Sect.~\\ref{atlas0} we present the notes for individual observed galaxies and discuss their common properties in Sect.~\\ref{disc}; a short  conclusion is drawn in Sect.~\\ref{concl}. ",
        "conclusions": "\\label{concl}  The  analysis of the morphology and kinematics of 13 double-barred galaxy candidates observed by means of  panoramic spectroscopy at the SAO RAS 6m telescope is presented in this paper. We tried to collect the maximum possible information from observed 2D distributions of the optical and NIR surface brightness and kinematic properties. For this goal, panoramic spectroscopy observations were carried out at the 6m telescope. Velocity fields and velocity dispersion fields of ionized gas and stars  were constructed from these observations. The rotation curves and radial variations of the dynamical axis position angle were calculated. The analysis of the morphology of the galaxies was done by means of isophote analysis of NIR and optical ground-based and HST archival images.   We suggest that candidate double-barred galaxies are, in fact, galaxies with \\textit{very different circumnuclear structure}. It is necessary to note that the majority of the observed  morphological and kinematic features in our sample galaxies may be explained without the secondary bar hypothesis. Three cases of inner polar disks, one counter-rotating  gaseous disk and seven nuclear disks (with and  without mini-spirals) nested in   large-scale bars were found in this work. However, we think that the observational data presented here, particularly the 2D data, will be useful to proponents as well as to opponents of the secondary bar hypothesis. Doubtless  the models of such galaxies must explain all their observed kinematic properties."
    },
    "0311/astro-ph0311536_arXiv.txt": {
        "abstract": "{Recent measurements of the galaxy clustering environments around intermediate-redshift radio sources have suggested a systematic environmental difference between radio galaxies and radio-loud quasars, in contradiction to the predictions of simple unified models for the two classes of object. I show that the apparent difference arises mainly as a result of the properties of low-excitation radio galaxies included in the radio-galaxy sample, which tend to lie in significantly richer environments. The environmental properties of high-excitation radio galaxies and quasars are statistically consistent in the redshift range $0.15 < z <0.4$, as unified models would predict. ",
        "introduction": "In unified models for powerful radio galaxies and radio-loud quasars (Scheuer 1987, Barthel 1987, 1989) isotropic properties of the two classes of objects, such as their clustering environments, should be statistically identical. However, the classical unified models work best at high redshift (there are few low-redshift objects classified as radio-loud quasars) and at these redshifts the clustering environments of the sources are hard to measure. Some attempts have been made to test the predictions of unified models using X-ray observations of extended emission, assumed to be cluster-related (Hardcastle \\& Worrall 1999) but sample sizes are generally small and {\\it ROSAT}-based work is generally not reliable because of the contamination of the X-ray emission by AGN- and radio-source-related components. {\\it Chandra} should provide significantly better constraints from the X-ray once sufficient objects are in the archive, though it will be important to take into account the significant evolution undergone by the ICM of a radio-source host cluster between $z \\sim 1$ and 0. In the meantime, there is still an important r\\^ole for the traditional method of calculating the amplitude of the galaxy-galaxy or quasar-galaxy spatial covariance function, hereafter referred to as $B_{\\rm gg}$ (e.g. Longair \\& Seldner 1979; Yee \\& Green 1987; Prestage \\& Peacock 1988; Yates, Miller \\& Peacock 1989; Ellingson, Yee \\& Green 1991; Hill \\& Lilly 1991; Yee \\& Ellingson 1993). Although it is recognised that selection in more than one optical filter is desirable to achive good statistics at high redshifts (Barr \\etal\\ 2003), at lower redshifts simple $B_{\\rm gg}$ determinations seem likely to be able to provide a statistical measure of the environments of different classes of source.  Recently Harvanek \\etal\\ (2001: hereafter HESR) have carried out a large systematic study of the clustering properties of 3CR (Spinrad \\etal\\ 1985) radio sources. By making new observations, and collating and cross-calibrating data from the literature, they were able to obtain a nearly complete set of $B_{\\rm gg}$ values in the redshift range $0.15 < z < 0.65$. 3CR objects in this redshift range are almost all luminous FRII-type radio galaxies or lobe- or core-dominated quasars. Yee \\& Green (1987) had shown that few low-redshift, radio-loud quasars are located in rich environments. HESR aimed to test the unified-model prediction that radio galaxies at these redshifts would also tend to lie in poor environments. They obtained the surprising result that the $B_{\\rm gg}$ values of their 3CR radio galaxies were not consistent with those of a sample of radio-loud quasars compiled by Yee \\& Ellingson (1993: hereafter YE93) in the redshift range $0.15 < z < 0.4$, in the sense that a number of the radio galaxies were found in significantly richer environments than any quasar. An immediate objection to this result is that the YE93 quasars were not in general as radio-luminous in extended emission as the 3CR comparison sample (otherwise they would all have been 3CR objects themselves on the basis of their low-frequency extended emission, whereas only a fraction of them are in fact in 3CR). The effects of this bias are hard to quantify, but since (for equal jet powers and source ages) simple radio-galaxy physics suggests that a more luminous object will lie in a richer environment, it might tend to produce a difference in the sense of the one observed by HESR. In what follows I shall ignore this potential problem and concentrate instead on the effects on HESR's conclusions of using a more sophisticated variant of the unified model. ",
        "conclusions": "\\label{discussion}  The key result of this work is that 3CR LERG lie in significantly richer environments than the HERG/quasar population at comparable redshifts. This strong difference between the environmental properties of LERG and HERG supports suggestions by Hardcastle \\etal\\ (1998) and Hardcastle \\& Worrall (1999) that the 3CR LERG population is physically different from the HERG. If we believe that the emission-line properties of the radio source reflect the power of the jet (e.g. Rawlings \\& Saunders 1991) then a plausible explanation is that at least some LERG are intrinsically low-jet-power radio galaxies which owe their comparatively high overall radio luminosity and their other properties (small size, dissipative jets, relaxed appearance -- cf. Hardcastle \\etal\\ 1998, Harvanek \\& Stocke 2002) to a rich environment. This recalls, and may be related to, the arguments of Barthel \\& Arnaud (1996).  The fact that the narrow-line objects and quasars have similar $B_{\\rm gg}$ distributions, when the LERG are excluded, supports the version of the orientation-based unified model for radio galaxies and quasars outlined in section \\ref{unif}."
    },
    "0311/astro-ph0311193_arXiv.txt": {
        "abstract": "We consider several mechanisms that possibly affect the evolution of disk galaxies in clusters using analytical models based on a hierarchical clustering scenario.  We especially focus on the evolution of disk galaxies in subclusters located around a main cluster. We show that ram-pressure stripping cannot be always ignored in subclusters, although their masses are much smaller than that of the main cluster. The star-formation rate of a galaxy may gradually decrease by the stripping of warm gas (`strangulation') in a main cluster. However, we find that ram-pressure stripping could start before the strangulation is completed, if a field galaxy directly falls into the main cluster.  Since this conflicts with some recent observations, many galaxies might have been affected by some environmental effects when they were in subclusters before they fell into the main cluster (`pre-processing'). We show that strangulation and evaporation of the cold gas by the surrounding hot ICM in subclusters are the possible candidates.  We also show that the observed morphological transformation of disk galaxies at $z\\lesssim 1$ is not chiefly due to galaxy mergers. ",
        "introduction": "Clusters of galaxies in the redshift range of $z\\gtrsim 0.2$ often exhibit an overabundance, relative to present-day clusters, of blue galaxies (Butcher, Oemler 1978, 1984). This star-formation activity is often called the Butcher--Oemler effect, and subsequent studies have confirmed this trend \\citep{cou87,rak95,lub96,mar00,ell01,got03a}. At the same time, various studies have reported a morphological evolution in cluster galaxies; they have shown that the fractions of S0 galaxies in clusters decrease toward a higher redshift (Dressler et al. 1997; Couch et al. 1994, 1998; Fasano et al. 2000). For many clusters, it has been found that their galaxy population gradually changes from a red, evolved, early type population in the inner part of the clusters to a progressively blue, later type population in the extensive outer envelope of the clusters \\citep{abr96,bal97,rak97,oem97,sma98,cou98,van98}. These are often interpreted as being the result that star-forming field disk galaxies infalling to a cluster are transformed to passively evolving cluster members.  Various mechanisms have been proposed for the transformation. In general, galaxies in the inner region of a cluster have lower star-formation rates both at $z\\sim 0$ \\citep{lew02,gav02,gra03} and at higher redshift (Couch, Sharples 1987; Balogh et al. 1997, 1999b; Poggianti et al. 1999; Postman et al. 2001; Couch et al. 2001; Balogh et al. 2002c). Ram-pressure stripping of cold gas in the disk of a galaxy is a possible candidate for the suppression of star formation in clusters because the removal of cold gas leads to the exhaustion of an ingredient of stars \\citep{gun72,bal94,fuj99a,aba99,qui00,mor00,vol01,bic02}, although ram-pressure may cause a mild increase of star formation of a galaxy \\citep{fuj99a} or a starburst with a short-time scale \\citep{bek03,gav03}. Ram-pressure stripping is likely to take place in the central regions of clusters because the density of intracluster medium (ICM) is high there. In fact, disk galaxies with a deficit of cold gas or a morphological sign of gas removal are often observed near cluster centers (Cayatte et al. 1990; Vollmer et al. 2000; Bravo-Alfaro et al. 2000, 2001; Solanes et al. 2001; Bureau, Carignan 2002; Vollmer 2003; Sofue et al. 2003; Yoshida et al. 2003).  Moreover, it is more effective in richer clusters because the typical velocity of galaxies is larger. Several authors actually indicated that the fractions of blue galaxies are systematically lower for rich clusters than for poor clusters at a given redshift \\citep{mar01,got03a}, which is consistent with a suppression of the star-formation activities by ram-pressure stripping. On the other hand, \\citet{bal02b} investigated clusters with small masses at $z\\sim 0.25$, and indicated that the striking similarity between the spectral and morphological properties of galaxies in these clusters and those of galaxies in more massive systems at similar redshifts implies that the physical processes responsible for truncating star formation in galaxies are not restricted to the rare, rich cluster environment, but are viable in much more common environments. Thus, they concluded that ram-pressure stripping cannot be solely responsible for the low star-formation rates in these systems. Moreover, several authors indicated that ram-pressure stripping cannot account for the suppression of star formation observed $\\sim$Mpc away from the center of a cluster \\citep{bal97,kod01c,lew02}, although we will show that it is not obvious.  \\citet{lar80} suggested that an infall from gaseous galactic halos might be important for sustaining star formation in spiral galaxies and that the gas in these halos might be stripped in the cluster environment, leading to the formation of passively evolving S0 galaxies. Although the gas in galactic halos, which we call `warm gas', has not been detected in nearby spiral galaxies \\citep{ben00}, models of galaxy formation based on the hierarchical clustering scenario have often assumed such an infall, and have been very successful in accounting for the properties of galaxies at optical and infrared wavelengths \\citep{col94,bau96,kau99,nag01}. If only warm gas is stripped, star formation may be allowed to continue by consuming the remaining cold disk gas, but without the infall to replenish this supply, star formation will die out on timescales of a few Gyr \\citep{lar80}. This is consistent with observations that indicate a slow decline of the star-formation rates of galaxies \\citep{bal99b,kod01a}. We call this scenario `strangulation' from now on. In this paper, we also consider the evaporation of cold gas in disk galaxies by energy transfer via thermal conduction from the surrounding hot ICM as another mechanism for the slow decline of the star-formation rates. As far as we know, this is the first time that the redshift evolution of the evaporation effect has been studied.  The above mechanisms, that is, ram-pressure stripping, strangulation, and evaporation, mainly have an influence on the star-formation activities of galaxies. Contrary to them, galaxy mergers are expected to significantly affect the morphology of galaxies. Thus, mergers may have played an important role in the observed morphological transformation in clusters. Mergers between galaxies with comparable masses (major mergers) could create elliptical galaxies \\citep{too72}. However, major mergers alone cannot account for the observed fractions of galaxies with intermediate bulge-to-disk luminosity ratios such as S0 galaxies in clusters \\citep{oka01,dia01}. Okamoto and Nagashima (2003) indicated that mergers between galaxies with significantly different masses (minor mergers) may produce the galaxies with intermediate bulge-to-disk luminosity ratios because these mergers do not disrupt galactic disks completely.  In this paper, we consider possible mechanisms responsible for the truncation of star formation and morphological transformation of disk galaxies in clusters. In particular, we investigate the evolution of disk galaxies in subclusters. Recent observations show that galaxies have already been affected by some environmental effects in subclusters \\citep{kod01c,tre03,got03b}. We define subclusters as clusters located around a larger cluster (main cluster) at relatively high redshift; these subclusters are to be merged with the main cluster until $z=0$. Thus, at high redshift, we can say that these subclusters are progenitors of the main cluster at $z=0$. On the other hand, at low redshift, they become subhalos included in the main cluster. Since high-resolution $N$-body simulations showed that a subcluster that has been included in the main cluster as a subhalo is not completely destroyed \\citep{moo99,oka99,ghi00,fuk01}, we do not explicitly discriminate a subcluster from a subhalo. However, we can estimate the time when subclusters are included by the main cluster (see section \\ref{sec:result}).  We consider the effect of cosmological evolution of clusters since the inner structure of clusters at high redshift is not the same as that at low redshifts, which makes differences of environmental effects on cluster galaxies at between high and low redshift. We take an analytical approach to study the environmental effects because at the moment it is difficult to study their redshift evolutions by genuine numerical simulations. For example, it is virtually impossible to study ram-pressure stripping of galaxies in clusters by cosmological numerical simulations including both dark matter and gas with a resolution that is sufficient to treat the interaction between ICM and interstellar gas in each galaxy correctly. Even if we introduce some assumptions about the stripping of the interstellar gas and the distribution of the ICM, as is done by Okamoto and Nagashima (2003), we need to find the average cluster and galaxy evolutions at least in terms of their dark halos. However, it is time-consuming to create many clusters by performing high-resolution $N$-body simulations to follow the average cluster and galaxy evolutions. Moreover, there is another merit for analytical approaches. Even in the future, when full numerical simulations are enabled, the simulations must include many physical processes, and it would be difficult to divide them when the results are analyzed. The division is easier for analytical studies, and the results of analytical studies would be very useful to be compared with those of numerical simulations.  We also emphasize that our model is different from previous semi-analytic models. Our model takes account of spatial correlations among initial density fluctuations, and can specifically predict the evolution of subclusters around the main cluster. On the other hand, previous models cannot discriminate between subclusters and the main cluster.  We mainly consider massive disk galaxies since they can be observed in detail even at high redshift. Thus, we do not consider the cumulative effect of high-speed encounters between galaxies (`galaxy harassment') because it influences disk galaxies with small masses \\citep{moo96}. We mostly study the redshift evolution of each environmental effect and discuss the relative strength of the effect at high and low redshift. We also compare results with observations qualitatively.  This paper is organized as follows. In section~\\ref{sec:model}, we summarize our models. In section~\\ref{sec:result} we give the results of our calculations, and compare them with observations in section~\\ref{sec:disc}. Summary and conclusions are given in section~\\ref{sec:sum}. Readers who are interested only in the results of calculations may skip section~\\ref{sec:model}. ",
        "conclusions": "\\label{sec:sum}  Using analytical models based on a hierarchical clustering scenario, we have investigated the evolution of massive disk galaxies in clusters. We have considered the effects of cosmological evolution of clusters on galaxy evolution. In particular, we have explored the galaxy evolution in subclusters located around the main cluster. The main conclusions from this work may be summarized as follows:  (i) As long as the spatial distribution of intracluster medium (ICM) follows that of dark matter in a cluster or a subcluster, galaxies are influenced by ram-pressure stripping of the cold disk gas. For a given massive cluster at $z=0$, ram-pressure stripping in the cluster progenitors is more effective at higher redshift although the masses of the cluster progenitors decrease with redshift. This is because the typical mass density of the progenitors increases with the redshift. Even in subclusters, ram-pressure stripping is effective.  (ii) If field galaxies directly fall into a main cluster, the exhaustion of the cold disk gas via star formation, after stripping of the warm diffuse gas in their galactic halos (`strangulation'), would not be completed before ram-pressure stripping of the cold disk gas becomes effective. This limits the time scale in which star-formation rates of the galaxies decline. The maximum time scale is $\\lesssim$~Gyr, which is inconsistent with some observations.  (iii) The conflict between the theoretical prediction and the observations may be avoided if the star-formation rates of galaxies had already dropped in subclusters before the subclusters plunged into the main cluster (`pre-processing'). This is consistent with recent observations indicating that the star-formation rates of galaxies are decreased at radii well larger than the virial radii of main clusters.  (iv) The star-formation rates of galaxies in subclusters might have gradually decreased via strangulation and/or evaporation of the cold gas by the surrounding hot ICM, if non-gravitational heating of ICM (`preheating') had changed the ICM distributions of the subclusters, and tidal force from the main cluster prevents the galaxies from being affected by ram-pressure stripping. These may be the mechanism of the pre-processing. In particular, the evaporation scenario is free from the problem that warm gas in galactic halos has not been detected.  (v) It is not obvious whether ram-pressure stripping takes place in subclusters (or main clusters at high redshift) or not. Even if future observations show that ram-pressure stripping does not occur in subclusters (or main clusters at high redshift), their small masses cannot be the only reason. It requires additional reasons, such as the preheating and/or the tidal acceleration.  (vi) Mergers between a galaxy at the center of a cluster progenitor and that not at the center had mostly finished at $z\\gtrsim 1$--$2$ because the velocities of galaxies increase as the host clusters grow. Thus, galaxy mergers do not appear to be the main cause of the observed morphological transformation from spiral galaxies to S0 galaxies at $z\\lesssim 1$.  (vii) The star-formation activity of a galaxy may resume after the galaxy has passed the central region of a cluster and the ram-pressure from the ICM has dropped. This may affect the fraction of blue galaxies in clusters.  \\vspace{10mm}  I am grateful to M. Nagashima, T. Okamoto, T. Kodama, N. Yoshida, and T. Goto for useful discussion.  I was supported in part by Grants-in-Aid from the Ministry of Education, Culture, Sports, Science and Technology (14740175)."
    },
    "0311/astro-ph0311470_arXiv.txt": {
        "abstract": "{The effect of the hyperon softening of the equation of state (EOS) of dense matter on the spin evolution of isolated neutron stars is studied for a broad set of hyperonic EOSs. We use a multidomain 2-D code based on a spectral method, and show how important the precision of solving the equations of stationary motion is for the stability analysis. For some EOSs, the hyperon softening leads to spin-up by the angular momentum loss, in the form of the back bending phenomenon, for a rather broad range of stellar baryon mass. We show that large segments of the evolutionary tracks exhibiting  the back bending behaviour in the moment-of-inertia - rotation-frequency plane, are actually unstable and therefore not astrophysically relevant. We show also that during the spin-up - angular-momentum-loss epoch, an isolated neutron star (e.g., a radio pulsar) can lose a sizable part of its initial angular momentum without significantly changing its rotation period.  We propose also simple arguments and criteria allowing  to connect the presence of a back bending epoch  with the mass-radius relations  and the stiffness and/or softness of the nucleon and hyperon EOSs of the neutron star core. ",
        "introduction": "\\label{sect:introduction} Some theories of dense matter predict a softening of the equation of state (EOS) at densities exceeding normal nuclear density $\\rho_0=2.7\\times 10^{14}~{\\rm g~cm^{-3}}$, implied by a phase transition to a new ``exotic'' phase. Several ``exotic'' high-density phases were proposed in the past, e.g., deconfined quark plasma, and pion-condensed or kaon-condensed  hadronic matter. The softening of the EOS could be due to a transition into a pure ``exotic'' phase, or to a mixture of an ``exotic'' phase with normal phase of dense matter. It has been suggested that the softening of the EOS due to a phase transition could lead to characteristic back bending phenomenon in the timing of spinning-down pulsars (Glendenning et al. 1997) or produce characteristic period clustering in spinning-up neutron stars in low-mass X-ray binaries (Glendenning 2001, Glendenning and Weber 2002). An interesting conclusion of these papers was that both back bending in spinning down pulsar and spin clustering in accreting millisecond pulsars stars is an {\\it evidence} of a phase transition taking place at the center of a spinning-down pulsar or a spinning-up accreting neutron star.  As we show in the present paper, using several EOS of dense matter and very precise code for the calculation of the rotating stellar configurations,  observation of the back bending in the timing behavior of isolated pulsars, or of the period  clustering,  is not an  unambiguous  evidence for an ``exotic'' phase in dense matter. Namely, these phenomena can also be implied  by the presence of hyperons in dense matter, a feature which is in no way ``exotic'' and which was predicted more than forty years ago (Cameron 1959, Salpeter 1960, Ambartsumyan and Saakyan 1960). A possibility of ``spin-up by the angular momentum loss'' for a normal sequence (baryon mass smaller than the maximum baryon mass of static configurations) of spinning-down neutron stars with hyperonic cores was previously noted by Balberg et al. (1999). However, as we will demonstrate in the present paper, a complete study of back bending requires a very precise ``exact'' code for calculating stationary configurations of rotating stars and a simultaneous careful checking of secular stability of these configurations. Such conditions  were typically not satisfied in previous calculations.  In previous works the  back bending phenomenon has been considered as a feature of the $I(\\Omega)$ dependence, where $I$ is the moment of inertia of the star and $\\Omega$ is the angular frequency of rotation (\\cite{gpw97}, \\cite{cyz02} and \\cite{ss02}). In the present  paper we clarify some statements about the back bending for rotating neutron stars, and we formulate some rules which are useful for searching for the back bending in rotational stellar sequences. As we demonstrate using high-precision evolutionary sequences, the whole analysis should be performed using different pair of variables: total stellar angular momentum $J$ versus $\\Omega$ instead of $I(\\Omega)$. High precision is particularly important because it is needed to reliably check the secular stability of rotating configurations; only the stable ones are interesting and observationally relevant.  The paper is organized in the following way. Softening of the EOS due to the presence of hyperons is discussed in Sect.\\ \\ref{sect:EOS}. The method allowing for a  high precision of the 2-D calculations of the equilibrium configurations of rotating neutron stars is briefly described in Sect.\\ \\ref{sect:NumCalc}. Different formulations of the stability criteria for rotating configurations are briefly summarized in Sect.\\ \\ref{sect:I.J.stability}. In Sect.\\ \\ref{sect:MR.minima} we propose a method of checking for the occurrence of the back bending phenomenon by inspecting the baryon-mass -- equatorial-radius relations at fixed values of rotation frequencies. We apply this method  to several  EOSs with a hyperon softening. The interplay between the back bending and stability is discussed in Sect.\\ \\ref{sect:Jf.BB.stab}, where we study neutron-star evolution tracks in the angular-momentum -- rotation-frequency plane. In Sect.\\ \\ref{sect:BB.EOS} we study the dependence of the back bending phenomenon on the EOS of the hyperonic matter. Final Sect.\\ \\ref{sect:DiscConcl} contains discussion of our results, including their possible observational aspects, and ends with  concluding remarks. ",
        "conclusions": "\\label{sect:DiscConcl} The presence of hyperons in neutron-star cores  can strongly affect the spin evolution of a solitary neutron star (e.g., an isolated pulsar). As we have shown, such a neutron star can live  a long epoch  of ``spin down by the angular momentum loss'', and this could occur for a broad range of baryon mass of neutron star. The epoch with $\\dot{P}<0$ despite $\\dot{J}<0$ could terminate by an instability or by a stable continuous transition to a ``standard'' spin-down epoch.  We paid particular attention to the $(\\dot{P}<0$, $\\dot{J}<0,M_{\\rm B}=const)$ epoch with a stable termination. It is  represented by an S-shaped segment of the spin-evolution track  in the  $J-f$ and $I-f$ planes,  and was baptized ``back bending phenomenon'' in the previous literature. Various regimes of the spin evolution were shown to be correlated with the behaviour of the $M_{\\rm B}(R_{\\rm eq})_f$ curves at fixed $f$. In particular, we pointed out importance of the location of the inflection points of the $M_{\\rm B}(R_{\\rm eq})$ curves for the existence of the back bending phenomenon.  We were also inclined to leave the name of back bending only to the evolution-track segments with a stable termination (${\\rm BB_{st}}$ in our terminology).  Epochs with back bending for normal rotating configurations were found for two of the four EOS of Balberg \\& Gal (1997). On the other hand, we found that the back bending phenomenon for normal sequences was absent for five hyperonic EOSs of Glendenning (1985). This illustrates the uncertainties in the hyperonic EOS, stemming for a high degree of ignorance concerning the nucleon-hyperon and particularly hyperon-hyperon interactions in dense matter.  Throughout this paper we stressed that the back bending  in the $I-f$ plane, considered up to now in the literature, should be accompanied by stability analysis; only stable back-bending configurations are astrophysically relevant. We performed such a stability analysis, and we found that very often  dominant back-bending segments of the $I(f)_{\\rm M_{\\rm B}}$  are {\\it unstable} with respect to the axisymmetric perturbations, and therefore do not exist in the Universe.  \\vskip 3mm  As it has  been mentioned by Spyrou \\& Stergioulas (2002), in such kind of calculations it is extremely important to assure the thermodynamical consistency of the EOS (the first law of thermodynamics has to be strictly fulfilled). It is well known that a rough treatment of this condition can lead to an inaccurate determination of the maximum mass and  stability conditions  (for example the configurations corresponding to the maximum of  $M$ and $M_{\\rm B}$ do not coincide - this is an evidence for  a lack of the thermodynamical consistency in the EOS).  The precision of the code is also very important and in our case the proper division of the star into two computational domains (at the threshold for appearance of hyperons) allows us to obtain high precision results without a large increase of the number of grid points.  \\vskip 3mm  An isolated pulsar, born in a SN II explosion, could have an interesting and nonstandard past due to a hyperon softening of the EOS. As we showed, such a pulsar could lose some half  of its initial angular momentum without changing much its  rotation period. Therefore, if one observes a rapid pulsar with a characteristic age $\\tau_{\\rm PSR}$ significantly longer than the age of a supernova remnant where this pulsar is born, $\\tau_{\\rm SNR}$, this might be due to some back bending episode, resulting from a  (hyperon? phase-transition?) softening of the EOS of its core. Such a possibility of explaining a seeming contradiction between  $\\tau_{\\rm SNR}$ and $\\tau_{\\rm PSR}$ has been already noticed, in the context  of a  mixed-phase EOS softening, by \\cite{ss02}. Clearly, the observational pulsar-timing signatures  of the EOS softening due to the hyperon or phase-transition softening of the EOS deserve further studies, and we are planning to continue such studies using our high-precision 2-D code.  Another interesting consequence of the hyperon softening of the EOS could be a ``period clustering'' of rotating neutron stars  powered by accretion in the long living low-mass X-ray binaries. This problem is now being studied."
    },
    "0311/astro-ph0311120_arXiv.txt": {
        "abstract": "We present an estimate of the bolometric X-ray luminosity - velocity dispersion ($L_x - \\sigma_v$) relation measured from a new, large and homogeneous sample of 171 low redshift, X-ray selected galaxy clusters. The linear fitting of $\\log(L_x) - \\log(\\sigma_v)$ gives $ L_x = 10^{32.72 \\pm 0.08} \\sigma^{4.1 \\pm 0.3}_v$ erg s$^{-1}$ h$^{-2}_{50}$. Furthermore, a study of 54 clusters, for which the X-ray temperature of the intracluster medium T is available, allows us to explore two other scaling relations, $L_x -T$ and $\\sigma_v -T$. From this sample we obtain $L_x \\propto T^{3.1 \\pm 0.2}$ and $\\sigma_v \\propto T^{1.00 \\pm 0.16}$, which are fully consistent with the above result for the $L_x$-$\\sigma_v$. The slopes of $L_x -T$ and $\\sigma_v -T$ are incompatible with the values predicted by self-similarity ($L_x \\propto T^{2} \\propto \\sigma_v^4$), thus suggesting the presence of non-gravitational energy sources heating up the intracluster medium, in addition to the gravitational collapse, in the early stages of cluster formation. On the other hand, the result on $\\log(L_x) - \\log(\\sigma_v)$ supports the self-similar model. ",
        "introduction": "Galaxy clusters are dark matter haloes which form mainly through the process of gravitational collapse, and so they are expected to be self-similar: small clusters or galaxy groups must be scaled-down versions of the more massive systems. Self-similarity predicts a series of relationships between different cluster observables such as X-ray luminosity $L_x$, mass $M$, temperature $T$ and galaxy velocity dispersion $\\sigma_v$, which are the same in all systems (Kaiser 1991).  The radial velocity dispersion $\\sigma_v$ of the galaxies in a cluster probes the depth and shape of the potential well, assuming that the luminous matter traces reasonably well the dark matter in clusters. Moreover, the intracluster gas emits X rays through a process of thermal bremsstrahlung and its bolometric luminosity $L_x$ is found to be strongly correlated with $\\sigma_v$, as both gas and galaxies share a common potential well (Solinger \\& Tucker 1972). Also, since the X-ray emission can be well modelled by thermal emission of a hot, optically thin plasma, $L_x$ and the plasma temperature $T$ must be correlated. Finally, because galaxies are embedded in the intracluster medium, a correlation is also expected between the gas temperature $T$ and the galaxies velocity dispersion $\\sigma_v$.   Self-similar models consider that the only energy source in the cluster comes from the gravitational collapse, predicting the following scaling relations: $L_x \\propto T^{2} \\propto \\sigma_v^{4}$. Whereas there seems to be a general agreement between different measurements that $L_x \\propto T^{ \\sim 3}$, the measurement of $L_x - \\sigma_v$ has given somewhat contradictory results so far. Some authors have found that $L_x \\propto \\sigma^{4}$ indeed, although with quite large measurement errors or rather small data samples, while others find slopes larger than 4 (see Table \\ref{literature}). Part of the differences in the results could come from different ways of selecting the sample, with a preference for regular clusters in some of these surveys.  It has also been suggested that clusters and groups do not follow the same $L_x- \\sigma_v$ scaling relation, the latter being flatter than the former (e.g. Mahdavi et al. 2000, Xue \\& Wu 2000, Dell'Antonio et al. 1994). But there are other measurements contrary to that conclusion (Mahdavi \\& Geller 2001, Mulchaey \\& Zabludoff 1998).  For more distant clusters ($z$ between $ \\sim 0.15$ and $\\sim 0.6$) there is some evidence that the slope is also $> 4$ (Borgani et al. 1999, Girardi \\& Mezetti 2001), although only small samples are available at the moment, and more data are needed to reduce the error bars.  There is a clear need of new measurements performed on large samples (more than 100 clusters), selected from a homogeneous dataset. The largest cluster (not group) samples available up to date, on which this kind of study has been performed, are literature compilations combining data from different authors and/or instruments (for example Xue \\& Wu 2000, Mahdavi \\& Geller 2001). The present work comes to partially fill this gap, as we have a large homogeneous subsample of clusters selected from REFLEX, a flux-limited X-ray selected cluster catalogue built under well defined selection criteria.  The structure of this paper is as follows. In Section \\ref{thedata} we give a brief introduction to the REFLEX catalogue and the way in which the subsample in this paper has been selected. Section \\ref{measurement} illustrates in detail the process of the $L_x$ and $\\sigma_v$ measurements, which lead to the the fitting of the $\\log(L_x) - \\log(\\sigma_v)$ in Section \\ref{lxsigma}. In Section \\ref{nosubst} we propose the removal of those clusters in the sample which are multiple systems. In Section \\ref{discussion} we analyse possible biases due to the nature of the selected sample. A study on the $L_x - T$ and $\\sigma_v - T$ relation follows in Section \\ref{lxt}, and we finish with a summary in Section \\ref{summary}. We are assuming H$_0=50$ km s$^{-1}$ Mpc$^{-1}$ in a flat universe with $\\Omega _{\\Lambda}=0.7$. ",
        "conclusions": "\\label{summary}  We present an estimate of the $L_x - \\sigma_v $ relation computed with a new large and homogeneous sample of clusters from the REFLEX survey. All studies on the $L_x - \\sigma_v$ relation in large samples (with more than $\\approx 100$ clusters) have been performed so far on compilations from the literature requiring the use of models to combine different data from different authors and/or instruments (e.g. Xue \\& Wu 2000, Mahdavi \\& Geller 2001).  From the REFLEX catalogue we construct a sample of 171 clusters for which we can derive reliable estimates of $\\sigma_v$, and then fit the $ \\log(L_x)-\\log(\\sigma_v) $ relation, finding  \\begin{equation} \\log(L_{45})= (-1.34 \\pm 0.08)+ (4.1 \\pm 0.3) \\cdot \\log ( \\sigma_{500})\\ \\rmn{erg} \\ \\rmn{s}^{-1} \\rmn{h}^{-2}_{50} \\end{equation}  in agreement with what is expected from self-similar models of cluster formation.  The large size and homogeneity of the sample allows us to make an estimate of the intrinsic dispersion of the $L_x - \\sigma_v$ relation. We find that $\\Sigma_{int}= 195$ km s$^{-1}$.  Measurements of $L_x$ and $\\sigma_v$ in clusters with multiple components may not be reliable. The identification in the literature of those clusters allows us to build a new sample where we remove those with significant substructure in their X-ray emission in Schuecker et al. (2001b). This procedure leaves us with a sample of 123 clusters, on which we fit again the $ \\log(L_x)-\\log(\\sigma_v) $ relation, finding  \\begin{equation} \\log(L_{45})= (-1.28 \\pm 0.10)+ (4.2 \\pm 0.4) \\cdot \\log ( \\sigma_{500} )\\ \\rmn{erg} \\ \\rmn{s}^{-1} \\rmn{h}^{-2}_{50} \\end{equation}  \\noindent compatible at one-sigma confidence level with the slope and intercept that we find when using the 171 cluster sample, suggesting that the sample is not noticeably affected by clusters with multiple components. The fact that we are considering only galaxies close to the cluster centre may be a reason for that.  We have also investigated the $L_x- \\sigma_v $ relation in a volume-limited sample, finding $\\log(L_{45})= (-1.20 \\pm 0.09) + (3.2 \\pm 0.3) \\times \\log(\\sigma_{500})$. This significantly flatter slope is an indication that some bias may be present on the whole sample due to the flux limit imposed on the data.  The $L_x - T$ and $\\sigma_v - T $ scaling relations have also been explored, finding that $L_x \\propto T^{3.1 \\pm 0.2}$ and $\\sigma_v \\propto T^{1.00 \\pm 0.16}$, in a subsample of 54 clusters for which $T$ was available from the literature. These results are consistent (within the error bars) with the one obtained for the $L_x - \\sigma_v$, as we find $ L_x \\propto \\sigma_v^{3.1 \\pm 0.7}$ when using these $L_x - T$ and $\\sigma_v - T$ to derive $L_x - \\sigma_v$.  The slope found in the $ L_x - T $ relationship is steeper than the value predicted by a purely gravitational collapse model of cluster formation, and is in good agreement with previous measurements performed on other cluster samples. As for the $ \\sigma_v - T $, the slope we find is only compatible with the self-similar value at the three-sigma confidence level, and slightly larger than the value found by other authors. Also, if we consider a volume-limited sample, the slope found is then in good agreement with previous studies, yet incompatible with self-similarity.  The fact that this model fails to reproduce the measured slope is attributed to the contribution to the ICM energy budget from other non-gravitational heating sources (see e.g. Bialek, Evrard \\& Mohr 2002), like AGNs, supernovae, gas cooling (Pearce et al. 2000), the presence of cool cores (Allen \\& Fabian 1998), or the effect of shocks in cluster formation (Cavaliere et al. 1997).  Here we are faced with a paradoxical result: the slope of $L_x - \\sigma$ supports the self-similar scenario, whereas $ L_x - T $ and $\\sigma_v - T$ do not, and all three relations are consistent with each other within one sigma confidence level. Clearly, more data are required to clarify the situation and reduce the error bars. This is a reflection of the general situation in this topic in the literature: so far there is a general agrement on $ L_x - T $ and $ \\sigma_v - T $, but there is not such an agreement on the $L_x - \\sigma_v$."
    },
    "0311/astro-ph0311316_arXiv.txt": {
        "abstract": "{We present results from the spectral analysis of the Seyfert 1.9 galaxy MCG~--5-23-16, based on $ASCA$, BeppoSAX, $Chandra$ and XMM-$Newton$ observations. The spectrum of this object shows a complex iron K$\\alpha$ emission line, which is best modeled by a superposition of a narrow and a broad (possibly relativistic) iron line, together with a Compton reflection component. Comparing results from all (six) available observations, we do not find any significant variation in the flux of both line components. The moderate flux continuum variability (about 25\\% difference between the brightest and faintest states), however, does not permit us to infer much about the location of the line-emitting material. The amount of Compton reflection is lower than expected from the total iron line EW, implying either an iron overabundance or that one of the two line components (most likely the narrow one) originates in Compton-thin matter. ",
        "introduction": "X--ray spectra of type 1 Active Galactic Nuclei (AGNs) have been extensively studied by many satellites in the last 15 years. After $Ginga$ it has become clear that typical X--ray broad band spectra of Seyfert galaxies result from the superposition of different components. The most common features arising over the primary power law emission are two reprocessed components: the so-called Compton reflection component at high energies (above $\\sim$10 keV) and the Fe K$\\alpha$ emission line at 6.4 keV from neutral iron. Both are due to reflection of the primary radiation by optically thick, neutral or low ionized, circumnuclear matter. At low energies, typically below about 1 keV, a further component may arise (soft excess), the origin of which is still rather unclear.  The study of the reprocessed components and in particular of the iron line profile provides important information about the region from which they originate. If the line is produced in the innermost regions of the accretion disc, its profile must be broad, asymmetric and double-peaked due to kinematic and relativistic effects \\citep[and references therein]{Fabian00}, while a narrow profile indicates an origin far from the nucleus, either in the torus envisaged in the Unification model \\citep{Antonucci93} or in the Broad Line Regions.  While ASCA observations indicate that a relativistic line is a common feature in Seyfert 1s \\citep{Nandra97}, XMM-$Newton$ have so far found unambiguous evidence for relativistic lines in an handful of objects only \\citep{Wilms01, Turner02}, while finding routinely a narrow line \\citep{Reeves02, Bianchi03}.  It is therefore very important for understanding the properties of the innermost regions of the accretion disc to search for more relativistic lines, to assess their frequency and strength. The fact that the relativistic line is broad, and therefore not easy to be separated from the continuum, and the confusing presence of the narrow component, implies that, unless the line equivalent width (EW) is very large, it should necessarily be searched for in bright sources.  MCG~--5-23-16 is a nearby (z~=~0.0083) X--ray bright narrow emission line galaxy. Its 2--10 keV flux has varied by a factor of 4 within about 10 years: from $\\sim8\\times10^{-11}$ erg cm$^{-2}$ s$^{-2}$ in 1978 \\citep{tennant83}, to $\\sim2\\times10^{-11}$ erg cm$^{-2}$ s$^{-2}$ in 1989 \\citep{np94}, and again to a high state of $\\sim9\\times10^{-11}$ erg cm$^{-2}$ s$^{-2}$ in 1996 \\citep{weaver98}. The nucleus is obscured by neutral matter with a column density of $\\sim10^{22}$ cm$^{-2}$; no signature of a warm absorber is detected. A soft excess is present in the spectrum below about 1 keV. $RossiXTE$ detected a Compton reflection signature in this Seyfert galaxy for the first time \\citep{weaver98}; $ASCA$ showed a strong (EW$\\sim300$ eV) and broad Fe K$\\alpha$ line with a possible complex profile \\citep{weaver97}; BeppoSAX confirmed the presence of the Compton reflection component and measured an equivalent width of $\\sim100$ eV for the iron line when fitted with a single gaussian \\citep{per02}. More recently $Chandra$ HETGS spectra, briefly described in \\citet{weaver01}, showed a narrow (FWHM $<$4\\,000 km s$^{-1}$) line with EW$\\sim90$ eV, while XMM-$Newton$ revealed a complex Fe K$\\alpha$ line constituted by a narrow unresolved component (EW$\\sim40$ eV) and a broad (FWHM $\\sim40\\,000$ km s$^{-1}$) component with an equivalent width of $\\sim120$ eV \\citep{dew03}.  In this paper we present a reanalysis of the $ASCA$, BeppoSAX, $Chandra$ and XMM-$Newton$ observations based on a model which includes Compton reflection and two distinct components, a narrow and a broad one, for the iron line. We will exploit the different capabilities of the various satellites (energy resolution for $Chandra$, sensitivity for XMM-$Newton$, hard X--ray coverage for BeppoSAX) to obtain informations on the location and physical properties of the emitting regions.  \\begin{table}[t]  \\caption{\\label{log}The log of all analysed observations, and exposure times. Values in brackets show net exposure times after filtering for particle induced flares.}  \\begin{center}  \\begin{tabular}{llll} \\hline \\hline \\textbf{Date} & \\textbf{Mission} & \\textbf{Instr.} & \\textbf{T$\\mathrm{_{exp}}$ (ks)} \\\\ \\hline 11/05/1994 & $ASCA$ & \\textsc{sis0-1} & 34 \\\\ \\hline 29/11/1996 & $ASCA$ & \\textsc{sis0-1} & 27 \\\\ \\hline & & \\textsc{lecs} & 36 \\\\ 24/04/1998 & BeppoSAX & \\textsc{mecs} & 77 \\\\ & & \\textsc{pds} & 33 \\\\ \\hline 14/11/2000 & $Chandra$ & \\textsc{acis-s hetg} & 76 \\\\ \\hline 13/05/2001 & XMM-$Newton$ & \\textsc{epic-pn} & 38 (6.0)\\\\ \\hline 01/12/2001 & XMM-$Newton$ & \\textsc{epic-pn} & 25 (19.1)\\\\ \\hline \\end{tabular}  \\end{center}  \\end{table} ",
        "conclusions": "Discussion}  We have reanalysed the XMM-$Newton$ observations of MCG~--5-23-16, as well as previous ASCA, BeppoSAX and $Chandra$ ones, confirming the presence of a significantly broad line already reported by \\citet{dew03}. A fit with a relativistic disc profile is somewhat better than that with a broad Gaussian. Moreover, a so broad line ($\\sigma=0.28$ keV) is hard to form other than in a relativistic disc \\citep[see e.g.][]{Fabian00}, which we therefore consider the most likely explanation. A narrow line is also present, as clearly shown by $Chandra$, as well as a Compton reflection component, measured by BeppoSAX. Exploiting the modest long term flux variability, we could confidently fit the XMM-$Newton$ spectrum adding these two further components (with the parameter $R$ describing the relative amount of reflection component kept fixed to the value found by BeppoSAX).  The detection of a broad, likely relativistic line in this source is important. Predicted on theoretical grounds, relativistic lines were found to be common in ASCA spectra of AGN \\citep{Nandra97}. However, \\citet{Lub01} showed that at least part of the lines could be ascribed to a narrow component from distant matter, also to be expected on the basis of Unification Models \\citep[e.g.][]{Ghi94}. $Chandra$ and XMM-$Newton$ have indeed confirmed the almost ubiquity of the narrow component, while relativistic lines have been unambiguously detected so far only in a few cases \\citep[e.g.][]{Wilms01, Turner02}. The discovery of a relativistic line in MCG~--5-23-16 adds therefore one more case to a still short list. It is worth remarking that the relativistic lines have been observed so far in sources in which either the EW of the line is very large, as in the case of MCG~--6-30-15 and NGC~3516 \\citep{Turner02}, or that are very bright, like MCG~--5-23-16. The presence of the narrow component, in fact, could make difficult to detect the broad component in not very bright sources, if its EW is modest. Even if in some cases, e.g. NGC~5506 \\citep{Matt01, Bianchi03}, NGC~7213 \\citep{b03}, NGC~5548 \\citep{Pounds03} and NGC~4151 \\citep{Schurch03}, relativistic lines with significant equivalent widths are definitely absent in high signal-to-noise XMM-$Newton$ spectra, the actual frequency of relativistic lines is still to be assessed.  The amount of reflection component, $R=0.45$, is low if compared with the total EW of the lines, which is about 170 eV \\citep[e.g.][]{matt91, gf91}. One possible explanation is iron overabundance \\citep{Matt97}; alternatively, one of the two line components may originate in Compton-thin matter. The obvious candidate is the narrow component (the broad one being likely emitted by the accretion disc), which therefore could be associated with the BLR rather than the molecular torus. (It is worth noting that the column density of the absorber is too low to make it a good candidate for producing the line \\citep[e.g.][]{Matt03}). This hypothesis cannot be directly tested, because MCG~--5-23-16 is a NELG and therefore the BLR is supposed to be obscured in the optical band, and therefore no BLR line width, to be compared with the upper limit derived from $Chandra$, is available. Alternatively, the line can be emitted in a Compton-thin torus. If indeed the narrow line is emitted in Compton-thin matter, no evidence for a Compton-thick torus is present in this source, similar to the case of NGC~7213 \\citep{b03}.  \\acknowledgement We thank G.C. Perola for useful discussions on BeppoSAX data analysis, Kazushi Iwasawa for his help on reducing $ASCA$ data and the anonymous referee for helpful comments and suggestions. We acknowledge financial support from ASI."
    },
    "0311/astro-ph0311585_arXiv.txt": {
        "abstract": "In a detailed analysis of the three point correlation function (3PCF) for the 2dF Galaxy Redshift Survey we have accurately measured the 3PCF for galaxies of different luminosity. The 3PCF amplitudes [$\\Qsu$ or $\\Qrpu$] of the galaxies generally decrease with increasing triangle size and increase with the shape parameter $v$, in qualitative agreement with the predictions for the clustering of dark matter in popular hierarchical CDM models. The 2dFGRS results agree well with the results of Jing \\& B\\\"orner for the Las Camapanas Redshift Survey (LCRS), though the measurement accuracy is greatly improved in the present study because the 2dFGRS survey is much larger in size than the LCRS survey.  The dependence of the 3PCF on luminosity is not significant, but there seems to be a trend for the brightest galaxy sample to have a lower amplitude than the fainter ones.  Comparing the measured 3PCF amplitudes [$\\Qsu$ or $\\Qrpu$] to the prediction of a WMAP concordance model, we find that the measured values are consistently lower than the predicted ones for dark matter. This is most pronounced for the brightest galaxies (Sample I), for which about one-half of the predicted $Q$ value provides a good description of $\\Qrpu$ for the 2dFGRS data.  For the less luminous sample (Sample II), the $Q$ values are also smaller than in the dark matter model on small scales, but on scales larger than $s=8 \\mpc$ and $r_p=3.25 \\mpc$ they reach the model values.  Therefore, the galaxies of sample II are unbiased tracers on linear scales, but the bright galaxies (sample I) have a linear bias factor of $\\sim 1.5$.  As for the LCRS data, we may state that the best fit DM model gives higher values for the 3PCF than observed. This indicates that the simple DM models must be refined, either by using more sophisticated bias models, or a more sophisticated combination of model parameters. ",
        "introduction": "We select data for our analysis from the 100k public release \\footnote{Available at http://www.mso.anu.edu.au/2dFGRS} of the 2dFGRS \\citep[][; hereafter C01]{colless01}. The survey covers two declination strips, one in the Southern Galactic Pole (SGP) and other in the Northern Galactic Pole(NGP), and 99 random fields in the southern galactic cap. In this paper, only galaxies in the two strips are considered. Further criteria for the inclusion of galaxies in our analysis are that they are within the redshift range of $0.02<z<0.25$, have the redshift measurement quality $Q\\ge 3$, and are in regions with redshift sampling completeness $R({\\mathbf \\theta})$ better than 0.1 (where ${\\mathbf \\theta}$ is a sky position). The redshift range restriction is imposed so that the clustering statistics are less affected by the galaxies in the local supercluster, and by the sparse sampling at high redshift. The redshift quality restriction is imposed so that only galaxies with reliable redshifts are used in our analysis. An additional reason is that the redshift completeness mask provided by the survey team, which is used in our analysis, is constructed for the redshift catalog of $Q\\ge 3$. The last restriction is imposed in order to eliminate galaxies in the fields for which the field redshift completeness $c_{F}$ is less than 70 percent (see C01 about the difference between $R({\\mathbf \\theta})$ and $c_F$). These fields are (or will be) re-observed,and have not been included in computing the redshift mask map $R({\\mathbf \\theta})$. Finally, there are a total of 69655 galaxies satisfying our selection criteria, 30447 in the NGP strip and 39208 in the SGP strip.  It is well known that the two-point clustering of galaxies depends on the luminosity \\citep{xia87, boerner91, loveday95,norberg02a}, and the luminosity dependence is an important constraint on galaxy formation models \\citep{kauffmann97,kauffmann99,benson00, yang03a}. We take advantage of the size of the 2dFGRS to carry out a first study of the luminosity dependence of the three point correlation function. The galaxies are divided into three classes; luminous galaxies with absolute magnitude $M_b \\le M_b^{\\star}=-19.66+5\\log h$, faint galaxies with $M_b>-18.5+5\\log h$, and typical galaxies with luminosity in between, where $ M_b^{\\star}$ is the characteristic luminosity of the Schechter function in the $b_J$ band \\citep{norberg02b}, and $h$ is the Hubble constant in units of $100\\kms{\\rm Mpc}^{-1}$. We will also do the analysis for galaxies with $M_b\\le -18.5+5\\log h$ in order to compare the results with the previous study of the Las Campanas Redshift Survey \\citep{jb98}.  The details of the subsamples studied in this paper are given in Table 1. For computing the absolute magnitude, we have used the k-correction and luminosity evolution model of \\citet[][${\\rm k+e}$ model]{norberg02b}, i.e., the absolute magnitude is in the rest frame $b_j$ band at $z=0$. We assume a cosmological model with the density parameter $\\Omega_0=0.3$ and the cosmological constant $\\lambda_0=0.7$ throughout this paper.  A detailed account for the observational selection effects has been released with the catalog by the survey team (C01). The limiting magnitude changes slightly across the survey region due to further magnitude calibrations that were carried out after the target galaxies had been selected for the redshift measurement. This observational effect is documented in the magnitude limit mask $b_J^{\\rm lim} (\\mathbf \\theta) $ (C01). The redshift sampling is far from uniform within the survey region, and this selection effect is given by the redshift completeness mask $R (\\mathbf \\theta)$. The redshift measurement success rate also depends on the brightness of galaxies, making fainter galaxies more incomplete in the redshift measurement. The $\\mu ( \\mathbf \\theta)$ mask provided by the survey team is aimed to account for the brightness-dependent incompleteness.  These observational effects can be corrected in our analysis for the three-point correlation function through properly generating random samples. To construct the random samples, we first select a spatial volume that is sufficiently large to contain the survey sample. Then, we randomly distribute points within the volume, and eliminate the points that are out of the survey boundary. Adopting $15.0<b_J\\le b_J^{\\rm lim} (\\mathbf \\theta)$ for the magnitude limits\\footnote{We assume that the brighter magnitude limit for the survey is 15.0. This is a reasonable value for the survey, but our results are insensitive to the choice of this value.} of the survey in the direction ${\\mathbf \\theta}$, we select random points according to the luminosity function of the 2dFGRS and the $k+e$ model for the k-correction and luminosity evolution \\citep{norberg02a}, and assign to each point an apparent magnitude (and an absolute magnitude). This unclustered sample is a random sample for the 2dFGRS photometric catalog. Then we implement the magnitude-dependent redshift selection effect according to C01. We keep random points of magnitude $m$ in the direction ${\\mathbf \\theta}$ at a sampling rate $S({\\mathbf \\theta},m)$ which reads as [Eq.(11) of C01], \\beq S({\\mathbf \\theta},m)=\\frac{N_p({\\mathbf \\theta})}{N_e({\\mathbf \\theta})} R({\\mathbf \\theta}) c_z[m,\\mu({\\mathbf\\theta})] \\eeq where $N_p({\\mathbf \\theta})$ is the number of parent catalog galaxies in the sector ${\\mathbf \\theta}$ and $N_e({\\mathbf \\theta})$ is the number of galaxies which are expected  to have measured redshifts for given $\\mu({\\mathbf \\theta})$. The ratio $N_e({\\mathbf\\theta})/N_p({\\mathbf\\theta})$ is actually the field completeness $c_F(\\mathbf \\theta)$ defined by C01 which we compute according to their Eq.(7) \\citep[see also ][]{norberg02b}. The function $R({\\mathbf \\theta})$ is given by the redshift completeness mask and $c_z[m,\\mu({\\mathbf\\theta})]$ can be easily computed from the $\\mu$ mask [eq.(5) of C01]. We have used the corrected value $0.5\\ln (10)$ for the $\\alpha$ parameter in the power-law galaxy count model according to the Web page of the 2dFGRS.  We have checked the random samples carefully by reproducing the angular distribution, mean redshift distribution, and especially the two-point statistics of clustering of the observed catalog. It is known that the two-point correlation function measured from galaxy catalogs on large scale is sensitive to the details of corrections for the above selection effects. We have estimated the redshift and projected correlation functions by the same method as in JMB98 for the Las Campanas Redshift Survey. The two-point correlation functions are shown in Figure \\ref{xis} and Figure \\ref{wrpfig}, and can be compared with the results of the 2dFGRS team for the clustering of galaxies \\citep[e.g.][]{hawkins02,norberg02a}. In addition to the broad agreement with their results, even the subtle difference between the north and south caps (the clustering on large scales is slightly larger in the southern cap than in the northern cap), and the luminosity dependence of the clustering, is well reproduced in our analysis.  We did not take into account in our analysis the fiber collision effect that two galaxies closer than $\\sim 30$ arcsec cannot be assigned fibers simultaneously in one spectroscopic observation. Thus one of them will not have a redshift observation if no re-observation is arranged. This effect reduces the real space (or projected) two-point correlation function at small separations. With LCRS, JMB98 estimated the effect to lead to a 15 percent reduction in the two-point correlation function at projected separations $100 \\kpc$ and to a less than 5 percent reduction at separations larger than $400\\kpc$. This effect is smaller (10 percent reduction at separations $100\\kpc$) in the 2dFGRS \\citep{hawkins02}, because the limiting fiber separation is slightly smaller (30 arc sec in the 2dFGRS vs 55 arcsec in the LCRS), and one field may be observed more than once in the 2dFGRS observation strategy.  JB98 have examined the fiber collision effect on their measurement of the three-point correlation function of the LCRS. They found that the effect reduces the real space (projected) three-point correlation function at small separation, but changes little the normalized three-point correlations functions $Q$ that we will measure in this paper, because the effects on the two-point CF and three-point CF are canceled out when $Q$ is measured. Since the effect is slightly smaller in 2dFGRS in terms of the two-point clustering, we believe that only a negligible effect on our measurement of the normalized three-point correlations would result. ",
        "conclusions": "In a detailed analysis of the 3PCF for the 2dFGRS survey we have accurately measured the 3PCF for galaxies of different luminosity. The 3PCF amplitudes ($\\Qsu$ or $\\Qrpu$) of galaxies generally decrease with the increase of the triangle size and increase with the increase of $v$, qualitatively in agreement with the predictions for the dark matter clustering in popular hierarchical CDM models.  Some dependence on luminosity is found, but not a strong effect, except for the brightest galaxy sample which seems to have lower amplitudes of up to $50 \\%$.  Comparing with the previous study on the LCRS galaxies (JB98), we find good agreement between the two studies, though the results from the 2dFGRS are more accurate, since the 2dFGRS survey is much larger than the LCRS survey.  The amplitudes in redshift space $\\Qsu$ are very similar, but the projected ones $\\Qrpu$ show some difference. It seems that the projected 3PCF from the LCRS is systematically underestimated for $r_p$ in the range of a few $\\mpc$, because of the thin slice geometry of that survey. The dependence of $\\Qrpu$ on $r_p$ is much weaker in the 2dFGRS survey than in the LCRS survey.  Comparing the measured 3PCF amplitudes ($\\Qsu$ or $\\Qrpu$) to the prediction of a WMAP concordance model, we find that the measured values are consistently lower than the predicted ones for dark matter. As in the case of the LCRS about one-half of the predicted $Q$ value provides a good description of $\\Qrpu$ for the 2dFGRS data. As in JB98 for the LCRS data, we may state that the best fit DM model gives higher values for the 3PCF than observed. This indicates that the simple DM models must be refined, either by using more sophisticated bias models, or a more sophisticated combination of model parameters.  The division of galaxies into luminosity classes reveals that the brightest galaxies are biased even on large scales, while the galaxies of sample II show a nonlinear bias on small scale, but appear unbiased on linear scales."
    },
    "0311/astro-ph0311066_arXiv.txt": {
        "abstract": "Interstellar grains span a wide range of sizes from a few angstroms to a few micrometers. The presence of nanometer-sized or smaller particles in the interstellar medium is indicated directly by the interstellar far ultraviolet (UV) extinction, the ubiquitous 3.3, 6.2, 7.7, 8.6, and 11.3$\\mum$ polycyclic aromatic hydrocarbon (PAH) emission features, the near and mid infrared broadband emission seen in the IRAS 12 and 25$\\mum$ bands and the COBE-DIRBE 3.5, 4.9, 12 and 25$\\mum$ bands, the 10--100$\\GHz$ Galactic foreground microwave emission, and indirectly by the heating of interstellar gas. For nanoparticles under interstellar conditions, UV/visible photon absorption is the dominant excitation process. With a heat capacity smaller than or comparable to the energy of an energetic stellar photon, nanoparticles are subject to single-photon heating, followed by vibrational relaxation, photoionization, and photodestruction. With excited electrons spatially confined, semiconductor nanoparticles are expected to luminesce efficiently. This review focuses on the photophysics of nanoparticles with emphasis on the stochastic heating and the vibrational excitation of interstellar PAH molecules, and the excitation of photoluminescence with special attention given to silicon nanoparticles. ",
        "introduction": "When the existence of solid particles in interstellar space was first convincingly demonstrated by Trumpler (1930) through the discovery of color excess between the photographic (with an effective wavelength $\\lambdaeff$\\,$\\approx$\\,$4300\\Angstrom$) and visual ($\\lambdaeff$\\,$\\approx$\\,$5500\\Angstrom$) magnitudes, the grains were estimated to be $\\gtsim$$10^{-19}\\g$, corresponding to radii of $\\gtsim$$20\\Angstrom$ (Trumpler 1930). By the end of the 1930s, a $\\lambda^{-1}$ extinction law (i.e., the interstellar extinction varied approximately inversely with wavelength $\\lambda$) in the wavelength range 1--3$\\mum^{-1}$ had been well established (Hall 1937; Stebbins, Huffer, \\& Whitford 1939). At that time, metallic grains with a dominant size of $\\simali$0.05$\\mum$ (Schal\\'{e}n 1936) or a power-law size distribution $dn(a)/da$\\,$\\simali$$a^{-3.6}$ in the size range $80\\Angstrom$\\,$\\ltsim$\\,$a$\\,$\\ltsim$\\,$1\\cm$ (Greenstein 1938) were proposed to explain the $\\lambda^{-1}$ extinction law, partly because meteoritic particles (predominantly metallic) and interstellar grains were then thought to have the same origin.  In view of the spatial correlation between gas concentration and dust extinction, Lindblad (1935) suggested that interstellar grains were formed by condensation from the interstellar gas through random accretion of gas atoms. However, it was found later that in typical interstellar conditions, the Lindblad condensation theory would result in a complete disappearance of all condensable gases and the grains would grow to sizes ($\\simali$10$\\mum$) well beyond those which would account for the interstellar extinction (see Oort \\& van de Hulst 1946 and references therein). By introducing a grain destruction process caused by grain-grain collisions as a consequence of interstellar cloud encounters, Oort \\& van de Hulst (1946) further developed the interstellar condensation idea and led to the ``dirty ice'' model consisting of saturated molecules such as H$_2$O, CH$_4$, and NH$_3$ with an equilibrium size distribution which could be roughly approximated by a functional form $dn(a)/da\\,$$\\simali$$\\exp\\left[-5 \\left(a/0.5\\,\\mu {\\rm m}\\right)^3\\right]$ and an average size of $\\simali$0.15$\\mum$.  The discovery of interstellar polarization (Hall 1949; Hiltner 1949) cast doubts on the ``dirty ice'' model since ice grains are very inefficient polarizers. This stimulated Cayrel \\& Schatzman (1954) to consider graphite grains as an interstellar dust component since aligned graphite particles would be a very efficient interstellar polarizer because of their strong optical anisotropy. The ``dirty ice'' model was also challenged by nondetection of the 3.1$\\mum$ feature of H$_2$O ice outside of molecular clouds (Danielson, Woolf, \\& Gaustad 1965; Knacke, Cudaback, \\& Gaustad 1969). This gave the late J. Mayo Greenberg (1922--2001) the incentive to perform the early experiments on the ultraviolet (UV) photoprocessing of low temperature mixtures of volatile molecules simulating the ``original'' dirty ice grains (Greenberg et al.\\ 1972) to understand how and why the predicted H$_2$O was not clearly present. From such experiments was predicted a new component of interstellar dust in the form of complex organics known as ``organic refractories'', containing a mixture of aliphatic and aromatic carbonaceous molecules (Greenberg et al.\\ 2000).  As an alternative to the interstellar condensation process, Hoyle \\& Wickramasinghe (1962) proposed that graphite particles with radii a few times 0.01$\\mum$ could form in the atmospheres of cool N-type carbon stars, and subsequently be ejected into interstellar space by the stellar radiation pressure, and be responsible for (part of) the interstellar extinction and polarization. The graphite proposal further gained its strength from the prominent 2175$\\Angstrom$ interstellar extinction hump detected by the {\\it Aerobee} rocket observation (Stecher 1965), which was generally attributed to small graphitic grains (Stecher \\& Donn 1965), although its exact nature still remains unidentified (Draine 1989, 2003a).  Similarly, Kamijo (1963) first proposed that SiO$_2$ whose size is of the order of $\\simali$20$\\Angstrom$, condensed in the atmospheres of cool M-type stars and expelled into interstellar space, could serve as condensation nuclei for the formation of ``dirty ices''. It was later shown by Gilman (1969) that grains around oxygen-rich cool giants are mainly silicates such as Al$_2$SiO$_3$ and Mg$_2$SiO$_4$. Interstellar silicates were first detected in emission in the Trapezium region of the Orion Nebula (Stein \\& Gillett 1969); in absorption toward the Galactic Center (Hackwell, Gehrz, \\& Woolf 1970), and toward the Becklin-Neugebauer object and the Kleinmann-Low Nebula (Gillett \\& Forrest 1973). Silicates are now known to be an ubiquitous interstellar dust component (see \\S1 in Li \\& Draine 2001a for a review).  In the 1960s and early 1970s, the extension of the extinction curve toward the middle and far UV ($\\lambda$$\\gtsim$$3\\mum^{-1}$) was made possible by rocket and satellite observations, including the rocket-based photoelectric photometry at $\\lambda=2600\\Angstrom$ and 2200$\\Angstrom$ (Boggess \\& Borgman 1964); the {\\it Aerobee} Rocket spectrophotometry at 1200$\\Angstrom$\\,$\\ltsim$\\,$\\lambda$\\,$\\ltsim$3000$\\Angstrom$ (Stecher 1965); the {\\it Orbiting Astronomical Satellite} (OAO-2) spectrophotometry at 1100$\\Angstrom$\\,$\\ltsim$\\,$\\lambda$\\,$\\ltsim$\\,$3600\\Angstrom$ (Bless \\& Savage 1972); and the {\\it Copernicus} Satellite spectrophotometry at 1000$\\Angstrom$\\,$\\ltsim$\\,$\\lambda$\\,$\\ltsim$\\,$1200\\Angstrom$ (York et al.\\ 1973). By 1973, the interstellar extinction curve had been determined over the whole wavelength range from 0.2$\\mum^{-1}$ to 10$\\mum^{-1}$. The fact that the extinction continues to increase in the far UV (e.g., see York et al.\\ 1973) implies that no single grain type with either a single size or a continuous size distribution could account for the observed optical to far-UV interstellar extinction (Greenberg 1973). This led to the abandonment of any one-component grain models and stimulated the emergence of various kinds of models consisting of multiple dust constituents, including silicate, silicon carbide, iron, iron oxide, graphite, dirty ice, solid H$_2$, etc.\\\\[1mm]  In 1956 Platt first suggested that very small grains or large molecules of less than 10$\\Angstrom$ in radius grown by random accretion from the interstellar gas could be responsible for the observed interstellar extinction and polarization. Platt (1956) postulated these ``Platt'' particles as quantum-mechanical particles containing many ions and free radicals with unfilled electronic energy bands. Donn (1968) further proposed that polycyclic aromatic hydrocarbon-like ``Platt particles'' may be responsible for the UV interstellar extinction.  Greenberg (1968) first pointed out that very small grains with a heat content smaller than or comparable to the energy of a single stellar photon, cannot be characterized by a steady-state temperature but rather are subject to substantial temporal fluctuations in temperature. Under interstellar conditions, grains larger than a few hundred angstroms in radius would attain an equilibrium temperature in the range of 15--25$\\K$ (Greenberg 1968, 1971; Mathis, Mezger, \\& Panagia 1983; Draine \\& Lee 1984; Li \\& Draine 2001b). This temperature was long thought to be $\\simali$3$\\K$ before van de Hulst (1949) made the first realistic estimation of 10--20$\\K$ for both metallic and dielectric grains.  In a study of the scattering properties of interstellar dust (albedo and phase function) determined from the OAO-2 observations at 1500$\\Angstrom$\\,$\\ltsim$\\,$\\lambda$\\,$\\ltsim$\\,$4250\\Angstrom$ of the diffuse Galactic light (Witt \\& Lillie 1973), Witt (1973) first explicitly suggested a bi-modal size distribution for interstellar grains: large grains with radii $\\gtsim$$2500\\Angstrom$ would provide extinction in the visible region including scattering which is strongly forward directed, and small particles with radii $\\ltsim$$250\\Angstrom$ would dominate the UV region and contribute nearly isotopic scattering.  The presence of ``Platt'' particles in a dust cloud near M17 was first suggested by Andriesse (1978) based on an analysis of its IR spectral energy distribution which was shown to be characterized by a combination of widely different color temperatures (e.g., the far-IR and near-IR color temperatures are $\\simali$38$\\K$ and $\\simali$150$\\K$, respectively) and that the 8--20$\\mum$ emission spectrum is similar over a distance of $\\simali$2$^{\\prime}$ through the source. Andriesse (1978) argued that the color temperature differences cannot be easily ascribed to a spatial variation of the dust temperature in this cloud, because the cloud profile appears to be independent of the wavelength; instead, ``Platt'' particles with sizes of $\\simali$10$\\Angstrom$ exhibiting temperature fluctuations up to $\\simali$150$\\K$ plus large grains with an equilibrium temperature of $\\simali$36$\\K$ can explain  the observed IR spectrum of M17.  In a near-IR (1.25--4.8$\\mum$) photometric and spectrophotometric study of three visual reflection nebulae NGC 7023, 2023, and 2068, Sellgren, Werner, \\& Dinerstein (1983) discovered that each nebula has extended near-IR emission consisting of emission features at 3.3 and 3.4$\\mum$ and a smooth continuum which can be described by a greybody with a color temperature $\\simali$1000$\\K$. They found that the emission spectrum (i.e., the color temperature for the continuum and the 3.3$\\mum$ feature) shows very little variation from source to source and within a given source with distance from the central star. Sellgren (1984) argued that this emission could not be explained by thermal emission from dust in radiative equilibrium with the central star since otherwise the color temperature of this emission should fall off rapidly with distance from the illuminating star; instead, she proposed that this emission is emitted by ultrasmall grains, $\\simali$10$\\Angstrom$ in radius, which undergo large excursions in temperature due to stochastic heating by single stellar photons.  The presence of a population of ultrasmall grains in the diffuse ISM was explicitly indicated by the 12$\\mum$ and 25$\\mum$ ``cirrus'' emission detected by the {\\it Infrared Astronomical Satellite} (IRAS) (Boulanger \\& P\\'{e}rault 1988), which is far in excess (by several orders of magnitude) of what would be expected from large grains of 15--25$\\K$ in thermal equilibrium with the general interstellar radiation field. Subsequent measurements by the {\\it Diffuse Infrared Background Experiment} (DIRBE) instrument on the {\\it Cosmic Background Explorer} (COBE) satellite confirmed this and detected additional broadband emission at 3.5$\\mum$ and 4.9$\\mum$ (Arendt et al.\\ 1998).  More recently, spectrometers aboard the {\\it Infrared Telescope in Space} (IRTS) (Onaka et al.\\ 1996; Tanaka et al.\\ 1996) and the {\\it Infrared Space Observatory} (ISO) (Mattila et al.\\ 1996) have shown that the diffuse ISM radiates strongly in emission features at 3.3, 6.2, 7.7, 8.6, and 11.3$\\mum$. These emission features, first seen in the spectra of the planetary nebulae NGC 7027 and BD+30$^{\\rm o}$3639 (Gillett, Forrest, \\& Merrill 1973), have been observed in a wide range of astronomical environments including planetary nebulae, protoplanetary nebulae, reflection nebulae, HII regions, circumstellar envelopes, and external galaxies (see Tielens et al.\\ 2000 for a review). Often referred to as ``unidentified infrared'' (UIR) bands, these emission features are now usually attributed to free-flying PAH molecules which are vibrationally excited upon absorption of a single UV/visible photon (L\\'{e}ger \\& Puget 1984; Allamandola, Tielens, \\& Barker 1985; Allamandola, Hudgins, \\& Sandford 1999; Draine \\& Li 2001; Li \\& Draine 2001b) although other carriers have also been proposed such as carbonaceous grains with a partly aromatic character (e.g., HAC [Duley \\& Williams 1981, Duley, Jones, \\& Williams 1990], quenched carbonaceous composite [QCC; Sakata et al.\\ 1992], coal [Papoular et al.\\ 1996], fullerenes [C$_{60}$; Webster 1993]), and surface-graphitized nanodiamonds (Jones \\& d'Hendecourt 2000).  \\vspace{3mm} Since the late 1970s, various modern interstellar dust models have been developed (see Li \\& Greenberg 2003, Draine 2004, Dwek et al.\\ 2004 for recent reviews). In general, these models fall into three broad categories: \\vspace{-1.0mm} \\begin{itemize} \\item \\underline{The {\\bf silicate-graphite} model} -- this model consists of two (physically) separate dust components each with a power-law size distribution $dn(a)/da \\sim a^{-3.5}$ in the size range $50\\Angstrom$\\,$\\ltsim$\\,$a$\\,$\\ltsim$\\,$0.25\\mum$ (Mathis, Rumpl, \\& Nordsieck 1977 [hereafter MRN]; Draine \\& Lee 1984). Modifications to this model were later made by Draine \\& Anderson (1985), Weiland et al.\\ (1986), Sorrell (1990), Siebenmorgen \\& Kr\\\"{u}gel (1992), Rowan-Robinson (1992), Kim, Martin, \\& Hendry (1994), Dwek et al.\\ (1997), Clayton et al.\\ (2003), and Zubko, Dwek, \\& Arendt (2003) by including new dust components (e.g., amorphous carbon, carbonaceous organic refractory, and PAHs) and adjusting dust sizes (e.g., deriving dust size distributions using the ``Maximum Entropy Method'' or the ``Method of Regularization'' rather than presuming a certain functional form).  Recent developments were made by Draine and his coworkers (Li \\& Draine 2001b, 2002a, 2002b; Weingartner \\& Draine 2001a) who have extended the silicate-graphite grain model to explicitly include a PAH component as the small-size end of the carbonaceous grain population. The PAH component, containing $\\approx 45\\times 10^{-6}$ of C relative to H, is represented by a log-normal size distribution $dn(a)/d\\ln a \\sim \\exp\\left\\{-\\left[\\ln\\left(a/a_0\\right)\\right]^2/ \\left(2\\sigma^2\\right)\\right\\}$ with $a_0\\approx 3.5\\Angstrom$ and $\\sigma\\approx 0.4$ for $a$$\\gtsim$$3.5\\Angstrom$ (see Li \\& Draine 2001b). Note that with the PAH component added, the size distribution of the larger grains is changed from the simple MRN power-law (see Weingartner \\& Draine 2001a; Zubko et al.\\ 2003). \\vspace{-1.0mm} \\item \\underline{The {\\bf silicate core-carbonaceous mantle} model} -- originally proposed by Greenberg (1978) in the context of a complex cyclic evolutionary scenario (see Greenberg \\& Li 1999), this model consists of larger silicate grains coated by a layer of carbonaceous organic refractory material, produced by UV photolysis of ice mixtures which attempts to simulate the physical and chemical processes occurring in interstellar space. The most recent development of this model was that of Li \\& Greenberg (1997), who modeled the core-mantle grains as 2:1 (the ratio of the length to the diameter) finite cylinders (to account for the interstellar polarization) with a Gaussian size distribution for the mantle $dn(a)/da \\sim \\exp\\left[-5\\left(a-a_{\\rm c}\\right)^2/a_{\\rm i}^2\\right]$ and a single silicate core radius $a_{\\rm c}$$\\approx$$0.07\\mum$ and a ``cut-off'' size $a_{\\rm i}$$\\approx$$0.066\\mum$. In addition, a PAH component and a population of small graphitic grains are added respectively to account for the far-UV extinction rise plus the ``UIR'' emission bands and the 2175$\\Angstrom$ extinction hump.  Again, modifications to this model were also made by considering different coating materials (e.g., amorphous carbon, HAC), including new dust type (e.g., iron, small bare silicates), and varying dust size distributions (Chlewicki \\& Laureijs 1988; Duley, Jones, \\& Williams 1989; D\\'{e}sert, Boulanger, \\& Puget 1990; Li \\& Greenberg 1998; Zubko 1999a). The Duley et al.\\ (1990) model exclusively consists of silicate core-HAC mantle grains with a bimodal size distributions: (1) $\\simali$43\\% of the total dust mass in grains $\\ltsim$$100\\Angstrom$ which produce the far-UV extinction rise and the 2175$\\Angstrom$ hump through an electronic transition of the OH$^{-}$ ions in low-coordination sites on or within silicate grains, and (2) the remaining dust mass in grains 0.05$\\ltsim$$a$$\\ltsim$$0.25\\mum$ with a power-law $dn(a)/da \\sim a^{-3.5}$ size distribution. The ``UIR'' bands were attributed to the thermally isolated aromatic ``islands'' of HAC material -- Duley et al.\\ (1989) postulated that an absorbed stellar photon might remain localized in a single aromatic unit of the grain long enough for the unit to cool by IR vibrational relaxation rather than by transferring the photon energy to the phonon spectrum of the grain. \\item \\underline{The {\\bf composite grain} model} -- realizing that grain shattering due to grain-grain collisions and subsequent reassembly through agglomeration of grain fragments may be important in the ISM, Mathis \\& Whiffen (1989) modeled the interstellar grains as composite collections of small silicates, vacuum ($\\approx 80\\%$ in volume), and carbon of various kinds (amorphous carbon, HAC, organic refractories) with a power-law size distribution $dn(a)/da$$\\simali$$a^{-3.7}$ in the size range $0.03\\mum$\\,$\\ltsim$\\,$a$\\,$\\ltsim$\\,$0.9\\mum$. In addition, a separate small graphite component containing $\\approx 59\\times 10^{-6}$ of C is needed to account for the 2175$\\Angstrom$ extinction hump. Wright (1987) argued that a fractal structure would be expected for interstellar grains formed through the coagulation of small grain fragments created from grain disruption caused by supernova shock waves. In view that the relative abundances of refractory elements in the ISM may be as low as 65\\% of solar (Snow \\& Witt 1996), Mathis (1996) revised the composite model in order to satisfy the tighter abundance constraints. To optimize the use of the heavy elements, Mathis (1996) derived a vacuum fraction of $\\simali$45\\%. The ``UIR'' emission, which remained unaccounted for in the models of Mathis \\& Whiffen (1989) and Mathis (1996), was taken into account in Zubko et al.\\ (2003) by including a population of PAH molecules. \\end{itemize}  It is fair to say that, as can be seen from the above summarizing description of the key contemporary dust models, a consensus on the interstellar dust compositions and sizes is now approaching among various grain models, although the debate in details is still going on and is not expected to disappear in the near future (except that the IR spectropolarimetric observation at 3.4$\\mum$ and 9.7$\\mum$ may allow a direct test of the core-mantle model for interstellar dust; see Li \\& Greenberg 2002) --- \\begin{itemize} \\item Regarding the grain chemical composition, the most generally accepted view is that interstellar grains consist of amorphous silicates and some form of carbonaceous materials; the former is inferred from the 9.7$\\mum$ Si--O stretching mode and 18$\\mum$ O-Si-O bending mode absorption features in interstellar regions as well as the fact that the cosmically abundant heavy elements such as Si, Fe, Mg are highly depleted; the latter is mainly inferred from the 2175$\\Angstrom$ extinction hump (and the ubiquitous 3.4$\\mum$ C--H stretching vibrational band) and the fact that silicates alone are not able to provide enough extinction. For the carbonaceous component, a wide range of dust materials have been suggested including amorphous carbon, coal, C$_{60}$, diamond, graphite, HAC, PAHs, organic refractory, and QCC (see Pendleton 2004 for a review of the carbonaceous component). \\item It is also generally accepted that, through the analysis of the wavelength-dependent interstellar extinction and polarization curves as well as the near and mid IR emission, interstellar grain sizes may be separated into two domains -- (1) the ``large'' grain component (with radii $a$$>$$0.025\\mum$; including the ``classical'' grains [with $a$$\\gtsim$$0.1\\mum$]) which is primarily responsible for the extinction, polarization and scattering at visible wavelengths and the IR emission at $\\lambda$$\\gtsim$$60\\mum$; and (2) the ``very small grain'' component (with $a$$<$$0.025\\mum$) which contributes importantly to the extinction in the vacuum-UV and emit strongly in the near and mid IR at $\\lambda$$\\ltsim$$60\\mum$ (see Figs.\\,8,16 of Li \\& Draine 2001b) when transiently heated to high temperatures during quantized absorption events.  While the size distribution for the ``classical grain'' component is relatively well constrained by fitting the observed interstellar extinction curve for an assumed dust composition, our knowledge of the size distribution $dn/da$ for the ``very small grain'' component is better constrained by the interstellar near and mid IR emission due to the fact that, for $\\lambda$$\\gtsim$$0.1\\mum$, these very small grains are in the Rayleigh limit ($x \\equiv 2\\pi a/\\lambda$$\\ll$1) and their extinction cross sections $C_{\\rm ext}(a,\\lambda)$ per unit volume $V$ are independent of size, so that the observed UV/far-UV extinction curve $A(\\lambda)/\\NH$ (${\\rm mag\\,cm^{-2}\\,H^{-1}}$) only constrains the total volume $V_{\\rm tot}$ of this grain component: $A(\\lambda)/\\NH = 1.086\\,\\NH^{-1} \\int C_{\\rm ext}(a,\\lambda) \\left(dn/da\\right)da = 1.086\\,\\NH^{-1} V_{\\rm tot} \\left(C_{\\rm ext}/V\\right)$, where $\\NH$ is the hydrogen column density. In contrast, the near and mid IR intensity $I_\\lambda$ is sensitive to the grain heat capacity ($\\propto$\\,$a^3$) which determines the maximum temperature to which the grain can reach when illuminated by a radiation field: $I_\\lambda = \\NH^{-1} \\int C_{\\rm abs}(a,\\lambda) \\left(dn/da\\right)da \\int B_\\lambda(T[a])\\left(dP/dT\\right)dT$ where $C_{\\rm abs}(a,\\lambda)$ is the absorption cross section for a grain of radius $a$ at wavelength $\\lambda$ (for grains in the Rayleigh limit, $C_{\\rm abs}$\\,$\\approx$\\,$C_{\\rm ext}$), $B_\\lambda$ is the Planck function, $dP/dT$ is the dust temperature distribution function. Evidently, $dP/dT$ is a sensitive function of grain size $a$ (see Draine \\& Li 2001, Li \\& Draine 2001b; and \\S3, Fig.\\,6 of this paper). \\end{itemize}  In the following sections of this review, we will focus on the ultrasmall grain component ($a$$\\ltsim$$25\\nm$), with particular emphasis on the photophysical processes including the stochastic heating and the vibrational excitation of PAH molecules, and the excitation of the photoluminescence of silicon nanoparticles. In astrophysical literature, one frequently encounters terms like ``ultrasmall grains'', ``very small grains'', ``large molecules'', ``tiny grains'', and ``nanoparticles''; to avoid confusion, we emphasize here that they are synonymous. In the following we will use the term ``nanoparticles'' -- by ``nanoparticles'' we mean grains of a few angstroms to a few tens of nanometers in radius.  In \\S2 we summarize the direct and indirect evidence for the existence of nanoparticles in the ISM. The vibrational excitation of nanoparticles is detailed in \\S3. In \\S4 we discuss the physics regarding the excitation of photoluminescence. In \\S5 we present an overview of the nanoparticle species known or proposed to exist in interstellar space. Concluding remarks are given in \\S6.  \\vspace{-3.0mm} ",
        "conclusions": "\\vspace*{-0.5em} Interstellar nanoparticles are important as emitters of near and mid IR radiation, as absorbers of far-UV radiation, as heating agents of the interstellar gas, and possibly as luminescing agents of red light. Our understanding of interstellar nanoparticles obviously remains incomplete. Besides the long-standing problems regarding the carriers of the 2175$\\Angstrom$ extinction hump and the ERE, there are many unanswered questions which will be demanding close attention in the future: \\begin{itemize} \\vspace*{-0.5em} \\item The origin and evolution of interstellar PAHs are not very clear. Suggested sources for interstellar PAHs include (1) injection (into the ISM) of PAHs formed in carbon star outflows (Latter 1991); (2) shattering of carbonaceous interstellar dust or of photoprocessed interstellar dust organic mantles (Greenberg et al.\\ 2000) by grain-grain collisions in interstellar shocks (Jones, Tielens, \\& Hollenbach 1996); (3) {\\it in-situ} formation through ion-molecule reactions (Herbst 1991). Although interstellar PAHs containing more than $\\simali$20--30 carbon atoms can survive the UV radiation field (Guhathakurta \\& Draine 1989; Jochims et al.\\ 1994; Allain, Leach, \\& Sedlmayr 1996; Le Page, Snow, \\& Bierbaum 2003), they are still subject to destruction by sputtering in interstellar shock waves (but also see Tielens et al.\\ [2000] who argued that destruction of PAHs by sputtering is unimportant, except in extreme environments such as very young supernova remnants with $\\gtsim$200$\\kms$ shocks, because they couple very well dynamically to the gas which cools down rapidly), chemical attack by atomic oxygen, and coagulation in dense regions (Draine 1994). Clearly, a detailed study of the evolution of interstellar PAHs would be very valuable. \\item The physics and astrophysics of the radiative electronic transitions of interstellar PAHs are not fully understood.  Small neutral PAHs with $\\ltsim$40 C atoms are expected to emit near-UV and blue photons through fluorescence, phosphorescence, and perhaps also through the recurrent Poincar\\'e fluorescence. Laboratory studies show that the fluorescence quantum yield can be quite high for isolated molecules (e.g., fluorescence quantum yields in the range of 10\\%--45\\% have been measured for gas-phase, collision-free naphthalene C$_{10}$H$_{8}$ [Reyl\\'e \\& Br\\'echignac 2000]). However, no PL shortward of $\\lambda$=$5000\\Angstrom$ has been seen in the ISM (Rush \\& Witt 1975; Vijh, Witt, \\& Gordon 2003). Although PAH ions do not exhibit strong ordinary fluorescence or phosphorescence (see the caption of Fig.\\,2), they can undergo Poincar\\'e fluorescence which may have a quantum yield larger than one. Therefore, it is difficult to explain nondetection of blue PL in the ISM in terms of PAH photoionization. Moreover, for PAHs containing $\\ltsim$40 C atoms in the diffuse ISM, the probability of finding them in a nonzero charge state is smaller than $\\simali$30\\% (see Fig.\\,7 of Li \\& Draine 2001b). One may argue that interstellar PAHs are larger (which is true, as shown in Li \\& Draine 2001b, the average [diffuse ISM] PAH size is $\\simali$6$\\Angstrom$, corresponding to $\\simali$100 C atoms) so that their fluorescence mainly occurs at longer wavelengths, say, in the wavelength range over which the ERE is observed.  However, if interstellar PAHs indeed luminesce in the ERE band with the required quantum efficiency ($\\simali$100\\%; see \\S2), the ``UIR'' emission bands expected for these luminescing PAHs would be strongly suppressed because a considerable fraction of the excitation energy is released in the form of PL photons (e.g., in the diffuse ISM with a mean energy $\\approx$\\,2.1$\\eV$ for ERE photons [Szomoru \\& Guhathakurta 1998], only $\\simali$60\\% of the original excitation energy is available as heat for PAHs which have a mean absorbed photon energy $\\approx$5.2$\\eV$ [Draine \\& Li 2001]; this could be even worse since the Poincar\\'e fluorescence may have a quantum yield higher than one). One consequence of this would be that even more carbon atoms in the form of PAHs would be needed to explain the intensity of the ``UIR'' bands, deteriorating the already tightened carbon budget ``crisis'' (Snow \\& Witt 1996). \\vspace*{-0.4em} \\item It is unclear how different the optical and thermal properties of nanometer-sized materials are compared with their bulk counterparts, although it is generally believed that they may be very different. \\begin{itemize} \\vspace*{-0.4em} \\item For a small metallic grain, the imaginary part $\\ei$ of its dielectric function $\\epsilon(\\lambda) = \\er +i\\,\\ei$ is expected to be larger compared to that of its bulk counterpart, as a consequence of the so-called {\\it electron mean free path limitation} effect (e.g. see Bohren \\& Huffman 1983). This is easier to understand if we decompose $\\epsilon$ into two components $\\eb$ and $\\ef$, contributed by bound charges (``interband transitions'') and free electrons, respectively. The free electron component is well described by the Drude theory $\\epsilon = 1 - \\frac{\\omega_p^2}{\\omega^2+i\\,\\gamma\\omega}$ with an imaginary part $\\ei = \\frac{\\gamma\\omega_p^2} {\\omega\\left(\\omega^2+\\gamma^2\\right)}$, where the plasma frequency $\\omega_p$ is related to the free electron density $n_e$: $\\omega_p^2 = n_e e^2/\\meeff$, the damping constant $\\gamma$ is related to the average time $\\tau$ between collisions: $\\gamma$=$1/\\tau$. In bulk materials, $\\gamma$ is mainly determined by the scattering of the electrons with phonons (lattice vibrations), and to a lesser degree, with electrons, lattice defects, or impurities. However, for particles in the nanometer size domain, $\\gamma$ is increased because of the additional collisions of the conducting electrons with the grain boundary: $\\gamma = \\gamma_{\\rm bulk} + \\frac{v_F}{\\beta a}$, where $\\gamma_{\\rm bulk}$ is the bulk metal damping constant, $v_F$ is the electron velocity at the Fermi surface, $\\beta$ is a dimensionless constant of order unity which depends on the character of the scattering at the boundary: $\\beta$=1 for classic isotropic scattering, $\\beta$=4/3 for classic diffusive scattering, $\\beta$=1.16 or $\\beta$=1.33 for scattering based on the quantum particle-in-a-box model (see Coronado \\& Schatz 2003 and references therein). Since $\\omega^2\\gg \\gamma^2$ in metals near the plasma frequency, $\\ei$ can be written as $\\ei = \\ei_{\\rm bulk} + \\frac{v_F \\omega_p^2}{\\beta a \\omega^3}$. Clearly, for a metallic grain $\\ei$ increases as the grain becomes smaller.\\\\[1mm] In contrast, our knowledge regarding the size dependence of the dielectric function of {\\it dielectric} materials is controversial. For example, some theoretical and experimental studies have concluded that the dielectric function of Si nanoparticles is significantly reduced relative to the bulk value (Koshida et al.\\ 1993; Wang \\& Zunger 1994; Tsu, Babic, \\& Ioriatti 1997; Amans et al.\\ 2003), which is generally attributed to the quantum confinement effect. But we find that the absorption and reflectivity measurements for porous silicon appear to be consistent with bulk Si together with voids and SiO$_2$ (see Li \\& Draine 2002a for details). \\vspace*{-0.5em} \\item It has been reported that the specific heat of small metal particles is strongly enhanced over the bulk value (see  Halperin 1986, Meyer et al.\\ 2003 and references therein). For example, a progressive decreasing of Debye temperature $\\Theta$ with the decrease of grain size has been observed for palladium (Pd): $\\Theta\\approx 273, 226, 193, 175\\K$ for bulk Pd and Pd particles of radius $a=42,33,15\\Angstrom$, respectively (Chen et al.\\ 1995 and references therein). This has been attributed to quantum effects on the vibrational spectrum in small particles: as a grain becomes smaller, a larger fraction of atoms occupy surface sites which are weakly bound to the grain; therefore, small grains are expected to have a larger low-frequency mode density due to the weaker bonds of the surface atoms.  However, although an enhancement of the vibrational specific heat is also expected for dielectric grains, the degree of enhancement is unclear due to the differences between the binding and structural properties of dielectrics with metals. For example, based on a lattice dynamical calculation, Hu et al.\\ (2001) have shown that the difference in vibrational specific heat between Si nanocrystals and the bulk is just about a few per cent. The specific heat of nanocrystalline diamond measured by Moelle et al.\\ (1998) shows a close agreement with that of {\\it bulk} diamond. It is worth noting that the heat capacity of PAHs can be well described by the two 2-dimensional Debye models of {\\it bulk} graphite (together with contributions from the C--H vibrational modes) (see Fig.\\,2 of Draine \\& Li 2001). Clearly, laboratory studies of the optical and thermal properties of interstellar nanoparticle analogues would be very valuable. \\end{itemize} \\end{itemize}"
    },
    "0311/astro-ph0311299_arXiv.txt": {
        "abstract": "% Supernova can be polarized by an asymmetry in the explosion process, an off-center source of illumination, scattering in an envelope distorted by rotation or a binary companion, or scattering by the circumstellar dust. Careful polarimetry should thus provide insights to the nature of supernovae. Spectropolarimetry is the most powerful tool to study the 3-D geometry of supernovae. A deep understanding of the 3-D geometry of SNe is critical in using them for calibrated distance indicators. ",
        "introduction": "Polarimetry of supernovae (SNe) reveals the intrinsic ejecta asymmetries (Shapiro \\& Sutherland 1982, McCall 1984, H\\\"oflich 1991, 1996). SN~1987A represent a breakthrough in this area, by providing the first detailed record of the spectropolarimetric evolution (e.g. Mendez et al. 1988; Cropper et al. 1988). SN~1993J also provided a wealth of data (Trammell, Hines, \\& Wheeler 1993; Tran et al. 1997). Most of the theoretical interpretations of the polarimetry data are based on oblate or prolate spheroid geometries. A very different picture of SN polarization is discussed in Wang \\& Wheeler (1996) where time-dependent dust scattering is shown to be a potential mechanism. New attempts are made with more complicated geometrical structures (Kasen et al. 2003; 2004).   We started a systematic program of supernova spectropolarimetry in 1995 using the 2.1 meter telescope of the McDonald Observatory which nearly doubled the number of SNe with polarimetry measurements in the first year of the program. The early qualitative conclusion was that core-collapse SNe are in general polarized at about 1\\% level and that Type Ia are much less polarized ($\\le 0.3\\%$) (Wang et al. 1996). We have also observed several SNe for their spectropolarimetry using the Image Grism Polarimeter (IGP) at the 2.1 meter telescope and a polarimeter at 2.7 meter telescope. The spectropolarimetry data show that normal SN Ia are not highly polarized, but a hint of intrinsic polarization is detected at levels around 0.3\\% for SN~1996X at around optical maximum. We found also that a peculiar, subluminous SN~Ia 1999by was clearly polarized at a much higher level ($\\sim$ 0.7\\%, Howell et al. 2001) than normal SN~Ia. The polarization of core-collapse SNe evolve with time (Wang et al. 2001; Leonard et al. 2001), with the general trend being that the polarization is larger post optical maximum than before or around optical maximum. Our observations indicate further that for core-collapse SNe the degree of polarization is larger for those with less massive envelopes (Wang et al 2001).  The ESO-VLT, with its extremely flexible ToO capability, makes it possible to study in great detail SN polarimetry. Not only can the SNe be follow to late nebular stages, but also they can be observed early enough so that the most energetic outer layers can be probed. We have obtained high quality multi-epoch spectropolarimetry data of a normal Type Ia SN~2001el. The data allow us to probe both the mechanisms of SN Ia explosions, and the effect of asymmetry on the use of SN Ia as distance indicators (Wang et al. 2003a).  For core-collapse events, the inferred degree of asymmetry is in general of the order of 10-30\\% if modeled in terms of oblate/prolate spheroids. The asphericity of the photosphere of SN Ia is perhaps not larger than 10-15\\%.  SN polarimetry is still a rapidly growing field in terms of both observations and theories. Spectropolarimetry of SNe has been a driving force for 3-D models of SN Ia and in turn the advance of 3-D models of SN explosion will help us to understand better the observed phenomena. Here I will concentrate on the observational efforts of our SN polarimetry program. ",
        "conclusions": ""
    },
    "0311/astro-ph0311250_arXiv.txt": {
        "abstract": "We discuss the generalisation of the so-called traditional approximation, well known in geophysics, to general relativity. We show that the approximation is applicable for rotating relativistic stars provided that one focuses on relatively thin radial shells. This means that the framework can be used to study waves in neutron star oceans. We demonstrate that, once the effects of the relativistic frame-dragging are accounted for, the angular problem reduces to Laplace's tidal equation. We  derive the dispersion relation for various classes of waves in a neutron star ocean and show that the combined effects of the frame-dragging and the gravitational redshift typically lower the frequency of a mode by about 20\\%. ",
        "introduction": "Accreting neutron stars exhibit a variety of quasiperiodic phenomena, with frequencies ranging from a few Hz to well above 1~kHz. Many of the observed oscillations are likely associated with processes in the accretion disk, but it is also probable that some of the observed features are linked to waves in the neutron star ocean \\citep{bc,heyl,stroh}. For example, it has been observed that neutron stars that accrete matter at rates $\\gtrsim 10^{-9} M_{\\odot} \\mathrm{yr}^{-1} $(the  so-called Z-sources) show quasi-periodic oscillations with frequencies of $\\sim 6$ Hz. It has also been argued that the Z-sources are covered by massive oceans comprising a mass of $\\sim 10^{-6} M_{\\odot}$ composed of a degenerate liquid of light elements (C, O, Ne, Mg,..). This led \\citet{bc} to argue that the observed oscillations may correspond to ocean g-modes.  To analyse this suggestion in more detail \\citet{buc} investigated the nature of low-frequency waves in the ocean of a rotating neutron star. They considered slowly rotating stars (far below the breakup limit) in Newtonian gravity, and made use of the assumption that $\\Omega\\gg\\sigma$, where $\\Omega$ is the rotation  frequency and $\\sigma$ is the mode frequency. This is a reasonable assumption since neutron stars accreting at a high rate are expected to spin much faster than a few Hz. The studies of \\citet{bc} and \\citet{buc} both made use of the so-called ``traditional'' approximation, which has its origins in geophysics. This approximation leads to a separation of variables for the linearised Euler equations, where the angular problem is described in terms of Laplace's tidal equation. The solutions to this equation have been discussed in detail by many authors \\citep{mkb,mil,buc,ls,town}.  The purpose of this paper is to extend the ``traditional approximation'' to general relativity. The main motivation for this is that we know from the studies of \\citet{bc,buc} \\citet{town} that the approximation provides a relatively simple way of analysing ocean waves in rotating Newtonian stars. Yet there are several reasons for studying this problem in the framework of general relativity. First of all, neutron stars are highly relativistic objects with $GM/Rc^2\\approx 0.15-0.2$. This means that one must account for the gravitational redshift when comparing observed oscillations to theoretical models. A second relativistic effect that needs to be accounted for is the rotational frame-dragging. Finally, the traditional approximation scheme is interesting from a conceptual point of  view. Especially since it provides an alternative approach to existing work on the low-frequency oscillations of rotating relativistic stars (see \\citet{lock} for a discussion and references to the relevant literature). ",
        "conclusions": "We have shown how the so-called traditional approximation can be generalised to general relativity. This leads to a set of equations that ought to prove valuable for future studies of waves in the oceans of rotating neutron stars. The derivation of these equations required several assumptions in addition to those made in the Newtonian case. First of all, we have to make use of the Cowling approximation. That is, we neglect all metric perturbations. This may seem a drastic simplification, but it is a reasonable approximation for waves in the relatively low-density neutron star ocean. After all, such waves are not expected to lead to appreciable gravitational-wave emission. Secondly, we have to assume that the frame-dragging $\\tilde{\\omega}$ is constant, which restricts us to consider thin shells. Again, this should be a good approximation for neutron star oceans. Finally, and perhaps most interestingly, we have to neglect some terms of post-Newtonian order from the perturbed relativistic Euler equations, an approximation which is valid for low-frequency oscillations. Having made these various assumptions we can separate the eigenfunctions into a radial and an angular part. The angular eigenvalue problem is formally identical to the Newtonian one: it corresponds to solving Laplace's tidal equation, which now depends also on the dragging of inertial frames. It should be emphasised that, since the two problems are identical we can use the results for the Newtonian angular problem also in the general relativistic case.  Our calculation prepares the ground for detailed calculations based on sophisticated models of the neutron star ocean. Given the eigenvalues to Laplace's tidal equation, we need to solve the two radial differential equations (\\ref{final1}) and (\\ref{final2}), together with the relevant boundary conditions. Once one provides a description of the physics of the ocean, the solution of these equations should be straightforward. Here we opted to discuss the effect that general relativity has on various classes of ocean modes in a less quantitative way. We simplified our equations to the conditions that should prevail near the surface of a neutron star. This allowed us to derive a dispersion relation which we used to deduce to what extent the various classes of modes are affected by the gravitational redshift and the rotational frame-dragging. Thus we have shown that the mode frequencies are typically lowered by roughly 20\\%. This is not particularly surprising, but it should be noted that the overall effect is due to different combinations of the frame-dragging and the redshift factors for the different kinds of modes. We also discussed (briefly) the effect that general relativity has on the ocean r-modes.  Our estimates provide a strong argument in favour of using the derived formulae in more detailed (numerical) studies. In particular, we believe it is clear that one must  account for the relativistic effects if the aim is to identify the individual ocean modes present in observed data."
    },
    "0311/astro-ph0311558_arXiv.txt": {
        "abstract": "AA Dor is an eclipsing, close, post common-envelope binary (PCEB) consisting of a sdOB primary star and an unseen secondary with an extraordinary small mass ($M_2 \\approx 0.066\\,\\mathrm{M_\\odot}$) -- formally a brown dwarf. The brown dwarf may have been a former planet which survived a common envelope (CE) phase and has even gained mass.  A recent determination of the components' masses from results of NLTE spectral analysis and subsequent comparison to evolutionary tracks shows a discrepancy to masses derived from radial-velocity and the eclipse curves. Phase-resolved high-resolution and high-SN spectroscopy was carried out in order to investigate on this problem.  We present results of a NLTE spectral analysis of the primary, an analysis of its orbital parameters, and discuss possible evolutionary scenarios. ",
        "introduction": "AA\\,Dor (LB\\,3459, $\\alpha_{2000} = 05^\\mathrm{h} 31^{\\mathrm m} 40.349^\\mathrm{s}$, $\\delta_{2000} = -69^\\mathrm{o} 53' 02.18''$) is an eclipsing, single-lined binary system (Kilkenny et al\\@. 1978). It is a relatively bright ($m_{\\mathrm V} = 11.138$), extremely blue foreground object ($d \\approx 400\\,\\mathrm{pc}$, Rauch 2000) of the Large Magellanic Cloud (Feast et al\\@. 1960). Kilkenny et al\\@. (1979) measured a period of $P = 0.2615\\,\\mathrm{d}$ and found an inclination of $i = 90.00^\\mathrm{o}\\pm 0.02^\\mathrm{o}$. The duration of the primary eclipse is $\\approx 22\\,\\mathrm{min}$ with $\\Delta m_{\\mathrm V} \\approx 0.4\\,\\mathrm{mag}$. There is a reflection effect (cf\\@. Hilditch et al\\@. 1996) with $\\Delta m_{\\mathrm V} \\approx -0.05\\,\\mathrm{mag}$ and a secondary eclipse with $\\Delta m_{\\mathrm V} \\approx 0.06\\,\\mathrm{mag}$.  Photometric investigations and light-curve analyses (Kilkenny et al\\@. 1978, 1979, 1981) as well as evolutionary models (Paczynski 1980) have shown that AA\\,Dor consists of a sdOB primary star and an unseen, nearly degenerate, hydrogen-rich dwarf star ($M_2 \\approx 0.054\\,\\mathrm{M_\\odot}$) of low temperature ($T_\\mathrm{eff} \\approx 3\\,\\mathrm{kK}$). One hemisphere of the secondary is heated by the primary (up to 15 - 20\\,kK) and responsible for the reflection effect.  Regarding the evolutionary times of both components of AA\\,Dor, Paczynski (1980) considered the primary to have only a hydrogen-burning shell and $M_1 \\approx 0.36\\,\\mathrm{M_\\odot}$ to be more likely than to have double-shell burning and $M_1 \\approx 0.54\\,\\mathrm{M_\\odot}$. ",
        "conclusions": "The evolutionary scenario of the PCEB AA\\,Dor is still unclear, although Eggleton \\& Kiseleva-Eggleton (2002) have little doubt to be able to adjust their model (see Section~\\ref{evo}) to obtain good agreement with AA\\,Dor. However, little is known about the secondary. Thus, it would be a valuable approach to hunt for weak spectral signatures of the secondary (cf\\@. Rucinski 2002) in high-resolution and high-SN spectra or to try to measure its radial-velocity curve using the reflected light.  A disagreement remains between the results of spectral analysis and light curve and radial-velocity curve analysis (see Section~\\ref{prob}). The validity of the latter has been verified recently by Hilditch et al\\@. (2003) and thus, it appears likely that the spectral analysis yields --- by an unknown reason so far --- a too-low value of $g$. Since the decrement of the hydrogen Balmer series is very sensitive to variation of $g$ (e.g\\@. Rauch et al\\@. 1998), the analysis of medium-resolution and very high-SN optical spectra, covering the wavelength range down to the Balmer edge, would be a suitable way to attack this problem. \\vspace{-2mm}"
    },
    "0311/astro-ph0311072_arXiv.txt": {
        "abstract": " ",
        "introduction": "Dark matter on 100 \\hmpc\\ scales is accessible to only a couple of techniques. Cosmic Microwave Background (CMB) fluctuations probe dark matter on these scales at $z \\sim 1100$, but the anisotropy spectrum is sensitive to other parameters as well as the dark matter power spectrum.  Peculiar velocities offer a complementary approach, and have the advantage that the dark matter distribution can be compared directly to the distribution of galaxies in the nearby Universe. The bulk flow statistic measures the mean motion of a sample with respect to the CMB, and thus gives an indication of the level of mass density fluctuations on scales larger than the sample size. Recent large-scale peculiar surveys have measured bulk flows which, \\emph{at face value}, appear to be in conflict. The purpose of this paper is to quantify the effect of sparse sampling on the bulk flow statistic and to determine whether recent results are consistent. ",
        "conclusions": "These results suggest substantial contributions to the Local Group's motion from large scales.  A natural question is whether we can identify the structures responsible.  Part of the motion may be due to the Shapley Concentration; the peculiar velocity data favor a substantial mass for this supercluster complex.  If X-ray emitting clusters are used to trace the gravity field, one predicts a very strong contribution to the LG's motion from a distance of $\\sim 150 \\hmpc$ (Kocevski \\emph{et al.}\\,\\cite{kem}). One prominent supercluster at this distance in the right direction on the sky is the Shapley Concentration.  On the other hand, Saunders and collaborators\\cite{sau} have extended the IRAS PSCz survey closer to the Galactic Plane. They find that approximately 50\\% of the Local Group's motion, or $\\sim 300 \\kms$ in the direction of the negative Galactic Y (\\lb{270}{0}) arises from $\\sim 200 \\hmpc$. The structure(s) responsible for this gravity have not been completely identified.  We expect to resolve some of these issues with a new peculiar velocity survey, the NOAO Fundamental Plane Survey (NFPS).  The NFPS is a survey of 93 X-ray selected clusters within 200 \\hmpc.  For each cluster, we obtain 20 -- 70 FP distances per cluster, with a total of 4000 FP cluster galaxies.  The NFPS is therefore 4 times the size of the SMAC and EFAR surveys combined.  We expect random, systematic and sampling errors each $\\lesssim 100 \\kms$ and so expect to resolve the question of the amplitude of the large-scale flow. A description of this survey is given by Smith \\emph{et al.}\\,\\cite{nfps}"
    },
    "0311/astro-ph0311591_arXiv.txt": {
        "abstract": "{ We have observed the continuum emission of the interacting galaxy NGC\\,4631 at $\\lambda\\lambda$\\,870\\,$\\mu$m and 1.23\\,mm using the Heinrich-Hertz-Telescope on Mt. Graham and the IRAM 30-m telescope on Pico Veleta. We have obtained fully sampled maps which cover the optical emission out to a radius of about $7'$ at both wavelengths. For a detailed analysis, we carefully subtracted the line contributions and synchrotron and free-free emission from the data, which added up to 6\\,\\% at 1.23\\,mm and 10\\,\\% at 0.87\\,mm. We combined the flux densities with FIR data to obtain dust spectra and calculate dust temperatures, absorption cross sections, and masses. Assuming a ``standard'' dust model, which consists of two populations of big grains at moderate and warm temperatures, we obtained temperatures of 18\\,K and 50\\,K for the both components. However, such a model suffers from an excess of the radiation at $\\lambda\\,1.23$\\,mm, and the dust absorption cross section seems to be enhanced by a factor 3 compared to previous results and theoretical expectations. At large galactocentric radii, where the galaxy shows disturbances as a result of gravitational interaction, this effect seems to be even stronger. Some possibilities to resolve these problems are discussed. The data could be explained by a very cold dust component at a temperature of 4\\,--\\,6\\,K, an increased abundance of very small grains, or a component of grains with unusual optical properties. We favour the latter possibility, since the first two lead to inconsistencies. ",
        "introduction": "The efficiency and time-scale of star-formation processes in galaxies depend strongly on the amount, distribution, and composition of the interstellar matter. One component of the ISM whose properties are still not well known is interstellar dust.  Since most of the dust in galaxies is cold, with temperatures of $< 20\\,{\\rm K}$, infrared observatories like IRAS are almost blind for the gross amount of dust, and can only observe the warm dust component. The major fraction of the interstellar dust radiates mainly in the sub-mm range, at wavelengths $> 100\\,\\mu{\\rm m}$. This emission was hardly accessible in the past, due to the lack of good sub-mm telescopes at sites with sufficiently good atmospheric conditions. This situation has improved in the last few years: Instruments like the bolometer arrays installed at the 30-m telescope on Pico Veleta or - more recently - SCUBA at the JCMT and the 19-channel array at the Heinrich-Hertz-Telescope have produced important results for the investigation of this cold dust component. Also, the ISO satellite could routinely measure the peak of the cold dust emission from grains with temperatures ranging down to 10~K (e.g.\\ Popescu et al.~\\cite{popescu+02}).  However, dust properties in external galaxies are still not well determined because of gaps in the spectral coverage. Another difficulty in the analysis and interpretation of continuum data is the possible contribution of molecular lines to the total flux measured with the (broad-band) bolometer arrays, especially in the usually used atmospheric windows around $\\lambda\\lambda$\\,870\\,$\\mu$m and 1.23\\,mm which contain the strong CO(3--2) and (2--1) lines. This situation currently improves as more and more objects are mapped even in the higher CO transitions.  This current paper is the first in a series that reports on observations of the thermal emission of cold dust in selected galaxies. NGC\\,4631 is a nearby ($D = 7.5\\,{\\rm Mpc}$, e.g.\\ Golla \\& Wielebinski \\cite{golla+rw94}) galaxy which is close to edge-on. While the position angle of the large-scale (cm-) radio continuum emission is about $86^{\\circ}$, that of the inner disk is closer to $82^{\\circ}$. A value of $84^{\\circ}$, which we assume throughout this paper, seems to be most appropriate for the mm/sub-mm emission out to a radius of a few arcminutes. NGC\\,4631 is embedded in a small gravitationally interacting group with two neighboring galaxies. The dwarf elliptical galaxy NGC\\,4627 is located only $3'$ northwest of the nucleus of NGC\\,4631, and another edge-on spiral, NGC\\,4656, is located about $30'$ to the southeast. This interaction has created several prominent \\ion{H}{i} bridges and spurs (Rand \\cite{rand94}). Since galaxy interactions play an important role in the evolution of galaxy systems and the ISM content in galaxies, the investigation of this object can provide useful information about many processes in the ISM.  NGC\\,4631, which was classified as ``mild starburst'' by Golla \\& Wielebinski (\\cite{golla+rw94}), may be in a late stage of its interaction, where the central star formation (triggered by molecular inflow due to the gravitational forces) has already ceased and an energetic outflow as well as a huge radio halo have been formed (Ekers \\& Sancisi \\cite{ekers+sancisi77}; Rand \\cite{rand00}).  The observations of the dust component of NGC\\,4631 started with the IRAS satellite (Rice et al.\\ \\cite{rice+88}; Young et al.\\ \\cite{young+89}). While the IRAS observations were sensitive mainly to the warm dust present in this galaxy, the $\\lambda$\\,1.3\\,mm map of Braine et al.\\ (\\cite{braine+95}) proved the existence of a significant amount of cold dust in the central area. The inner $2'\\!\\!.5$ were also observed at $\\lambda\\,850\\,\\mu{\\rm m}$ by Alton et al.\\ (\\cite{alton+99}). Neininger \\& Dumke (\\cite{neininger+dumke99}) presented a more extended map at $\\lambda$\\,1.23\\,mm (which was obtained from the same data as shown in this paper) and detected intergalactic cold dust which was pulled out of the disk by the gravitational interaction of NGC\\,4631 with its neighbours. Here we present a more extended map at $\\lambda\\,870\\,\\mu{\\rm m}$, covering the disk out to a radius of $7'$.  Some basic parameters of NGC\\,4631 are compiled in Table \\ref{tab:basic}.  \\begin{table} \\caption[]{Some basic parameters of NGC\\,4631 as obtained from the literature.} \\label{tab:basic} \\begin{tabular}{lll} \\hline Type & Sd & (de Vaucouleurs et al.\\ \\cite{rc3}) \\\\ \\multicolumn{2}{l}{Position:} & (Young et al.\\ \\cite{young+95}) \\\\ RA[2000] & $12^{\\rm h}42^{\\rm m}07^{\\rm s}\\!\\!.65$ & \\\\ Dec[2000] & $32^{\\circ}32'27'\\!\\!.9$ & \\\\ Distance & 7.5\\,Mpc & (Golla \\& Wielebinski \\cite{golla+rw94}) \\\\ Pos.\\ Angle & $84^{\\circ}$ & \\\\ Inclination & $86^{\\circ}$ & \\\\ $M (\\ion{H}{i})$ & $5 \\times 10^9\\,M_{\\odot}$ & (Rand \\cite{rand94}) \\\\ $M ({\\rm H}_2)$ & $1.2 \\times 10^9\\,M_{\\odot}$ & (Golla \\& Wielebinski \\cite{golla+rw94}) \\\\ \\multicolumn{2}{l}{IRAS flux densities:} & (Young et al.\\ \\cite{young+89}) \\\\ $12\\,\\mu{\\rm m}$ & $  5.1 \\pm  1.1$ & \\\\ $25\\,\\mu{\\rm m}$ & $  8.8 \\pm  1.8$ & \\\\ $60\\,\\mu{\\rm m}$ & $ 90.7 \\pm 18.2$ & \\\\ $100\\,\\mu{\\rm m}$& $170.4 \\pm 34.1$ & \\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "\\label{section:summary}  We have observed the nearby interacting galaxy NGC\\,4631 in the radio continuum emission at $\\lambda\\lambda$\\,0.87\\,mm and 1.23\\,mm. The emission is concentrated on the galactic plane, although some halo emission could also be detected, especially at 1.23\\,mm where intergalactic cold dust seems to follow the previously detected \\ion{H}{i} spurs. The two maxima visible in both maps, located in the plane symmetrically around the nucleus at a radius of about $30''$, resemble a ring-like structure, similar to the distribution of the CO molecular line emission. In the outer disk, where no CO is detected, the mm/sub-mm emission follows the \\ion{H}{i} distribution, similar to other edge-on galaxies.  After subtracting non-dust contributions (line, synchrotron, and free-free emission) from the bolometer flux, the remaining flux densities of $S_{\\rm 0.87mm} = 3.41 \\pm 0.58$\\,mJy and $S_{\\rm 1.23mm} = 2.04 \\pm 0.23$\\,mJy can be attributed to the thermal emission of dust. In order to estimate the temperature of this dust, we fitted a two-component modified Planck spectrum to the FIR-to-mm data and found a temperature of 50\\,K for the warm component and of 18\\,K for the cold component, which is responsible for most of the emission at $\\lambda > 200\\,\\mu$m.  This two-component model suffers from the fact that the observed dust spectrum is too flat in the sub-mm/mm range. Furthermore it leads to absorption cross sections too high compared with theoretical expectations or previous results on other galaxies, especially in the outer part of the disk which are disturbed because of gravitational interaction.  We suggest several possibilites to resolve these inconsistencies, the most likely being grains with unusual optical properties which can account for the measured mm-excess as well as the high absorption cross sections.  These observations have shown that results of mm and sub-mm observations have to be combined to investigate the cold dust component in external galaxies."
    },
    "0311/astro-ph0311134_arXiv.txt": {
        "abstract": "Collapsars may be a source for the ``long'' Gamma Ray Bursts (GRBs) in the BATSE catalog. Collapsars may radiate gamma rays anisotropically by beamed jet emission close to the observer's line of sight. These jets must penetrate the radiation-dominated medium of their collapsar progenitor and break through its atmosphere in order to produce a GRB. We present a study of previously published, axisymmetric, collapsar jet simulations. Here we use the linearized, relativistic fluid equations to find the Kelvin-Helmholtz modes that are triggered by recollimation shocks within the jet.  The modes will grow as they propagate with the jet. These are of interest since the light curves of GRBs show evidence of a variable flow from the GRB engine. We also evaluate effects of grid scaling in the numerical simulation and show that short wavelength modes are suppressed by grid scaling before the jet breaks out of the Helium shell. ",
        "introduction": "\\label{sec:1}  The close association between GRB030329 and SN2003dh proves that a subset of ``long'' GRBs is associated with core collapse supernovae [1,2]. The light curve of GRB030329 steepens over time. Price et al. [3] interpret this as a broadening of a jet-like outf\\mbox{}low from the GRB source. Multiwavelength observations by Uemura et al. [4] along with those from GRB021004 [5] strengthen the case for a variable jet outf\\mbox{}low from the central engine.\\\\  These observations support the prompt (Type I) collapsar model of GRBs [6,7]. In this model, a progenitor star with mass $>$ 30 $\\mathrm{M}_\\odot$\\space f\\mbox{}ails to eject its envelope after the core collapses into a black hole. Accretion of free f\\mbox{}alling material may drive a relativistic jet of gas that may break through the stellar envelope. At issue is whether hydrodynamic instabilities within the jet will grow so that the radial profile of jet velocity becomes time and space dependent. Such information could of\\mbox{}fer insight into the production of gamma rays as well as prompt optical afterglows of GRBs.\\\\ ",
        "conclusions": ""
    },
    "0311/astro-ph0311187_arXiv.txt": {
        "abstract": "We present a detailed abundance analysis for 21 elements based on high dispersion, high spectral resolution Keck spectra for four members of the outer halo ``young'' Galactic globular cluster Pal 12. All four stars show identical abundance distributions with no credible indication of any star-to-star scatter. However, the abundance ratios of the Pal 12 stars are very peculiar.  There is no detected enhancement of the $\\alpha$-elements; the mean of [Si/Fe], [Ca/Fe] and [Ti/Fe] is $-0.07\\pm0.05$ dex, O/Fe is also Solar, while Na is very deficient.  The distribution among the heavy elements shows anomalies as well. These are inconsistent with those of almost all Galactic globular clusters or of field stars in the Galaxy.  The peculiarities shown by the Pal 12 stars are, however, in good general agreement with the trends established by Smecker-Hane \\& McWilliam and by Bonifacio \\etal\\ for stars in the Sgr dSph galaxy evaluated at the [Fe/H] of Pal 12.  This reinforces earlier suggestions that Pal 12 originally was a cluster in the Sgr dSph galaxy which during the process of accretion of this  galaxy by our own was tidally stripped from the Sgr galaxy to become part of the extended Sgr stream. ",
        "introduction": "Pal 12 is a sparse globular cluster (henceforth GC) located in the outer halo of the Milky Way.  The CMD studies of \\cite{gratton88} and of \\cite{stetson89} suggested that Pal 12 is probably somewhat younger than the vast majority of Galactic GCs. However, until there was a measurement of the metallicity of this GC, its age could not be determined robustly due to degeneracies in the CMD between age and metallicity.  \\cite{dacosta91} provided the crucial datum; they found that that Pal 12, in spite of its large galactocentric distance, is quite metal rich, obtaining [Fe/H] = $-0.6$ dex\\footnote{The standard nomenclature is adopted; the abundance of element $X$ is given by $\\epsilon(X) = N(X)/N(H)$ on a scale where $N(H) = 10^{12}$ H atoms.  Then [X/H] = log$_{10}$[N(X)/N(H)] $-$ log$_{10}$[N(X)/N(H)]\\subsun, and similarly for [X/Fe].} from low resolution spectroscopy.  More recently, \\cite{brown97} carried out a high resolution spectroscopic study of two stars in Pal 12, which yielded [Fe/H] = $-1.0$ dex.  The combination of deep photometry and an estimate of [Fe/H] led to the verification that Pal 12 is indeed a young cluster. \\cite{rosenberg98}, as part of their recent study of the age dispersion within the Galactic GC system \\citep{rosenberg99}, suggest an age for Pal 12 of roughly 70\\% that of the majority of the halo GCs.  Assuming the latter group to be 12 Gyr old, they then infer an age of 8.4 Gyr for Pal 12, consistent with that of the earlier CMD studies of \\cite{gratton88} and \\cite{stetson89}.  After the discovery of the Sgr dSph galaxy by \\cite{ibata94}, several groups, including \\cite{layden00}, noted that four Galactic GCs (M54, Arp 2, Terzan 7 and Terzan 8), based on their positions on the sky, appear to form a stream extending from the Sgr dSph, and suggested that these GCs had been tidally stripped away from the Sgr galaxy. The GC M54 was postulated to be the original nucleus of the Sgr galaxy.  \\cite{irwin99} was the first to suggest that the GC Pal 12 had also been tidally captured from the Sgr dSph galaxy by our Galaxy. \\cite{dinescu00} measured the proper motion of this GC and calculated its orbit to find that Pal 12's tidal capture from the Sgr dSph took place about 1.7 Gyr ago.  \\cite{ibata01} demonstrated the existence of an extended tidal stream of debris from the Sgr dSph galaxy in the Galactic halo, which has recently been detected in the 2MASS  \\citep{majewski03} and in the SDSS \\citep{ivezic03} databases. Deep optical imaging over wide fields around Pal 12 by \\cite{martinez02} and by \\cite{bellazzini03} show that this GC is embedded in the extended debris stream of stars torn from the Sgr dSph.   The only existing high dispersion spectroscopic study of Pal 12 \\citep*[that of][]{brown97} was hitting the limits of what was observationally possible with a 4-m telescope.  They analyzed the spectra of only the two brightest probable members. Their primary result was that Pal 12 did not appear to show the enhancement of the $\\alpha$-process elements seen in almost all GC stars.   Although this is extremely interesting, the accuracy of their analysis was limited by the quality of their spectra and some key elements were not included.  The purpose of the present paper is to provide a firmer foundation for their results, to extend them as possible, and to compare the properties of the Pal 12 stars with the recently published abundance distributions for stars in the Sgr dSph galaxy of \\cite{bonifacio00}, \\cite{smecker02}, \\cite{mcwilliam03} and of \\cite{bonifacio03}. ",
        "conclusions": "We present a detailed abundance analysis for 21 elements based on high dispersion, high spectral resolution Keck/HIRES spectra for four members of the outer halo ``young'' Galactic GC Pal 12 which has an age of $\\sim$8 Gyr, $\\sim30\\%$ younger than almost all GCs \\citep{rosenberg98}, and is known to have a rather high metallicity for a GC in the outer halo \\citep{brown97}.  Since the discovery of the Sgr dSph galaxy, the Galactic GCs M54, Arp 2, Terzan 7 and Terzan 8 have been believed to be associated with the this  galaxy, which is currently being accreted by the Milky Way.  \\cite{irwin99} recently suggested that Pal 12 is also a Galactic GC which has been tidally stripped from the Sgr galaxy.  All four stars in our sample in Pal 12 show identical abundance distributions with no credible indication of any star-to-star scatter. However, the abundance ratios of the Pal 12 stars are very peculiar.  With $[{\\rm{Fe/H}}] \\sim -0.8$ dex, there is no detected enhancement of the $\\alpha$-elements; the mean of [Si/Fe], [Ca/Fe] and [Ti/Fe] is $-0.07\\pm0.05$ dex, O/Fe is also Solar, while Na is very deficient.  The distribution among the heavy  elements shows anomalies as well, with a shift from the first $s$-process peak (Y and Zr) to the second peak (Ba and La). (Eu is highly enhanced, as is typical in Galactic GCs of similar metallicity.) We show that these abundance peculiarities are not seen among other stellar populations in the Galaxy, including almost all Galactic GCs and field stars in the disk, halo and bulge of the Galaxy.  The abundance anomalies shown by the Pal 12 stars are, however, in good general agreement with the trends established by \\cite{smecker02} and \\cite{bonifacio00} for stars in the Sgr dSph galaxy, when evaluated at the [Fe/H] of Pal 12.  It is interesting to note that the trend of low/no $\\alpha$-enhancement has been found in all high metallicity stars in all the dSph satellites of the Galaxy studied to date. The abundance peculiarities exhibited by Pal 12 can in general be explained by initiating star formation in a metal poor system, and having an extended period of star formation lasting at least a few Gyr.   Our abundance analysis of a sample of four stars in Pal 12 thus reinforces earlier suggestions that this GC originally was a cluster in the Sgr dSph galaxy which during the process of accretion of this  galaxy by our own was tidally stripped from the Sgr galaxy to become part of the extended Sgr stream. Further searches for such anomalies are perhaps best concentrated among metal rich outer halo Galactic GCs; in the metal poor ones the duration of the major epoch of star formation may have been too brief for significant differences to develop in star formation history and hence in abundance distributions among the GCs."
    },
    "0311/astro-ph0311522_arXiv.txt": {
        "abstract": "2MASS has provided a three-dimensional map of the $>360^{\\circ}$, wrapped tidal tails of the Sagittarius (Sgr) dwarf spheroidal galaxy, as traced by M giant stars.  With the inclusion of radial velocity data for stars along these tails, strong constraints exist for dynamical models of the Milky Way-Sgr interaction.  N-body simulations of Sgr disruption with model parameters spanning a range of initial conditions (e.g., Sgr mass and orbit, Galactic rotation curve, halo flattening) are used to find parameterizations that match almost every extant observational constraint of the Sgr system. We discuss the implications of the Sgr data and models for the orbit, mass and $M/L$ of the Sgr bound core as well as the strength, flattening, and lumpiness of the Milky Way potential. ",
        "introduction": "The relatively nearby Sagittarius (Sgr) dwarf galaxy offers the opportunity to explore in exquisite detail the interaction of a satellite with its parent galaxy.  Moreover, the extensive Sgr tidal tail system gives sensitive leverage on the properties of the Galactic potential, much as polar ring galaxies have been exploited to determine the properties of extragalactic systems (e.g., Sparke 2002).  Since Sgr's discovery by Ibata, Gilmore \\& Irwin (1994), this archetype of a dwarf galaxy merger has remained difficult to study because its core (centered at $[l,b]=[6,-14]^{\\circ}$), lies obscured by foreground dust and stars of the Galactic disk and bulge.  However, the Two Micron All Sky Survey (2MASS) has opened a new window on Sgr because of the reduced extinction by dust in the near infrared (NIR) and because Sgr contains a significant population of M giant stars, which are bright in the NIR.  To improve the contrast of Sgr, these M giants can be separated from foreground M dwarfs using the combination of 2MASS $JHK_s$ photometry because of differential, surface gravity-sensitive opacity effects near 1.6$\\mu$ (Bessell \\& Brett 1988). Majewski et al. (2003a; MSWO hereafter) used this method to make all-sky maps that reveal Sgr to be the primary source ($>80\\%$) of M giants in the high halo ($|Z_{GC}|>20$ kpc), excluding those bound in the LMC and SMC (which actually show no M giant tidal tails themselves).  The Sgr M giants lie along a great circle tipped only 13.5$^{\\circ}$ from a truly polar ring.  A three dimensional analysis is made possible by converting $K_s$ magnitudes into photometric parallaxes using the red giant branch relation defined by the Sgr core (for which we adopt $[m-M]=16.90$). The resulting three dimensional distribution reveals a flattened, planar alignment with a $<2$ kpc RMS spread, and clear trailing (predominantly south of the Galactic Plane) and leading (predominantly north of the Galactic Plane) tidal arms. Even without benefit of modeling, the debris arms show clearly that Sgr is orbiting within a non-precessing plane in an elliptical (approximately 3:1 or 4:1 apo:peri-Galacticon) orbit of $\\sim$14 kpc peri-Galacticon. Together the tidal arms wrap at least 360$^{\\circ}$ around the sky, and overlap at even greater length.  To constrain the dynamics of the Sgr-Milky Way interaction further, we (e.g., Majewski et al. 2003b, Paper II hereafter) have been collecting radial velocities of Sgr M giants with the Swope 1-m, KPNO 2.1-m, Yalo 1.5-m and Bok 2.3-m telescopes. As of July 2003, our sample of M giant velocities, distributed around the Sgr plane, is approaching 900 M giants.  The $\\sim 5$ km s$^{-1}$ precision velocities provide critical input to Sgr models, and help compensate for remaining vagaries in the M giant spatial distribution due to contamination by disk/bulge M giants and random (and perhaps systematic) errors in the distance scale. ",
        "conclusions": ""
    },
    "0311/astro-ph0311008_arXiv.txt": {
        "abstract": "We investigate the constraints on models of supermassive black hole (SMBH) and quasar formation obtainable from two recent observational developments: the discovery of luminous quasars at $z\\sim6$, and estimates of the local mass density of SMBHs. If $\\sim90$ per cent of this mass was accreted at redshifts $z\\la 3$, as suggested by the observed quasar luminosity functions, these joint constraints pose a challenge for models, which must account for the observed luminous quasar population at $z\\sim 6$ within a very limited `mass budget'. We investigate a class of models based within the hierarchical structure formation scenario, in which major mergers lead to black hole formation and fuelling, and the resulting quasars shine at their Eddington-limited rate until their fuel is exhausted. We show that the simplest such model, in which a constant fraction of the gas within the halo is accreted in each major merger, cannot satisfy both constraints simultaneously.  When this model is normalized to reproduce the number density of luminous quasars at $z\\sim6$, the mass budget is grossly exceeded due to an overabundance of lower mass SMBHs. We explore a range of modifications to the simple model designed to overcome this problem. We show that both constraints can be satisfied if the gas accretion fraction scales as a function of the halo virial velocity. Similar scalings have been proposed in order to reproduce the local $M_{\\bullet}-\\sigma$ relation. Successful models can also be constructed by restricting the formation of seed black holes to redshifts above $z_{\\rm crit} \\sim 11.5$ or to haloes above a velocity threshold $v_{\\rm crit} \\sim 55 \\; {\\rm km} \\, {\\rm s}^{-1}$, or assuming that only a fraction of major mergers result in formation of a seed SMBH.  We also briefly discuss the issue of trying to assume a `universal $M_{\\bullet}-\\sigma$ relation' within the framework of simple Press--Schechter models, and further show that a fixed universal relation between SMBH mass and host halo mass is unlikely. ",
        "introduction": "In recent times, a growing wealth of observations has revealed a strong correlation between the mass of supermassive black holes (SMBHs) and the luminosity of the stellar bulges in their host galaxies, and an even stronger correlation between SMBH mass and the bulge velocity dispersion (Magorrian et al.\\ 1998; van der Marel 1999; Gebhardt et al.\\ 2000; Ferrarese \\& Merritt 2000; Tremaine et al.\\ 2002). This has in turn produced increasingly accurate estimates of the overall local SMBH mass density.  If these SMBHs are, as now widely believed, the relics of earlier quasar activity (e.g.\\ Rees 1984 and references therein) then these estimates can place strong constraints on the nature and evolution of the quasar population as they effectively limit the `mass budget' available to account for quasar activity.  The last few years have also seen dramatic progress in the numbers of observed quasars as well as in the range of luminosities and redshifts probed. The 2dF Quasar redshift survey (2QZ; Boyle et al.\\ 2000) and the earlier Large Bright Quasar Survey (LBQS; Hewett, Foltz \\& Chaffee 1995) have identified over 23000 quasars at redshifts $0.3 \\la z \\la 3.0$. The Sloan Digital Sky Survey (SDSS; Gunn et al.\\ 1998) has observed over 650 confirmed quasars in the redshift range $3.0 \\la z < 5.7$ (Schneider et al.\\ 2003) and a small number at $z\\ga 5.7$ which have been sufficient to constrain for the first time the density of bright quasars at $z\\sim6$ (Fan et al.\\ 2001a, 2003). These objects must host SMBHs with masses on the order of several $10^9\\;{\\rm M}_{\\sun}$ (Fan et al.\\ 2001a, 2003; Willott, McLure \\& Jarvis 2003). Moreover, evidence has been found even in these highest redshift objects for molecular gas (Bertoldi et al.\\ 2003a; Walter et al.\\ 2003), dust (Bertoldi et al.\\ 2003b), and metals (Freudling, Corbin \\& Korista 2003), underlining the fact that even at these early epochs quasars are associated with galaxies developed enough to have experienced significant star formation.  In addition to the optical surveys already mentioned, observations of quasars and other active galactic nuclei (AGN) in other wavebands have also been steadily improving.  This includes substantial work done determining the soft X-ray quasar/AGN luminosity function (Miyaji, Hasinger \\& Schmidt 2000, 2001) and also several imaging surveys of AGN selected in the hard X-ray band (e.g.\\ Mushotzky et al.\\ 2000; Giacconi et al.\\ 2001; Hasinger et al.\\ 2001; Alexander et al.\\ 2003; Barger et al.\\ 2003) where samples are unbiased by any line-of-sight absorption which may be present.  Thus through multiwaveband observations progress is being made in piecing together the history of AGN and their SMBH.  It has long been established that quasars, and indeed all forms of AGN, are the result of gaseous matter accreting on to a SMBH at the centre of a galaxy (Zel'dovich \\& Novikov 1964; Salpeter 1964; Lynden-Bell 1969; Bardeen 1970) and hence must somehow be linked to galaxy formation.  Since that time a great number of models have been published exploring quasar formation and evolution within the context of the modern Cold Dark Matter (CDM) paradigm of structure formation (including e.g.\\ Efstathiou \\& Rees 1988; Haehnelt \\& Rees 1993; Haehnelt, Natarajan \\& Rees 1998; Cattaneo, Haehnelt \\& Rees 1999; Kauffmann \\& Haehnelt 2000; Cavaliere \\& Vittorini 2000; Monaco, Salucci \\& Danese 2000; Nulsen \\& Fabian 2000; Haiman \\& Loeb 2001; Wyithe \\& Loeb 2002a, 2003; Hatziminaoglou et al.\\ 2003; Granato et al.\\ 2003; for a brief overview of several of these models see Haehnelt 2003). However, none of the above works have concentrated on successfully reproducing the highest redshift population specifically within the constraints of the limited `mass budget' implied by the estimated local SMBH mass density.  Any potential model of quasar formation finds itself faced with three major unknowns. Firstly: where, how and with what mass do the initial `seed' BHs form? Secondly: what events trigger their subsequent fuelling and growth, how much fuel do they supply and how efficiently is it converted into radiative energy?  Thirdly: how does feedback (from star formation or the AGN activity itself) regulate and possibly even check their growth? These uncertainties translate inevitably to a relatively large number of free parameters in models.  In this Paper we aim to reduce this number by concentrating on the high redshift quasar population, allowing us to concentrate on just the first two issues: formation and fuelling.  In order to investigate very luminous quasars such as those discovered at $z\\sim6$, a very large volume must be simulated so as to capture these extremely rare objects. For this reason, we rely, as in the majority of the work cited above, on the analytic Press--Schechter (Press \\& Schechter 1974) and extended Press--Schechter (Lacey \\& Cole 1993) formalism to predict the merger history of dark matter haloes. We consider a class of models based on the basic assumption that seed black holes are created in major mergers, and that these events trigger an episode of gas accretion in galaxies with a pre-existing SMBH. We further assume that SMBHs accrete and radiate at their Eddington-limited rate until their fuel supply is exhausted, and that a fixed fraction of the accreted mass is converted to radiation.  To rigorously explore this problem would require detailed modelling of the star formation, gas cooling and feedback processes which alter and compete for the intergalactic gas thought to fuel quasars. In this Paper, we avoid these complications by focusing on constraints at very high redshift $z\\sim6$, when cooling times are short and only a small fraction of the baryons is locked up in stars, so these competing processes should be relatively unimportant. As well, we do not attempt to model the complex interconnection of galaxy and BH/quasar formation through feedback. Instead, we consider a suite of models with simple parametrized recipes for SMBH formation and fuelling, which point the way for more sophisticated and physical modelling. As we discuss in the next section, attempts to combine data for observed AGN in X-ray and optical wavebands now appear to be able to account for the observed present day SMBH mass density, and moreover seem to suggest that most of it was assembled at fairly low redshifts $z\\la 3$.  This enables us to place rough limits on the `mass budget' available at higher redshifts. Combining this with the SDSS observations of the quasar number density at $z \\sim 6$ allows us to identify successful models within the broad class we have considered, and to rule out other model variants.  Our results will in some sense reflect a `best case scenario' as inclusion of star formation and feedback would presumably only act to make fuel less plentiful for quasars, and thus strengthen our constraints.  As a result however it is difficult to place strong constraints on the lower redshift behaviour of the models as the characteristic upper peak and subsequent decline of the quasar population number density at low redshifts is likely to be strongly related to fuel depletion due to star formation (see e.g.\\ Kauffmann \\& Haehnelt 2000; Di Matteo et al.\\ 2003).  Nevertheless we can demand that the bright quasar population we track in our models never decreases in number below that of the observed population at lower redshifts. In practice this constitutes a relatively weak constraint.  In Section~\\ref{sec:dens} we detail our arguments limiting the SMBH mass density at high redshift and present the criteria we shall use to judge which of our quasar models are successful. In Section~\\ref{sec:framework} we present the basis for our family of quasar models and the methods of our implementation.  We then present in Section~\\ref{sec:simple} a realization of the simplest member of this family of models.  We show why such a model is not workable and use it to demonstrate the challenges a successful model must face before going on to explore other models within the family in Section~\\ref{sec:main}. In Section~\\ref{sec:others} we briefly compare our work to other recent results in the literature, and discuss the implications of imposing a fixed scaling between black hole mass and halo virial mass or velocity at all redshifts. We summarize and conclude in Section~\\ref{sec:conc}. Throughout this paper we assume a `standard' $\\Lambda$CDM\\ cosmology with matter density $\\Omega_0=0.3$, vacuum density $\\Omega_{\\Lambda}=0.7$, Hubble parameter $H_0=70 \\; \\rm{km}\\, \\rm{s}^{-1}\\, \\rm{Mpc}^{-1}$, and normalization $\\sigma_8=1$, although our conclusions are not significantly altered for a lower value of $\\sigma_8=0.9$ (the predicted high redshift SMBH mass densities would typically change by 5--10 per cent for most of the models considered here, and in no cases would change so much as to alter our conclusions as to the `success' of a model). ",
        "conclusions": "\\label{sec:conc}  We have investigated a basic scenario for the formation and fuelling of SMBHs within hierarchical dark matter merger trees. In this scenario, SMBH formation and fuelling are triggered by major mergers, and once `activated' SMBHs accrete at their Eddington rate until their fuel is consumed.  We have not attempted to model the depletion of fuel arising from star formation and related feedback processes, but instead have concentrated upon meeting the very high redshift constraints on the quasar population as a minimum criterion for success.  To do this we have combined the number density of observed high redshift quasars at $z \\sim 6$ with an estimated limit on the high redshift `SMBH mass budget' inferred from local measurements of the SMBH mass density and the observed quasar luminosity function at $z\\la3$.  In our simplest model, a constant fraction of the gas in the halo falls to the centre after every major merger where it forms a seed SMBH if there is no pre-existing SMBH, or else accretes on to the existing SMBH. We have shown that such a model, when the single free parameter (the accreted fraction of gas $f_m$) is normalized to reproduce the number density of high redshift quasars, badly exceeds the SMBH mass budget at $z=6$ placing too much mass in smaller SMBHs. This is a direct result of the steep slope of the dark matter halo mass function in CDM models, and echoes similar problems with reproducing the galaxy luminosity function is such models (e.g. White \\& Frenk 1991).  We further showed that the seed SMBHs formed in this scenario can tend to unphysically large values in large mass haloes. We found that we could counter this by imposing a reduced efficiency of seed formation, in which only about ten per cent of the gas freed by the major merger goes into forming a SMBH (the rest being lost back to the galaxy), and a maximum seed mass of $10^6 \\; {\\rm M}_{\\sun}$. We also found reducing the seed formation efficiency to much smaller values $\\sim 10^{-4}$ allowed us to bring the total SMBH mass density down to a reasonable level, but at the expense of probably overpredicting the number of lower redshift quasars.  Though within the current framework of our model such a low efficiency would seem perhaps hard to justify physically.  Including now a reduced seed formation efficiency of $f_{\\rm seed}=0.1$ and upper mass limit $M_{\\rm seed,max}=10^6 \\; {\\rm M}_{\\sun}$, we then presented a broad range of models involving simple modifications to the the formation of seed SMBHs or to their growth via accretion.  These modifications were designed to bias SMBH growth toward high mass haloes, reducing the relative overabundance of small SMBHs seen in the original model and thus allowing us to reproduce the observed number density of luminous quasars without exceeding the imposed mass budget. We have shown that several of these recipes do indeed prove successful; our conclusion would not be changed by relaxing the maximum seed mass to $10^7 \\; {\\rm M}_{\\sun}$, indeed most of the successful variants are affected only slightly even by a complete removal of the limit. Our findings for these model variants are summarized below:  \\begin{enumerate}  \\item A model in which the fraction of baryons accreted by the SMBH following a merger, $f_m$, was scaled as $(1+z)^n$ was not successful. We suspect that any non-contrived scaling with redshift will not solve the problem at hand.  \\item Introducing a strong scaling of $f_m$ with halo virial velocity was very successful. We found very good results with $f_m \\propto v_{\\rm vir}^2$ and slightly inferior, but still acceptable, results with the relation suggested by Kauffmann \\& Haehnelt (2000) and given in Eqn.\\ \\ref{eqn:kh}. The latter was advocated by these authors in order to reproduce the local $M_{\\bullet}-\\sigma$-type scaling relations.  \\item Decreasing the number of seed SMBHs able to form, either through a redshift cut-off, a virial velocity cut-off or a purely stochastic effect, also had the desired effect. Of these, we feel that the model with the velocity cut-off was the least preferable as it tends to produce only very large mass seeds.  \\item Allowing accretion to take place only when SMBHs exceed $10^6 \\; {\\rm M}_{\\sun}$ can also allow our constraints to be met, but is perhaps less effective than the other models. Its success is dependent upon the reduced seed formation efficiency (i.e.\\ it requires $f_{\\rm seed} \\leq 0.1$).  \\end{enumerate}  The mechanisms adopted in the successful models identified here may all be associated with physical processes.  For example, if reionization can prevent gas from collapsing in haloes as large as $55 \\; {\\rm km} \\, {\\rm s}^{-1}$, which is not implausible, this would be sufficient to make our models successful, as seen in Model F. As we have noted before, our conclusions will likely also be pertinent to other possible BH formation scenarios.  In particular if the dominant mechanism for forming seed BHs is core collapse of massive $\\sim 200 \\; {\\rm M}_{\\sun}$ Pop III stars, or even of supermassive ($\\sim 10^6 \\; {\\rm M}_{\\sun}$) stars in early metal-poor dwarf galaxies (Bromm \\& Loeb 2003), then it is quite reasonable to expect that this mechanism would shut off at a redshift around $z_{\\rm crit} \\sim 10-12$, as the IGM becomes polluted with metals and proto-stellar gas clouds are more likely to fragment into stars with a `normal' IMF (e.g. Bromm \\& Clarke 2002; Bromm \\& Loeb 2003). This scenario is similar to our successful Model E. The strong scaling of accretion fraction with halo circular velocity seen in Models B and C could arise from stellar or AGN-driven feedback (e.g. Silk \\& Rees 1998). While a stochastic element to seed SMBH formation or fuelling in major mergers, as in Model G, may be motivated if we consider that the efficiency of merger-triggered inflows on larger (kpc) scales seems to depend on many variables such as impact parameter, orbital inclination, and host galaxy morphology (e.g. Mihos \\& Hernquist 1994).  Lastly, we have drawn attention to the difficulties faced when trying to assume a redshift invariant `universal $M_{\\bullet}-\\sigma$ relation' in the context of these simple Press--Schechter based models; and also we have shown that if a relation between SMBH mass and halo virial mass is taken to be redshift invariant, then the inferred total SMBH mass density at $z\\sim6$ is so low that producing enough luminous quasars within this budget will be extremely difficult in any CDM model.  While the proliferation of parameters and modelling uncertainties needed to attempt a full self-consistent treatment of galaxy and quasar formation is somewhat daunting, this work suggests that it is necessary if we are to understand the important population of quasars and AGN and how they are related to galaxies. We hope to use the insights gained here to build a more physically motivated, joint model of quasar and galaxy formation in the near future."
    },
    "0311/hep-ph0311367_arXiv.txt": {
        "abstract": "\\singleandabitspaced ",
        "introduction": "There exists an enormous body of evidence in the favor of cold dark matter. This evidence includes observations of galactic clusters and large scale structure \\cite{structure}, supernovae \\cite{supernovae} and the cosmic microwave background (CMB) anisotropies \\cite{cmb,wmap}. Recently, WMAP has provided the most detailed information on the CMB to date, quoting a total matter density of $\\Omega_{\\rm{m}} h^2 = 0.135^{+0.008}_{-0.009}$ \\cite{wmap}. Furthermore, data from WMAP and other prior experiments indicate a considerably smaller quantity of baryonic matter \\cite{wmap,nonbaryonic}. At the 2-$\\sigma$ confidence level, the density of non-baryonic, and cold dark matter is now known to be $\\Omega_{\\rm{CDM}} h^2 = 0.113^{+0.016}_{-0.018}$ \\cite{wmap}.  Many methods have been proposed to search for evidence of particle dark matter. In addition to accelerator searches, direct and indirect dark matter detection experiments have been performed. Direct dark matter searches attempt to measure the recoil of dark matter particles scattering elastically off of the detector material \\cite{direct}. Indirect dark matter searches have been proposed \\cite{ss} to observe the products of dark matter annihilation including gamma-rays \\cite{indirectgamma,subgamma}, positrons \\cite{positrons}, anti-protons \\cite{antiprotons} and neutrinos \\cite{indirectneutrino}.  Each of these methods has its advantages and disadvantages. Direct detection is most suited for WIMPs with relatively large elastic scattering cross sections with nucleons. Indirect detection using gamma-rays is often studied in the context of observing annihiltions in the galactic center region, which depends dramatically on the dark matter profile used. Alternatively, dark matter substructure could be studied with this method. This would still depend on poorly known dark matter distributions, however \\cite{subgamma}. Measurements of the local anti-proton or positron cosmic-ray spectrum depends critically on the location of, and amount of, dark matter substructure within tens of kiloparsecs and a few kiloparsecs, respectively. Also, positrons and anti-protons do not point towards their sources making an unambiguous separation from backgrounds very difficult.  WIMPs which scatter elastically in the Sun or Earth may become gravitationally bound in these gravitational wells. Over the age of the solar system, they may accumulate in substantial numbers in these objects, greatly enhancing their annihilation rate. Although gamma-rays, positrons and anti-protons produced in these annihilations do not escape the Sun or Earth, neutrinos often can. Also, unlike other indirect dark matter searches this method does not depend strongly on our galaxy's dark matter halo profile or on the distribution of dark matter substructure. In this article, we review this method of indirect dark matter detection, and discuss the prospects for such observations in existing and planned experiments. ",
        "conclusions": ""
    },
    "0311/astro-ph0311604_arXiv.txt": {
        "abstract": "The unresolved transition array (UTA) of iron M-shell ions is a prominent absorption feature in the X-ray spectrum of many active galactic nuclei (AGNs). Modeling photoionized plasmas in attempt to match the observed silicon and oxygen lines fail to predict the level of ionization of iron as inferred by this feature. It is suggested that the discrepancy is due to underestimation of the low-temperature dielectronic recombination rates for iron M-shell ions. Modified ionization balance calculations, based on new (guessed) atomic data, support this idea. The results are shown and compared to the global properties of several observed UTAs. Implications for AGN absorbing gas are discussed including an analysis of the ionization parameter distribution in such sources. The need for real calculations of such atomic data is stressed. ",
        "introduction": "\\chandra\\ and \\xmm\\ spectroscopy of  active galactic nuclei (AGNs) show a rich spectrum of X-ray absorption lines superimposed on a power-law continuum. Common features in such spectra are H-like and He-like lines of the  astronomically abundant elements and inner-shell lines of silicon, sulphur, iron and other elements (e.g. Kaspi et al, 2002, Behar and Netzer 2002). Among these, the n=2--3 inner-shell transitions of \\ion{Fe}{1}-\\ion{Fe}{16} (so called iron M-shell ions) are particularly strong. They have been observed in IRAS\\,13349+2438 (Sako et al. 2001), NGC\\,3783 (Kaspi et al. 2001), NGC\\,5548 (Steenbrugge et al. 2003) and in other sources showing warm absorber (WA) features. The lines are hard to resolve spectroscopically  and form a broad unresolved transition array (UTA) around 16--17\\AA. A detailed investigation of these transitions is given in Behar, Sako \\& Kahn (2001). These authors comment on the early UTA observations, compute synthetic absorption spectra  and provide abbreviated set of atomic parameters for these lines.  The  UTA feature has been discussed in a couple of recent papers analyzing the 900 ks \\chandra\\ spectrum of NGC\\,3783. Krongold et al. (2003) used the feature to argue for a very low ionization absorber in this source while Netzer et al. (2003, hereafter N03)  noted a disagreement between the predicted and observed wavelength of the feature. The equivalent width (EW) and the central wavelength of the UTA are potentially important diagnostics of AGN absorbers and it is important to resolve the uncertainties in the calculation of the feature. This paper provides a new look at recently observed AGN spectra and suggests a major revision in iron dielectronic recombination (DR) rates that can resolve the UTA problem. \\S2  describes the available data and the old calculations. In \\S3  I present new calculations of the spectral distribution of the UTA based on ad hoc increases in the DR rates and \\S4 summarizes the more important new findings of this work. ",
        "conclusions": "Detailed analysis and modeling of  iron UTA features in AGN spectra suggest that the iron level of ionization is inconsistent with the ionization of other elements. It is suggested that the reasons are the inaccurate DR rates used for the iron M-shell ions and new \"assumed\" rates are used to support this claim. Theoretical spectra using the new atomic data give much better fit to the observed spectrum of NGC\\,3783. It is also shown that UTA observations by themselves cannot be used to distinguish between multi-component and single component low ionization AGN absorbers. New measurements and calculations of DR rates for M-shell lines are badly needed to produce more realistic models to the X-ray spectrum of AGN."
    },
    "0311/astro-ph0311432_arXiv.txt": {
        "abstract": "We have analysed a sample of satellite and primary galaxies in the 2dF galaxy redshift survey. In our study we find a strong statistical evidence of the Holmberg effect (that is a tendency for satellites to avoid regions along the line defined by the primary plane) within 500 kpc of projected distance to the primary. This effect is present only when we restrict to objects with radial velocity relative to the primary $|\\Delta v|<160$ km/s which correspond approximately to the mean of the distribution. We explore the dependence of this anisotropy on spectral type ($\\eta$), colours, and luminosities of both primaries and satellites, finding that objects with a low present-day star formation rate present the  most significant effect. ",
        "introduction": "The origin and evolution of structures in the universe can be suitably understood within the hierarchical scenario. In such schemes, present-day systems are the result of the aggregation of several sub-structures, and in particular, satellite galaxies can contribute significantly as building blocks of galaxy formation. However, present-day systems of satellites have evolved since their formation; and therefore the observations of such systems can provide invaluable information on the formation and evolution of galaxies.  There are several works that explore the characteristics of systems of satellites. Holmberg (1969) found that the distribution of galaxy satellites presents a lack of objects within $30^\\circ$ of the primary disk plane ($\\phi<30^\\circ$). He searched for companions of nearby bright disk galaxies (the primaries) within 50 kpc of projected separation ($r_p$). His sample was composed of 218 satellites including both physical systems and interlopers that contaminated his photometrically selected sample. He found 45 objects within $\\phi<30^\\circ$ while 173 had $\\phi>30^\\circ$. If the position angles relative to the primary plane are homogeneously distributed, then the number expected within $\\phi<30^\\circ$ would be $\\sim$ 73. After removing contamination by presumed interlopers (which are expected to be distributed isotropically) he concluded that there were no physical objects with $\\phi<30^\\circ$.  The results remained controversial since Valtonen et al. (1978) found an excess of compact satellites with small angles relatives to the primary disk. Based on this work, Byrd et al. (1987) argued that the anisotropy of the Holmberg sample was mainly due to selection effects for he may have not considered these compact dwarf objects. The authors proposed a dynamical stripping of the gas content of the satellite as a possible explanation for the excess of such objects near the primary plane.  Some years later, the problem was investigated again in a statistical way by Zaritsky et al. (1997) using redshift data. They selected a sample of 115 satellites of 69 nearby primaries, these satellites have a maximum projected separation $r_p<500$ kpc. These authors could not confirm the Holmberg effect in the inner 50 kpc since they had few objects at this distance , but they found a similar anisotropy effect with reliable statistical significance at larger separations, $r_p>250$ kpc.  One simple starting point to approach this problem is to assume an isotropic initial distribution of satellites and that evolution generates the observed anisotropies. This scenario has been explored by Quinn $\\&$ Goodman (1986), who found that anisotropy in the distribution of satellite positions is not efficiently produced by dynamical evolution, even at small separation, $r_p<50$ kpc.  On the other hand, numerical simulations like those carried out by Abadi et al. (2002) indicate that, in a hierarchical scenario, the formation of disks involve substantial accretion of satellites which strongly influence the angular momentum vector and therefore the primary plane. Thus, observed anisotropies in the distribution of the remaining satellites could be a direct consequence of this process. In addition to this work, Pe\\~narrubia et al. (2002) also based on numerical simulations, established that the time interval before the satellite disruption strongly depends on the dark matter halo shape, the initial orbital eccentricities and inclinations, and on the mass fraction of satellites relative to the primary.  Our work aims to search for anisotropies of the angular distribution of satellites relative to the primary plane as well as exploring the possible dependence on satellites properties. Our analysis is based on the study of an unprecedentedly large statistical sample extracted from the 2dF galaxy redshift survey. In section 2 we present the data and selection criteria, in section 3 we show the main results, and in section 4 we summarise our conclusions. ",
        "conclusions": "We have carried out an analysis of the anisotropy of satellite angular positions with respect to the primary plane. The sample was selected in redshift space from the 2dFGRS and is large enough for studying possible dependencies of anisotropies on both primary and satellite properties.  We confirm the existence of the effect first claimed by Holmberg, that is, the preference of the satellites to lie near the pole of the primary. We note however that this effect is present up to relative projected separations $r_p<500$ kpc.  We use an anisotropy parameter A suitable to measure the departure from isotropy in the different subsamples analysed. Nevertheless, this parameter indicates a lower limit to the real anisotropy, given the contamination by false satellites expected to be uniformly distributed with respect to the primary plane.  The deviation from isotropy is detected at least at $\\sim3\\sigma$ level for objects with low values of relative velocity $|\\Delta v|<160$ km/s while the high velocity subsample is consistent with isotropy. Although in principle this effect could be related to different levels of contamination by interlopers, Zaritsky 1992 test provides comparable values of contamination in both high and low relative velocity subsamples.  By inspection to the luminosity of the primaries, we conclude that the anisotropy of the distribution of satellites around faint primaries shows a high statistical significance ($\\simeq3\\sigma$).  The most important signal was detected in the sample of satellites corresponding to passive star forming primaries, where the distribution of $\\phi$ angles departs from isotropy at the $\\simeq6\\sigma$ level; while the opposite behaviour is present in primaries with $\\eta>-1.4$, for which the A parameter results consistent with zero. Moreover, we also notice that the signal is even higher when we raise the eccentricity threshold of the primaries. We also have considered the different spectral types of the satellites, finding the strongest anisotropy for satellites with $\\eta<1.1$ (objects with little present day star formation activity). We notice that primaries with a low value of $\\eta$ in our sample also have satellites with poor star formation activity.  The isolation criteria imposed to our sample lowers the probability of including a significant number of elliptical galaxies. Thus, most primaries in our sample are expected to be of late type morphology. Due to the morphology density relation we expect only 10$\\%$ of our isolated sample of primaries to be ellipticals (Whitmore et al. 1993).  We could in principle interpret the Holmberg effect as being mainly caused by evolutionary processes. Systems dominated by old stellar population objects have formed early, so the physical processes that may have removed  satellites near the primary plane have been present for a large time. By contrast, in younger primary and satellite systems, we would expect a larger fraction of objects that have not been accreted yet on to the primary so that they would show an isotropic distribution of satellites. However, as discussed by Quinn $\\&$ Goodman (1986), it may be difficult to produce anisotropies in the distribution of satellites from pure dynamical evolution, specially at large scale  $r_p<500$ kpc. More likely, we suggest that according to Abadi et al. (2002) results, such anisotropies could be the remnants of the formation of the disks of the primaries, where substantial orbital momentum is transfered into primary rotation during the early stages of galaxy formation. Thus, the anisotropy effects we detect at large projected separations, could be the relics of such mechanism.  We also argue that the significant dependence of the anisotropy effect on the relative velocity between the satellite and the primary could be an important point of disagreement between the results of previous works."
    },
    "0311/astro-ph0311118_arXiv.txt": {
        "abstract": "We analyze samples of nearby clusters taken from the Abell catalog and the X-ray Sample of Bright Clusters (De Grandi et al 1999) including a wide range of X-ray luminosities. Using the usually adopted background subtraction procedures, we find that galaxies in clusters selected by means of their X-ray emission show a flat luminosity function (faint end slope $\\alpha \\simeq -1.1$) consistent with that derived for galaxies in the field and groups.  By contrast, the sample of Abell clusters that do not have an X-ray counterpart shows a galaxy luminosity function with a steep faint end ($\\alpha \\simeq -1.6$).  We investigate the possibility that cD halos could be formed by the disruption of galaxies in rich relaxed clusters that show an apparently flat faint end galaxy luminosity function (Lopez-Cruz et al 1997). We find that clusters dominated by a central cD galaxy (Bautz-Morgan classes I and II) show the same systematic trend: X-ray selected clusters have flatter faint end slopes than those clusters with no detected X-ray emission. Thus, it is likely the X-ray selection and not the cluster domination by central galaxies what correlates with background decontamination estimates of the galaxy luminosity function. Moreover, no significant correlation between X-ray luminosity and the galaxy LF faint end slope is found. These results do not support a scenario where flat faint end slopes are a consequence of cD formation via the disruption of faint galaxies. We argue that the clusters without X-ray emission are strongly affected by projection effects which give rise to spurious faint end slopes estimated using background subtraction procedures (Valotto et al 2001). ",
        "introduction": "Determining the galaxy luminosity function (hereafter LF) down to faint magnitudes in rich clusters has been the subject of many studies in the last years ({\\it e.g.} Sandage, Binggeli \\& Tammann 1985; Driver et al. 1994, De Propris et al. 1995; Lobo et al.  1997; Lumsden et al. 1997; Lopez-Cruz et al. 1997; Valotto et el. 1997; Wilson et al. 1997; Smith et al. 1997; Trentham 1998a; Trentham 1998b; Driver et al.  1998; Garilli et al. 1999).  Most of these authors show that galaxy clusters are dominated by a large population of high surface brightness dwarf galaxies, corresponding to a steep faint end slope of the luminosity function, $\\alpha$ in the range -1.4 to -2.0.  The origin and evolution of this faint galaxy population has important consequences for our current understanding of galaxy formation and evolution in dense environments.  For instance, Lopez-Cruz et al (1997) have suggested that the galaxy LF is flat in dynamically evolved clusters characterized by the presence of a dominant cD galaxy, high richness, symmetrical single-peaked X-ray emission, and high gas mass. On the other hand, steep faint-end slopes ($-2.0 \\leq \\alpha \\leq -1.4$) are detected in poorer clusters. It is worth noticing the fact that the galaxy luminosity function of groups selected in redshift space is flat at faint magnitudes (Muriel, Valotto \\& Lambas, 1998) which shows the lack of a universal trend of the parameter $\\alpha$ with  system richness. More recently, Martinez et al. (2002), have obtained reliable determinations of galaxy LF in groups obtained from the 2dF galaxy redsfhit survey. Their results are consistent with a nearly universal galaxy LF with little dependence on environment which is consistent with these previous findings. Lopez-Cruz et al. (1997) suggest that cD galaxies are formed from the disruption of many faint galaxies in the cluster cores, thus resulting in a globally flat faint end slope. Indeed, dynamical processes operating in relaxed clusters are, in general, destructive.  Ram pressure stripping (Gunn \\& Gott 1972, Abadi, Moore \\& Bower 1999) and gravitational tides/galaxy harassment ({\\it e.g.} Moore et al. 1996) both tend to fade galaxies by removing gas or stripping stars. These process are most effective for smaller, less bound galaxies, and would cause a flattening of the faint end slope. However, mergers and galaxy interactions, when the cluster environment then was more like a group environment could have contributed in the opposite direction at early evolutionary stages.  Although large deep redshift surveys of clusters are not presently available, Ferguson \\&  Sandage (1991), and Zabludoff \\& Mulchaey (2000) have analyzed cluster fields with redshift information and obtained faint end slopes $\\alpha \\sim -1.25$ to $-1.3$, which are not significantly steeper than field and group measurements in the 2dF galaxy redshift survey, Norberg et al. 2001, Martinez et al. (2002) both consistent with $\\alpha \\sim -1.13$. Most other estimates of galaxy luminosity functions in clusters rely critically on background subtraction and are in general consistent with significantly steeper faint end slopes ($\\alpha \\le -1.5$). The accuracy of this procedure depends on the statistical assumption that galaxy clusters correspond to density enhancements unbiased with respect to the distribution of foreground or background galaxies.  However, significant projection effects are found in Abell clusters (e.g. Lucey 1983, Sutherland 1988, Frenk et al 1990) that can systematically bias the observed correlation function and the mass function of these systems.  In fact, van Haarlem et al (1998) claim that one third of Abell clusters are not real physically bound systems but simply projections of galaxies and groups along the line of sight.  Valotto, Moore and Lambas (2001) have analyzed several sources of systematic effects present in observational determinations of the galaxy luminosity function in clusters. They used mock catalogues derived from numerical simulations of a hierarchical universe to identify clusters of galaxies in two dimensions in a similar fashion to Abell 1958 and Abell et al. 1989.  Applying standard background subtraction procedures to these data gave rise to artificially steep faint end slopes since many of the clusters do not have significant counterparts in physical space. These projection effects result almost entirely from the large scale structure behind the cluster, a result that was also concluded by Adami et al (2000) from measuring $\\sim 100$ redshifts for faint galaxies thought to lie in the Coma cluster. Color information (eg. b-r) is useful to improve the signal to noise in the process of background decontamination by efficiently removing red background galaxies at $z\\geq0.5$. However, very unlikely contaminating structures at significantly lower redshifts would be eliminated by the use of colors. See for instance Adami et al. (2000) where a significant number of background galaxies in the field of Coma Clusters are at $0.02 < z < 0.3$.  The X-ray emission of the hot intracluster gas provides a confirmation of the presence of a bound cluster of galaxies.  Thus, estimates of the galaxy luminosity function in clusters by means of background decontamination techniques restricted to an X-ray selected sample may provide a useful insight on the issues previously mentioned. In this paper we explore the nature of the faint galaxy population in clusters obtained by background decontamination techniques in X-ray and optically identified clusters of galaxies. By considering a subsample restricted to clusters dominated by a central cD galaxy we can also explore the disruption hypothesis suggested by Lopez Cruz et al (1997). ",
        "conclusions": ""
    },
    "0311/astro-ph0311268_arXiv.txt": {
        "abstract": "We perform an analytic investigation of the stability of line-driven disk winds, independent of hydrodynamic simulations. Our motive is to determine whether or not line-driven disk winds can account for the wide/broad UV resonance absorption lines seen in cataclysmic variables (CVs) and quasi-stellar objects (QSOs). In both CVs and QSOs observations generally indicate that the absorption arising in the outflowing winds has a steady velocity structure on time scales exceeding years (for CVs) and decades (for QSOs). However, published results from hydrodynamic simulations of line-driven disk winds are mixed, with some researchers claiming that the models are inherently unsteady, while other models produce steady winds. The analytic investigation presented here shows that if the accretion disk is steady, then the line-driven disk wind emanating from it can also be steady. In particular, we show that a gravitational force initially increasing along the wind streamline, which is characteristic of disk winds, does not imply an unsteady wind. The steady nature of line-driven disk winds is consistent with the 1D streamline disk-wind models of Murray and collaborators and the 2.5D time-dependent models of Pereyra and collaborators. This paper emphasizes the underlying physics behind the steady nature of line-driven disk winds using mathematically simple models that mimic the disk environment. ",
        "introduction": "Accretion disks are commonly believed to be present in both cataclysmic variables (CVs) and quasi-stellar objects and active galactic nuclei (QSOs/AGN). In both types of object blue-shifted absorption troughs in UV resonance lines are sometimes present. Therefore, a popular scenario put forth to explain them involves outflowing winds emanating from an accretion disk, with the winds giving rise to absorption troughs in the objects' spectra when they are viewed along preferential sight-lines. However, while this scenario is well motivated, the driving mechanism for the winds is still debatable.  In CVs, the P-Cygni type line profiles are observed only at inclination angles close to the pole (i.e., at $\\lesssim 65^\\circ$, where $0^\\circ$ corresponds to the disk rotation axis) when there are high inferred mass accretion rates \\citep[e.g.,][]{war95}. Given these requirements, the obvious source of the wind material is the disk itself. The terminal velocities of the blue-shifted absorption troughs lie in the range $3,000 - 5,000 \\; {\\rm km \\; s}^{-1}$, which is comparable to escape velocities in the inner disk regions surrounding the white dwarf.  In QSOs, broad absorption lines (BALs) from an outflowing wind are observed in $>10\\%$ of QSOs \\citep[e.g.,][]{hf03}. Outflow velocities up to $20,000 - 30,000 \\; {\\rm km \\; s}^{-1}$ are common. However, it is not clear to what degree the detection of BALs is an orientation-angle effect \\citep[e.g.,][]{elv00}, as is the case for CVs, or if instead BAL QSOs are a distinct QSO type with some special intrinsic or evolutionary properties \\citep[e.g.,][]{bor02}. Nevertheless, if accretion disks power QSOs/AGN, some viewing angle effects are likely to be present. \\citet{tur84} cited evidence for and speculated that the outflowing BAL gas may result from material being radially driven off an inner rotating disk, however the 1D model of \\citet{mur95} was the first serious attempt to explain QSO/AGN observations using an accretion disk wind.  Another property that CVs and QSOs have in common is the existence of persistent velocity structure in their absorption troughs (when present) over significantly long time scales. This steady velocity structure, with changes $<10$ ${\\rm km \\; s}^{-1}$, persists at least up to years and decades, for CVs and QSOs, respectively. However, changes in the depths of the absorption have been observed in both classes of objects. For some relevant observations showing the behaviors see \\citet{fron01} and \\citet{har02} for results on CVs and \\citet{fol87} and \\citet{bar92} for results on QSOs. The behavior of the absorption in Q1303+308 \\citep[][]{fol87} is possibly very illustrative, with more recent data showing that the absorption depth has steadily increased over 20 years while the underlying velocity structure persists (Foltz, private communication).  To explain CVs \\citet{per97a} developed 2.5D dynamical models of line-driven disk winds \\citep[see also][]{per97a,per97b,per00,per03}. They concluded that steady wind solutions do exist and that wind terminal velocities are approximately independent of disk luminosity, similar to line-driven winds in early-type stars. They found that the mass-loss rate increased with disk luminosity and that rotational forces were important, causing velocity streamlines to collide and reduce speed, giving rise to an enhanced density region where the strongest absorption occurs. Although the approximations of single scattering and constant ionization were made, these models are generally consistent with observations and they show a strong dependence on inclination angle.  \\citet{pro98} also developed 2.5D models for CVs \\citep[see also][]{pro99,pro02}. Some of their results, for example the magnitude of the wind velocities, were similar to those obtained by \\cite{per97a}. However, a significant difference between the two groups was that \\citet{pro98,pro99} found unsteady flows characterized by large amplitude fluctuations in velocity and density on very short time scales.  For QSOs/AGN the 1D models of \\citet{mur95} indicate that, with an appropriate x-ray shielding mechanism, a suitable accretion disk wind can be driven off by line radiation pressure. Unlike the case for CVs, the wind streamlines are approximately parallel to the disk at high velocities (i.e., close to inclination angles near $90^\\circ$). The models of \\citet{mur95} are able to account for the observed outflow velocities seen in BAL QSOs, the presence of detached absorption troughs seen in some BAL QSOs, and the approximate fraction of QSOs observed to have BALs. But the 2.5D QSO/AGN line-driven wind models of \\citet{pro00} challenge this result. Similar to their earlier CV disk-wind models, \\citet{pro00} and \\cite{pro02} report finding intrinsically unsteady winds.  Proga and collaborators have suggested that since their unsteady line-driven disk-wind models are inconsistent with observational results, ``that a factor other than line-driving is much more likely to be decisive in powering these outflows'' \\citep[][]{har02} \\footnote{ In line-driven disk wind models the mass loss rate is expected to increase/decrease with increasing/decreasing disk luminosity. The observations presented by \\citet{har02} indicate that observables like the CIV$\\lambda1549$ absorption equivalent width do not scale with disk luminosity. However, this argument should not be used to invalidate line-driven disk wind models since the CIV$\\lambda1549$ absorption equivalent width may not be a direct measure of mass loss rate, e.g., due to ionization changes. } and that ``indeed radiation pressure alone does not suffice to drive the observed hypersonic flow'' \\citep[][]{pro03}. However, the line-driven disk-wind models of Pereyra and collaborators \\citep[see][]{hil02} continue to find steady wind solutions, similar to the earlier CV models of \\citet{per97a} and \\citet{per00}.  Thus, an impasse of sorts has developed with regard to the applicability of line-driven accretion disk wind models to CVs and QSOs/AGN. Clearly the observations indicate that steady winds do exist in these objects. However, while one group believes that steady line-driven winds can be achieved, the other group has advocated either abandoning this approach or adopting one that incorporates additional physics (e.g., magnetic fields) into the problem. Therefore, the objective of this work is to develop a series of semi-analytical models, independent of previous 2.5D dynamical efforts, to test for the existence/nonexistence of steady disk winds. {\\it We find that steady winds can exist} and here we emphasize the physics behind these solutions by employing mathematically simple models designed to mimic a disk environment. In a subsequent paper we will present detailed numerical calculations.  In \\S\\ref{sec_general_equations} we present the notation and general equations used in this paper. We define the nozzle function in \\S\\ref{sec_nozzle} and discuss its relationship to steady wind solutions. In \\S4 we analyze the FSH02 model \\citep{fel02}. This analysis clearly demonstrates that an increase in gravity along a streamline, which is characteristic of disk winds, does not necessarily imply an unsteady wind. In \\S5 we present the standard simple models which demonstrate steady winds can exist. The summary and conclusions are presented in \\S6. ",
        "conclusions": "The objective of this work is to determine, in a manner independent of from results of previous numerically-intensive 2.5D hydrodynamic simulations, whether steady line-driven disk wind solutions exist or not. The motive behind this objective is to, in turn, determine whether line-driven disk winds could potentially account for the wide/broad UV resonance absorption lines seen in CVs and QSOs. In both types of objects, it is observationally inferred that the associated absorption troughs have steady velocity structure.  Our main conclusion is that if the accretion disk is steady, then the  corresponding line-driven disk winds emanating from it can also be steady. We have confirmed this conclusion with more realistic (and mathematically more elaborate) models that we will present in a subsequent paper which implements the exact flux distribution above a standard Shakura-Sunyaev disk. This paper, in particular, has sought to emphasize on the underlying physics behind the steady nature of line-driven disk winds through mathematically simple models that mimic the disk environment.  The local disk wind mass loss rates and local disk wind velocity laws represented by the standard simple models of this work are a consequence of balancing the gravitational forces and the radiation pressure forces. Although gas pressure is present, we find that in our models inclusion of gas pressure gives rise to only minor corrections to the overall results.  The balance between gravitational and radiation pressure forces is represented quantitatively through a nozzle function. The spatial dependence of the nozzle function is a key issue in determining the steady/unsteady nature of supersonic wind solutions. In the case of a steady solution which neglects gas pressure effects, the minimum of the nozzle function determines the corresponding wind mass loss rate and the position of the minimum determines the critical point.  In the vicinity of the disk, gas pressure effects only produce minor corrections to the nozzle function, namely the critical point is shifted slightly to the right of the nozzle minimum, where the nozzle function has a positive derivative. However, in cases where the nozzle function is monotonically decreasing (as in the S~model), these minor corrections generate a minimum in the nozzle function at a finite distance from the disk surface, i.e., shifting the critical point from infinity to a finite distance.  The steady nature of line-driven disk winds found in this paper is consistent with the steady nature of the streamline disk wind models of Murray and collaborators \\citep{mur95,mur96,chi96,mur98}, and it is also consistent with the 2.5D time-dependent models of Pereyra and collaborators \\citep{per97a,per00,hil02,per03}."
    },
    "0311/astro-ph0311097_arXiv.txt": {
        "abstract": "{We measure the variance in the distribution of off-plane ($|b|>20^\\circ $) galaxies with $m_K<13.5$ from the 2MASS K-band survey in circles of diameter between 0.344$^\\circ $ and 57.2$^\\circ $. The use of a near-infrared survey makes negligible the contribution of Galactic extinction to these fluctuations. We calculate these variances within the standard $\\Lambda $CDM model assuming that the sources are distributed like halos of the corresponding mass, and it reproduces qualitatively the galaxy counts variance. Therefore, we test that the counts can be basically explained in terms only of the large scale structure. A second result of this paper is a new method to determine the two point correlation function obtained by forcing agreement between model and data. This method does not need the knowledge of the two-point angular correlation function, allows an estimation of the errors (which are low with this method), and can be used even with incomplete surveys.  Using this method we get $\\xi (z=0, r<10\\ h^{-1}{\\rm Mpc})=(29.8\\pm 0.3) (r/h^{-1}{\\rm Mpc})^{-1.79\\pm 0.02}$, which is the first measure of the amplitude of $\\xi $ in the local Universe for the K-band. It is more or less in agreement with those obtained through red optical filters selected samples, but it is larger than the amplitude obtained for blue optical filters selected samples. ",
        "introduction": "The distribution of the number of galaxies or clusters of galaxies in a certain volume $V$ can be used to test the large-scale structure. The calculation of probabilities of different structures in a clustering model are sometimes amenable to analytical techniques: for instance, for voids (Betancort-Rijo 1990), density of Abel clusters (Betancort-Rijo 1995), number of galaxies in a randomly placed sphere and cluster density profiles (Betancort-Rijo \\& L\\'opez-Corredoira 1996), etc. The calculation of variance of the number of galaxies within the cones corresponding to the circles in the sky, or another 2-D figure, (i.e., galaxy counts) is another possibility, and even easier to compare with the observations since we do not need any information about the redshifts of the individual galaxies. This is precisely the purpose of this paper: the calculation from a model of the variance of the number of galaxies in the different regions of the sky, and to see how well they fit real data obtained directly from a sky survey. This will allow us to determine some parameters of the two-point correlation function independent of any other method.  One disadvantage of the galaxy counts is that they are affected by Galactic extinction, and it is difficult to separate the fluctuations due to this extinction from the real fluctuation in the galaxies distribution. However, the arrival of near-infrared galaxy surveys such as DeNIS or 2MASS provides us a way to avoid this problem, since the extinction in K-band, the one to be used in this paper, is 10 times lower than in V-band. In V-band, the dispersion of galaxy counts due only to the fluctuations of extinction in scales around 1 degree is around 30\\% [see eq. (\\ref{sigmaext}) and the parameters given in the paragraph before eq. (\\ref{sigmaext})] at galactic latitudes $|b|>20^\\circ$, while in K-band is only $\\sim 3$\\%. Since the fluctuations we measure due to the large-scale structure are 25-80\\% in the explored scales, it is clear that the calculation of the structure parameters are very sensitive to the exact knowledge of the galactic extinction in visible bands; however, in near-infrared, the calculation of the parameters is much less sensitive to the extinction and we can be sure that the measured fluctuations correspond to the real distribution of galaxies. Moreover, in the K-band we have an important advantage: the K-correction is nearly independent of galaxy type (Mannucci et al. 2001), and this will allow us to make the K-corrections without knowing the galaxy type of the sources in the survey.  Maps of galaxy counts have been available for quite a long time and have been used for several purposes. Zwicky (1957, p. 84), for instance, produced them from the available catalogues in visible bands, with a smooth correction of galactic extinction as a function of $b$, in order to explore the intergalactic extinction. Rather than obtaining the Galactic or intergalactic extinction, the purpose of this paper is to show that the count variances are in qualitative good agreement with those obtained assuming that the source correlation function is equal to that obtained in the standard $\\Lambda $CDM model. Cumulative galaxy counts (in optical) versus the limiting magnitude were used to constrain the departure from homogeneity---fractal distribution---at large scales (Sandage et al. 1972). A $\\log N(m)\\propto 0.6 m$ points out total homogeneity and this is more or less observed in near infrared counts too: with 2MASS data for low redshift (Schneider et al. 1998, fig. 5) or other surveys (e.x., K\\\"ummel \\& Wagner 2000). Here, we will not explore further the homogeneity at very large scales but the clustering at small scales for the local Universe ($z<0.4$). The use of galaxy counts to derive the two-point angular correlation function was already used by Porciani \\& Giavalisco (2002); however, we will recover directly (under some assumptions) the two-point correlation function of the sources, and, therefore, the relative biasing with respect to the mentioned halos. We may accurately determine it, by considering a slightly modified power law (the correlation function is a power law in the most relevant region) and determining the value of the amplitude and the exponent which lead to the best agreement.  The paper is divided as follows: \\S \\ref{.galcounts} describes the observational data and the way to measure the variance in the distribution of galaxies; \\S \\ref{.model} explains how to calculate this variance within a model for the large scale structure; and \\S \\ref{.comparison} makes the comparison between data and model predictions and derive the parameters of the two-point correlation function necessary to get an agreement between model and data. ",
        "conclusions": ""
    },
    "0311/astro-ph0311574_arXiv.txt": {
        "abstract": "Foreseeing the era of high spatial resolution measurements of the Sunyaev--Zel'dovich effect (SZE) in clusters of galaxies, we present a prototype analysis of this sort combined with Chandra X-ray data. It is applied specifically to RX~J1347--1145 at $z=0.451$, the most X-ray-luminous galaxy cluster known, for which the highest resolution SZE and X-ray images are currently available. We demonstrate that the combined analysis yields a unique probe of complex structures in the intracluster medium, offering determinations of their temperature, density, and line-of-sight extent.  For a subclump in RX~J1347--1145, previously discovered in our SZE map, the temperature inferred after removing the foreground and background components is well in excess of 20 keV, indicating that the cluster has recently undergone a violent merger.  Excluding the region around this subclump, the SZE signals in submillimeter to centimeter bands (350, 150, and 21~GHz) are all consistent with those expected from Chandra X-ray observations.  We further present a temperature deprojection technique based on the SZE and X-ray images, without any knowledge of spatially resolved X-ray spectroscopy. The methodology presented here will be applicable to a statistical sample of clusters available in the future SZE surveys. ",
        "introduction": "In the era of precision cosmology represented by the first-year WMAP data (Bennett et al. 2003; Spergel et al. 2003) among others, the scientific goals of studies on galaxy clusters are gradually shifting from the exploration of cosmology to investigations of cluster physics itself.  For cosmological uses, clusters of galaxies are often assumed to be idealized spherical and homogeneous systems. Recent X-ray observations by Chandra and XMM-Newton, however, have revealed complex structures in a number of clusters, whose origins are still to be clarified (e.g. Markevitch et al. 2000, 2002; Peterson et al. 2003; Fabian et al. 2003).  There has hence been an increasing need for accurate and reliable measurements of clusters at multi-wavelengths to further pursue studies on cluster physics.  The Sunyaev--Zel'dovich effect (SZE), the inverse Compton scattering of the cosmic microwave background by hot electrons, serves as yet another diagnosis of the intracluster medium (Sunyaev, Zel'dovich 1970, 1972; for reviews see Rephaeli 1995; Birkinshaw 1999; Carlstrom et al. 2002). The intensity of the thermal SZE is proportional to the line-of-sight integral of the electron density times the temperature, $\\int n_{\\rm e} T_{\\rm e} dl$, whereas that of the X-ray thermal Bremsstrahlung emission is proportional to $\\int n_{\\rm e}^2 T_{\\rm e}^{1/2} dl$.  The combination of the two thus yields independent measures of the density and temperature distributions, including the departure from spherical symmetry (Zaroubi et al. 1998; Hughes, Birkinshaw 1998; Yoshikawa, Suto 1999; Donahue, et al. 2003).  In addition, it is still not easy to explore the nature of very hot ($> 10$~keV) gas with X-ray facilities alone, as the energy band of X-ray spectrometers with sub-arcminute spatial resolution is currently limited to below 10~keV.  The SZE has an advantage in resolving such hot gas even in distant clusters, given its linear dependence on $T_{\\rm e}$ and its redshift independence (e.g. Rephaeli 1995).  This is of particular importance in depicting violent activities in clusters, such as mergers and associated heating of the gas.  As of today, the spatial resolution of the majority of the SZE imaging observations is limited to arcminute scales, considerably lower than current X-ray observations ($6''$ for XMM-Newton, $0''.5$ for Chandra). This is largely due to the limited number of radio telescopes with sub-arcminute resolution in the range of wavelengths relevant to the SZE measurements.  Controlling the systematics of such large telescopes to the sensitivity level required for the SZE imaging is also challenging.  Interferometers are stable with carefully controlled systematics and have allowed high signal-to-noise ratio (S/N) imaging of the SZE (e.g., Jones et al. 1993; Carlstrom et al. 1996); the effective resolution ranges from $15''$ to a few arcmin, with more S/N at the largest angular scales. A single-dish telescope with an array of detectors working at millimeter or submillimeter wavelengths is also a promising tool of the SZE imaging observations. It is highly plausible that an array with hundreds to thousands of bolometers, such as BOLOCAM (Glenn et al. 2003) and SCUBA-2 (Holland et al. 2003), will achieve an angular resolution of $\\sim 10''$ while retaining a wide field-of-view of $\\sim 10'$.  Given the recent progress and numerous on-going plans of the SZE surveys, it is likely that future cluster data, especially at high redshifts, will be dominated by the SZE images and supplemented by the X-ray and optical follow-up observations. Therefore, we explore uses of the SZE data beyond the standard analysis done so far.  An ideal and unique template for this purpose is provided by RX~J1347--1145 at redshift $z=0.451$, the most X-ray-luminous cluster of galaxies known (Schindler et al. 1995, 1997).  This cluster is an exceptional target of high-resolution SZE observations for two reasons: 1) the central SZE intensity, characterized by the central $y$-parameter of $y_0 \\sim 10^{-3}$, is more than twice as large as any other clusters observed to date, and 2) its compact size, characterized by an angular core radius of $\\theta_{\\rm c} \\sim 7''$, enables mapping within a relatively narrow field-of-view of existing facilities.  A number of measurements have thus been carried out for this cluster (Komatsu et al. 1999, 2001; Pointecouteau et al. 1999, 2001; Reese et al. 2002).  We currently have SZE maps of RX~J1347--1145 at three bands: an SZE {\\it increment} map at 350~GHz with a spatial resolution of $15''$ (Komatsu et al. 1999), and {\\it decrement} maps at 150~GHz and 21~GHz with the resolutions $13''$ and $76''$, respectively (Komatsu et al. 2001).  In particular, our 150~GHz observation has achieved the highest resolution of the SZE images so far and revealed a complex substructure with enhanced decrement signals at $\\sim 20''$ off the X-ray center.  The substructure was not visible in the previous ROSAT X-ray image (Schindler et al. 1997), indicating that it has a very high temperature that is far beyond the energy band of the ROSAT/HRI ($<2$~keV).  This finding was subsequently confirmed by the Chandra observations with an emission-weighted temperature toward the substructure of $18.0^{+2.7}_{-2.3}$~keV, much higher than the ambient temperature $12.7 \\pm 1$~keV (Allen et al. 2002).  Based on their Keck spectroscopy of 47 member galaxies, Cohen and Kneib (2002) suggested that this cluster is undergoing a major merger.  Indeed, this is a good demonstration that the SZE imaging observations provide an independent and powerful diagnosis of cluster physics.  In this paper, we combine our multi-band, high-resolution SZE images and the Chandra X-ray data of RX~J1347--1145 to present a proto-type analysis that will continue to prove useful in the future. More specifically, we take into account radial temperature profile and non-sphericity of the intracluster gas to constrain the gas distribution and the cluster bulk motion.  We also attempt to deproject temperature profile of the cluster on the basis of the SZE and X-ray surface brightness, without a knowledge of spatially resolved X-ray spectroscopy.  Finally, we explore the physical nature of the high-temperature substructure and discuss its implications on the merger history and dynamical activities of the cluster.  Throughout the paper, we assume a standard set of cosmological parameters: $\\Omega_m=0.27$, $\\Omega_\\Lambda=0.73$, and $h=0.71$ (Spergel et al. 2003). In this cosmology, an angular size of 1$''$ corresponds to a physical size of 5.74 kpc at the cluster redshift $z=0.451$. Unless stated otherwise, errors are given in 1-$\\sigma$.   \\begin{table*} \\caption{Summary of the SZE observations and the data used in the analysis. } \\label{tab:obs} \\begin{center} \\begin{tabular}{cccc} \\hline \\\\[-9pt] \\hline \\\\[-6pt] Frequency & 350 GHz& 150 GHz& 21 GHz\\\\[4pt]\\hline \\\\[-6pt] Facility & JCMT/SCUBA & Nobeyama/NOBA& Nobeyama/HEMT \\\\ Beam-size (FWHM)  &$15''$ & $13''$ & $76''$ \\\\ Field-of-view & $2'.8$ diameter  & $1'.9 \\times 1'.9$ & $6'.0 \\times 6'.0$\\\\ Average 1-$\\sigma$ noise & 5.3 mJy~beam$^{-1}$& 1.6 mJy~beam$^{-1}$& 0.9 mJy~beam$^{-1}$\\\\ Radial bins [arcsec] & 10--20, 20--30, 30--40, & 10--20, 20--30, 30--40, & 60--110, 110--140, \\\\ & 40--50, 50--65, 65--85&  40--50, 50--60, 60--75& 140--170, 170--240 \\\\[4pt] \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "We have performed combined analyses of the highest resolution SZE and X-ray images currently available for RX~J1347--1145.  The data at both 150~GHz and 350~GHz show clear signatures of enhanced SZE signals at $\\sim 150$~kpc southeast from the center of this cluster.  We have separately determined the temperature, density, and the line-of-sight extent of this excess component, after removing the foreground and background signals.  The presence of high-temperature gas with $28 \\pm 7$~keV strongly indicates that the cluster has recently experienced a major merger.  Excluding the regions around the southeast enhancement and the radio sources, the SZE signals at 350, 150, and 21~GHz agree well with those expected from the electron density and temperature profiles of Chandra X-ray observations.  Our multi-band SZE data indicate marginal evidence for the line-of-sight elongation of the gas distribution or non-zero peculiar velocity.  The central $y$-parameter extrapolated from radii greater than $60$~kpc is $y_0 \\simeq 1.0 \\times 10^{-3}$.  We further present a temperature deprojection technique using the SZE and X-ray imaging data. The radial temperature profile so obtained between 60 and 350~kpc are all consistent with those based on the X-ray spectral deprojection analysis of Allen et al. (2002).  Our results demonstrate the power of high-resolution SZE mapping observations. In particular, they provide a unique probe of the temperature structure in distant clusters for which spatially resolved X-ray spectroscopy is unavailable.  With its unlimited sensitivity to high temperature, the SZE has an advantage over X-rays in detecting the gas with temperatures well above 10 keV. As demonstrated for RX~J1347--1145, an angular resolution of $\\sim 10''$ is essential in resolving irregular morphology of merger-heated gas in distant clusters. Large detector arrays, such as BOLOCAM on the Large Millimeter Telescope (Glenn et al. 2003) and SCUBA-2 on the James-Clerk-Maxwell Telescope (Holland et al. 2003), will indeed be suitable for such observations. Though performed on a single cluster so far, the methodology presented in this paper will be applicable to a statistical sample of clusters obtained in the future SZE surveys.  \\bigskip  We thank Steve Allen for providing deprojected temperature data and for useful discussions, Naoki Itoh and Satoshi Nozawa for supplying subroutines of their relativistic correction formulae, and Erik Reese, Motokazu Takizawa, Ichi Tanaka and the referee for helpful comments and discussions. NO and KY acknowledge support from Research Fellowships of the Japan Society for the Promotion of Science for Young Scientists. This work is supported in part by the Grant-in-Aid by the Ministry of Education, Culture, Sports, Science and Technology (14740133, TK)."
    },
    "0311/astro-ph0311024_arXiv.txt": {
        "abstract": "We have measured mid-infrared (MIR) fluxes of 3C sources on images taken with ISOCAM onboard the {\\it Infrared Space Observatory (ISO)}. The photometric data were combined with existing photometry at other wavelengths to assemble the spectral energy distribution (SED) for each galaxy from infrared to radio wavelengths. In addition, we used ISOPHOT spectra to compute average MIR spectra for different types of Active Galactic Nuclei (AGNs).  The MIR emission of \\ndust\\ of our sources shows evidence of dust.  We find a clear correlation between the SEDs, emission from polycyclic aromatic hydrocarbons (PAHs), and the type of AGN. Specifically, we find that in broad emission line radio galaxies (BLRGs), the dust emission peaks at $\\lambda\\approx40\\mu$m and PAH bands are weak, whereas emission in comparable narrow line radio galaxies (NLRGs) peaks at a longer wavelength of $\\lambda\\approx100\\mu$m and the PAH bands are much stronger. Although less pronounced, a similar trend is seen when comparing quasars (QSOs) with high-luminosity NLRGs.  We used dust radiative transfer models with a central heating source to describe the SEDs. The difference in the dust emission appears to be the effect of heating by radiation from the central engine.  In type~1 AGNs (BLRGs and QSOs), the broad line region (BLR) is directly visible and therefore hot ($T\\gtrsim300K$) dust dominates the emission. In this region, PAHs are destroyed which explains the weak PAH emission. On the other hand, the BLR is hidden in type~2 AGNs (NLRGs). Their SEDs are therefore dominated by cooler dust including PAHs. ",
        "introduction": "The widely accepted paradigm for the power source of AGNs is a massive black hole. In this picture, accretion of dust and gas powers the AGN. Beamed radiation escapes from the central engine and affects the surrounding gas and dust.  The effect on gas can directly be observed in close-by AGNs through the spatial separation of ionized and neutral gas (e.g. Tadhunter and Tsvetanov 1989, Prieto and Freudling 1993). The impact of this radiation on dust, if present, is expected to be similarly dramatic. The intense radiation field in the immediate vicinity of the central engine prevents dust from forming and evaporates any dust approaching that region. In this region, PAHs cannot exist. At somewhat larger distances from the nucleus, the radiation field heats the dust to just below its evaporation temperature. Dust at such distances will have temperatures of about 1500K. Cold dust ($T\\lesssim50K$) should exist only at distances greater than several hundred pc to a few kpc. At intermediate distances, a region should exist were large quantities of dust are significantly heated by the AGN to temperatures of a few hundred Kelvin and should therefore be detectable at MIR wavelengths.  Unfortunately, detailed mapping of dust temperatures in AGN host galaxies is not feasible at present. Our knowledge of dust in 3C sources is therefore based on low spatial resolution photometry. Observations with the {\\it Infrared Astronomical Satellite (IRAS)} showed that 3C sources which are bright at $\\lambda=60\\mu$m are typically much fainter or not detectable at $\\lambda=12\\mu$m \\citep[][and references therein]{heck}. This signature of cold dust is also seen in the more recent data presented by \\cite{Meise}. \\citet{heck} find that the NLRGs are about a factor of 4 to 5 fainter in the MIR than BLRGs, but are similar to BLRGs at $\\lambda=60\\mu$m. These differences are qualitatively consistent with unification schemes since hotter dust is visible in BLRGs due to their viewing angle.  ISOCAM was more than an order of magnitude more sensitive than {\\sl IRAS} at MIR wavelengths. It was able to detect the hot dust component in 3C sources at much higher redshifts and therefore expand comparison of type~1 and~2 AGNs to a much wider range of luminosities. In this letter, we report first results from a program to investigate the hot dust component in 3C sources based on ISOCAM observations. A detailed discussion of our data reduction and modelling procedures as well as a full presentation of our data are deferred to a later paper \\citep[][hereafter SFKH]{mainpaper}. ",
        "conclusions": "We have detected hot dust in ISOCAM images of 3C galaxies out to redshift z$\\approx$\\zmax.  We found a clear correlation between the type of the AGN and the overall SEDs in the sense that SEDs of BLRG and QSOs are dominated by emission from dust which is about a factor of 2 hotter than the one in comparable NRLGs.  MIR spectra show that PAH emission bands in spectra of type~1 AGNs are significantly weaker than the ones in spectra of type~2 AGNs with the same total dust luminosity. These trends are naturally predicted if dust in AGNs is predominantly heated by hard radiation from a central engine which destroys PAHs close to the nucleus and heats larger dust grains at intermediate distances."
    },
    "0311/astro-ph0311448_arXiv.txt": {
        "abstract": "{ In a recent paper, Edvardsson \\etal (\\cite{edvardsson}) propose a new solution for the spin evolution of the Earth and Mars. Their results  differ significantly with respect to previous studies, as they found a large contribution on the precession of the planet axis from the tidal effects of Phobos and Deimos. In fact, this probably results from the omission by the authors of the torques exerted on the satellites orbits by the planet's equatorial bulge, as otherwise the average torque exerted by the satellites on the planet is null. }  \\begin{document} ",
        "introduction": "In a recent paper, Edvardsson \\etal (\\cite{edvardsson}) propose a new solution for the spin evolution of the Earth and Mars. Considering the absence of a precise evaluation of the errors due to the integrator and the absence of relativity in the model, one could discuss the use of the word 'accurate' in the  title of the paper, but unfortunately  there are some  more important flaws in this paper. In their integration of the spin of Mars, the authors found that the integration of the spin changes in a large amount when Phobos and Deimos are taken into account (see their Fig.12). They also notice that their \" curve without the moons is very similar to the curve given by Bouquillon \\& Souchay (\\cite{bouquillon}) \", who included the moons. In fact,  the  paper of Edvardsson \\etal (\\cite{edvardsson}) is largely in error on this point. ",
        "conclusions": "Thus, contrarily to what is found by Edvardsson \\etal (2002), the averaged torque exerted by Phobos and Deimos on Mars should be null (at first order). The origin of the authors' error can be traced by looking to their figures 1 and 2. The inclination of the satellites with respect to the equator start with a value close to zero at the origin, and then increase up to about 50 degrees and then back to zero with a period of about 6000 and 1500 years  for Phobos and  Deimos, while a back on the enveloppe calculation gives respectively 5400 and 1400 years for the precession of the node of these satellites submitted uniquely to solar perturbation. The issue is then clear. By neglecting the torque exerted on the satellites orbits by the equatorial bulge of Mars, the authors found that instead of remaining close to the equatorial plane of Mars, as demonstrated by Goldreich (1965), the satellites were precessing along the orbital plane. Instead of finding a zero value for the mean precession torque (11), they obtain an effect following Eq. (10) which gives a contribution of about $0.5$\"/year to the precession of mars axis, in agreement with the difference observed by the authors in figure 12. It is surprising that the authors did not wonder how Phobos and Deimos could both be at present with an inclination of less than 2 degrees with respect to the equator if they were precessing along the orbit.  Goldreich indeed realised more that this `would amount to an unbelievable coincidence'."
    },
    "0311/astro-ph0311162_arXiv.txt": {
        "abstract": "In this contribution, I review the current observational evidence for the existence of filaments of Warm-Hot Intergalactic Medium (WHIM). In particular, I first focus on the controversial issue of the identification of the $z\\sim 0$ highly ionized far-ultraviolet (i.e. OVI) and X-ray (i.e. OVII, OVIII and NeIX) absorbers with either a very tenuous and diffuse WHIM filament, or with much denser condensations of material at large distances in the Galactic halo. I then present our recent detection (confidence level $> 3\\sigma$) of the OVII WHIM at $z > 0$ and derive an estimate of the total number of baryons contained in this hard-to-detect phase of the IGM. ",
        "introduction": "Despite recent progress in cosmology, very little is still known about the location of baryons in the local Universe. The extraordinary WMAP ({\\em Wilkinson Microwave Anisotropy Probe}) measurements of the Cosmic Microwave Background (CMB) anisotropies favor a $\\Lambda$-CDM scenario in which the baryon density in the Universe amounts to $\\Omega_b = (4.4 \\pm 0.4) h_{70}^{-2}$ \\% of the total matter-energy density, with a baryon-to-dark-matter ratio of $0.17 \\pm 0.1$ (Bennet et al., 2003). This number agrees well with predictions by the standard 'big bang nucleosynthesis' when combined with light element ratios ($\\Omega_b = (3.9 \\pm 0.5) h_{70}^{-2}$ \\%: Burles \\& Tytler, 1998), and also with the actual number of baryons at redshifts larger than 2 in the 'trees' of the Ly$\\alpha$ Forest: $\\Omega_b \\ge 3.5 h_{70}^{-2}$ \\% (Rauch, 1998; Weinberg et al., 1997). This concordance cosmology represents a major advance. However, the number of baryons detected in the ``virialized Universe'' at $z \\ls 2$ (i.e. stars, neutral hydrogen associated with galaxies, and X-ray emitting gas in clusters) is more than a factor of 2 smaller than the concordance value, $\\Omega_b < 2 h_{70}^{-2}$ \\% (2$\\sigma$; e.g. Fukugita, Hogan, \\& Peebles, 1998). It has also been known since 1959 (Kahn \\& Woltjer, 1959) that, locally, more than $1.5 \\times 10^{12}$ M$_{\\odot}$ are needed to dynamically stabilize our own Local Group. So, one fundamental question arises: where have the baryons been hiding for the last $\\sim 10$ Gyrs of life of the Universe?  \\medskip Hydrodynamical simulations for the formation of structures in the Universe, predict that in the present epoch ($z \\ls 1-2$), half of the normal baryonic matter in the Universe (the ``missing baryons'') is in a tenuous (overdensities $\\delta \\simeq 5-50$, relative to the mean baryon density in the Universe)``warm phase'' (the so called 'Warm Hot Intergalactic Medium': WHIM), shock-heated to temperatures of $10^5-10^7$ K during the continuous process of structure formation (e.g. Hellsten et al., 1998; Cen, \\& Ostriker, 1999; Dav\\'e et al., 2001). These hot filaments are now detectable with current instruments. ",
        "conclusions": ""
    },
    "0311/astro-ph0311176_arXiv.txt": {
        "abstract": "{ The imager on board INTEGRAL (IBIS) presently provides the most detailed sky images ever obtained at energies above 30 keV. The telescope is based on a coded aperture imaging system which allows to obtain sky images in a large field of view (29$^{\\circ}$~$\\times$~29$^{\\circ}$) with an angular resolution of 12$'$. The System Point Spread Function of the telescope and its detailed characteristics are here described along with the specific analysis algorithms used to derive the accurate point-like source locations. The derived location accuracy is studied using the first in-flight calibration data on strong sources for the IBIS/ISGRI system. The dependence of the calibrated location accuracy with the signal to noise ratio of the sources is presented. These preliminary studies demonstrate that the IBIS/ISGRI telescope and the standard scientific analysis software allow source localizations with accuracy at 90$\\%$ confidence level better than 1$'$ for sources with signal to noise ratios $>$ 30 over the whole field of view, in agreement with the expected performances of the instrument. ",
        "introduction": "The IBIS telescope (Imager on Board of the INTEGRAL Satellite) (Ubertini et al. 2003), launched onboard the major ESA gamma-ray space mission INTEGRAL (Winkler et al. 2003) on October 2002, is a hard X-ray / soft $\\gamma$-ray telescope based on a coded aperture imaging system (Goldwurm et al. 2001). The IBIS imaging system and the IBIS scientific data analysis are described in Goldwurm et al. (2003). Here we discuss more in detail the specific image data analysis procedures used to evaluate the point source locations and we present the characteristics of the System Point Spread Function obtained with the implemented analysis. We also provide preliminary results on the source location accuracy obtained from the first in-flight calibration data of ISGRI, the low energy (15-1000 keV) detector of IBIS (\\cite{FL03}). ",
        "conclusions": ""
    },
    "0311/astro-ph0311495_arXiv.txt": {
        "abstract": "Axion-photon mixing has been proposed as an alternative to acceleration as the explanation for supernovae dimming. We point out that the loss of photons due to this mixing will induce a strong asymmetry between the luminosity, $d_L(z)$, and angular-diameter distance, $d_A(z)$, since the latter is unaffected by mixing. In a first search for such an asymmetry we introduce a dimensionless mixing amplitude $\\lambda$ so that $\\lambda=0$ if no photons are lost and $\\lambda\\approx1$ if axion-photon mixing occurs. The best-fit to  SNIa and radio galaxy data is  $\\lambda = -0.3^{+0.6}_{-0.4}$ (95\\% CL), corresponding to an unphysical, negative, mixing length. This same argument limits the attenuation of light from supernovae due to dust. We show that future $d_L$ and $d_A$ data from SNAP and galaxy surveys such as DEEP2 and KAOS will detect or rule out mixing at more than 5$\\sigma$, almost independently of the dark energy dynamics. Finally we discuss the constraints from the near maximal polarisation of the gamma-ray burst (GRB) GRB021206. Since mixing reduces the polarisation of distant sources, future observations of high redshift GRBs may provide orthogonal constraints on axion-photon mixing and related scenarios. ",
        "introduction": "The reciprocity relation is a wonderfully powerful result valid for any metric theory of gravity where photons travel on null geodesics, as long as photon number is conserved (Etherington 1933; Ellis 1971).  It ensures that the luminosity distance, $d_L(z)$ is exactly the same as the angular-diameter distance, $d_A(z)$, up to a factor of $(1 + z)^2$. In this paper we turn the reciprocity relation around and use it to probe alternatives to cosmic acceleration.  The accumulating evidence for recent cosmic acceleration (Barris {\\em et al} 2003; Knop {\\em et al}, 2003) leaves us with the familiar coincidence problem - why do we live at such a special time? An attractive alternative is that the acceleration is a mirage and not a real feature of the dynamics of our Universe.  Although such a non-accelerating cosmology can be made reasonably compatible with cosmic microwave background (CMB) and large scale structure (LSS) data (Blanchard {\\em et al.} 2003), the dilemma then is to explain, without acceleration, the dimming of distant Type Ia supernovae (SNIa) and the observed $\\sim 3\\sigma$ correlation between the CMB and LSS (Boughn \\& Crittenden 2003; Nolta \\etal 2003; Scranton \\etal 2003; Fosalba \\etal 2003). The latter can perhaps be explained by negative spatial curvature, while the dimming of supernovae can be explained by axion-photon mixing (Cs\\'aki {\\em et al} 2002), meaning that the evidence for acceleration is not yet overwhelming.  The basic idea of axion-photon mixing is simple. On average, and on large scales relative to the mixing length, $1/3$ of photons in the visible would be lost through conversion to a light axion state, $a$, in the presence of the cosmic magnetic field $\\overrightarrow{B}$. This proceeds through the axion interaction term $\\frac{a}{4M}\\overrightarrow{E}\\cdot\\overrightarrow{B}$. Cs\\'aki {\\em et al.} (2002) argue that an axion mass scale of $M \\simeq 4\\times 10^{11} GeV$ would provide a good fit to SNIa luminosity data as a function of redshift, (quantified in Erlich \\& Grojean 2001), without the need for cosmic acceleration, while still being (marginally) consistent with other constraints (see Deffayet {\\em et al} 2002; Mortsell {\\em et al.} 2002; Christensson and Fairbairn 2003; Mortsell and Goobar, 2003), especially if non-flat FLRW models are considered.  Intriguingly, axion-photon (AP) mixing, with a similar mass scale, can explain the existence of super-GZK cosmic rays (Cs\\'aki {\\em et al.} 2003) if the primaries are taken to be photons, since they can travel most of the way as axions, then oscillate back into photons before reaching earth. Axion-photon mixing can therefore provide a simultaneous solution to the super-GZK and coincidence problems and is thus worth further detailed study.  We make four points about the axion-photon mixing scenario: \\begin{enumerate} \\item Axion-photon mixing should induce a violation of the reciprocity relation and a fundamental disagreement between the dimensionless coordinate distance, $y(z)$, inferred from the luminosity, $d_L(z)$, and angular-diameter distances, $d_A(z)$. \\item The observed $d_A(z)$ should correspond to a decelerating universe if axion-photon mixing is the source of supernovae dimming. However, current estimates of $d_A(z)$,  favour an accelerating Universe. \\item Future data from SNAP and KAOS will allow for constraints beyond the 5$\\sigma$ level. Tests of number counts from the on-going DEEP2 survey will provide further tight constraints. \\item Mixing leads to depletion of the polarisation levels of extra-galactic sources (Cs\\'aki {\\em et al.} 2002; Mortsell and Goobar 2003). The near maximal polarisation seen in the gamma-ray burster (GRB) GRB021206 (Coburn and Boggs, 2003) suggests that GRB's may provide powerful constraints on the mixing scenario when more data is available. \\end{enumerate} In this paper we will assume a flat FLRW model consisting of dust ($\\Omega_M$) and dark energy $X$ ($\\Omega_X=1-\\Omega_M$) with a constant equation of state $w = p_X/\\rho_X$. ",
        "conclusions": "The dimming of distant supernovae (SNIa) remains the most direct evidence for cosmic acceleration. Nevertheless  alternative explanations exist, such as axion-photon mixing in which roughly one third of all photons from distant SNIa are lost into axion states. We have pointed out that such mixing will not affect the angular-diameter distance $d_A(z)$ and hence will cause a fundamental asymmetry between measurements of the luminosity distance, $d_L(z)$, and $d_A(z)$ that can be searched for.  In a first search for such asymmetry we have undertaken a joint analysis of high-redshift SNIa ($d_L(z)$) and radio galaxy data ($d_A(z)$). The results do not favour the loss of photons and hence disfavour mixing. Future data will improve the limits, and be able to test very generally for an asymmetry between $d_A(z)$ and $d_L(z)$. Number counts versus redshift are a promising test while estimates of $d_L(z)$ from SNAP and $d_A(z)$ from a large 2nd generation galaxy survey such as KAOS will allow axion-photon mixing to be detected or ruled out at more than $5\\sigma$, almost independently of the dynamics of the dark energy, showing the power in constraining non-standard physics implicit in combining $d_L(z)$ and $d_A(z)$ data."
    },
    "0311/astro-ph0311389_arXiv.txt": {
        "abstract": "The Milagro telescope monitors the northern sky for 100 GeV to 100 TeV transient emission through continuous very high energy wide-field observations.  The large effective area and $\\sim$100 GeV energy threshold of Milagro allow it to detect very high energy (VHE) gamma-ray burst emission with much higher sensitivity than previous instruments, and a fluence sensitivity at VHE energies comparable to that of dedicated gamma-ray burst satellites at keV to MeV energies.  Even in the absence of a positive detection, VHE observations can place important constraints on  gamma-ray burst (GRB) progenitor and emission models. We present limits on the VHE flux of 40 s -- 3 h duration transients nearby to earth, as well as sensitivity distributions which have been corrected for gamma-ray absorption by extragalactic background light and cosmological effects.  The sensitivity distributions suggest that the typical intrinsic VHE fluence of GRBs is similar or weaker than the keV -- MeV emission, and we demonstrate how these sensitivity distributions may be used to place observational constraints on the absolute VHE luminosity of gamma-ray bursts for any GRB emission and progenitor model. ",
        "introduction": "Very high energy gamma-ray burst (GRB) observations have the potential to constrain theoretical models of both the prompt and extended phases of GRB emission. Models based on both internal and external shocks have predicted VHE fluence comparable to, or in certain situations stronger than, the keV -- MeV radiation, with durations ranging from shorter than the keV -- MeV burst to extended TeV afterglows \\citep{Dermer:VHEGRB,PillaLoeb:VHEGRB,Bing}.  GRB emission above 100 GeV is particularly sensitive to the Lorentz factor and the photon density of the emitting material --- and thus the distance of the radiating shock from the source --- due to $\\gamma\\gamma \\rightarrow e^+e^-$ absorption in the emission region.  EGRET observed several GRBs at GeV energies and observed no evidence of a high energy rollover in the GRB spectrum \\citep{EGRET}, and recent results by \\citet{Gonzalez} indicate that the spectrum of some GRBs contains a very hard, luminous, long duration component.  At energies above a few 100 GeV several observation attempts have been made, including: coincidence observations between the Tibet air-shower array and BATSE \\citep{TibetGRB}, rapid follow-up observations by the Whipple Air-Cherenkov telescope \\citep{WhippleGRB}, and coincident and monitoring studies by AIROBICC \\citep{Padilla}, Whipple \\citep{WhippleGRB1sec}, and the Milagro prototype \\citep{MilagritoGRB, Leonor, McCullough}.  While several of these studies have reported evidence for VHE GRB emission, no clear picture of the VHE emission from gamma-ray bursts has emerged.   In recent years the rate of possible GRB coincident detections has been relatively low due to the absence of sensitive, large field of view GRB detectors like BATSE. However, Milagro is a very sensitive detector with a low $\\sim$100 GeV energy threshold and wide $\\gtrsim$ 2 sr field of view, and is fully capable of autonomous identification of a VHE burst of emission. The low energy threshold is particularly important because of the reduced attenuation by extragalactic background light near 100 GeV \\citep{Jelley,Primack:TeVAbsorption,Stecker:EBL}, which dramatically increases the volume of space observed. This paper details the monitoring of the northern sky for VHE transients of 40 s -- 3 hours duration with the Milagro observatory. ",
        "conclusions": ""
    },
    "0311/astro-ph0311340_arXiv.txt": {
        "abstract": "The Guitar nebula is a spectacular example of an \\Halpha\\ bow shock nebula produced by the interaction of a neutron star with its environment.  The radio pulsar B2224+65 is traveling at $\\sim$800--1600~km~s$^{-1}$ (for a distance of 1--2~kpc), placing it on the high-velocity tail of the pulsar velocity distribution.  Here we report time evolution in the shape of the Guitar nebula, the first such observations for a bow shock nebula, as seen in \\Halpha\\ imaging with the Hubble Space Telescope.  The morphology of the nebula provides no evidence for anisotropy in the pulsar wind, nor for fluctuations in the pulsar wind luminosity.  The nebula shows morphological changes over two epochs spaced by seven years that imply the existence of significant gradients and inhomogeneities in the ambient interstellar medium.  These observations offer astrophysically unique, {\\em in situ} probes of length scales between $5\\times10^{-4}$~pc and $0.012$~pc.  Model fitting suggests that the nebula axis --- and thus the three-dimensional velocity vector --- lies within 20\\arcdeg of the plane of the sky, and also jointly constrains the distance to the neutron star and the ambient density. ",
        "introduction": "While the steady decay of spin rates observed in radio pulsars provides a good estimate of the rate of energy loss \\Edot, only a small fraction of the energy output of a neutron star is typically converted to directly detectable electromagnetic radiation. Most of the spindown energy loss is carried away by a relativistic wind, the properties of which are largely unknown.  The best constraints on the relativistic wind derive from its interaction with the interstellar medium (ISM).  This interaction has been observed, for example, in synchrotron nebulae such as the Crab Nebula \\citep{KC84,GA94,MM96}, where wisp structures are observed moving at $\\sim 0.5c$ \\citep{HMB+02}, as well as bow shocks observed in \\Halpha\\ emission from shock-excited neutral gas around PSR~B1957+20 \\citep{KH88} and various other pulsars \\citep{CC02,GJS02}.  The Guitar nebula was discovered in deep \\Halpha\\ imaging observations with the 5-m Hale telescope at Palomar Observatory \\citep{CRL93}. Using a model for the Galactic electron density distribution \\citep{CL02}, the dispersion measure (DM $= \\int_0^Dn_e\\,ds = $ 35.3 pc~cm$^{-3}$) of PSR~B2224+65 implies a distance of 1.9~kpc (hereafter parameterized as D$_{1.9}$).  While the pulsar has a modest \\Edot = 10$^{33.1}$ erg~s$^{-1}$, the large space velocity of the pulsar ($\\sim 1640$ D$_{1.9}$ km~s$^{-1}$) provides the ram pressure needed to create a detectable bow shock nebula.  The existence of an observable bow shock nebula also implies the presence of a significant neutral hydrogen component in the ISM near the pulsar.  Here we describe Hubble Space Telescope (HST) observations of the time evolution of the nebula, and discuss the implications of the observations for our understanding of neutron star (NS) relativistic winds and the ISM. ",
        "conclusions": "The wavenumber spectrum for electron density fluctuations in the general ISM has been delineated through a variety of measurements, including radio scintillation, scattering, pulse time-of-arrival, and Faraday-rotation, as illustrated in Figure~\\ref{fig-spec}. These measurements constrain the power levels in different wavenumber intervals, and suggest an overall consistency with a Kolmogorov turbulence process \\citep{ARS95}, although this consistency is only coarse and may be an illusion associated with the large amplitude scale of the diagram.  Departures from the Kolmogorov spectrum are also inferred from pulsar scintillation measurements.  The HST measurements of the Guitar Nebula provide constraints on wavenumbers not easily accessible by other probes and therefore provide a new tool for investigating fine structure in the interstellar density.  \\begin{figure}[ht] \\epsscale{1.0} \\plotone{f3.eps} \\caption{Power spectrum for electron density variations in the ISM, defined such that the integral over all wavenumbers is the mean-square electron density.  The region labeled `GN' designates the constraints derived from HST observations of the Guitar Nebula described here. The wavenumber range corresponds to the reciprocals of the length scales probed on angular scales between $\\sim 0.05$ and $1.3$ arcsec, assuming a distance of 1.9 kpc and a combined filling factor and ionized fraction ranging from 0.02 to 0.5.  The dashed line labelled `HI' is an estimate using the length scales of standard HI clouds \\citep{MO77}, assumed to be partially ionized.  The line labelled `MSP' represents variations in dispersion measure for the millisecond pulsar B1937+21 \\citep{CWD+90,KTR94} The observational constraints at the highest wavenumbers result from measurements of interstellar scintillation and scattering.  The point for Sgr~A* represents the power level for scattering of the Galactic center source, corrected for the proximity of the scattering screen to the source \\citep{LC98}. The hatched regions incorporate constraints from scintillation and scattering observed for several hundred sources \\citep[NE2001;][and references therein]{CL02}.  The light solid line designates a power spectrum with power-law index equal to the Kolmogorov value ($-11/3$). \\label{fig-spec} } \\end{figure}  Along with the properties of the ISM, the Guitar nebula also provides a probe of the pulsar's relativistic wind.  The properties of pulsar winds, including the magnetization parameter $\\sigma$, the ratio of the Poynting flux to the kinetic energy flux, have been inferred primarily from the Crab nebula \\citep{KC84,GA94,MM96}.  The existence of bow shock nebulae such as the Guitar requires collisional excitation of the neutral interstellar medium through interactions with shocked electrons and protons. These charged particles can be ejected from the NS itself, or originate from interstellar atoms through photoionization or magnetic reconnection outside the pulsar light cylinder radius ($R_{LC} = c P/2\\pi$). The process by which the NS Poynting flux is converted to particle kinetic energy is not well understood: overviews of different processes are provided by \\citet{BL92}, \\citet[ion loading]{conf-GVKA02} and \\citet[striped pulsar winds]{LK01}. At the stand-off radius ($R_0 \\gg R_{LC}$), pressure balance requires that the spindown energy of the NS be carried by the particle flux ($\\sigma\\ll 1$), but the shape of the nebula may encode information about the (rotation-averaged) shape of the NS wind.  Along with the Guitar nebula, future observations of the evolution of bow shock nebulae will be possible for the nearby radio-quiet neutron star RX~J1856.5$-$3754 \\citep{VK01}, only $\\sim$120 pc away \\citep{KVA02,WL02}.  Even at its relatively low speed ($< 200$~km~s$^{-1}$), the neutron star travels $\\sim 1$\\arcsec\\ in a year, and evolution of the nebula should be evident on roughly this time scale.  PSR~J2124$-$3358, a nearby millisecond pulsar with a complex bow shock nebula \\citep{GJS02}, is also promising in this regard, while the discovery of a bow shock nebula powered by the Poynting flux from a magnetar would allow investigation of the relativistic wind in the presence of an ultra-strong magnetic field."
    },
    "0311/astro-ph0311030_arXiv.txt": {
        "abstract": "{ We discuss the selection and observations of a large sample of nearby galaxies, which we are using to quantify the star formation activity in the local Universe.  The sample consists of 334 galaxies across all Hubble types from S0/a to Im and with recession velocities of between 0 and 3000 km~s$^{-1}$.  The basic data for each galaxy are narrow band H$\\alpha +$[N{\\sc ii}] and $R$-band imaging, from which we derive star formation rates, H$\\alpha +$[N{\\sc ii}] equivalent widths and surface brightnesses, and $R$-band total magnitudes.  A strong correlation is found between total star formation rate and Hubble type, with the strongest star formation in isolated galaxies occurring in Sc and Sbc types.  More surprisingly, no significant trend is found between H$\\alpha +$[N{\\sc ii}] equivalent width and galaxy $R$-band luminosity.  More detailed analyses of the data set presented here will be described in subsequent papers. ",
        "introduction": "Knowledge of the star formation histories of both the Universe and of individual galaxies provides the foundation for our understanding of the evolution of the Universe we see today.  Substantial advances have been made in our understanding of high-redshift star formation, (e.g. Madau et al. \\cite{madau96}; Madau et al. \\cite{madau98}; Steidel et al. \\cite{stei}; Hopkins et al. \\cite{hopk}; Lanzetta et al. \\cite{lanz}). This has resulted in an anomalous situation in which the star formation history of the Universe appears to be more fully quantified at high redshift than it is locally.  This results at least partly from the relative ease with which a representative volume of the Universe can be observed at high redshift, since only a small area of the sky need be observed. Star formation rates have been determined for many local galaxies using a range of different techniques, e.g. emission-line fluxes (e.g., Kennicutt \\& Kent \\cite{ken83}; Young et al. \\cite{young}; Gallego et al. \\cite{gallego95}; Gallego et al. \\cite{gallego96}), far-infrared luminosities (e.g., Kennicutt et al. \\cite{ken87}), radio luminosities (e.g., Condon \\cite{condon}, Cram et al. \\cite{cram}) or direct ultraviolet emission from hot stars (e.g. Bell \\& Kennicutt \\cite{bell}). However, in the main these studies have looked at the brightest and most rapidly star-forming galaxies, and {\\bf most} spectroscopic studies are biased towards galaxies with large equivalent width (EW) in the line concerned.  In this study, one focus will be on the star formation properties across the full range of the numerically dominant dwarf galaxies, and the global star formation rate from this population. A further consequence of the observational limitations on areal coverage is that galaxy clusters are easier to study than the field environment, and hence local field galaxy star formation rates are relatively poorly constrained compared with their cluster counterparts.  In the present study, we attempt to quantify star formation activity across the full range of star-forming galaxies, down to the faintest dwarf irregular types, by using narrow-band imaging through filters centred on the redshifted Balmer H$\\alpha$ line.  The advantages of this technique are that it is sensitive to low levels of star formation even in faint, low surface brightness galaxies, and that it traces high mass stars, and hence recent star formation.  Given suitable assumptions, principally about extinction and the stellar initial mass function (IMF), it yields quantitative measurements of the star formation rate.  The method can be applied to a sensitive level using relatively modest integration times on small telescopes, which is an important consideration given that this project must be done one object at a time, to ensure that each target galaxy is observed with the correct filter.  The resulting data not only give estimates of the total star formation rate for each galaxy, but also detailed information on the star formation distribution, enabling, for example, the separation of nuclear and disk activity in spiral galaxies. Finally, the large format of current CCDs gives a reasonable probability of detecting nearby star-forming companion galaxies which lie in the same field as target galaxies, and which have similar recession velocities.  The principal drawbacks of narrow-band H$\\alpha$ imaging are the need for large and uncertain extinction corrections; the need to assume an IMF to extrapolate from the quantity of high mass stars, responsible for the ionising flux, to the total mass of the young stellar population; contamination in most of the filters used by the [N{\\sc ii}] 6548 and 6584~\\AA\\ lines, necessitating a further correction to the derived star formation rates; and possible contributions to the line emission by central active galactic nuclei (AGN). All of these assumptions and corrections have been investigated in detail, principally by Kennicutt and his collaborators (Kennicutt \\& Kent \\cite{ken83}; Kennicutt \\cite{ken98}); we will re-examine some of the corrections they derived in Paper II of this series (James et al. \\cite{jame}).  The final disadvantage of using narrow-band filters is that a comprehensive survey of a contiguous area (c.f. Sloan, IRAS) cannot be completed in a reasonable amount of time, due to the small recession velocity coverage of each of the narrow-band filters, and hence a pre-existing galaxy catalogue must be used to provide a target list for specific pointed observations with the appropriate filters.  Any discussion of previous work in this area must first acknowledge the extensive studies undertaken by Kennicutt and collaborators in defining the techniques for H$\\alpha$ measurement and deriving star formation rates from such measurements (Kennicutt \\& Kent \\cite{ken83}), and in applying these measurements to studies of bright spiral and irregular galaxies (e.g., Kennicutt et al. \\cite{ken94}) and interacting galaxies (Kennicutt et al. \\cite{ken87}). All of this work was comprehensively reviewed by Kennicutt (\\cite{ken98}). Ryder \\& Dopita made a detailed study of 34 nearby southern spirals in H$\\alpha$, $V$ and $I$ band emission, demonstrating that star formation can be much more asymmetric than the underlying stellar distribution (Ryder \\& Dopita \\cite{ryde93}) and that the H$\\alpha$ scale length tends to be longer than that of the stellar light distribution (Ryder \\& Dopita \\cite{ryde94}).  These observations were used to derive a form for the star formation law in spiral galaxies (Dopita \\& Ryder \\cite{dopi94}). Young et al. (\\cite{young}) undertook a major study of 120 spiral galaxies using similar techniques to those used in the present work, to look at trends in H$\\alpha$ surface brightness with Hubble type and at the dependence of star formation rate on gas mass and interactions. Koopman et al. (\\cite{koopman}) looked at the H$\\alpha$ total emission and light profiles for 63 bright spiral galaxies in the Virgo cluster, again using narrow-band imaging techniques.  The Universidad Complutense de Madrid (UCM) emission-line survey (Gallego et al. \\cite{gallego95}, Gallego et al. \\cite{gallego96}) adopted a different and in many ways complementary strategy of searching areas of sky for all emission-line galaxies over a wide range of recession velocities, using the objective prism technique to detect candidates, and follow-up spectroscopy for confirmation and derivation of detailed spectroscopic parameters.  Gallego et al. (\\cite {gallego96}) present H$\\alpha$ and H$\\beta$ fluxes and EWs for over 200 galaxies detected in this way, and Gallego et al.  (\\cite{gallego95}) use this dataset to determine the total star formation rate in the Local Universe.  One of the aims of the present study is to rederive this parameter using a sample with very different selection criteria which are not dependent on H$\\alpha$ line strengths.  This will be presented in a later paper in this series.  Finally, an important recent study with which the current work should be compared is that of Charlot and collaborators (e.g. Charlot et al. \\cite{char02}), who used spectroscopic data from the Stromlo-APM survey to study a representative sample of star-forming galaxies with H$\\alpha$ EW down to 0.2~nm, compared with a limit of 1.0~nm for Gallego et al. (\\cite{gallego95}). The corresponding limit for the current survey is about 0.4~nm EW, although a few detections of line emission below this level are made.  The sample selection for the H$\\alpha$ Galaxy Survey is described in Sect. 2 of this paper.  The observational strategy and the data reduction process are detailed in Sections 3 and 4 respectively. Sect. 5 contains the data for the 334 galaxies, and Sect. 6 a discussion of some of the first results of this study. Sect. 7 contains our conclusions and a summary of planned future publications. ",
        "conclusions": "We have presented new narrow-band H$\\alpha +$[N{\\sc ii}] line and $R$-band imaging data photometry, of a sample of 334 galaxies carefully selected to span all star-forming types in the local Universe.  The galaxies were taken from the UGC, and hence are subject to the selection criteria of that catalogue, and in addition we have required a measured recession velocity less than 3000~km~s$^{-1}$, D$_{25}$ diameters between 1\\farcm7 and 6\\farcm0, and Hubble types between S0/a and Im.  The new data have been used to quantify the distributions of star formation activity as a function of galaxy type and luminosity.  We find a strong correlation between galaxy Hubble type and total star formation rate, in the expected sense that the most strongly star-forming galaxies are those of Hubble types Sbc and Sc, although even within these types there is a wide range in star formation rates.  H$\\alpha +$[N{\\sc ii}] EW values, averaged over entire galaxies, give a measure of star formation activity normalised by the stellar luminosity of the galaxy; these also show a trend with galaxy type, with early-type spiral galaxies having low EW values, and irregular galaxies the highest EWs, at least in the mean.  However, the distribution of EW as a function of galaxy luminosity is surprisingly flat, both in terms of median EW and the scatter in this parameter.  In this paper we have only skimmed the surface of the science to be undertaken in this project.  A key area for further study concerns the corrections which need to be applied to H$\\alpha$ fluxes derived from narrow-band imaging to account for [N{\\sc ii}] line contamination and extinction internal to the galaxies under study, since these dominate the uncertainties in the derived star formation rates (see also the detailed examination of this question by Charlot et al. \\cite{char02}, and work by Buat et al. \\cite{buat02} and Rosa-Gonz\\'alez et al. \\cite{gonz02}).  The spatial dependence of the [N{\\sc ii}] correction will be investigated, since the [N{\\sc ii}] line strength depends on metallicity and requires a harder radiation field than H$\\alpha$, suggesting that the relative strengths of the two lines may be very different in the disk and circumnuclear environments. We will also look at the type dependence of the H$\\alpha$ extinction correction in detail, building on the work of Kennicutt \\& Kent (\\cite{ken83}), since our sample contains many late-type and low-luminosity galaxies which were not well represented in their study.  Sect.  6.1 of the present paper illustrates that the H$\\alpha$ images contain substantial information on the distribution of line-emitting gas {\\em within} galaxies, as well as on the total quantity.  We will use this information to compare the distributions of star formation regions with the old stellar populations (as represented by $R$-band light) in galaxy disks and in irregular galaxies, with a key question being whether new stars are produced with a distribution similar to that of the old stars, or whether there is any evidence for disks and irregular galaxies being constructed `inside-out' or `outside-in'. The distributions of both new and old stellar populations will be quantified for this analysis by concentration indices.  Whilst the present H$\\alpha$ sample has been constructed to sample the field population, i.e. avoiding the centres of rich galaxy clusters, it still samples a range of galaxy environments, including completely isolated galaxies, galaxies in known pairs or groups, and some interacting systems.  We will thus study the effects of environment (local galaxy density) and interactions with very close neighbours on star formation rate, surface density and distribution in spiral and irregular galaxies, for comparison with results obtained by Kennicutt et al. (\\cite{ken87}), Loveday et al. (\\cite{love99}), Carter et al. (\\cite{cart01}) and Koopman et al. (\\cite{koopman}).  A final aim for this study is to determine the total star formation rate in the local Universe, by correcting the observed galaxy sample for the known selection effects to represent a volume-limited sample.  This will yield a galaxy luminosity function, from which we can estimate the total star formation rate per unit volume of the local Universe, and the contribution to this total from different galaxy luminosities and Hubble types.  We believe this to be a complementary approach to previous studies in this area, e.g. those by Gallego et al. (\\cite{gallego95}) who looked at 176 star-forming galaxies out to $z\\sim$0.045, and Tresse \\& Maddox (\\cite{tres98}), 138 galaxies to $z\\sim$0.3. Our study is much more local than these, raising concerns about how representative a volume it may sample, but intensively samples that volume given the large sample size. The proximity of the galaxies also enables us to include the faintest star-forming dwarfs, and gives us excellent spatial resolution within the galaxies, thus disentangling nuclear and disc contributions to the total star formation rate."
    },
    "0311/astro-ph0311206_arXiv.txt": {
        "abstract": "We present the intrinsic spectral energy distribution (SED) of the Narrow-Line Seyfert 1 galaxy (NLS1) \\objectname[MGC +05-53-012]{Ark~564}, constructed with contemporaneous data obtained during a multi-wavelength, multi-satellite observing campaign in 2000 and 2001. We compare this SED with that of the NLS1 Ton~S180 and with those obtained for Broad-Line Seyfert 1s to infer how the relative accretion rates vary among the Seyfert 1 population. Although the peak of the SED is not well constrained, in our parameterization most of the energy of this object is emitted in the 10--100\\,eV regime, constituting roughly half of the emitted energy in the optical/X-ray ranges. This is consistent with a primary spectral component peaking in the extreme UV/soft X-ray band, and with disk--corona models, hence high accretion rates. Indeed, we  estimate that  $\\dot{m} \\approx 1$. We also address the issue of the energy budget in this source by examining the emission lines observed in its spectrum, and we constrain the physical properties of the line-emitting gas through photoionization modeling. The available data suggest that the line-emitting gas is characterized by $\\log n \\approx 11$ and $\\log U \\approx 0$, and is stratified around $\\log U \\approx 0$. Our estimate of the radius of the H$\\beta$-emitting region $R_{\\rm BLR}^{{\\rm H}\\beta} \\approx 10 \\pm 2$ lt-days is consistent with the $R_{\\rm BLR}^{{\\rm H}\\beta}$--luminosity relationships found for Sy1 galaxies, which indicates that the narrowness of the emission lines in this NLS1 is not due to the Broad-Line Region  being relatively further away from the central mass than in BLS1s of comparable luminosity. We also find evidence for super-solar metallicity in this NLS1. We show that the emission lines are not good diagnostics for the underlying SEDs and that the absorption line studies offer a far more powerful tool to determine the ionizing continuum of AGNs, especially if comparing the lower- and higher-ionization lines. ",
        "introduction": "%  The population of Seyfert 1 galaxies has a widely used sub-classification into Narrow-Line Seyfert 1 galaxies (NLS1s) and Broad-Line Seyfert 1 galaxies (BLS1s). While this classification appears to make an arbitrary distinction based on the widths of the optical emission lines (NLS1s having FWHM(H$\\beta) \\lesssim 2000$ ${\\rm km\\ s^{-1}}$, \\citealt{Goodrich89}), this is in fact an extremely useful scheme since the X-ray properties of the two subclasses are systematically different. As a class, NLS1s show rapid and large-amplitude X-ray variability \\citep[][hereafter BBF96; \\citealt{Turnerea99b}]{BBF96}, with the excess variance \\citep{Nandraea97a} typically an order of magnitude larger than that observed for samples of BLS1s  with the same luminosity distribution \\citep{Turnerea99b,Leighly99I}. Analogously, the spectral properties also vary across the Seyfert population with NLS1s showing systematically steeper spectra than those of BLS1s in both the soft and hard X-ray bands \\citep{BBF96,BME97,TGN98,Leighly99II,Vaughan99b}.  One increasingly popular hypothesis to explain the differences in X-ray properties across the Seyfert population is that NLS1s have relatively low masses for the central black hole compared to BLS1s with similar luminosities. Smaller black-hole masses naturally explain both the narrowness of the optical emission lines, which are generated in gas that has relatively small Keplerian velocities, and the extreme X-ray variability, since the primary emission would originate in a smaller region around the central engine \\citep[e.g.,][]{Laorea97}. Given that NLS1s have comparable luminosity to that of the BLS1s, \\citet{PDO95} suggested that they must be emitting at higher fractions of their Eddington luminosity, hence higher fractional accretion rates ($\\dot{m} = \\dot{M}/\\dot{M}_{\\mbox{\\scriptsize Edd}}$) are also required. The closer the luminosity is to the Eddington limit (and the lower the black-hole mass), the greater the fraction of the energy emitted by the accretion disk in the soft X-rays \\citep{RFM92}. Thus NLS1s might be expected to show disk components which peak at higher energies than for BLS1s. \\citet{PDO95} and \\citet{MH97} noted that soft photons from the disk may Compton-cool hard X-rays from the corona, and cause the observed steep spectra.  Given the dependence of the peak energy of the disk emission on the accretion rate, the spectral energy distribution (SED) of an active galactic nucleus (AGN) will provide critical information about accretion rates and conditions close to the disk. In particular, examination of the SED of a NLS1, and comparison with that obtained for BLS1s will offer insight into the relative accretion rates across the Seyfert population. However, measuring the SED requires observations taken simultaneously over a long wavelength baseline stretching from infrared to hard X-ray wavelengths, and the determination of the UV/X-ray continuum in AGN is still difficult due to the severe attenuation by even small amounts of Galactic interstellar gas along the line-of-sight. Another complication comes from intrinsic reddening, i.e., reddening associated with the active nucleus itself. Indeed, the steeper UV/blue continua observed in NLS1s (when compared to AGN spectrum composites) can be attributed at least in part to reddening, though the ionization of the absorbing material and its location with respect to the accretion source is still not well determined \\citep{Constantinea03}.  Arakelian~564 (Ark~564, IRAS 22403+2927, MGC +05-53-012) is a bright, nearby NLS1 galaxy, with $z = 0.02467$ and  $V = 14.6$ mag \\citep{rc3.9catalogue}, and a mean 2--10\\,keV luminosity $L_{\\mbox{\\scriptsize 2--10{} keV}} \\approx 2.4 \\times 10^{43}$ \\ergsec{} with flux variations of a factor of a few in a few thousand seconds \\citep[][hereafter Paper~I]{Akn564I}. It was the object of an intense multiwavelength monitoring campaign that included simultaneous observations from {\\it ASCA} (2000 June 1 to July 6, Paper~I; \\citealt{Poundsea01,Edelsonea02}), {\\it XMM-Newton} (2000 June 17, \\citealt{Vignaliea03}), {\\it Chandra} (2000 June 17, \\citealt{Matsumotoea02}), {\\it HST} (2000 May 9 to July 8, \\citealt{Collierea01}, Paper~II; \\citealt{Crenshawea02}, Paper~IV), {\\it FUSE} (2001 June 29--30, \\citealt{Romanoea02b}, Paper V), and from many ground-based observatories as part of an AGN Watch\\footnote{\\anchor{http://www.astronomy.ohio-state.edu/~agnwatch} {All publicly available data and complete references to published AGN Watch papers can be found at http://www.astronomy.ohio-state.edu/$\\sim$agnwatch.}} project \\citep[1998 November to 2001 January, ][ Paper~III]{Shemmerea01}. Ark~564 has shown a strong associated UV absorber (\\citealt{Crenshawea99}, Paper~II; Paper~IV; Paper V). There are indications that it also possesses a warm X-ray absorber, as seen by the narrow absorption lines of \\ion{O}{7} and \\ion{O}{8} detected in a {\\it Chandra} spectrum \\citep{Matsumotoea02}, and that the UV and X-ray absorbers in Arakelian~564 are physically related, possibly identical, and may be spatially extended along the line of sight (Paper~V).  In this paper we present a contemporaneous SED of Ark~564, based on the extensive monitoring of 2000. In \\S\\ref{sa:allobs} we describe the observations and data reduction, summarize the main results of the monitoring campaign, and describe the adopted method for reddening correction. In \\S\\ref{sa:sed} we present the SED of Ark~564. In \\S\\ref{sa:pioniza} we constrain the mean physical properties of the line-emitting gas through photoionization modeling. In \\S\\ref{sa:discuss} we discuss some implications of our investigation. Finally, our results are summarized in \\S\\ref{sa:summary}. ",
        "conclusions": "}  A non-simultaneous optical, UV and X-ray spectral energy distribution of Ark~564 was presented by \\citet{Comastriea01} who found that it peaks in the soft-X-ray band. Here we present a spectral energy distribution which is obtained from contemporaneous data covering almost 5 decades in energy\\footnote{The Lick spectrum, which was obtained in 1980 is merely used to fill in a small gap in the optical data, and the 2001 {\\it FUSE} spectrum shows a continuum slope consistent with the one obtained from the {\\it HST} spectrum. This latter fact suggests that, although the flux level changed between 2000 and 2001, the overall shape of the optical/FUV spectral energy distribution did not change. }. Simultaneity is particularly important for NLS1s, since, as a class, they are extremely variable in time, although Ark~564 has shown only weak variability in the optical/UV bands (Paper II-III).  We report some relevant data from the spectral energy distribution in Table~\\ref{sa:sedvalues}. These were derived using the simple parameterization of the combination of an optical/FUV power law with spectral index $\\alpha =0.42$ breaking at $\\sim 50$\\,eV to $\\alpha_{ASCA} = 1.538$ (SEDA in Figure~\\ref{sa:sed}), as well as the {\\it IRAS} and {\\it IRTF} data points (\\S\\ref{sa:dataredux}). A more conservative spectral energy distribution, instead, connects the FUV and the soft X-ray data with a simple power law ($\\alpha = 1.08$, the solid line in Figure~\\ref{sa:sedcorr}), and is a combination of an optical/FUV power law with $\\alpha =0.42$ breaking at 1000\\AA{} to $\\alpha =1.08 1$, then again breaking at $\\approx 0.8$\\,keV to $\\alpha_{ASCA} = 1.538$ (SEDB in Figure~\\ref{sa:sed}). The ambiguity in the shape of the spectral energy distribution in the 900\\AA--0.8\\,keV region is rather unfortunate, since a considerable portion of the energy of Ark~564 might be output  in this range. Previous {\\it ROSAT}  \\citep{Brandtea94} and {\\it BeppoSAX} \\citep{Comastriea01} data showed a flattening of the soft excess toward the lowest X-ray energies available, and the {\\it XMM} spectrum obtained during the monitoring campaign of 2000 shows a definite curvature in the soft excess \\citep{Vignaliea03}.  An interesting issue is how Ark~564 compares to other NLS1 galaxies and with BLS1 in terms of its broad-band properties, as they can be quantified by spectral indices. Table~\\ref{sa:aox} reports the intrinsic spectral indices calculated between different wavelength bands for Ark~564 and, as a comparison, the corresponding values for the NLS1 Ton~S180 and BLS1s. These show that while the inter-band properties of Ton~S180 are not significantly different from the ones observed in BLS1s \\citep{Turnerea02}, this may not be the case for Ark~564. Table~\\ref{sa:aox} indicates that the two NLS1s have steeper X-ray slopes than BLS1s, Ark~564 more so than Ton~S180 ($\\alpha_{\\rm x} = 1.57$ and 1.44, respectively, compared to 0.91 for BLS1s), which is consistent with the general characteristics of NLS1s. The optical/X-ray spectral index ($\\alpha_{\\rm ox}$), on the other hand, is lower in Ark~564 than in BLS1s and Ton~S180 with $\\alpha_{\\rm ox}=1.11$ for Ark~564\\footnote{The observed $\\alpha_{\\rm ox}$ is 0.94 in Ark~564.}, 1.52 for Ton~S180. This reflects the fact that Ark~564 is relatively more X-ray bright, or that the optical continuum is suppressed, compared to other AGNs. The other indices also reflect the X-ray brightness of Ark~564 when compared to BLS1s. We cannot exclude the possibility that intrinsic reddening in excess of the  \\centerline{} \\vspace{-1.5cm}  \\centerline{\\includegraphics[width=10.5cm,height=10.0cm]{figure8.ps}}  \\figcaption{Comparison of the SEDs of Ark~564 with the mean SED for radio-quiet quasars (RQ; \\citealt{Elvisea94}) and LZ \\citep{Laorea97,Zhengea97}, one Seyfert 1 galaxy (NGC~4151,  \\citealt{kraemerea00}), one Seyfert 1.5 galaxy (NGC~5548, \\citealt{KCFP98}, and the NLS1 Ton~S180 \\citep{Turnerea02}. [{\\it A color version of this plot is available in the electronic edition of the Journal.}] \\label{sa:compseds}}  \\centerline{} \\vspace{0.3cm}  \\noindent one considered here may be present that could be suppressing the optical, but we argue that it would require the continuum to be more reddened than the emission lines (e.g.\\ \\ion{He}{2} $\\lambda 1640/\\lambda 4686$ ratio; Paper IV) show. If these differences, however,  are not due to intrinsic reddening, then they represent a true difference in the energy balance of this NLS1 and that the X-ray emission is therefore consistent with the predictions of the slim disk models, indicating high accretion rates.  We estimated the luminosity in different energy bands using our parameterizations of the spectral energy distribution described above (SEDA and SEDB). Though with a $\\sim 20$\\% systematic error in the 0.3--0.6\\,keV band flux (C.\\ Vignali 2003, private communication), the {\\it XMM} spectrum provides a way to extend the energy range below the $\\sim$ 0.8\\,keV limit of the {\\it ASCA} spectrum. Therefore, we also estimated the luminosity for the best-fit model to this spectrum. Table~\\ref{sa:enbudget} reports these luminosities, and indicates that more energy is emitted in the 10--100\\,eV band than in any other decade, constituting roughly half of the emitted energy in the opt/X-ray ranges. This implies that the primary spectral component peaks in the extreme UV--soft X-ray band and this is generally consistent with disk-corona models \\citep{hm91}. The integrated luminosity between $10^{-5}$ and 10\\,keV gives a lower limit on the bolometric luminosity of $L_{\\rm bol} \\gtsim 10^{45}$\\,\\ergsec.  Growing evidence has been gathered that NLS1s are super-Eddington accretors. \\citet{CH01} and  \\citet{Collinea02} studied accretion rates in a sample of AGNs using the \\citet{Kaspiea00} relationship, and found that not only half of their sample accretes close to or above the Eddington rate, but also that the largest Eddington ratios are found in NLS1s. Indeed, they showed that NLS1s are at the extreme of a well-defined sequence relating the Eddington ratio to the line widths. Furthermore, \\citet{WN03} show that a model of extreme slim disk (which is responsible for the soft X-ray excess, or hump,  seen in most NLS1s) and a hot corona (contributing to the hard X-ray emission) can also naturally explain the X-ray spectral variability  characteristics observed in Ark~564 such as simultaneous variations of the soft hump and the hard X-ray without a significant time lag \\citep{Akn564I}. Through a comparison of the X-ray variability of Ark~564 and the BLS1 NGC~3516, \\citet{Poundsea01} estimate that the mass of the central black hole in Ark~564 is $\\sim 10^{7}\\,\\Msun$, implying an accretion rate in the range $\\dot{m} \\approx 0.2$--1. With a bolometric luminosity in the order of $10^{45}$\\,\\ergsec{} and an Eddington luminosity $L_{\\rm Edd} \\approx 10^{45}$\\,\\ergsec, we also infer that $\\dot{m} \\approx 1$. \\citet{WN03} also provide a means of estimating the black hole mass (if the accretion is super-critical) which is independent on the accretion rate itself. For Ark~564, we obtain $M \\approx 2 \\times 10^{6} \\Msun$, which is within a factor of a few from the \\citet{Collierea01} estimate.  Figure~\\ref{sa:compseds} compares the SED of Ark~564 with the mean SED for radio-quiet quasars \\citep{Elvisea94}, the LZ SED \\citep{Laorea97,Zhengea97}, the Seyfert 1.5 galaxies NGC~5548 \\citep{KCFP98} and NGC~4151 \\citep{kraemerea00}, and the NLS1 Ton~S180 \\citep{Turnerea02}. There are significant differences in the intrinsic shape of the SED across the AGN population (see, also, \\citealt{Turnerea02}), the most evident being the energy of the peak and the presence (or lack of) of the big blue bump (BBB), the signature of the emission from the accretion disk. The radial dependence of the temperature for an optically thick, geometrically thin accretion disk \\citep{SS73} is, $T(R) \\sim 6.3 \\times 10^{5} \\; (\\dot{m})^{1/4} \\; M_{8}^{-1/4} \\; (R/R_{\\rm S})^{-3/4} {\\rm K}$ \\citep{P2000}, where $M_8$ is the mass in units of $10^8 \\Msun$, $R$ is the radius, and $R_{\\rm S}$ is the Schwarzschild radius. For Ark~564, using $\\dot{m}\\approx 1$, as we derived above, and $M \\approx 8 \\times 10^{6}\\,\\Msun$ (Paper~II), the peak temperature is $\\sim 125$\\,eV; this is within a factor of 3 from the peak of our less conservative parameterization, SEDA (defined extrapolating the optical power-law continuum to meet the extrapolation of the X-ray power law) which peaks at 50\\,eV. The true SED probably peaks somewhere between these two values. For comparison, the NLS1s RE~J1034$+$396 \\citep{Puchnaea01} and Ton~S180 \\citep{Turnerea02} peak at $\\approx 250$\\,eV and $\\ltsim 100$\\,eV, respectively. Therefore, even among the NLS1s, differences in the shape of the SED are observed. However, in none of these NLS1s there is an indication of the presence of optical/UV BBB, and  a strong soft X-ray excess is seen, instead. In this light, Ark~564, is also consistent with the paradigm that the accretion disk is so hot in NLS1s that the BBB is shifted in the EUV--soft X-rays.  We also note that Ark~564 is rather FIR bright (with respect to the optical), compared to the sample of radio-quiet quasars and the LZ sample, as can be seen in Figure~\\ref{sa:compsedsIR}, where the SEDs have been normalized to match their optical/UV slope (as opposed to the 2\\,keV flux in Figure~\\ref{sa:compseds}). Indeed, if we use the definition of the IR flux as a function of the {\\it IRAS} fluxes given by \\citet{SandersMirabel96}\\footnote{$F(8-1000\\mu$m$) = 1.8 \\times 10^{-14} \\{13.48 f_{12} +5.16 f_{25} + 2.58 f_{60} +f_{100}\\}$ W m$^{-2}$, where $f_{12}, f_{25}, f_{60}$, and $f_{100}$ are the {\\it IRAS} flux densities in Jy at 12, 25, 60 and 100\\,$\\mu$m.}, we obtain that $L(8-1000\\mu$ m$) \\approx 10^{11} \\Lsun$, which makes Ark~564 a luminous IR galaxy. The shape of the IR/optical SED of Ark~564 also resembles the shapes observed in the {\\it IRAS} Bright Galaxy Survey (see Figure~2 in \\citealt{SandersMirabel96}, for intermediate values of $f_{60}$). Using SEDA, we  \\centerline{} \\vspace{-1cm}  \\centerline{\\includegraphics[width=10.5cm,height=10.0cm]{figure9.ps}}  \\centerline{} \\vspace{-0.3cm}  \\figcaption{Comparison of the Ark~564 SED  with the mean SED for radio-quiet quasars (RQ; \\citealt{Elvisea94}) , LZ \\citep{Laorea97,Zhengea97}, and the ``table agn'' model in Cloudy (AGN). The SEDs are normalized so that they match in the optical/UV part. \\label{sa:compsedsIR}}  \\centerline{} \\vspace{0.3cm}  \\noindent estimate that the FIR luminosity is $L_{\\rm FIR} \\approx 1.3\\times 10^{44}$\\,\\ergsec, i.e., $\\sim 10\\,\\%$ of the total luminosity and $\\sim 20\\,\\%$ of the combined optical/UV/X-ray luminosity. \\citet{Crenshawea02} noted that the associated warm UV absorber is lukewarm and dusty. The IR emission observed in this object could then be thermal emission from the dust grains embedded in the absorber as they are heated by the strong UV/EUV continuum. Given the IR brightness in this object and in many NLS1s \\citep{Moranea96}, it is not unlikely that a contribution might be coming from the host galaxy, in the form of a nuclear starburst \\citep{Mathur00}. \\citet{Crenshawea02}, however, did not detect any extended emission in the two-dimensional STIS spectral images that could be due to a nuclear starburst. We compared the {\\it FUSE} spectrum, which was taken through a much larger aperture, and hence is more likely to show stellar absorption features, with a {\\it FUSE} `template' spectrum of the starburst galaxy NGC~7496; using the constraints from the Ly$\\gamma$--\\ion{C}{3}$] \\lambda$977 profiles, we obtained a rough upper limit of the starburst contribution at 1000\\AA{} of 50\\%. A further constraint on the IR starburst contribution would probably come from detection and measurements of the 3.3--3.4$\\mu$m PAH emission features, which have been found in starburst galaxies, luminous IR galaxies, and obscured AGNs \\citep{Moorwood86,Imanishi02}, and which have been successfully detected in the NLS1 NGC~4051 \\citep{rodriguezviegas03}.   Using our SEDA and SEDB as inputs to {\\tt Cloudy} we predicted the intensity of the strongest lines in the FUV/UV spectrum and compared them with the observed values in order to constrain the physical parameters of the line-emitting gas, namely, the density and ionization parameter. Figure~\\ref{sa:diagc3andn5} and \\ref{sa:otherdiags} show that SEDA and SEDB, because of their strong EUV to soft X-ray flux, for a given observed line ratio predict values of ionization parameter which are lower than those with standard AGN continuum. From  \\ion{C}{3}] $\\lambda$1909, \\ion{C}{3} $\\lambda$977, and \\ion{N}{5} $\\lambda$1240, we infer that $\\log n \\approx 11$. Two classes of solutions for $U$ are consistent with this density value, one with low $U$ values ($\\log U \\approx -1.5$) and one with high $U$ values ($\\log U \\approx 0$). We discarded the low-$U$ class (\\S\\ref{sa:pionizaresults}) on the basis that the predicted widths of the lines, derived from the velocity dispersion $V = (GM/R_{\\rm BLR}f)^{0.5}$, of 740--880 \\kms{} are too small with respect to the observed ones ($\\approx 2000$ \\kms; Paper~II). As expected, we find that the BLR is stratified around $\\log U \\approx 0$, with higher ionization lines originating from regions with higher $U$.  Column (6) of Table~\\ref{sa:elratios} shows that FWHM(Ly$\\alpha$)$=$ 2114 \\kms, FWHM(\\ion{N}{5})$=$ 2809 \\kms, and FWHM(\\ion{C}{4})$=$ 1934 \\kms, which indicate that the radii of the L$\\alpha$, \\ion{N}{5}, and \\ion{C}{4} broad-line emitting regions are $R_{\\rm BLR}^{{\\rm L}\\alpha} \\approx 4.3$ lt-days, $R_{\\rm BLR}^{\\rm N \\sc V} \\approx 2.5$ lt-days, and $R_{\\rm BLR}^{\\rm C \\sc IV } \\approx 5.2$ lt-days. Using the findings of previous monitoring programs on Seyfert 1s \\citep{NetzerPeterson97}, we can estimate the size of the H$\\beta$-emitting region from $R_{\\rm BLR}^{{\\rm L}\\alpha} \\approx 0.5 R_{\\rm BLR}^{{\\rm H}\\beta}$, and $R_{\\rm BLR}^{\\rm N \\sc V} \\approx 0.2 R_{\\rm BLR}^{{\\rm H}\\beta}$. For NGC~5548, furthermore, $R_{\\rm BLR}^{\\rm C \\sc IV} \\approx 0.5 R_{\\rm BLR}^{{\\rm H}\\beta}$ \\citep{Peterson93}. We can conclude that $R_{\\rm BLR}^{{\\rm H}\\beta} \\approx 10 \\pm 2$ lt-days, which is consistent with the $R_{\\rm BLR}^{{\\rm H}\\beta}$--luminosity relationships of \\citet{Kaspiea00} and \\citet{P2000}, when we assume a luminosity $\\lambda L_{\\lambda} (5100$ \\AA)$ \\approx 3.2 \\times 10^{43}$ \\ergsec{} (Table~\\ref{sa:sedvalues}). This indicates that the BLR radius of this NLS1 is consistent with the distribution of BLR radius in BLS1s, and that the narrowness of the emission lines is not due to the BLR being relatively further away from the central mass than in BLS1s of comparable luminosity.  Table~\\ref{sa:elratios} shows that some interesting differences in line ratios can be found with respect to the Sy1.5 NGC~5548. Indeed, \\ion{N}{5} $\\lambda$1240 is stronger in Ark~564 by a factor of $\\sim 2.3$, while \\ion{C}{4} $\\lambda$1550, \\ion{C}{3}] $\\lambda$1909, and Mg II $\\lambda$2800 are weaker in Ark~564 by a factor of $\\sim 4.4$, 2.5 and 2.8, respectively. While all line ratios in Ark~564 are statistically consistent with the ones measured for a mean QSO (given the large uncertainties on our measurements), Carbon lines are at the lower end and Nitrogen at the upper end of the QSO distribution, confirming this trend for weak Carbon and strong Nitrogen in Ark~564. Furthermore, \\ion{C}{3}] $\\lambda$977 would indicate that $-3.13 < \\log U < -0.05$ (Figure~\\ref{sa:diagc3andn5}), and for this range, the observed \\ion{N}{5}$/$\\ion{C}{4} ratio is larger than the model predictions by a factor of $\\sim 8$. This may imply super-solar metallicity in this NLS1 as suggested by \\citet{Mathur00} and is consistent with the finding of \\citet{Shemmerea02} that NLS1s have higher metallicities than BL AGNs for a given luminosity.  An interesting question is how sensitive the emission lines are to the true shape of the ionizing continuum. Given the difficulty related to deblending the BEL and NEL components of the lines and consequent large errors on the observed line ratios, the present emission line data do not allow us to discriminate between SEDA and SEDB. These SEDs differ in the range 1000\\AA--750\\,eV (corresponding to the gap in the data between the {\\it FUSE} and {\\it ASCA} spectra), with the maximum difference at around 50\\,eV. This difference should show the most for the high ionization lines of \\ion{C}{4}, \\ion{N}{5} and \\ion{O}{6}. However, the predicted strength of their {\\it emission} lines is very similar for the two SEDs, for a wide range of ionization parameters of interest (Figures~\\ref{sa:diagc3andn5} and \\ref{sa:otherdiags}). Perhaps, this is the reason why the emission line spectra of most AGNs look so very similar, over a wide range of luminosities. This demonstrates how unsuitable emission lines are as diagnostics for the underlying SEDs.  The absorption lines, on the other hand, are sensitive to the  \\centerline{} \\vspace{-1cm}  \\centerline{\\includegraphics[width=10.5cm,height=10.0cm]{figure10.ps}}  \\centerline{} \\vspace{-0.3cm}  \\figcaption{Top panel shows the ratio of the column densities relative to SEDA and SEDB as a function of $U$ for \\ion{C}{4}, \\ion{N}{5}, \\ion{O}{6}, and \\ion{O}{7}. Bottom panel shows the ratio of the column densities of \\ion{O}{7} to \\ion{C}{4} as a function of $U$ for solar abundances, for SEDA and SEDB. \\label{sa:fions} }  \\centerline{} \\vspace{0.3cm}  \\noindent input SED \\citep{MWEF94}. Column densities of different ions can be inferred from the the fractional abundances $f_{\\rm ion}$, the total column density $N_{\\rm H}$, and the assumed abundances $N_{\\rm X}$, $N_{\\rm ion} = N_{\\rm H}\\, N_{\\rm X}\\,N_{\\rm ion}$, through photoionization calculations (assuming an SED). Figure~\\ref{sa:fions} (Top) shows the fractional column densities for the same ion calculated with the different SEDs [$N_{\\rm CIV}$(SEDA) /$N_{\\rm CIV}$(SEDB), $N_{\\rm NV}$(SEDA) /$N_{\\rm NV}$(SEDB), $N_{\\rm OVI}$(SEDA) /$N_{\\rm OVI}$(SEDB), and $N_{\\rm OVII}$(SEDA) /$N_{\\rm OVII}$(SEDB)], and illustrates how the  predicted column densities of high ionization lines differ, with \\ion{O}{6} showing the largest variations, with over two orders of magnitude difference. Figure~\\ref{sa:fions} (Bottom) compares the column density for \\ion{O}{7} and \\ion{C}{4} calculated with the same SED and solar abundances [$N_{\\rm OVII}$(SEDA) / $N_{\\rm CIV}$(SEDA) and $N_{\\rm OVII}$(SEDB)  / $N_{\\rm CIV}$(SEDB)], which turn out to be good diagnostics even at lower ionization parameters. Thus, in principle, absorption line studies offer a far more powerful tool to determine the ionizing continuum of AGNs, compared to the emission lines. This also underscores the power of multiwavelength observations, as the maximum discriminator comes from comparing the lower- and higher-ionization lines.   }  We presented the intrinsic spectral energy distribution of Ark~564, constructed with quasi-simultaneous data obtained during 2000 and 2001. We compared this SED with that of Ton~S180 and with those obtained for Broad-Line Seyfert 1s to infer how the relative accretion rates vary among the Seyfert 1 population. The peak of the SED is not well constrained; however, in our parameterization most of the energy of this object is emitted in the 10--100\\,eV regime, and constitutes roughly half of the emitted energy in the optical/X-ray ranges. This is consistent with a primary spectral component peaking in the soft X-ray band, therefore with the predictions of the slim disk models, hence high accretion rates. Indeed, we estimate that  $\\dot{m} \\approx 1$.  We constrained the mean physical conditions in the BELR of this AGN, by examining the emission lines observed in its spectrum, and deduced the physical properties of the line-emitting gas through photoionization modeling. We concluded that the line-emitting gas is characterized by $\\log n \\approx 11$ and $\\log U \\approx 0$, and is stratified around $\\log U \\approx 0$. Our estimate of the radius of the H$\\beta$ emitting region $R_{\\rm BLR}^{{\\rm H}\\beta} \\approx 10 \\pm 2$ lt-days, is consistent with the $R_{\\rm BLR}^{{\\rm H}\\beta}$-luminosity relationships of \\citet{Kaspiea00} and \\citet{P2000}. This indicates that the narrowness of the emission lines is not due to the BLR being relatively further away from the central mass than in BLS1s of comparable luminosity. We also find evidence for super-solar metallicity in this NLS1, based on the low \\ion{C}{4}/\\ion{N}{5} observed line ratio. While the emission lines turn out to be unsuitable as diagnostics for the underlying SEDs, we showed that absorption line studies offer a far more powerful tool to determine the ionizing continuum of AGNs, especially if comparing the lower- and higher-ionization lines. This underlines the power of multiwavelength observations."
    },
    "0311/astro-ph0311560_arXiv.txt": {
        "abstract": "\\noindent We present the analysis of a 35 ksec long {\\it Chandra} observation of the neutron star soft X--ray transient (SXT) SAX~J1810.8--2609. We detect three sources in the field of view. The position of one of them is consistent with the location of the ROSAT error circle of SAX~J1810.8--2609. The accurate {\\it Chandra} position of that source coincides with the position of the proposed optical counterpart, strengthening the identification as the counterpart. We detected the neutron star SXT system in quiescence at an unabsorbed luminosity of $\\sim1\\times10^{32}$ erg s$^{-1}$ (assuming a distance of 4.9 kpc). This luminosity is at the low--end of quiescent luminosities found in other neutron star SXTs. This renders support to the existence of a group of faint soft X--ray transients of which the accreting millisecond X--ray pulsar SAX~J1808.4--3658 is the most prominent member. The quiescent spectrum of SAX~J1810.8--2609 is well--fit with an absorbed power law with photon index of $3.3\\pm0.5$.  With a value of 3.3$\\times10^{21}$ cm$^{-2}$ the Galactic absorption is consistent with the value derived in outburst. Since the spectra of quiescent neutron star SXTs are often fit with an absorbed blackbody or neutron star atmosphere plus power--law model we also fitted the spectrum using those fit functions. Both models provide a good fit to the data. If cooling of the neutron star core and/or crust is responsible for the soft part of the spectrum the time averaged mass accretion rate must have been very low ($\\sim$5.7$\\times10^{-13}$ M$_\\odot$ yr$^{-1}$; assuming standard core cooling only) or the neutron star must be massive. We also discuss the possibility that the thermal spectral component in neutron stars in quiescence is produced by residual accretion. ",
        "introduction": "\\label{intro}  Low--mass X--ray binaries are binary systems in which a $\\approxlt$1\\,$M_{\\odot}$ star transfers matter to a neutron star or a black hole. A large fraction of the low--mass X--ray binaries is transient; these systems form the so called soft X--ray transients (SXTs). The characterising property of SXTs is that the accretion rate drops several orders of magnitude when the source returns to quiescence (\\pcite{1984heta.conf...49V}).  SXTs have been and are being studied extensively with previous and present X--ray satellites; for an overview of SXTs see \\scite{1997ApJ...491..312C}. Using data obtained with the Beppo{\\it SAX} satellite \\scite{1999ApL&C..38..297H} and \\scite{2001ESASP.459..463I} suggested that $\\sim$10 bursting neutron stars form a separate class of {\\it faint} SXTs; the outburst peak luminosity is low (typically $\\sim10^{36.5}$ erg s$^{-1}$). Later, \\scite{2000MNRAS.315L..33K} argued that a class of faint SXT could be explained evolutionarily. \\scite{2000MNRAS.315L..33K} argues that faint SXTs should have evolved beyond the period minimum of $\\sim$80 minutes to orbital periods of 80--120 minutes.  Using the {\\it Chandra} and {\\it XMM/Newton} satellites it is possible to study these systems even when they are in quiescence (cf.~\\pcite{2001ApJ...560L.159W}; \\pcite{2002ApJ...580..413R}; \\pcite{2002ApJ...575L..15C}). A large fraction of the neutron star SXT quiescent X--ray luminosity is often ascribed to cooling of the hot neutron star core and/or crust (cf.~\\pcite{1987A&A...182...47V}; \\pcite{1998ApJ...504L..95B}).  Assuming that the thermal spectral component is coming from the neutron star surface, in principle neutron star parameters such as the mass and radius can be derived when using a neutron star atmosphere model to fit the spectrum (\\pcite{2001ApJ...551..921R}; \\pcite{2001ApJ...559.1054R}; \\pcite{2003ApJ...588..452H}).  The quiescent luminosity of the neutron star SXT and accreting millisecond X--ray pulsar SAX~J1808.4--3658 was found to be very low (L$_X=5\\times10^{31}$ erg s$^{-1}$; \\pcite{2002ApJ...575L..15C}). This is nearly as low as the luminosities found for several quiescent BHC SXTs with a short orbital period (\\pcite{2002ApJ...570..277K}).  Such a low neutron star luminosity could result if the neutron star core and crust are relatively cool (\\pcite{2001ApJ...548L.175C}). This in turn could hint at a massive neutron star (M$_{NS}>1.7$ M$_\\odot$; \\pcite{2001ApJ...548L.175C}), or a very low time--averaged mass accretion rate (\\pcite{1998ApJ...504L..95B}).  Several other emission mechanisms have been proposed to explain the quiescent luminosity of SXTs as well. First, in quiescence mass accretion may be ongoing at a low level (\\pcite{1995ApJ...439..849Z}, possibly via an advection dominated accretion flow, e.g.~\\pcite{1999ApJ...520..276M}). Second, shocks formed by the propeller mechanism preventing accretion onto the neutron star may produce X--rays (\\pcite{1975A&A....39..185I}; \\pcite{1986ApJ...308..669S}). Third, the switch--on of a radio pulsar mechanism may produce the observed X--rays (\\pcite{2000ApJ...541..849C}).  Finally, leaking of matter through the magnetospere especially at high latitudes in an ADAF--like spherical flow may produce X--rays (e.g.~\\pcite{1998ApJ...494L..71Z}).  The quiescent source luminosity seems to be a characterising difference between the systems containing a neutron star and those containing a black hole candidate (BHC) compact object.  When plotted versus the binary orbital period the neutron star SXTs seem to be systematically brighter by a factor of $\\sim$10 than the BHC SXTs (cf.~\\pcite{2001ApJ...553L..47G}). A possible and exciting explanation is that in the BHC SXTs energy is advected across the event horizon and not thermalised and radiated away like in neutron star SXTs (\\pcite{1997ApJ...482..448N}; \\pcite{2001ApJ...553L..47G}).  However, \\scite{2003MNRAS.343L..99F} suggested that the difference between the quiescent X--ray luminosities from neutron stars and BHC SXTs could also be explained by taking into account the jet power.   SAX~J1810.8--2609 was discovered with the Wide Field Cameras on board the Beppo{\\it SAX} satellite (\\pcite{1998IAUC.6838....1U}). A type~I X--ray burst was discovered establishing that the compact object is a neutron star (\\pcite{1998IAUC.6838....1U}; \\pcite{1999ApL&C..38..133C}). The burst showed evidence for radius expansion which led to a distance estimate of $\\sim4.9$ kpc (\\pcite{2000ApJ...536..891N}). \\scite{1999MNRAS.308L..17G} obtained ROSAT HRI follow--up observations and reported a candidate optical and near--infrared counterpart of SAX~J1810.8--2609.  In this Paper we present the analysis of our 35 ksec {\\it Chandra} observation of the faint neutron star SXT SAX~J1810.8--2609 while the source is in quiescence. ",
        "conclusions": "We detected X--ray emission from three sources using data obtained with the {\\it Chandra} satellite. The coordinates of only one of the sources are consistent with the ROSAT position of the neutron star SXT SAX~J1810.8--2609. We conclude that we detected SAX~J1810.8--2609 in quiescence. The source position is also fully consistent with the proposed optical--infrared counterpart of SAX~J1810.8--2609 (\\pcite{1999MNRAS.308L..17G}) strengthening the identification. The source spectrum is well described with an absorbed neutron star atmosphere, or blackbody, plus power--law model. With a value of 0.14 or 0.07 keV, the temperature of the neutron star as fitted with the blackbody and the NSA model, respectively, is low compared with many other neutron star SXTs in quiescence where values around 0.2--0.4 keV for the blackbody or 0.1--0.2 keV for the NSA model are usual (cf.~\\pcite{1996PASJ...48..257A}; \\pcite{2001ApJ...559.1054R}, but see also \\pcite{2001ApJ...551..921R}).  A single absorbed power law also gave a good fit. However, with a power--law index of $\\sim$3.3 it was rather soft.  The unabsorbed 0.5--10 keV source flux and luminosity are $\\sim4\\times10^{-14}$ erg cm$^{-2}$ s$^{-1}$ and $\\sim1\\times10^{32}$ erg s$^{-1}$, respectively.  In calculating the luminosity we used the source distance of 4.9 kpc as determined by \\scite{2000ApJ...536..891N}. The outburst peak luminosity of SAX~J1810.8--2609 was $\\sim2\\times10^{36}$ erg s$^{-1}$ (\\pcite{2000ApJ...536..891N}).  The outburst and quiescent luminosity are at the low end of typical outburst and quiescent luminosities of neutron star SXTs of $\\sim2\\times10^{38}$ erg s$^{-1}$ and several times $\\sim10^{32-33}$ erg s$^{-1}$, respectively. Furthermore, they are close to those of the accreting millisecond X--ray pulsar SAX~J1808.4--3658 ($\\sim2\\times10^{36}$ erg s$^{-1}$, and $\\sim5\\times10^{31}$ erg s$^{-1}$, for a distance of 2.5 kpc; \\pcite{1998A&A...331L..25I}; \\pcite{1998Natur.394..344W}; \\pcite{2001A&A...372..916I}; \\pcite{2002ApJ...575L..15C}). Both systems were proposed to be members of the class of recently recognised sub--group of {\\it faint} neutron star SXTs (\\pcite{1999ApL&C..38..297H}; \\pcite{2001ESASP.459..463I}). Obviously, if one would underestimate the source distance of the faint systems by a factor of a few one could derive anomalously low values for the luminosity. However, in both these neutron stars, radius expansion bursts have been found (see references above). As was shown by \\scite{2003A&A...399..663K} using a sample of neutron stars in Globular Clusters exhibiting radius expansion bursts, the typical error in the distance estimate as derived from the radius expansion burst is 15 per cent. This is insufficient to explain the difference in outburst and quiescent luminosity between the faint and 'normal' neutron star SXTs. The low luminosity of SAX~J1810.8--2609 in quiescence hence strengthens the case for the existence of the class of faint (neutron star) SXTs.  The quiescent luminosity of neutron star SXTs has been ascribed to thermal cooling of the hot neutron star crust/core (\\pcite{1998ApJ...504L..95B}). However, this interpretation is not without problems. As was noted by several authors (among which \\pcite{1998ApJ...504L..95B} themselves, see also \\pcite{2000ApJ...541..849C}), thermal cooling cannot explain the power--law spectral component. As mentioned in the introduction, a low quiescent luminosity could arise if the neutron star is more massive than $\\sim1.7$ M$_\\odot$ since in that case enhanced core cooling due to the direct URCA process could take place (\\pcite{2001ApJ...548L.175C}). However, the neutron star core could also be cold if the time--averaged mass accretion rate onto the neutron star is very low (\\pcite{1998ApJ...504L..95B}). In case of SAX~J1810.8--2609 the time average mass accretion rate necessary to explain the flux coming from the thermal spectral component would be 5.7$\\times10^{-13}$ M$_\\odot$ yr$^{-1}$ assuming that the amount of heat deposited in the crust per accreted nucleon is 1.45 MeV (assuming that only standard core cooling processes are at work; see \\pcite{1998ApJ...504L..95B}).  Initially, the relatively large variability found in the temperature associated with the soft component in Cen~X--4 and Aql~X--1 was regarded as a problem for the cooling neutron star model (\\pcite{1999ApJ...514..945R}; \\pcite{2001ApJ...559.1054R}) and only if the neutron stars are massive could the variability be explained (\\pcite{2001MNRAS.325.1157U}).  However, \\scite{2002ApJ...574..920B} found that since the residual hydrogen and helium mass in the neutron star atmosphere varies from one outburst to the next the thermal quiescent flux can vary between outbursts by a factor of 2--3. However, \\scite{2002ApJ...577..346R} found that the quiescent thermal flux of Aql~X--1 after an outburst first decreased by up to 50 per cent and then increased again by 35 per cent on timescales of a month. It also varied by nearly 40 per cent on short time scales (ksec). \\scite{2002ApJ...577..346R} concluded that accretion ongoing at a low level is the most likely explanation (however, it is unclear why the the luminosity agrees well with the predictions from deep crustal heating. See also \\pcite{2003ApJ...597..474C} for an alternative explanation).  \\scite{1987A&A...182...47V} also considered the possibility that mass accretion is ongoing at a low rate when the systems are in quiescence.  \\scite{1995ApJ...439..849Z} showed that emergent spectra of neutron stars accreting at a low rate would resemble a thermal spectrum albeit overall hardened.  We investigated the observational findings of neutron star SXTs in quiescence.  There is evidence that in sources which have a low quiescent source luminosity the contribution of the power--law component to the spectrum is large compared with more luminous neutron star sources in quiescence. The source with the lowest luminosity, SAX~J1808.4--3658 has a spectrum which is well--fitted with just a power--law (\\pcite{2002ApJ...575L..15C}). Other sources like Cen~X--4 and SAX~J1810.8--2609 with luminosity, L, 10$^{32}\\approxlt$L$\\approxlt5\\times10^{32}$ erg s$^{-1}$ show a power--law component which contributes up to about 50 per cent of the source luminosity (cf.~\\pcite{1996PASJ...48..257A}; \\pcite{2001ApJ...551..921R}; this work). Sources with luminosity near 10$^{33}$ erg s$^{-1}$ like X~5 and X~7 in the Globular Cluster 47~Tucanae (\\pcite{2003ApJ...588..452H}) are well--fit without the power--law component. Most quiescent neutron star SXTs in Globular Clusters fit this trend (the addition of a power--law component to the fit was never necessary to obtain a good fit, \\pcite{2003ApJSubmittedHetal}). However, the neutron star SXT in Terzan 5 and Aql~X--1 complicate the picture.  Their quiescent spectrum is dominated by or in the case of Aql~X--1 has a power--law component even though the quiescence luminosity is high $\\sim$2--4$\\times$10$^{33}$ erg s$^{-1}$ (for Terzan 5 see \\pcite{wijnsubter5}; for Aql~X--1 see e.g.~\\pcite{2001ApJ...559.1054R}; \\pcite{2002ApJ...577..346R}). Perhaps the power--law spectral component at source luminosities higher than 10$^{33}$ erg s$^{-1}$ has a different origin than those at lower source luminosities.  Finally, as was noted before (cf.~\\pcite{2003ApJ...594..952W}), we like to point out that assuming the observed black body radiation is caused by residual accretion and not a cooling neutron star crust the power--law spectral component at the lowest source luminosities (i.e.~those found in SAX~J1808.4--3658) could be similar in nature to the one found for black hole candidate SXTs in quiescence. This could tie in with the neutron stars entering a jet--dominated state when the accreting rate drops (\\pcite{2003MNRAS.343L..99F})."
    },
    "0311/astro-ph0311083_arXiv.txt": {
        "abstract": "Numerous numerical studies indicate that dark matter halos show an almost universal radial density profile. The origin of the profile is still under debate. We investigate this topic and pay particular attention to the velocity dispersion profile. To this end we have performed high-resolution simulations with two independent codes, ART and {\\sc Gadget}. The radial velocity dispersion can be approximated as function of the potential by $\\sigma_r^2 = a (\\Phi / \\Phi_{\\rm{out}} )^\\kappa ( \\Phi_{\\rm{out}} - \\Phi )$, where $\\Phi_{\\rm{out}}$ is the outer potential of the halo. For the parameters $a$ and $\\kappa$ we find $a=0.29\\pm0.04$ and $\\kappa=0.41\\pm0.03$. We find that the power-law asymptote $\\sigma^2 \\propto \\Phi^\\kappa$ is valid out to much larger distances from the halo center than any power asymptote for the density profile $\\rho \\propto r^{-n}$. The asymptotic slope $n(r \\to 0 )$ of the density profile is related to the exponent $\\kappa$ via $n=2\\kappa/(1+\\kappa)$. Thus the value obtained for $\\kappa$ from the available simulation data can be used to obtain an estimate of the density profile below presently resolved scales. We predict a continuously decreasing $n$ towards the halo center with the asymptotic value $n \\lesssim 0.58$ at $r=0$. ",
        "introduction": "The formation of structures by purely gravitationally interacting cold dark matter (CDM) is one of the fundamental paradigms in cosmology.  The luminous baryonic matter is embedded in dark matter halos.  Numerical studies of structure formation, which allow only for gravitational interactions, predict the distribution of galaxies and clusters of galaxies in excellent agreement with observations.  Earlier studies had indicated that not only the distribution of halos but also their density profiles depend on the underlying cosmological model \\citep{quinn:86,frenk:88,dubinski:91,crone:94}. However, subsequent numerical investigations revealed that the halo density profiles are almost universal and depend neither on the mass of the halo \\citep{navarro:96,navarro:97}, the initial fluctuation spectrum, nor the underlying cosmological model. This similarity of dark matter halos is complemented by the discovery of a universal angular momentum profile \\citep{bullock:01}.   With the dramatically increased resolution of simulations the properties of the profiles can now be studied over several orders of magnitude in radius. This allows a more detailed study of the central slope of the density profile. \\citet{navarro:96,navarro:97} first gave an analytic  approximation  for the density profiles. They obtained $\\rho \\propto r^{-n}$, where the slope depends on the radius according to $n(r) = (1+3r/r_s)/(1+r/r_s)$, $r_s$ denotes the scale radius of a given halo. This NFW-profile implies an asymptotic slope of $n = 1$ for the central density profile. Other numerical studies produced a significantly steeper inner slope, $n \\approx 1.5$, \\citep{moore:98,moore:99,ghigna:98,ghigna:00,fukushige:01,klypin:01}.  More recently \\citet{power:02} found no asymptotic slope at all but a continuously decreasing slope with values down to $\\approx 1.2$ approaching the innermost radius resolved. In order to determine the inner slope of dark matter halos in numerical studies more reliably, more efforts are needed to increase the resolution of simulations.   In galaxies the motion of the stars is governed by the dark matter halo. Measuring rotation curves in the very center of a galaxy allows in principle the determination of the central profile of the underlying density distribution. Low surface brightness galaxies are almost unbiased by any gas and are thus promising candidates for the determination of the innermost density profile. A `cuspy', $n \\gtrsim 1$, dark matter core is ruled out by many studies \\citep{gaugh:98,blok:01} including models of the inner rotation curve of the Milky Way \\citep{binney:01}. \\citet{salucci:00} have inferred a constant central density from the kinematics of a large sample of spiral galaxies. However, some recent investigations of dwarf galaxies suggest that an innermost slope of $n \\approx 1$ is consistent with rotation curve data \\citep{bosch:01}.  Lensing studies of clusters of galaxies indicate that there is a shallow core, with some preference of an isothermal profile in the very center \\citep{seitz:98,gavazzi:03}. Thus, present observations favor an inner slope of $n \\approx 1$ at most and in many cases a considerably shallower one.  Some mechanisms have been proposed which are able to erase a central cusp of dark matter halos.  \\citet{el-zant:01}, e.g., suggest that a core may be eliminated by dynamical friction of an initially clumpy gas distribution. Alternatively the occurrence of bar instabilities in the center have been discussed \\citep{weinberg:02,sellwood:03}. \\citet{dekel:03} have argued that a central cusp can be avoided by puffing up infalling halos, which could occur due to baryonic feedback processes.   The uncertainty of the central slopes reflects that up to now dark matter halo profiles have only been obtained on a empirical basis. Efforts are made to calculate the profiles analytically and semi-numerical models arew used to calculate the profile using the growth rate of halos. Since the mass accretion depends on the environment those models may indicate that the profile should also vary with the supposed cosmology \\citep{syer:98,nusser:99,lokas:00}. An alternative explanation was given by \\citet{taylor:01,navarro:01}. They argued that recurrent merging results in a phase-space density profile which decreases according to a power-law. As a result they found that the central slope should be as low as $n=0.75$. One may summarize the results about dark matter density profiles by stating the open questions. Why do halos show a profile close to the NFW-profile, why are the profiles almost universal, and what is the central slope?    In this paper we investigate the velocity dispersion profile of dark matter halos that were simulated with high resolution and using two independent codes. In Sec.~\\ref{sec-vdpr} we present a suitable formula that approximates the dispersion profile as a function of the potential. The probability distribution of the particle velocities within shells of different radii is investigated in Sec.~\\ref{sec-pdf}. This allows us to interpret the individual factors which are present in the suggested formula. The analysis of the possible asymptotic behavior of the profiles in Sec.~\\ref{sec-asymp} demonstrates that the relation between velocity dispersion and potential is strongly restricted by physical requirements. Moreover, limits for the inner slope of the density profile can be set. For consistency we show in Sec.~\\ref{sec-jeans} that the numerical density profiles can be recoverd by integrating the basic equations.    \\begin{figure}[t] \\centering \\resizebox{1.02\\hsize}{!}{\\includegraphics[angle=0]{f1.eps}} \\caption{ The radial profiles with respect to different quantities for two halos (galaxy sized and cluster sized).  Upper panel: The density $\\rho/\\rho_0$ (open squares, scaling given by left-hand y-axis), the potential $\\sqrt{\\Phi}$ (crosses, scaling by right-hand y-axis), the radial velocity dispersion $\\sigma_r$ (solid squares, scaling by right-hand y-axis) and the anisotropy $\\beta$ (triangles, scaling by right-hand y-axis) are plotted. Potential and dispersion are given in units of $[10^3\\Unitkmpers]$ what allows us to compare immediately these quantities to the velocity p.d.f. Fig.~\\ref{fig-ehist}. The density is approximated by generalized NFW-profile Eq.~(\\ref{eq-nfw-profile}) (solid line) with parameters $r_s = 0.42 \\UnitMpc$, $\\delta_c = 1.01 \\times 10^4$, $n = 1.54$ and $r_s = 0.034 \\UnitMpc$, $\\delta_c = 7.1 \\times 10^4$, $n = 1.11$ for the cluster-sized halo Cl6 and the galaxy-sized halo Gal, respectively. The anisotropy is approximated by $\\beta = 0.27 * \\Phi / \\Phi_{\\rm{out}}$. The values of $\\Phi_{\\rm{out}}$ are taken from Tab.~\\ref{table-halos}. Lower panel: Predictions by the left hand side of the Jeans equation (\\ref{eq-Jeans}) including the anisotropy term (crosses) and excluding it (open squares) match the circular velocity (solid squares). } \\label{fig-profiles} \\end{figure} ",
        "conclusions": "The numerous numerical studies of the structure of dark matter halos indicate that perfectly virialized systems show presumably a universal density profile. The latter can be approximated reasonably well by the NFW-profile or a generalized version of it. Similarly we have introduced the relation between velocity dispersion and potential Eq.~(\\ref{eq-of-state}) (VDPR): it characterizes the profile by very few parameters. In contrast to the approximations of the density profiles the VDPR is implicit in the sense that it as a function of the potential instead of the radius. However, by introducing the VDPR we can put restrictions even on the inner asymptotic slope of, both, the velocity dispersion profile and the density profile. The data can be approximated very well by the newly introduced VDPR. Almost all halos clearly show a power-law for $\\sigma_r^2(\\Phi) \\propto \\Phi^\\kappa$ near the center, the best resolved halos do so over one order of magnitude. We suppose that the obtained exponent $\\kappa$ reflects also the inner asymptotic slope of $\\sigma_r^2(\\Phi)$. As a result, the innermost density profile power index $n$ is related to $\\kappa$ by $n=2\\kappa/(1+\\kappa)$. Using the mean value $\\kappa = 0.41$, determined from all presented simulations we get $n=0.58$. The obtained standard deviation $\\Delta \\kappa = 0.03$ permits only a negligible variation of the exponent $n$. Thus, our prediction for $n$ is below the theoretical prediction by \\citet{taylor:01}. They argued that the inner asymptotic slope is given  by $n=0.75$. Our result also indicates that the innermost slope of the density profiles is less steep than that obtained from the density profiles directly. This is not a discrepancy since the prediction $n=0.58$ is valid for scales smaller than those which are currently resolved by numerical simulations. Eq.~(\\ref{rho-exponent}) indicates that the density profile may steepened by an exponential term in a range where the velocity dispersion profile already follows a power-law.  The density profile only becomes a power-law if $(\\Phi/\\Phi_{\\rm{out}})^{1-\\kappa} \\ll a(1-\\kappa)$ is fulfilled. For instance, the impact of the exponential term becomes smaller than 10\\% if the potential is below $0.001\\times \\Phi_{\\rm out}$, using the obtained values for $\\kappa$ and $a$.     In a theoretical analysis of all possible asymptotes \\citet{muecket:03} found that $\\kappa$ should be in the range $0 < \\kappa < 1/3$. Therefore, the inner asymptotic slope of the density profile should be as low as  $n<0.5$. Our numerical results are close to this theoretical predictions. Possible reasons for the 20\\% higher value of $\\kappa$ could be too strong assumptions in the theoretical analysis or steepening of the central density profile during relaxation. In fact, it is still uncertain to which degree dark matter halos are actually relaxed. During the cosmological evolution halos merge frequently and show almost always substructures. Therefore, the present substructure may have impact on the resulting profiles. Nevertheless it is assumed that virialization has occured to a high degree and the halo configuration obeys the Jeans equation. The sample of halos we considered contains very relaxed halos, e.g. the halo  Gal, as well as halos with clear deviations from relaxation, e.g. the halo Cl6. \\citet{dekel:02} found that ongoing merging enforces a central cusp, consequently we may have overestimated the mean value of $\\kappa$ to some extent.     Our numerical results do not indicate any significant variation of the exponent $\\kappa$ towards the center, see Fig.~\\ref{fig-sig-phi}. Only in the case of the highest resolved halo Gal a trend of profile flattening can be noticed. Generally, a core with constant velocity dispersion might exist one a sub-resolution scale. The existence of such a core would weaken the restrictions for $\\kappa$. However, the apparent inner constancy of $\\kappa$ in the numerical simulations rather points to the existence of a power-law asymptote.    A further benefit of analyzing the VDPR is that this relation enables the determination of the leading processes responsible for the quite universal halo density profiles. The VDPR consists of two parts, the power-law $\\Phi^\\kappa$ and the relative potential  $\\Phi_{\\rm{out}} - \\Phi$. The latter can be understood as the gain in energy of the particles falling into the potential well. Utilizing the virial theorem $2T(r) = U(r)$, where $T(r)  = (1/2)\\sum\\nolimits_{r<r_i<r+\\Delta r}m_iv_i^2$ and $U(r)  =\\sum\\nolimits_{r<r_i<r+\\Delta r}m_i(\\Phi_{\\rm{out}}-\\Phi)$ and assuming isotropy for the velocity dispersion, $3\\sigma_r^2 = v_i^2$ we obtain $\\sigma_r^2 = (1/3)(\\Phi_{\\rm{out}}-\\Phi)$. The factor 1/3 is very close to the value $a=0.29$. Dark matter halos, therefore, can be considered to be locally virialized in their outer region. This leads to the observed, steep outer profiles. In the inner region of the halo the additional factor $\\Phi^\\kappa$ occurs, which may be attributed to the motion of relatively slow particles concentrated within the central potential well. Consequently, slow particles are more likely to be found in the center. This argument is scale-independent and should therefore result either in a power-law for $\\sigma_r^2(\\Phi)$ or in a dispersion core in the center of the halo.  The characteristic and universal density profiles are a result of these two mechanisms, which do neither depend on the history of the halo formation nor on the cosmological environment. For a perfectly relaxed system an universal density profile should be expected.   In summary, we have shown that the radial velocity dispersion can be approximated well by the VDPR as a function of the potential. We have provided physical arguments for the analytical form of the VDPR. The numerical data already obey the inner power-law over one order in magnitude. Considering the inner asymptotics, we have shown that the density profile becomes a power-law if the potential is below $\\approx 0.001\\times  \\Phi_{\\rm out}$. We conclude that the density profile can reach its asymptotic slope $n\\lesssim 0.58$ only at scales smaller than those presently resolved."
    },
    "0311/astro-ph0311426_arXiv.txt": {
        "abstract": "We investigate whether magnetic tension can define the masses of forming stars by holding up the subcritical envelope of a molecular cloud that suffers gravitational collapse of its supercritical core.  We perform an equilibrium analysis of the initial and final states assuming perfect field freezing, no rotation, isothermality, and a completely flattened configuration. The sheet geometry allows us to separate the magnetic tension into a qlevitation associated with the split monopole formed by the trapped flux of the central star and a suspension associated with curved field lines that thread the static pseudodisk and envelope of material external to the star. We find solutions where the eigenvalue for the stellar mass is a fixed multiple of the initial core mass of the cloud.  We verify the analytically derived result by an explicit numerical simulation of a closely related 3-D axisymmetric system.   However, with field freezing, the implied surface magnetic fields much exceed measured values for young stars.  If the pinch by the central split monopole were to be eliminated by magnetic reconnection, then magnetic suspension alone cannot keep the subcritical envelope (i.e., the entire model cloud) from falling onto the star. We argue that this answer has general validity, even if the initial state lacked any kind of symmetry, possessed rotation, and had a substantial level of turbulence. These findings strongly support a picture for the halt of infall that invokes dynamic levitation by YSO winds and jets, but the breakdown of ideal magnetohydrodynamics is required to allow the appearance in the problem of a rapidly rotating, centrifugally supported disk.  We use these results to calculate the initial mass function and star formation efficiency for the distributed and clustered modes of star formation. ",
        "introduction": "Despite the empirical evidence for the formation of stars under a wide variety of physical conditions in the cosmos, and despite more than a half-century of theoretical study, the question of what determines the masses of forming stars from gravitational collapse in large interstellar clouds remains an open one. Ideas range from hierarchical, opacity-limited, thermal fragmentation (Hoyle 1953, Lynden-Bell 1973, Rees 1976, Silk 1977, Bodenheimer 1978, Zinnecker 1984), to magnetically-limited fragmentation (Mestel 1965, 1985), to turbulence-induced fragmentation (Scalo 1985,1990; Larson 1995; Elmegreen \\& Efremov 1997; Truelove et al. 1998; Elmegreen 2002; Falgarone 2002; Klein et al. 2003), to wind-limited mass infall (Shu \\& Terebey 1984, Shu et al. 2000).  Shu (1977) pointed out the intrinsic difficulty of defining starlike masses in a large, unmagnetized, isothermal cloud that starts in equilibrium or near-equilibrium and is not bounded by an artificial surface pressure.  He commented on the difficulty of fragmentation in such an environment, where the condition of equilibrium of the initial state and the steep velocity gradients of the subsequent inside-out collapse effectively prevent gravitational fragmentation of the type originally contemplated by Hoyle from taking place (see also the debate between Hunter 1967 and Layzer 1964). These comments were subsequently given added force in numerous numerical simulations (see, e.g., Tohline 1980, 1981; Truelove et al. 1998).  Mestel (1965, 1985; see also Nakano \\& Nakamura 1978) made a crucial distinction between clouds that have dimensionless mass-to-flux ratios greater than unity (supercritical -- capable of gravitational collapse), and less than unity (subcritical -- incapable of gravitational collapse).  He argued that spherical magnetized clouds which are marginally supercritical cannot fragment gravitationally because any spherical subpiece will be subcritical.  He speculated, however, that supercritical isothermal clouds that collapse coherently to highly flattened states can fragment into supercritical pieces, each with size comparable to the vertical scale height.   These pieces might correspond to stellar masses.  Shu \\& Li (1997) cast doubt on this scenario in the case when field freezing may be assumed and the initial state corresponds again to a state of initial force balance with a significant density and field stratification. Self-similar, inside-out, collapses in such circumstances do indeed show the anticipated flattening, but the dynamically infalling pseudo-disks that result exhibit no tendency to fragment gravitationally.  They are prone to other kinds of numerical and perhaps physical instabilities in the presence of non-ideal magnetohydrodynamic (MHD) effects (Allen, Shu, \\& Li 2003; Allen, Li, \\& Shu 2003).  Indeed, to our knowledge, no reliable numerical simulation of initial non-uniform states of isothermal equilibria, with or without magnetization, has ever been able to demonstrate fragmentation in the subsequent collapse. \\footnote{If the cloud mass is higher than the equilibrium value and starts with a nearly homogeneous density distribution, so as to contain either initially or shortly thereafter more than one Jeans mass, then fragmentation is possible, as shown, for example, by Bodenheimer et al. (2000) or Matsumoto \\& Hanawa (2003).  But such initial states require explication how they arose, since their evolutionary time scales are generally short in comparison with the probable ages of the parent molecular clouds.}  Truelove et al. (1998) and Klein et al. (2003) find fragmentation possible in unmagnetized clouds that start with substantial supersonic turbulence.  However, the entire self-gravitating cloud cannot resist being turned into stars in the presence of strong shock dissipation (see, e.g., Klessen 2001). This then introduces a problem why such clouds are around today in a universe whose age is many dynamical crossings of molecular clouds (Zuckerman \\& Evans 1974).  Scenarios have been proposed where the star formation occurring in molecular clouds on all scales occupies only one dynamical crossing time before being dispersed (Elmegreen 2000, Hartmann 2003). But given the known overall inefficiency of star formation, these proposals then need to elucidate how the small fraction of cloud matter that came to have star-formation capability arrived at the critical state, and how the vast bulk of the molecular cloud material avoids gravitational collapse and fragmentation.\\footnote{Lada \\& Lada (2003) estimate that less than 10\\% of the mass of a giant molecular cloud participates in star formation, and the star formation efficiency in even the densest regions rarely exceeds 10--30\\%.}  The problem of rapid turbulent dissipation persists even in the presence of magnetization (Heitsch, MacLow, \\& Klessen 2001; Ostriker, Stone, \\& Gammie 2001), unless the magnetization is so strong as to make the clouds subcritical, in which case neither runaway gravitational collapse nor gravitational fragmentation will occur, even if the turbulence dissipates completely.  In any case, in prototypical low-mass cores of molecular clouds that form sunlike stars, the level of turbulence is low (Myers 1995, Evans 1999), too subsonic to satisfy the supersonic conditions found necessary by Truelove et al. (1998) and Klein et al. (2003) for turbulent fragmentation.  The theoretical solution for making the vast bulk of molecular cloud material have a low efficiency for star formation is therefore simple, at least in principle: assume that the bulk of molecular clouds,i.e., their envelopes, are subcritical (Shu, Adams, \\& Lizano 1987). The trick to gravitational collapse (and fragmentation) is then to invent mechanisms to get pieces (the cores) that are supercritical. Lizano \\& Shu (1989, see also Shu et al. 1999) outlined a bimodal process to accomplish this task: with ambipolar diffusion producing small, quiescent, supercritical cores in the isolated (or distributed) mode of star formation, and cloud-cloud collisions along field lines producing large, turbulent, supercritical cores in the cluster mode of star formation.  Molecular-cloud turbulence could also play a role in accelerating the rate of star formation in the distributed mode, both in concentrating matter during the dissipation of turbulence (Myers \\& Lazarian 1998, Myers 1999) and in enhancing the effective speed of ambipolar diffusion by the effects of fluctuations (Fatuzzo \\& Adams 2002, Zweibel 2002). We postpone further discussion of such effects until \\S 4.  Mouschovias (1976) made the intriguing suggestion that magnetic tension in a magnetized cloud might be able to hold up the envelope, preventing it from joining the mass of the cloud core that collapses to the center to form the central star. Despite initial enthusiasm for this idea to define stellar masses (e.g., Shu 1977), subsequent collapse calculations of supercritical clouds (or the parts of them that are supercritical) showed that the original proposal based on curvature arguments was not well-founded (Galli \\& Shu 1993a,b; Li \\& Shu 1997; Allen et al. 2003a,b).  However, a modified form of the question remains open for models where the cloud cores are supercritical but the envelopes are subcritical.  \\subsection{Goal of this Work}  The present paper is motivated by the desire to settle the last issue. We set the discussion in the context of a very specific theoretical model, but we shall argue later that the answers derived are generic and suggestive.  Although we start with the picture of a cloud as a well-defined entity, we actually have considerable sympathy for the view of the turbulence camp that such a concept has limited utility above the mass scales of molecular cloud cores (see, e.g., Elmegreen 1995 or Larson 1995). However, the discussion becomes more concrete if we save the debate on the role of interstellar turbulence until the end of the paper.  In our formal calculations, we are interested in initial mass distributions given by the singular isothermal sphere (SIS), \\begin{equation} \\rho(r) = {a^2\\over 2\\pi G r^2} , \\label{e-10} \\end{equation} where $a$ is the isothermal speed of sound in cosmic molecular gas, $G$ is the universal gravitational constant, and $r$ is the radial distance from the cloud center. Like Galli \\& Shu (1993a,b), we thread this cloud with a uniform magnetic field of strength $B_0$ in the $z$ direction.  Since the field is uniform, it exerts no MHD forces, and the force balance between self-gravity and isothermal pressure represented by the (unstable) equilibrium (\\ref{e-10}) is undisturbed.  In cylindrical coordinates $(\\varpi, \\varphi, z)$, the surface density corresponding to a vertical projection of the SIS volume density (\\ref{e-10}) along field lines to the equatorial plane $z=0$ is given by \\begin{equation} \\Sigma (\\varpi) \\equiv \\int_{-\\infty}^{\\infty} {a^2\\, dz\\over 2\\pi G(\\varpi^2+z^2)}= {a^2\\over 2G\\varpi} . \\label{e-9} \\end{equation} The corresponding mass enclosed in an infinite cylinder of radius $\\varpi$ is then \\begin{equation} M(\\varpi) = \\int_0^\\varpi \\Sigma(\\varpi) \\, 2\\pi \\varpi d\\varpi = {\\pi a^2 \\over G} \\varpi. \\label{e-8} \\end{equation} The magnetic flux contained in the same cylinder is \\begin{equation} \\Phi(\\varpi) = B_0 \\pi \\varpi^2 . \\label{e-7} \\end{equation} Thus, the differential mass-to-flux ratio reads \\begin{equation} {dM\\over d\\Phi} = {a^2 \\over 2 GB_0 \\varpi} = {\\pi a^4\\over 2G^2 B_0 M}. \\label{e-6} \\end{equation} Notice that the surface density $\\Sigma$ (projected mass per unit area) and the magnetic field $B_0$ (flux per unit area)  have the same ratio as the differential mass to flux, $dM/d\\Phi$.  Following Basu \\& Mouschovias (1994) and Shu \\& Li (1997), we nondimensionalize using $(2\\pi G^{1/2})^{-1}$ as the basic unit of mass-to-flux (see also Nakano \\& Nakamura 1978).  Thus, the dimensionless differential mass-to-flux ratio is given by \\begin{equation} \\lambda (\\varpi) \\equiv 2\\pi G^{1/2} {dM\\over d\\Phi} = {M_0\\over M}, \\label{e-5} \\end{equation} where \\begin{equation} M_0 \\equiv {\\pi^2 a^4\\over G^{3/2} B_0} , \\label{e-4} \\end{equation} is the fundamental mass scale in the problem. We denote the cylindrical radius $\\varpi$ corresponding to the critical value $\\lambda(\\varpi) = 1$ by $r_0$; it is computed from equation (\\ref{e-8}) when $M=M_0$ and $\\varpi=r_0$: \\begin{equation} r_0 = {\\pi a^2\\over G^{1/2} B_0} . \\label{e-3} \\end{equation} Because $M(\\varpi)$ scales directly as $\\varpi$ in our initial state, we may also equivalently write equation (\\ref{e-5}) as \\begin{equation} \\lambda (\\varpi) = {r_0\\over \\varpi}. \\label{e-2} \\end{equation} We heuristically refer to the regions with $\\varpi < r_0$, where the initial cloud is supercritical, $\\lambda > 1$, as the ``cloud core;'' and the regions with $\\varpi > r_0$, where the cloud is subcritical, $\\lambda < 1$, as the``cloud envelope.''  For $a=0.2$ km/s and $B_0=30 \\; \\mu$G believed to be typical of dense regions in cold molecular clouds, we have $r_0 = 0.05$ pc and $M_0= 1.5\\; M_\\odot$. These are suggestive values for the cores in the Taurus molecular cloud (Jijina, Myers, \\& Adams 1999, Evans 1999), as has been noted in other contexts by many authors. Low-mass cores in the most crowded regions of clustered star formation might begin to overlap, unless such highly pressured regions have, as likely, larger ambient values of $B_0$. Alternatively, the whole region of embedded cluster formation may be supercritical, and the magnetic difference between core and (common) envelope loses some of its distinction (see \\S 4.5). In any case, the mass scale (\\ref{e-4}) resembles a kind of magnetic Bonnor-Ebert mass (Ebert 1955, Bonnor 1956), with $B_0^2/8\\pi$ replacing the role of the external pressure to define a critical condition for gravitational collapse (see the discussion of Shu et al. 1999 concerning why such a concept only holds for the supercritical cores). But the analogy is not completely apt, for reasons to be made clear in this paper.  Galli \\& Shu (1993 a,b) showed that an initial state with frozen-in magnetic fields corresponding to the mass-to-flux distribution $dM/d\\Phi$ of the above configuration is unstable to inside-out collapse, with the formation of a point object at the center that grows in mass with time.  The corresponding flux trapped at the center fans out as a split monopole eventually to connect at large distances onto the straight and uniform field lines of the initial state.  Galli \\& Shu followed only the initial stages of the inside-out collapse involving the supercritical regions of the cloud, where the mass accumulation rate at the center has its SIS value, $\\dot M = 0.975 \\, a^3/G$ (Shu 1977), even though the infall now takes place through a pseudodisk rather than spherically symmetrically onto the center.  We wish to extend the problem considered by Galli \\& Shu (1993a,b) by asking the following questions. What is the ``final'' state of the collapse?  How much of the ``core'' or envelope will eventually end up inside the star?  If it is not the entire cloud, what is the multiple $m_*$ of the original supercritical core $M_0$ that becomes stellar material?  By what mechanism would the material beyond the amount $m_*M_0$ be prevented from falling into the central star? How is the transition from the initial state to final state made in time?  How sensitive are our answers to the assumption of perfect field freezing, i.e., to the preservation of the function $dM/d\\Phi$ from beginning to end?  \\subsection{Findings of this Paper}  We approach these questions by two kinds of calculations.  One is by numerical simulation of the time-dependent evolution.  This simulation is presented in \\S 3. The other approach is to attack directly the final-state equilibrium, with the value of $m_*$ to be obtained as an eigenvalue of the problem.  The mathematical formulation of the resultant problem as an integro-differential equation in a single variable and its numerical solution occupy \\S 2 of this paper.  Readers uninterested in this mathematical derivation or its counterpart in numerical simulation (\\S 3) may jump directly to \\S 4 from the end of \\S 1.  The direct attack for the final state is extremely informative and will be discussed first. The full 3-D (axisymmetric) problem of magnetostatic equilibrium with prescribed $dM/d\\Phi$ is an involved and difficult calculation (see Mouschovias 1976, Nakano 1979, Lizano \\& Shu 1989, Tomisaka, Ikeuchi, \\& Nakamura 1989), even without the complication of an eigenvalue search for a central point mass. We therefore attack a simplified version of the problem posed above.  This simplification is motivated by the expectation that the inner regions of the suspended cloud envelope will be highly flattened by the same anisotropic magnetic forces that give rise to the pseudodisks of the dynamical collapse calculations.  By adopting a gas pressure that acts only in the horizontal directions rather than isotropically in all three directions, even the initial state flattens completely (see below).  The calculation of magnetic forces simplify considerably in a completely flattened geometry (Basu \\& Mouschovias 1994). In particular, only tension forces remain, and they can be computed from the flattened distribution of currents that act as the source of the magnetic fields as ``action at a distance'' (Shu \\& Li 1997).  As an additional bonus, Newtonian gravity has the fortunate coincidence that the radial gravity of a singular isothermal sphere is identical to a completely-flattened singular isothermal disk (SID) if they have the same column density distributions. In other words, the self-gravity of a completely flattened distribution, \\begin{equation} \\rho(\\varpi, z) = \\Sigma(\\varpi) \\delta (z), \\label{e-1} \\end{equation} when $\\Sigma(\\varpi)$ is given by equation (\\ref{e-9}), can be exactly offset by the isothermal (2-D axisymmetric), negative gas-pressure gradient, $-a^2 d\\Sigma/d\\varpi$, of this SID. Thus, the SID can also be threaded by an initially uniform, vertical magnetic field of strength $B_0$ and remain in (unstable) equilibrium.  Such a SID is no longer magnetized {\\it isopedically} in the nomenclature of Li \\& Shu (1997), and it will undergo inside-out gravitational collapse in a non-self-similar manner.  Nevertheless, we shall make good use of the remaining correspondences between the non-self-similar, axisymmetric, 2-D and 3-D problems in what follows. In particular, written in terms of the variables defined previously, the local value of the dimensionless mass-to-flux ratio of the completely flattened configuration, \\begin{equation} {2\\pi G^{1/2}\\Sigma(\\varpi)\\over B_0} = \\left( {r_0 \\over \\varpi}\\right), \\label{e0} \\end{equation} recovers our identification that $\\varpi = r_0$ marks the cylindrical radius where our model SID makes a transition from being supercritical (cloud core) to subcritical (cloud envelope).  For $B_0 = 30 \\; \\mu$G, the transition between core and envelope is made at a surface density equal to $B_0/2\\pi G^{1/2} = 1.8\\times 10^{-2}$ g cm$^{-2}$, which corresponds to about 4 magnitudes of visual extinction, a suggestive value from the points of view of observations (Blitz \\&  Williams 1999) and magnetically limited star formation (McKee 1989).  The gravity and magnetic forces will be different for the intermediate regions in the time-dependent collapse of the axisymmetric 2-D and 3-D problems.  Near the origin, the gravity and magnetic forces are dominated by the central object, which is a point mass and a split monopole in both calculations. Because the pseudodisk is very thin in its innermost regions even for the formal 3-D problem, it may not be a surprise to find that the eigenvalue $m_*$ for the final state, assuming field freezing, is the same for both problems: $m_* = \\surd 2$ (see below). We should add the immediate caveat that the answer $m_* = \\surd 2$ is an {\\it exact} result of the 2-D axisymmetric calculation, whereas it is less precisely determined in the 3-D calculation.  But because the analysis shows that this result depends only on what happens close to the central source, and not at all on what happens in the outer envelope, we have reason to believe that the result holds accurately for both configurations.  The mechanism of holding up the envelope is counter-intuitive, as hinted upon by the description given just now that the eigenvalue $m_*$ for the mass of the central object is determined entirely by what happens near the star. Instead of the envelope being suspended by magnetic tension working against the gravity of the cloud core and star, the envelope is being levitated by the split monopole pushing against the background magnetic field in the vacuum regions above and below the pseudodisk.  However, in order to hold off the inflow by this process, the central star would need to trap fields of $\\sim 10$ megagauss. This value is a factor of about 5000 times larger than the fields measured in low-mass pre-main-sequence stars (Basri, Marcy, \\& Valenti 1992; Johns-Krull, Valenti, \\& Koresko 1999). Evidently, through flux leakage to the surroundings, or through flux destruction via magnetic reconnection or anti-dynamo action, young stars destroy much if not all of the interstellar fields brought in by the mass infall.  In \\S 4, we shall comment on the implications for the overall problem by these various possibilities.  The tentative answer to the question posed in the title of this paper is, therefore, an equivocal ``no.''  The equivocation arises because the time-dependent simulations show, as was intuitively expected before we began the actual calculations, a significant reduction in the rate $\\dot M$ of mass infall onto the central source as the outwardly propagating wave of infall expands into the subcritical envelope.  Although the rate can be reduced all the way to zero only with the help of the magnetic push extended by the central split monopole, the mechanism for the reduction resides partially in the strong fields (relative to gas pressure) in the cloud envelope and not only in the magnetic levitation provided by the central source.  The two effects interact in a complex fashion to produce the final result. Indeed, before the paper ends, we will have revised our answer to a qualified ``yes!''  In this regard, the reader should not be fooled by the usual expression for magnetic tension, $({\\bf B}\\cdot \\nabla){\\bf B}/4\\pi$, into thinking that this force has a purely local origin.  The local expression is useful only if we have other means to compute the field ${\\bf B}$ and its derivative along field lines. Those other means, i.e., the equations of magnetohydrodynamics, include effects that mimic ``action at a distance,'' particularly in the near-vacuum conditions above and below the pseudodisk and in the outer cloud envelope where Alfv\\'en waves can propagate almost with infinite speed relative to $a$. This ``action at a distance'' brings magnetic influences from throughout the system, including the origin and the outermost regions.  Although we reject magnetic levitation by a trapped split monopole at the center as a viable practical mechanism in the long run, its role in halting infall can be effectively replaced by the magnetized winds and jets seen in actual protostars (Lada 1985, Bachiller 1996, Reipurth \\& Bally 2001, Shang et al. 2002). These are believed to arise because the realistic problem includes rotation in addition to magnetic fields (K\\\"onigl \\& Pudritz 2000, Shu et al. 2000). Even if the winds shoot purely out of the plane of an idealized 2-D axisymmetric calculation, they would still have an effect in blowing away the gas in a thin disk or pseudodisk because of their action on the fields from above and below that thread the flattened distribution of matter.  We defer this discussion, however, to the concluding remarks at the end of this paper. ",
        "conclusions": "\\subsection{A Powerful Conjecture} \\label{Conjecture}  Equations (\\ref{e28b})--(\\ref{e31}) give the most salient features of the formal solution for the ``initial'' and ``final'' states of the posed collapse problem in axisymmetric 2-D if we insist on strict field freezing. It is interesting that the central masses that result from the dimensional formula, \\begin{equation} M_* = \\surd 2 M_0 = \\pi^2 \\surd 2 {a^4\\over G^{3/2} B_0}, \\label{finalmass} \\end{equation} can be made to correspond to stellar masses only if we choose values of $a=(kT/\\mu)^{1/2}$ and $B_0$ that correspond to interstellar values, yet the physical determination of the mass eigenvalue occurs close to the star, where the local temperature $T$ and $B_z$ in the realistic situation are likely to differ considerably from their interstellar counterparts.  How sensitive, therefore, are our numerical conclusions to the exact details of the model adopted to do the actual calculations, in particular, to the assumptions of isothermality (same temperature in initial and final state) and field freezing (same $dM/d\\Phi$) in initial and final states?  Apart from numerical factors of order unity, we believe that our conclusions are generic. The following is a ``proof'' that does not quite live up to the standards of rigorous mathematics, but is powerful by astronomical standards.  If there is to be a final equilibrium involving a highly flattened, nonrotating, configuration of matter immediately outside the star, the final state of equilibrium still has to satisfy equation (\\ref{e11}).  For realistic temperatures in the pseudodisk (i.e., not approaching values comparable to that at the center of the star), the pressure gradient term in equation (\\ref{e11}) is negligible in comparison the stellar gravity term $-M_*$ throughout the inner pseudodisk. Similarly, as long as the mass contained in the pseudodisk from the stellar surface out to radius $\\varpi$ is less than the stellar mass $M_*$ multiplied by $\\lambda(\\varpi)$ times $\\bar \\lambda$ where $\\bar \\lambda$ is some appropriate average value in the pseudodisk, the contribution from the integral is negligible in comparison with the stellar gravity term.  (This is always true for $\\varpi$ near the inner edge $R_*^+$.)  Thus, the only way to satisfy equation (\\ref{e11}) is to have the split-monopole levitate the inner pseudodisk against the stellar gravity, i.e., the mass-to-flux ratio in the inner pseudodisk must have an inverse relationship to the stellar mass-to-flux ratio: \\begin{equation} \\lambda_*\\lambda(R_*^+) = 1 . \\label{eLL} \\end{equation} This elegant equation is the generalization of our eigenvalue determination for the final-state problem.  Suppose now that the initial cloud that gave rise to the collapsed final state has a flux-to-mass distribution given by \\begin{equation} {d\\Phi\\over dM} = 2\\pi G^{1/2} f^\\prime \\left({M\\over M_0}\\right), \\label{dPhidM} \\end{equation} where $M_0$ is no longer necessarily given by equation (\\ref{e-4}) but is equal, say, to the observed value where the cloud turns from being supercritical (core) to subcritical (envelope).  We suppose therefore that the function $f^\\prime(m)$ in equation (\\ref{dPhidM}) is a monotonically increasing function of its argument with the properties $f(0)=0$ and $f^\\prime(1)=1$. Assuming strict field freezing in the transition from the initial to final state, and carrying out an integration of equation (\\ref{dPhidM}), we find that the flux to mass ratio of the final star is given by \\begin{equation} {\\Phi_*\\over M_*} = 2\\pi G^{1/2} f\\left({M_*\\over M_0}\\right){M_0\\over M_*}. \\label{PstarMstar} \\end{equation} Denoting $m_* \\equiv M_*/M_0$, and identifying $\\lambda(R_*^+)= [f^\\prime (m_*)]^{-1}$ and $\\lambda_* = m_*/f(m_*)$, equation (\\ref{eLL}) obtains the general solution for the multiple $m_*$ of $M_0$ that will fall into the star as the positive root of the equation: \\begin{equation} f(m_*)f^\\prime(m_*)-m_* = 0. \\label{generaleigen} \\end{equation} For the special case, $f(m)=m^2/2$, we recover the eigenvalue solution $m_*=\\surd 2$ of \\S 2.  Equation (\\ref{generaleigen}) is a remarkable result.  It claims the following.  Start with a magnetized cloud of arbitrary shape and size (much bigger than a star) that has a dimensionless flux-to-mass distribution $f(m)$, where $m$ is the mass $M$ in units of the core mass $M_0$ (the supercritical part of the cloud) within magnetic flux $f$, which is $\\Phi$ in units of $2\\pi G^{1/2}M_0$.\\footnote{To construct the function $f(m)$ in the presence of turbulence and lack of any exact symmetries, start with the dimensional flux-to-mass ratio $d\\Phi/dM$ computed by projecting the mass along each (wiggly) flux tube onto some arbitrary mid-plane at some instant $t=0$.  Define the central flux tube $\\Phi=0$ to be the one that contains the minimum value of $d\\Phi/dM$.  Plot as a contour diagram the values of $d\\Phi/dM$ surrounding this central flux tube.  The function $\\Phi(M)$ is obtained by the two-dimensional integration of this contour diagram from $\\Phi=0$ to various isocontour levels of $d\\Phi/dM$ that contain flux $\\Phi$ and mass $M$.  The core mass $M_0$ is defined by the isocontour level where $f'(m) \\equiv 2\\pi G^{1/2} d\\Phi/dM = 1$, with $m\\equiv M/M_0$.  With $f'(m)$ known, a simple integration, with $f(0)=0$, yields $f(m)$. Notice that there is no need to define a length scale $r_0$ in such a formalism, nor do we need to introduce the concepts of a typical isothermal sound speed $a$ or background field strength $B_0$.  Finally, there are no restrictions on the nature of the initial velocity field of the fluid, laminar or turbulent. We have assumed that there is only one local minimum to $d\\Phi/dM$, i.e., that all isocontours close around a single point.  The case when a common (subcritical) envelope surrounds many local minima of $d\\Phi/dM$ remains to be explored (see \\S\\S 4.3 \\& 4.4).}  Then, without doing any calculations other than finding the root to equation (\\ref{generaleigen}), we are able to predict accurately the final mass of a star from the ideal MHD collapse of this cloud.  The initial conditions could contain an arbitrary amount of initial rotation, whose elimination by magnetic braking when field freezing holds -- see Allen et al. (2003b) -- will ensure the assumption of a single rotating star as the final state. The initial state also needs not be a configuration of equilibrium or near-equilibrium. It could contain sizable turbulent fluctuations in the density and velocity fields -- particularly in the magnetically subcritical envelope --  these will dissipate without having any effect on the final state. Sufficiently large turbulent fluctuations in the supercritical core (which are not seen in the cores that give rise to low-mass star formation) could give rise to temporary fragments, but in the presence of trapped stellar or pseudodisk flux which will stir the surrounding cloud envelope, the fragments are likely to lose all their orbital angular momentum and merge into a single final object.\\footnote{For very weakly magnetized initial states, or special geometric configurations, the time scale for the Alfv\\'en radiation of the system angular momentum to the envelope of the cloud may exceed relevant YSO time scales.}  Finally, the thermal history can be arbitrarily complex.  This will also have no effect on the final state, provided that radiation pressure on dust grains does not come into play (true for low-mass stars) and the temperature in the inner pseudodisk does not approach virial values.  \\subsection{Speculations}  For realistic astronomical application, the weakest link in the above argument is, of course, the assumption of field freezing.  Indeed, nearly fifty years ago, Mestel \\& Spitzer (1956) recognized that field freezing would predict stars at birth with surface magnetic fields of many megaguass, contrary to the observational evidence. They also pointed out (Mestel 1965, 1985; Spitzer 1968) that conservation of angular momentum in the gravitational collapse of stars from interstellar clouds would lead to absurd rotational velocities of the final objects, again contrary to the observational evidence. They proposed that a solution be found to link these two problems via magnetic braking and field slippage.  We now propose to link the two classical problems of magnetic flux loss and angular momentum redistribution to the third great problem of star formation: What processes define the masses of forming stars?  We even believe that the breakdown of field freezing plays a fundamental role in the fourth great problem of star formation: What determines whether gravitational collapse produces binary (or multiple) stars versus single stars surrounded by planetary systems?  (See, e.g., the discussion of Galli et al. 2000 and the simulations of Nakamura \\& Li 2003.)  The formal answer given by equation (\\ref{generaleigen}) to what determines stellar masses is magnetic levitation.  But levitation by a stellar split monopole leads to numerically unacceptable surface magnetic fields for the central star, as well as to the wrong prediction that magnetic fields should be {\\it weaker} in the regions adjacent to newly formed stars than the surrounding molecular cloud (see Figs. \\ref{fig:sigmaB} and \\ref{fig:sigmaB3} as well as Fish \\& Reid 2003). Is it possible that magnetic reconnection of the type contemplated by Galli \\& Shu (1993b, see also Mestel \\& Strittmatter 1967) completely eliminates the stellar monopole (except for some small seed fields needed to start stellar dynamo action)? After all, the strong sheet current corresponding to the reversal of field lines across the midplane of a split monopole will dissipate in the presence of finite resistivity.  The dissipation of the sheet current corresponds to the annihilation of field lines just above and below the midplane. The magnetic pressure of remaining field lines presses toward the midplane to replace the lines that have been annihilated.  They are annihilated in turn.  The result is an ever weaker split monopole, which continues to lose strength by sheet current dissipation until the sheet current and the split monopole are both gone.  (In practice, this dissipation would occur throughout the process of central mass and flux accumulation and not just in the final state.)  A vanishingly small split monopole, $\\lambda_*\\rightarrow \\infty$ in equation (\\ref{eLL}), requires $\\lambda(R_*^+) \\rightarrow 0$ if magnetic levitation is still to do its work. The only material so heavily magnetized outside the star (if magnetic flux were preserved even approximately in the rest of the system where the fields are not so severely pinched) comes from far out in the original model cloud. In other words, with no split monopole at the center to help it, magnetic suspension of the cloud envelope alone cannot withstand the very large gravity just above the surface of the star, and the infall would continue until the star has accumulated the entire molecular cloud (or stellar evolution intervenes).  But this possibility also contradicts the observational evidence. Therefore, imagine next that there is no electrical resistivity, but there is ambipolar diffusion.  Unlike electrical resistivity, ambipolar diffusion cannot destroy magnetic flux, but it can redistribute it. Can ambipolar diffusion (or C shocks, see Li \\& McKee 1996) transfer the flux from material before it became incorporated into the star and redistribute it over a much wider area of a static pseudodisk, so that the resultant fields are not such an embarrassment relative to observations? On a larger scale the levitation provided by such fields could serve to prevent the infall of the cloud envelope.  The difficulty of such a scheme is that there is only suspension to keep the static pseudodisk from falling into the star, and we proved in this paper that suspension without levitation cannot do such a job.  Keplerian rotation provides an excellent barrier to infall, so can the conversion of magnetized pseudodisk into a magnetized Keplerian disk provide the needed buttress to prevent the inward advance of the cloud envelope?  The question is especially cogent now that Allen et al. (2003b) have demonstrated that centrifugally supported disks cannot arise in realistically magnetized models of rotating gravitational collapse unless the assumption of field freezing is abandoned in the pseudodisks and disks of star formation (see, also, Krasnopolsky \\& K\\\"onigl 2002 for examples of formation of centrifugally supported disks in cases of low, prescribed, efficiency of magnetic braking in the presence of ambipolar diffusion). But strongly magnetized, rapidly rotating disks will develop magnetocentrifugally driven disk winds (Blandford \\& Payne 1982, K\\\"onigl \\& Pudritz 2000), which will provide {\\it dynamic rather than static} levitation.  Thus, ultimately, the realistic solution to the problem of obtaining stellar masses from the gravitational collapse of large interstellar clouds will end up invoking wind-limited mass infall (Shu \\& Terebey 1984). Moreover, numerical simulations of Keplerian disks threaded by such a network of open field lines show them to quickly fall into the innermost regions of the disk (Gammie \\& Balbus 1994, Miller \\& Stone 1997), concentrating all the trapped flux into a narrow annulus that makes the subsequent outflow indistinguishable from an X-wind.  X-winds will also arise in the alternative extreme scenario of all electrical resistivity and no ambipolar diffusion, if Keplerian disks appear outside accreting protostars (Shu et al. 1988, 1994, 2000).  In this case, although the stellar split monopole will eventually disappear because of field annihilation (eliminating the magnetic braking by long lever arms that prevents the formation of Keplerian disks in the simulations of Allen, Li, \\& Shu 2003), the operation of a stellar dynamo can replace such fields with those that have a dipole or higher multipole structure at the stellar surface (Mohanty \\& Shu 2003).  The interaction of these dynamo-driven fields with the surrounding accretion disk, and the subsequent opening of the stellar field lines (or the simple pressing of the remnant interstellar field lines against the magnetopause) can then create an X-wind.  The attractiveness of either proposal relative to that of static levitation is one of efficiency.  Newly formed stars do not need 10 megaguass fields to halt the inflow of gravitational collapse. Surface fields of a few kilogauss, when combined with the rapid rotation from an adjoining accretion disk, suffice to provide dynamic levitation (a physical blowing away) of cloud envelopes.  In the real situation, electrical resistivity and ambipolar diffusion are probably equally important in the disks and pseudodisks around young stars.  Since we have argued for X-winds as a likely outcome for the two extreme assumptions, we believe it is also the likely outcome in all intermediate cases.  \\subsection{From Isolated Star Formation to Distributed Star Formation}  Imagine a large magnetized cloud which is overall subcritical, and therefore probably fairly flattened, but which contains many dense pockets (cloud cores) that are supercritical. (We are supposing that the discussion of \\S \\ref{Conjecture} can be generalized to more than a single core.) Each of the cores can undergo inside-out collapse. If they later generate a magnetocentrifugally driven wind, they can each dynamically levitate their neighboring subcritical envelopes and prevent star formation from being anywhere near 100\\% efficient. A general inefficiency for forming stars is consistent with the observational evidence, except possibly for situations involving the bound-cluster mode of star formation. The supersonic turbulence and other complex factors in the subcritical common envelope play no dynamical role once the cores go into collapse and generate their winds, but the winds may play a crucial role in sustaining the observed turbulence and overall filamentary morphology of the cloud (Allen \\& Shu 2000). In turn, as discussed below, the turbulence in the general envelope may ultimately help set the mass distribution of molecular cloud cores {\\it before} the onset of gravitational collapse, and therefore the initial-mass-function of forming stars.  How did the cloud get so many supercritical cores in a general environment that is subcritical?  One answer proposed by Shu et al. (1987, see also Lizano \\& Shu 1989) is ambipolar diffusion acting as a background process in all molecular clouds.  Their name for this process was ``the isolated mode of star formation,'' but in the present generalization, we refer to it as ``the distributed mode of star formation.''  If ambipolar diffusion is the mechanism that condenses many cores from large cloud envelopes, it represents a kind of {\\it magnetic fragmentation}.  Whether this fragmentation occurs slowly or not relative to dynamical timescales may depend on the contractive and diffusive enhancements one ascribes to turbulence dissipation and fluctuations (Myers \\& Lazarian 1998, Myers 1999, Adams \\& Fatuzzo 2002, Zweibel 2002).  In any case, the first supercritical core to appear from the gravitational contraction of any dense pocket of gas and dust has, by definition, a vanishingly small mass $M_0$, since the transition from subcritical to supercritical is made gradually. Unless dynamically compressed, such a core is not immediately unstable to inside-out gravitational collapse; it has to evolve by further ambipolar diffusion (and perhaps turbulent decay) to reach a state of sufficient central concentration. In the numerical simulations performed to date of this process, the exact state of runaway collapse depends on the outer boundary conditions that are imposed (see, e.g., Nakano 1979; Lizano \\& Shu 1989; Tomisaka, Ikeuchi, \\& Nakamura 1990; Basu \\& Mouschovias 1994; Li 1998).  When those boundary conditions are applied at infinity, and the evolution is approximated as quasi-static, our group has argued for pivotal states that resemble singular isothermal toroids (Li \\& Shu 1996, Allen et al. 2003a), which are close relatives in their density distributions (but not flux distributions) to the magnetized SISs studied in this paper.  The details for cloud core structure are not important for the present application.  Recently, Motte, Andr\\'e \\& Neri (1998), Testi \\& Sargent (1998), and Motte et al. (2001) have found that the mass distributions of cloud cores as mapped by dust emission in turbulent molecular clouds resemble the empirical initial mass function (IMF) of newly formed stars (Salpeter 1955, Scalo 1986).  The weighting of cloud cores toward small (typically solarlike) masses instead of toward large masses (typically thousands of $M_\\odot$) as in cloud clumps (sometimes called {\\it dense} cores) is particularly striking.\\footnote{We reserve the name ``cloud core\" for the supercritical entities that will form individual stars or binaries. Observers define such cores differently than we do, based on observable contrast over the background rather than on the criticality of the mass-to-flux ratio. For realistic core-envelope models, the two definitions are related, since being supercritical allows the self-gravity of a region to produce large column density contrasts compared with their surroundings, which are probably subcritical or only weakly supercritical.} Lada \\& Lada (2003; see also Lefloch et al. 1998, and Sandell \\& Knee 2001) make the interesting suggestion that outflows may transform the mass-spectrum of cloud clumps, which resembles the mass distribution of embedded clusters, to the mass-spectrum of cloud cores, which resembles the mass distribution of stars. In what follows, we suggest a mechanistic process for performing such a task in a statistically invariant way, both in the distributed and the clustered modes of star formation, if outflows ultimately provide a major source of the turbulence present in star-forming clouds.  \\subsection{Initial Mass Function}  As an example of how the processes described in this paper can determine the stellar IMF (see also Silk 1995 and Adams \\& Fatuzzo 1996), we consider the contribution lent by turbulent velocities $v$ to the support of molecular clouds.  At a heuristic level, we imagine that $v$ can be added in quadrature to the isothermal sound speed $a$ in formulae that involve the gravitational constant $G$ (e.g., Stahler, Shu, \\& Taam 1980; Shu et al. 1987; McKee 2003; Pudritz 2003).  For even greater simplicity, we use $v$ below to mean the turbulent velocity when the turbulence is supersonic, and we substitute $a$ for $v$ when that turbulence is subsonic.   From IRAM and CSO observations of molecular clouds, Falgarone \\& Phillips (1996) find that CO line-profiles have high-velocity wings indicative of supersonic turbulence.  This interpretation is supported by the study of Falgarone (2002) who argues that the structure of molecular clouds is characterized by a unique power-law in mass, except possibly for a lower scale defined by cloud cores (Blitz \\& Williams 1999).  We adopt such a description and assume that the mass of magnetized molecular-cloud material with turbulent velocity between $v$ and $v+dv$ in 1-D is given by the power-law distribution, \\begin{equation} {\\cal M}(v)\\, dv \\propto v^{-s} \\, dv, \\label{turbvspectrum} \\end{equation} where $s$ is a positive number.  To avoid divergence of the integrated mass at very small $v$, we cut off the distribution (\\ref{turbvspectrum}) when $v$ becomes smaller than the thermal line-width $a$, noting that the important applications occur at $v\\lsim 1$ km s$^{-1}$.  To ensure that there is no confusion over what we mean by ${\\cal M}(v)\\, dv$, we emphasize that we have in mind a statistical sampling of a whole giant molecular cloud (GMC), or even of an ensemble of GMCs in an entire galaxy.  As a practical way to measure ${\\cal M}(v)\\, dv$, at least at relatively large $v$'s (a few km s$^{-1}$), imagine getting the integrated spectrum of a GMC in the line wings of some optically thin species, after we remove the effects of a systematic velocity field, and determining the mass with random velocity between $v$ and $v+dv$ along the line of sight.  Then our derivation below holds assuming that the same functional form ${\\cal M}(v)$ extrapolates to $v \\lsim 1$ km s$^{-1}$ and to the material in molecular cloud cores.  With this understanding, we further suppose that the core mass $M_0$ can be approximated as (see eq.~[\\ref{e-4}]): \\begin{equation} M_0 = m_0{v^4\\over G^{3/2} B_0}, \\label{m0M} \\end{equation} where the coefficient $m_0$ has a probabilistic distribution centered about $\\pi^2$ with a variance of a factor of a few (see \\S 4.6). Values of $m_0$ smaller than the typical $\\pi^2$ arise because turbulence, if compressive rather than expansive, can act to decrease as well as to increase the stability of some fiducial statistical equilibrium. Suppose now that the final stellar mass is some fixed fraction of $M_0$ (say, 1/3; see \\S 4.6): \\begin{equation} M_* \\sim {1\\over 3} M_0 \\sim {m_0 v^4\\over 3G^{3/2}B_0}. \\label{finalmscaling} \\end{equation} With a stellar mass-accumulation rate behaving as (2/3 times the infall rate) $\\dot M_* \\sim 2v^3/3G$ (cf. Shu et al. 1987), we get a characteristic formation time, \\begin{equation} t_{\\rm sf} \\equiv {M_*\\over \\dot M_*} \\sim {m_0 v\\over 2 G^{1/2}B_0} = {m_0^{1/2} r_0\\over 2v}, \\label{tsf} \\end{equation} where we have adopted equation (\\ref{e-3}) with $a$ and $\\pi$ replaced, respectively by $v$ and $\\surd m_0$ to estimate the core radius $r_0$.  The infall rate could exceed $v^3/G$ by a factor of a few if static magnetic fields contribute to core support before collapse, and if pivotal states contain nonzero contraction velocities (Allen et al. 2003a, b). The characteristic time $t_{\\rm sf}$ would then typically measure $\\sim 10^5$ yr, rather than a few times this value.  We further suppose that $B_0$ scales with $v$ as \\begin{equation} B_0 \\propto v^p, \\end{equation} where $p$ is again a positive number. The mass of stars formed during any fixed interval of time $t_{\\rm sf}$ with stellar masses between $M_*$ and $M_*+dM_*$ is now given by \\begin{equation} M_*{\\cal N}(M_*)\\, dM_* =\\varepsilon {\\cal M}(v)\\, dv \\propto M_*^{-(s+3-p)/(4-p)} \\, dM_*. \\label{IMF} \\end{equation} In equation (\\ref{IMF}), $\\varepsilon$ is an overall efficiency of turning molecular clouds into stars and is given by \\begin{equation} \\varepsilon = {F\\over 3}, \\label{SFE} \\end{equation} where $F$ is the cumulative fraction of molecular cloud mass in the region that has existed as star-forming cores.  Of the currently observable cores, a fraction $t_{\\rm sf}/(t_{\\rm sf}+t_{\\rm AD})$ exists as cores with actively accreting embedded protostars, and a complementary fraction $t_{\\rm AD}/(t_{\\rm sf}+t_{\\rm AD})$ exists as starless cores.  In the latter expressions, $t_{\\rm AD}$ is the time required for ambipolar diffusion and the decay of turbulence to transform a core with a critical level of magnetization to one in a pivotal state at the onset of protostar formation.  Except for the normalization factor $\\varepsilon$, where we have adopted the observer's definition of star formation efficiency as the mass of formed stars divided by the mass of total starting material in the region, equation (\\ref{IMF}) corresponds to Salpeter's (1955) IMF for low- and intermediate-mass stars if \\begin{equation} {s+3-p\\over4-p} = 4/3, \\qquad i.e., \\; {\\rm if} \\qquad s={7-p\\over 3}. \\end{equation} In the above, we have taken the liberty of approximating Salpeter's 1.35 by $4/3$.  One attractive combination satisfying the 4/3 rule is $p=1$ and $s=2$. This case corresponds to equal $t_{\\rm sf}$ for stars of all masses (eq. [\\ref{tsf}]) and Lorentzian line wings (eq. [\\ref{turbvspectrum}]).  From observation (Masson \\& Chernin 1992) and theory (Li \\& Shu 1996), it is known that bipolar outflows produce swept-up mass distributions that satisfy equation (\\ref{turbvspectrum}), with $s$ close to, but perhaps slightly smaller than 2, except at the highest velocities, where $s$ can be significantly larger than 2 (Lada \\& Fich 1996). It is not known, however, whether the same distribution applies after the gas has decelerated from velocities of order 10 km s$^{-1}$ to those of order 1 km s$^{-1}$. The assumption that $B_0$ scales linearly with $v$ ($p=1$) agrees with the estimate by Myers \\& Fuller (1993) that stars of all masses take $\\sim$ (a few times) $10^5$ yr to form once gravitational collapse starts.  A positive correlation will exist between $B_0$ and $v$ if regions with larger turbulent velocities require supercritical cores to reach higher mean densities before they become self-gravitating, resulting in a greater compression of the mean magnetic field threading the core. The Salpeter slope in equation (\\ref{IMF}) does not depend sensitively on the exact choice for $p$ if $s=2$.  In this paper, to match the properties of Taurus cores, we have adopted a normalization of $B_0 = 30 \\; \\mu$G when $v=a=0.2$ km s$^{-1}$.  With this normalization, the minimum stellar mass to result from equation (\\ref{finalmscaling}) is nominally 0.5 $M_\\odot$ if we take $m_0$ to have a typical value of $\\pi^2$.  In actual practice, 0.5 $M_\\odot$ is close to the flattening of the observed Galactic IMF for $M_*{\\cal N}(M_*)$ (Scalo 1986), but significant numbers of stars form in the Orion Trapezium region down to and below the hydrogen-burning limit $0.08$ $M_\\odot$ (cf. Fig 10 of Lada \\& Lada 2003).  In the current physical picture, stars can form with masses smaller than 0.5 $M_\\odot$, because there is a distribution in the value of $m_0 a^4/B_0$ for different regions.  In this context, it may be meaningful that Taurus appears to be deficient in brown dwarfs in comparison with the Orion Trapezium (Luhman 2000), consistent with our earlier suggestion that the effective value of $B_0$ may be statistically lower in the former region. But it could be that a significant number of brown dwarfs are initially formed as companions to normal stars and that there are larger numbers of ``free-floating'' brown dwarfs in Orion than in Taurus because of the greater role of binary disruption in the cluster environment (Kroupa 1995). It could also be that a limited form of Hoyle's (1955) fragmentation picture can still apply to turbulent, magnetized, cores if the eventual decoupling of magnetic fields occur sufficiently rapidly in realistic non-ideal MHD (Galli et al. 2000, Nakamura \\& Li 2003).  Apart from the different behaviors at low stellar masses, departures from a $-4/3$ power law in the observed IMF appear also at the highest stellar masses. We believe that an exponent for $M_*{\\cal N}(M_*)$ steeper than $-4/3$ can arise at large stellar masses because radiation pressure acting on dust grains aids YSO winds to reduce $M_*$ as a fraction of the initial core mass $M_0$ (Wolfire \\& Cassinelli 1987, Jijina \\& Adams 1996). If this interpretation is correct, the non-power-law features in the stellar IMF contain clues on how stars help to limit their own masses by blowing away cloud material that might have otherwise fallen gravitationally into the stars (cf. Adams \\& Fatuzzo 1996).  Figure 7 depicts pictorially the ideas of this subsection.  \\subsection{Embedded Cluster Formation}  Our derivation of the IMF formally seems to hold only for the distributed mode of starbirth, where competitive accretion does not occur.  But apart from the deficit of brown dwarfs, the IMF of YSOs does not appear to be different in Taurus compared to clusters or the general field (Kenyon \\& Hartmann 1995). Moreover, the sometimes evoked picture of protostars growing by Bondi-Hoyle accretion as they move freely in a background of more-or-less smooth clump gas may be flawed.  After all, stars do not appear half formed from nowhere; they probably have to grow by gravitational collapse and infall of small dense cores that are themselves self-gravitating substructures in the cloud clump. The gravitational potential associated with the cores forms local minima that are sharper, although perhaps less deep in absolute terms, than the large bowl that represents the smoothed gravitational potential of the clump and associated embedded cluster.  The tidal forces of the latter will not rip asunder the small cores if their mean densities are appreciably larger than the mean density of the background clump/cluster gas (which is the so-called Roche criterion).  In this regard, it is informative to note that the mean density inside a {\\it sphere} of radius $r_0$ of a uniformly magnetized SIS is given by \\begin{equation} \\bar \\rho_{\\rm core} = {2a^2 r_0/G\\over 4\\pi r_0^3/3} = {3B_0^2\\over 2\\pi^3 a^2}. \\end{equation} For $a=0.2$ km s$^{-1}$ and $B_0 = 30 \\; \\mu$G, $\\bar \\rho_{\\rm core}$ has a numerical value $1.1\\times 10^{-19}$ g cm$^{-3}$ $\\approx$ $1.6\\times 10^3$ $M_\\odot$ pc$^{-3}$.  This core density is larger than the average density of any observed clump forming an embedded cluster in the local GMCs of our Galaxy, except perhaps for the central regions of the very densest clumps (Lada \\& Lada 2003, Table 1). In the central regions of the densest clumps, individual small cores may merge, yielding the large cores that give rise to massive stars.  (We presume that such merger processes and large cores are part of the ``turbulent'' spectrum in eq.~[\\ref{turbvspectrum}].) However, the actual part of the core that forms a star is on average 9 times denser yet than the value ${\\bar \\rho}_{\\rm core}$, so forming low-mass stars can survive the neighboring appearance of luminous high-mass stars, although the remnant envelopes, pseudodisks, and disks not truncated by the tides of the central cluster/clump, may be photoevaporated away by the ultraviolet radiation field of the H II regions in which they are embedded (Johnstone, Hollenbach, \\& Bally 1998).   In contrast, the low-mass stars being born at the centers of the small cores at the peripheries of the clump, are in no greater danger of wandering away from those centers than we risk falling off the Earth because it orbits the Sun with a far deeper gravitational potential.  We do need to worry about the disruption of the cores from hydrodynamic effects as they orbit inside the clump/cluster.  However, such orbital effects are probably mitigated by the magnetization of the clump and cores.  This magnetization {\\it dilutes} the effect of gravitation (if both clump and core are supercritical); and the net forces can actually turn repulsive if the clump (common envelope of the cores) is subcritical.  Thus, the orbital motions of cores are likely to be sub-virial with orbital times to cross the clump in excess of the time $t_{\\rm sf}$ required to form individual stars.  The individual cores may hold together better as entities than envisaged in the standard picture of a highly turbulent and chaotic clump.  As long as the growing protostar (which feels only gravitational forces if it completely destroys its accreted interstellar flux by magnetic reconnection) is trapped by the local potential minimum of its parent core, the situation will resemble, to zeroth order, the case of isolated star formation. Sooner or later, such a protostar will develop an X-wind and begin to blow away its placental core.  The fact that the common envelope in a cluster environment may be mildly supercritical, instead of subcritical, may not make too big a difference on our (admittedly crude) estimate that only 1/3 of the initial supercritical core will end up falling onto the star (see below).  Indeed, in the present context, the 1/3 figure takes on added significance.  As an extreme, we may think of a turbulent clump which is supercritical everywhere as {\\it completely} packed with cores, with no surrounding ``common envelope'' (packed like ``kernels'' [of corn on the cob] in the nomenclature of P. C. Myers).  In that case, the efficiency of wind-limited star formation for the clump is the same as the average efficiency for a typical core, which we have taken to be 1/3 (see eq.~[\\ref{finalmscaling}]). The same result, $\\varepsilon = 1/3$, can be obtained from equation (~\\ref{SFE}) by setting $F=1$. It is then interesting to note that the maximum star-formation efficiency $\\varepsilon$ deduced by Lada \\& Lada (2003, Table 2) for embedded star clusters is indeed 33\\%.  To zeroth order, therefore, the clustered mode of star formation under present-day levels of magnetization may be considered as an extreme form of the distributed mode of star formation, which is itself a generalization of the isolated mode of star formation. Otherwise, it becomes a total coincidence that the IMFs found in many different environs of star formation, crowded or dispersed, are similar, with departures, if any, only at the lower end of the mass spectrum (Lada \\& Lada 2003).  \\subsection{Conclusion}  For the mass scale $M_0$ of cloud cores to play its part in setting stellar masses, cloud envelopes need to be not highly supercritical.  Indeed, for highly supercritical (e.g., unmagnetized) and turbulent clouds it is difficult even to define exactly what one means by a ``core'' or ``clump'' except by Roche-lobe or ``sphere of influence'' arguments. Moreover, without outflows to halt infall, artificial numerical means to pump in local turbulence (Klessen 2000) in combination, perhaps, with the ``Jeans swindle'' (Gammie 2003) is needed to keep the star-formation efficiency of supercritical regions that are overall gravitationally bound from approaching unity, contrary to the observational evidence. To understand the overall low efficiency of star formation, it would be best if the bulk of molecular clouds, except for their cores, were subcritical or only marginally critical. Unfortunately, observers have found no regions of molecular clouds that are definitely subcritical (Crutcher \\& Troland 2000).  This may be because they are biased to looking at regions where star formation has already taken place, and such regions are necessarily supercritical.  Allowing for various projection corrections, Shu et al. (1999) argued that most Zeeman measurements to date are consistent with the presence of a near-critical level of magnetization in molecular clouds. This leaves us optimistic about the fundamental correctness of the basic program begun by Mestel and Spitzer (1956) nearly fifty years ago.  Simple numerical estimates suffice to make our case here.  On average, the envelopes of GMCs possess 4 magnitudes of visual extinction (e.g., Blitz \\& Williams 1999).  With the usual caliberation that a number-column density of hydrogen $N_{\\rm H} = 1.9\\times 10^{21}$ cm$^{-2}$ yields 1 mag of visual extinction and that a helium atom accompanies each 10 hydrogen atoms, the observed extinction through envelopes corresponds to a mass-column density equal to \\begin{equation} N_{\\rm H} (1.4)m_{\\rm H} = 0.018 {\\rm g} \\; {\\rm cm}^{-2}. \\label{GMCenv} \\end{equation} Zeeman measurements yield the component of the magnetic field along a typical line of sight through a GMC as $B_{\\rm los}$ = 10 $\\mu$G (Crutcher \\& Troland 1999). This combination gives an apparent dimensionless ratio of magnetic field to mass-column density along the line of sight equal to \\begin{equation} {B_{\\rm los}\\over 2\\pi G^{1/2} N_{\\rm H} (1.4 m_{\\rm H})} = 0.34 \\label{los} \\end{equation} in the envelopes of GMCs.  For a flat sheet of surface density $\\Sigma$ threaded perpendicularly by a magnetic field $B_z$, viewed at an angle $\\theta$ with respect to the normal to the sheet, $B_{\\rm los} = B_z\\cos \\theta$ systematically underestimates the true field strength $B_z$, and $1.4 m_{\\rm H} N_{\\rm H} = \\Sigma/\\cos\\theta$ sytematically overestimates the true surface density $\\Sigma$.  Since $\\cos^2\\theta$ averaged over a hemisphere equals 1/3, we see that equation (\\ref{los}) is close to the expectation value for a sheet which is exactly critically magnetized: \\begin{equation} {B_z\\over 2\\pi G^{1/2}\\Sigma} = 1. \\label{envelope} \\end{equation}  Without bothering to argue about whether better numerical values or better corrections might make the right-hand side of equation (\\ref{envelope}) greater or less than unity, we will merely reiterate Shu et al.'s (1999) observation that it is probably no coincidence that modern-day GMCs have near-critical levels of magnetization.  Highly subcritical clouds are too non-self-gravitating to become molecular clouds; they are present in galaxies as H I clouds (Heiles 2003).  On the other hand, highly supercritical clouds are too vulnerable to gravitational collapse; they have long since disappeared to become the interiors of stars. McKee's important point that the ultraviolet photons associated with the interstellar radiation field suffice to keep regions with less than 4 mag of visual extinction too highly ionized to permit appreciable ambipolar diffusion then leads naturally to the idea of magnetically regulated star-formation (McKee 1989).  As a consequence, ambipolar diffusion can act to produce distinctly supercritical regions only in the cores of molecular clouds.  Many simulations (e.g., Lizano \\& Shu 1989, Basu \\& Mouschovias 1994) show that such cores of mass $M_0$, equatorial radius $r_0$, and mean magnetic field $\\bar B$ become susceptible to gravitational collapse once the dimensionless mass-to-flux ratio acquires a value of about 2: \\begin{equation} \\bar \\lambda \\equiv {2\\pi G^{1/2} M_0 \\over \\pi r_0^2 \\bar B} \\approx 2. \\label{lambda2} \\end{equation} The dilution of gravity $(1-\\bar \\lambda^{-2})=3/4$ (Shu \\& Li 1997) associated with this mean level of magnetization is not severe, and if the same core is supported against gravity largely by turbulent velocity $v$ (as measured in any single dimension), then virial equilibrium requires \\begin{equation} 2\\cdot {3\\over 2} M_0 v^2 \\approx {GM_0^2\\over r_0}. \\label{virialequil} \\end{equation} Solution of equations (\\ref{lambda2}) and (\\ref{virialequil}) yields \\begin{equation} M_0 \\approx {9v^4\\over G^{3/2} \\bar B}, \\qquad r_0 \\approx {3v^2\\over G^{1/2} \\bar B}. \\label{approxMandR} \\end{equation} When we identify $\\bar B$ with $B_0$, equation (\\ref{approxMandR}) for $M_0$ and $r_0$ differ from our expressions (\\ref{e-4}) and (\\ref{e-3}) for the same quantities only in the replacement of $\\pi$ by 3.  If turbulent cores are barely bound rather than fully virialized -- as implied by some suggestions that star-forming cores are produced as dynamical objects rather than as \"equilibrium\" states -- then the factor of 2 should be removed from the left-hand side of equation (\\ref{virialequil}). In fact, since fluctuations can exist on either side of virial equilibrium, we estimate that $M_0$ given by an expression for star-forming, bound cores such as equation (\\ref{m0M}) can have a coefficient $m_0$ that might vary by a factor as much as 4 about a central value of 3 or $\\pi$.  This crude estimate is depicted schematically in Figure 7 showing the typical distribution of $m_0$.  The following reasoning underlies the central idea behind our rough estimate that the final mass for the formed star is typically one-third of the core mass, $M_* = M_0/3$.  As long as the infall encompasses only supercritical regions, the inflow of matter through the envelope and pseudodisk may be too strong for an incipient X-wind to reverse the infall completely (see Fig. \\ref{fig:Mdot}).  However, when the outwardly expanding wave of infall begins to encompass the subcritical envelope (which happens at a dimensionless time $\\tau = 1$ in Fig.~\\ref{fig:flow}) and the mass infall rate declines significantly below its standard value $\\sim (1+H_0)a^3/G$, with $H_0$ a measure of the static magnetic contribution to cloud-core support (Li \\& Shu 1996; Allen et al. 2003a,b), the X-wind may begin to blast its way to ever wider opening angles and push on the cloud magnetic fields that thread through the infalling pseudodisk to reduce the accretion rate further.  The interaction may be quite complex, as the wind outflow rate itself may be a fraction ($\\sim$ 1/3) of the mass inflow rate (Shu et al. 1988, 1994, 2000).  Nevertheless, once outward motions are imparted to a significant portion of the surrounding cloud envelope, inertia and the ``action at a distance'' that is part of the effect of magnetic tension, may succeed in blowing away the infalling pseudodisk, even when the instantaneous wind rate from the central star drops to negligible values. Although detailed calculations are needed to quantify the estimates, which we are in the process of performing, comparison with the shape of the base of the outflow cavities in Class I sources suggests that the configuration at $\\tau = 2.5$ in Figure~\\ref{fig:flow} may be amenable to complete reversal of the inflow.  As an approximate guess, therefore, we suppose that 1/2 of the core mass $M_0$ will have dropped into the central regions before an X-wind can effectively halt the infall.  Of this 1/2, only 2/3 has ended up on the star (1/3 has come out as an X-wind); thus, $M_* = (2/3)(1/2)M_0 = M_0/3$. The entire description would then lend physical content to the observational definition that the transition from Class 0 to Class I marks the end of the phase of major infall in protostellar evolution (Andr\\'e, Ward-Thompson, \\& Barsony 1993).  In summary, the answer to the question posed by the title of our paper is a qualified ``yes'' --  yes, magnetic levitation and suspension may help to define the masses of forming stars -- but only in a dynamical context that involves rapid rotation and protostellar outflows.  Knowing the relationship between stellar mass to core mass for any single region then reduces the problem of the stellar IMF to an explanation for the mass distribution of the cores of molecular clouds."
    },
    "0311/astro-ph0311610_arXiv.txt": {
        "abstract": "We systematically reanalyzed all Gamma-Ray Burst (GRB) afterglow data published through the end of 2002, in an attempt to detect the predicted supernova light component and to gain statistical insight on its phenomenological properties. We fit the observed photometric light curves as the sum of an afterglow, an underlying host galaxy, and a supernova component. The latter is modeled using published multi-color light curves of SN 1998bw as a template. The total sample of afterglows with established redshifts contains 21 bursts (GRB 970228 - GRB 021211). For nine of these GRBs a weak supernova excess (scaled to SN 1998bw) was found, what makes this  to one of the first samples of high-$z$ core collapse supernovae. Among this sample are all bursts with redshifts less than $\\sim$0.7. These results strongly support the notion that in fact all afterglows of long-duration GRBs contain light from an associated supernova. A statistics of the physical parameters of these GRB-supernovae shows that SN 1998bw was at the bright end of its class, while it was not special with respect to its light curve shape. Finally, we have searched for a potential correlation of the supernova luminosities with the properties of the corresponding  bursts and optical afterglows, but we have not found such a relation. ",
        "introduction": "}  Observational and theoretical evidence suggest that the majority of GRB progenitors are stars at the endpoint in stellar evolution (e.g., Fryer, Woosley, \\& Hartmann 1999; Heger et al. 2003). Since the discovery of a near-by Type Ic supernova (SN 1998bw) in the error circle of the X-ray afterglow for GRB 980425 (Galama et al. 1998; Kulkarni et al. 1998), evidence is accumulating that core collapse supernovae are physically related to long-duration GRBs. The supernova picture is further supported by the observation that all GRB hosts are star-forming, and in some cases even star-bursting galaxies (e.g., Frail et al. 2002; Sokolov et al. 2001). Evidence for host extinction by cosmic dust in GRB afterglows (e.g., Castro-Tirado et al. 1999; Klose et al. 2000), and the discovery of an ensemble  of optically 'dark bursts' (for a recent discussion, see Fynbo et al. 2001; Klose et al. 2003; Lazzati, Covino, \\& Ghisellini 2002) also  is consistent with the picture that GRB progenitors are young, massive stars (Groot et al. 1998; Paczy\\'nski 1998). Furthermore, for several GRB afterglows X-ray data suggest a period of nucleosynthesis preceding or accompanying the burst (e.g.,  Antonelli et al. 2000; Lazzati, Campana, \\& Ghisellini 1999; M\\'esz\\'aros \\& Rees 2001). The angular distribution of the afterglows with respect to their hosts also favors a physical relation of young, massive stars to GRBs (Bloom, Kulkarni, \\& Djorgovski 2002).  As a natural consequence of a physical relation between the explosion of massive stars and GRBs supernova light should contribute to the afterglow flux, and even  dominate under favorable conditions. The most convincing example is GRB 030329 (Peterson \\& Price 2003) at $z$=0.1685 (Greiner et al. 2003a) with spectral confirmation of supernova light in its afterglow (Hjorth et al. 2003; Kawabata et al. 2003; Matheson et al. 2003; Stanek et al. 2003). In contrast to this unique spectroscopic evidence, several cases of photometric evidence for extra light in GRB afterglows have been reported, starting with GRB 980326 (Bloom et al. 1999). Various groups successfully fit SN 1998bw templates to explain these late-time bumps, the most convincing case being that of GRB 011121 (Bloom et al. 2002; Garnavich et al. 2003; Greiner et al. 2003b).  The goal of the present paper is to analyze the supernova bumps in GRB afterglow light curves using a systematic approach. While this was done also for several bursts by Dado, Dar, \\& de R\\'ujula (2002a; and references therein) with respect to a verification of their cannonball model (Dar \\& de R\\'ujula 2003), we tackle this issue in an independent and different way. First, from the numerical site, we have developed our own computational approach. This includes a numerical procedure to redshift SN 1998bw light curves (see \\S~\\ref{SNfits})  and to fit afterglow data within the context of the fireball model. Second, from the observational site, when necessary and possible we have performed late-time observations of some GRB host galaxies (\\S~2). A considerable part of the data we  have included in our study is based on observing runs where we have been involved. Additional data have been collected from the literature and checked for photometric consistency. Nearly two dozen afterglows  had sufficient data quality, and a known redshift, to search for a late-time bump in the light curve (\\S~3). Third, we concentrate our attention on a statistical analysis of the phenomenological parameters for this class of GRB-SNe (\\S~4). In this respect  our investigation goes beyond the approaches undertaken by others to  explain late-time bumps in individual afterglow light curves. ",
        "conclusions": "}  \\subsection{The luminosity distribution of the GRB-SNe}  When plotting the parameter $k$ deduced for the $R$ band in the observer frame, the width of the  luminosity  distribution of the class of GRB-SNe (Fig.~\\ref{histogram1}) is similar to what is observed for other classes of core collapse SNe (Richardson et al. 2002). The mean of $k$ in the $R$ band is 0.7 independent of whether or not we fix the stretch parameter at $s=1$, while our template SN 1998bw is at the bright end of the GRB-SN luminosity distribution. The latter conclusion is supported when we plot $k$ for the same wavelength region in the SN host frame, which is a better indicator of the spread in SN luminosities. Most of the GRB-SNe we explored have a data point for the photometric band centered around 395$\\pm$10 nm in the corresponding host frame (Table~\\ref{res}). The distribution of $k$ now indicates again that SN 1998bw is among the most luminous GRB-SNe. It also indicates that in fact there is no peak around $k$=0.8 (Fig.~\\ref{histogram1}), but we may so far have only sampled the bright part of the GRB-SN luminosity function. Some caution is of course required, given the partly large error bars of the $k$ factors, which are not shown in Figs.~\\ref{histogram1} and ~\\ref{histogram2}. While the conclusion that SN 1998bw was among the most luminous members of its class seems to be robust, the shape of the GRB-SNe luminosity function is still less well-determined. Extinction by interstellar dust in the host galaxies could in principle also affect these results, although only for GRB 010921 was a significant host extinction (\\gr 1 mag) reported (Price et al. 2003).  As we have outlined before, for redshifts $z$\\kr0.7 all GRB afterglows show evidence for an underlying late-time bump. Within our context this means that we trace a complete set of GRB-SNe, i.e., not only the brightest members of this class. The width of this GRB-SNe luminosity distribution in the photometric band centered around 395$\\pm$10 nm in the SN frame is $\\sim$1 to 1.5 magnitudes. This wavelength region is roughly placed in the $B$ band, so that we can compare the corresponding  luminosities with those of other Type Ib/c supernovae (Richardson et al. 2002), i.e., those class of SNe, which is believed to include the progenitors of GRBs. It turns out that the GRB-SNe do fit into a region between approximately $M_B=-19.5$ and $M_B=-18$ in figure 6 of Richardson et al., where no data on Type Ib/c SNe are known. If all GRB-SNe are indeed of type Ib/c this would favor the conclusion that the luminosity function of Type Ib/c SNe is rather described by a broad  Gaussian than by a bimodal distribution (figures 6 and 7 in Richardson et al.).  \\begin{figure} \\plotone{f4.eps} \\caption{The distribution of the luminosity parameter $k$ (Eq.~\\ref{ot}) as measured in the $R$ band in the observer frame (Table~\\ref{res}, with $s$ being a free parameter). GRB 980703 is not included here because the fitting procedure did  not find a solution in this case. Note that the histogram does not include the 1$\\sigma$ error bars of the individual $k$ factors, which are relatively large. \\label{histogram1}} \\end{figure}  The non-detection of a supernova bump in more than half of the investigated GRB afterglow light curves may be accounted for by  several reasons, like a relatively bright host, or a faint supernova. In particular, finding a SN bump for high-$z$ bursts is an observational challenge. For $z\\gr0.7$  and $k=1$, this peak magnitude exceeds $R_c=24$, which poses a major challenge for 3-m class telescopes, given the usually very limited amount of target of opportunity time for such observations. It is thus no surprise that only for three of the GRBs above $z=0.7$ a supernova component was found (GRB 980703, GRB 000911, GRB 021211), even though we can not rule out that the SN 'drop-out' is due to some evolutionary effect of the underlying burster population and their environment.  \\begin{figure} \\plotone{f5.eps} \\caption{The same as Fig.~\\ref{histogram1}, but for the  photometric band centered around 395$\\pm$10 nm in the SN host frame (Table~\\ref{res}, with $s$ being a free parameter). Not included here are GRB 980703, 010921, and 021211 since there is no data point in this wavelength range. \\label{histogram2}} \\end{figure}  \\subsection{The SupraNova model}  Vietri \\& Stella (2000) argued that GRBs are the result of delayed black hole formation, which implies that the core collapse and its subsequent supernova may significantly precede the burst. The delay could be of order months to years (Vietri \\& Stella 2000), or perhaps as short as hours (Woosley, Zhang, \\& Heger 2002). While constraining the latter possibility can not be accomplished with the data at hand, the longer time scales are easily constrained. For only two of the SN light curves the fit indeed improved if we allowed for a shift in time between the onset of the burst and the onset of the SN (GRB 990712, 011121). The offsets never exceeded 5 days, and were both negative and positive. However, the uncertainties in this parameter are large, due to the poorly sampled shape of the underlying supernova (e.g., Garnavich et al. 2003). The average shift is basically consistent with zero. Presumably, these shifts are due to an underlying correlation between luminosity and light curve shape, as observed in other types of supernovae (e.g., Candia et al. 2003; Stritzinger et al. 2002). This is just what the parameter $s$ takes into account. On the other hand, its clear that we have no information about this issue in those cases where we have not found evidence for a SN bump at late times. While this still leads open the possibility of two populations of bursters (collapsars \\& SupraNovae), we emphasize again that we find a late time bump in all afterglows with redshift $z<0.7$.  \\subsection{X-ray lines and supernova bumps}  The identification of late-time bumps in afterglow light curves with SN light would benefit from observations in  the X-ray band (for the cannonball model, e.g., Dado, Dar, \\& de R\\'ujula 2003a). If the X-ray lines seen in some afterglows (e.g., Reeves et al. 2002) have their origin in the circumburster medium (e.g., Lazzati et al. 2001)  and not in the exploding star (e.g., M\\'esz\\'aros \\& Rees 2001) this would be difficult to reconcile with the interpretation of a late-time bump with an underlying SN component. Unfortunately, the majority of bursts with high-resolution \\it XMM-Newton \\rm or \\it Chandra \\rm spectroscopic X-ray follow-up observations have no well-observed optical light curves. Among the afterglows with a detected optical late-time bump listed in Table~2 only for GRB 020405 such observations exist (Mirabal, Paerels, \\& Halpern 2003); no evidence for X-ray lines has been found there. \\it BeppoSAX \\rm observed the afterglow of GRB 970228 (Frontera et al.  1998) with comparable low energy resolution, and no X-ray lines have been reported.  Among those low-$z$ bursts with well-defined optical light curves and no evidence for a late-time bump in the data, only for GRB 970508 \\it BeppoSAX \\rm X-ray follow-up observations have been published (Piro et al. 1999). Evidence for an iron line was found.  Although one might add to the list of bursts with well-observed late-time light curves and additional spectral information in the X-ray band GRBs 990123,  000926, 010222, 011211, and 021004, the redshift of these bursts was \\gr 1.5, making it more or less hopeless to find an underlying SN component in the available data base (an upper limit for GRB 010222 is reported in Fig.~\\ref{luminosity1}b; Henden et al. 2004, in preparation).  While it is very interesting that neither for GRB 020405 (Mirabal et al. 2003) nor for GRB 030329 (Tiengo et al. 2003) lines have been found in X-ray spectra of their afterglows, at present we cannot confirm a possible anti-correlation between the occurrence of X-ray lines and the appearance of SN light in GRB afterglows. This important issue remains to be investigated in the \\it Swift \\rm era.  \\begin{figure} \\plotone{f6.eps} \\caption{The correlation between the afterglow parameters $\\alpha_1$ and $\\alpha_2$ (Eq.~\\ref{AG}) for all afterglows with a break in their light curves. The dotted line is the theoretical prediction in the simplest model ($\\Theta_{\\rm jet}$=const; M\\'esz\\'aros \\& Rees 1999). Open diamonds indicate the four afterglows with a break and a detected underlying late-time bump, i.e., a SN component (GRB 980703, 011121, 020405, 021211). Note that several afterglows listed in Table~\\ref{allgrbs} showed no evidence for a break in their light curves, so that they are not included in this figure. \\label{a2a1e}} \\end{figure}  \\subsection{SN properties vs. afterglow parameters}  Of particular interest is whether the burst and afterglow properties are to some degree related to the existence of an underlying SN component.  For this reason we have investigated if the deduced SN luminosity is correlated with the corresponding energy release in the gamma-ray band (as given in Bloom, Frail, \\& Kulkarni 2003). No such correlation was found. We have also checked whether the afterglow parameters $\\alpha_1$ (pre-break decay slope), $\\alpha_2$ (post-break decay slope), and $t_b$ (break time; Eq.~\\ref{AG}) from those GRBs with detected SN component are different from  those without such a component. Again, no correlation was apparent, even though one should keep in mind that the data base is still very small. For illustration, in Fig.~\\ref{a2a1e} we display the relation between the afterglow parameters $\\alpha_1$ and $\\alpha_2$ for all GRB afterglows we have investigated. While there is a 'forbidden region' with $\\alpha_2 \\kr \\alpha_1 + 3/4$ apparent in the data, with the border line representing the simplest jet model with no sideways expansion of the jet (M\\'esz\\'aros \\& Rees 1999), no bimodality in this distribution is visible. A tendency in our data that GRB afterglows with detected underlying SN component seem to prefer pre-break slopes $\\alpha_1>1$ should not be overinterpreted, since the $\\alpha_1<1$ sample includes several bursts with  redshifts $z>1.2$, where the photometric  detection of an underlying SN component is difficult. On the other hand, at least one selection effect does occur here. If a bright SN component is apparent in the data then the parameter $\\alpha_2$ of the afterglow light curve is usually much more difficult to quantify, because the late-time evolution of the genuine afterglow is less well-defined. This problem is well seen in the afterglow light curve of GRB 011121 (e.g., Greiner et al. 2003) and GRB 030329 (e.g., Lipkin et al. 2004)."
    },
    "0311/nucl-th0311054_arXiv.txt": {
        "abstract": "Several nuclear physics issues essential to understanding the $r$-process are discussed. These include validity of the waiting-point approximation, strength of closed neutron shells in neutron-rich nuclei far from stability, and effects of neutrino interaction with such nuclei. The needs for nuclear data in resolving these issues are emphasized. ",
        "introduction": "The dominant mechanism for producing nuclei beyond the Fe group is neutron capture. In particular, the rapid neutron capture process, or the $r$-process, is responsible for approximately half the solar abundances of nuclei with mass numbers $A>70$. The relative abundances of $r$-process nuclei in the solar system are shown in Figure 1. This solar $r$-process abundance pattern ($r$-pattern) has two prominent peaks at $A\\sim 130$ and 195, respectively, and provides an important basis for studying the $r$-process.  \\begin{figure}[tb] \\begin{minipage}{1.5cm} \\makebox[1cm]{} \\end{minipage} \\begin{minipage}{8cm} \\epsfig{file=fig1.ps,scale=0.7} \\end{minipage} \\caption{Solar $r$-process abundance pattern as derived in \\cite{Kapp89}.} \\end{figure}  By definition, neutron capture occurs much more rapidly than $\\beta$ decay during the $r$-process. With many neutrons available for capture by each seed nucleus, the nuclear flow quickly moves toward the neutron-drip line. At some point, the separation energy of the next neutron to be captured becomes so small that it will be quickly disintegrated by the photons in the $r$-process environment. At this so-called ``waiting point,'' the nuclear flow is characterized by a tug of war between neutron capture and photo-disintegration. Heavier nuclei cannot be produced until the waiting-point nucleus $\\beta$ decays to its daughter with a higher proton number. Then the tug of war repeats at the next waiting point. Through such a series of neutron capture and $\\beta$ decay, very neutron-rich nuclei far from stability are produced. Clearly, at a given proton number $Z$, the abundance is concentrated in the corresponding waiting-point nucleus. The more slowly this nucleus $\\beta$ decays, the more abundant it is. As nuclei with magic neutron numbers $N=82$ and 126 have (relatively) very slow $\\beta$-decay rates, abundance peaks are produced at these nuclei. After neutron capture ceases, these nuclei successively $\\beta$ decay to stability and give rise to the peaks at $A\\sim 130$ and 195 in the solar $r$-pattern (see Fig. 1). ",
        "conclusions": "In conclusion, a full understanding of the $r$-process depends on a lot of nuclear physics input. Neutron separation energies and capture cross sections and $\\beta$-decay rates for some key nuclei such as $^{195}$Tm are needed to establish the validity of the waiting-point approximation. Accurate mass measurements for extremely neutron-rich nuclei are needed to quantify the effects of neutron shell strength on production of the nuclei below the $A\\sim 130$ and 195 peaks. Finally, branching ratios of neutron emission and fission and fission yields for highly-excited neutron-rich nuclei are needed to understand the effects of supernova neutrinos on the $r$-pattern."
    },
    "0311/hep-th0311011_arXiv.txt": {
        "abstract": "We show that the cosmological constant may be reduced by thermal production of membranes by the cosmological horizon, analogous to a particle ``going over the top of the potential barrier\", rather than tunneling through it. The membranes are endowed with charge associated with the gauge invariance of an antisymmetric gauge potential. In this new process, the membrane collapses into a black hole, thus the net effect is to produce black holes out of the vacuum energy associated with the cosmological constant. We study here the corresponding Euclidean configurations (``thermalons\"), and calculate the probability for the process in the leading semiclassical approximation. ",
        "introduction": "One of the outstanding open problems of theoretical physics is to reconcile the very small observational bound on the cosmological constant, $\\Lambda$, with the very large values that standard high energy physics theory predicts for it \\cite{cosmoproblem}. This challenge has led to consider the cosmological constant as a dynamical variable, whose evolution is governed by equations of motion. In that context one has looked for mechanisms that would enable $\\Lambda$ to relax from a large initial value to a small one during the course of the evolution of the universe. The simplest context in which this idea may be analyzed, is through the introduction of an antisymmetric gauge potential, a three--form \\cite{3form} in four spacetime dimensions. The three--form couples to the gravitational field with a term proportional to the square of the field strength. In the absence of sources, the field strength is constant in space and time and provides a contribution to the cosmological term, which then becomes a constant of the motion rather than a universal constant.  Changes in the cosmological constant occur when one brings in sources for the three--form potential. These sources are two--dimensional membranes which sweep a three dimensional history during their evolution (``domain walls\"). The membranes carry charge associated with the gauge invariance of the three--form field, and they divide spacetime into two regions with different values of the cosmological term.  The membranes may be produced spontaneously in two physically different ways. One way is by tunneling trough a potential barrier as it happens in pair production in two--dimensional spacetime. The other, is by a thermal excitation of the vacuum analogous to going ``over the top\" of the potential barrier rather than tunneling through it.  The tunneling process was originally studied in \\cite{BT}, and was further explored in \\cite{all}. The purpose of this article is to study the spontaneous decay of the cosmological constant through the other process, namely, the production of membranes due to the thermal effects of the cosmological horizon.   It is useful to visualize the decay process in terms of  its simplest context, which is a particle in a one dimensional potential barrier, as recalled in Fig. \\ref{potential}. When the barrier is in a thermal environment, the particle can go from one side of the barrier to the other by ``climbing over the top\" rather than tunneling.  There is  a probability given by the Boltzmann factor $e^{-\\beta E}$ for the particle to be in a state of energy $E$. If $E$ is greater than the height of the barrier the particle will move from one side to the other even classically. The effect is optimized when the energy is just enough for the particle to be at the top of the barrier and roll down to the other side. In this case, the Boltzmann factor is as large as possible while still allowing for the process without quantum mechanical tunneling. It turns out that,  in the leading approximation,  the probability is given by the exponential of the Euclidean action evaluated on  an appropriate classical solution, just as for tunneling \\cite{linde}. In the case of tunneling, the classical solution is called an instanton and it is time dependent \\cite{Coleman}. In the present case, the classical solution  corresponds to the configuration in which the particle sits at the top of the barrier, and thus it is time independent. Since the solution is unstable, when slightly perturbed the particle will fall half of the time to the left side and half of the time to the right side. \\begin{figure}[h] \\centering \\includegraphics[width=10cm]{fig1} \\caption{The figure shows a one-dimensional potential barrier. If the particle is initially at the minimum $x=0$ it may end up in the other side of the barrier,  reaching the lower minimum  $x=x_T$ by two different mechanisms. It can either (i) tunnel through the potential quantum mechanically or (ii) it can jump over it by a thermal fluctuation when the barrier is in a thermal environment.} \\label{potential} \\end{figure}  In more complex situations, ``sitting at the top of the barrier\" is replaced by a ``time independent classical solution with one instability mode\"\\cite{linde}.  In the context of gauge theories such solutions appear in the analysis of violation of baryon--number conservation and have been called ``sphalerons\"\\cite{Manton}. In the present case we will use the name ``thermalon\", to emphasize that the static solutions will be intimately connected with the intrinsic thermal properties of event horizons in gravitational theory.  We show below that there exists a thermalon which reduces the cosmological constant through production of membranes in de Sitter space due to the thermal properties of the cosmological horizon. In this thermalon, a membrane  is emitted by the cosmological horizon and collapses forming a black hole. ",
        "conclusions": "In this paper, we have exhibited a new process through which the cosmological constant can decay.  At the same time, a black hole is created. The classical solution describing the process is an unstable (``ready-to-fall\") static solution which we have called ``cosmological thermalon\"; it is an analog of the sphaleron of gauge theories. Gravity and, in particular,  the thermal effects of the cosmological horizon is essential for the existence of the solution, which disappears in the flat space limit.  The net effect of the process is thus to transform non localized dark energy into localized dark matter, thus providing a possible link between the small present value of the cosmological constant and the observed lack of matter in the universe. Of course, the emergence in this way of (nearly) flat space as a natural endpoint of dynamical evolution is most intriguing in view of cosmological observations."
    },
    "0311/astro-ph0311298.txt": {
        "abstract": "Accumulating observational evidence for a number of radio galaxies suggests an association between their jets and regions of active star formation.  The standard picture is that shocks generated by the jet propagate through an inhomogeneous medium and trigger the collapse of overdense clouds,  which then become active star-forming regions. In this contribution, we report on recent hydrodynamic simulations of radiative shock-cloud interactions using two different cooling models: an equilibrium cooling-curve model assuming solar metallicities and a non-equilibrium chemistry model appropriate for primordial gas clouds. We consider a range of initial cloud densities and shock speeds in order to quantify the role of cooling in the evolution.  Our results indicate that for moderate cloud densities ($\\gtrsim 1$ cm$^{-3}$) and shock Mach numbers ($\\lesssim 20$), cooling processes can be highly efficient and result in more than 50\\% of the initial cloud mass cooling to below 100 K. We also use our results to estimate the final H$_2$ mass fraction for the simulations that use the non-equilibrium chemistry package.  This is an important measurement, since H$_2$ is the dominant coolant for a primordial gas cloud.  We find peak H$_2$ mass fractions of $\\gtrsim 10^{-2}$ and total H$_2$ mass fractions of $\\gtrsim 10^{-5}$ for the cloud gas, consistent with cosmological simulations of first star formation. Finally, we compare our results with the observations of jet-induced star formation in ``Minkowski's Object,'' a small irregular starburst system associated with a radio jet in the nearby cluster of galaxies Abell 194. We conclude that its morphology, star formation rate ($\\sim 0.3M_\\odot$ yr$^{-1}$) and stellar mass ($\\sim 1.2 \\times 10^7 M_\\odot$) can be explained by the interaction of a $\\sim 9 \\times 10^4$ km s$^{-1}$ jet with an ensemble of moderately dense ($\\sim 10$ cm$^{-3}$), warm ($10^4$ K) intergalactic clouds in the vicinity of its associated radio galaxy at the center of the galaxy cluster. ",
        "introduction": "\\label{sec:intro}  Combined optical and radio observations of a number of extragalactic radio jets in both low-luminosity (FRI) and high-luminosity (FRII) galaxies have revealed interesting correlations between the jets and apparent regions of active star formation.  One of the first objects demonstrated to show such a correlation was the nearest radio galaxy, Centaurus A \\citep[e.g.][]{Blanco75}.  Other examples have been found as the sensitivity and spatial resolution of radio and optical telescopes has improved.  In the case of Centaurus A, recent observations with the Hubble Space Telescope, when compared to radio images obtained with the Very Large Array, have confirmed that there are about half a dozen young ($< 15$ Myr) OB associations near filaments of ionized gas located between the radio jet and a large \\ion{H}{1} cloud \\citep{Mould00,rej02}. Another nearby example is ``Minkowski's Object,'' a peculiar small starburst system at the end of a radio jet emanating from the elliptical galaxy NGC 541, located near the center of the cluster of galaxies Abell~194 \\citep{vB85}. Star forming regions associated with radio sources have also been found in cooling flow clusters \\citep{McNamara02}.  Correlations between radio and optical emissions have also been observed in the so-called ``alignment effect'' in distant ($z>0.6$) radio galaxies \\citep{Chambers87,McCarthy87}. The best studied example here is 4C41.17 at $z=3.8$, where deep spectroscopic observations have shown that the bright, spatially extended, rest-frame UV continuum emission aligned with the radio axis of this galaxy is unpolarized and shows P Cygni-like features similar to those seen in star-forming galaxies \\citep{Dey97}.  Collectively, these observations are most compellingly explained by models in which shocks generated by the radio jet propagate through an inhomogeneous medium and trigger gravitational collapse in relatively overdense regions \\citep{Begelman89,DeYoung89,Rees89}. A detailed analysis of the jet-induced star formation in 4C41.17 has been presented by \\citet{Bicknell00}. In that object, Hubble Space Telescope images showed a bimodal optical continuum structure parallel to the radio jet \\citep{vB99}, strongly supporting the idea that the star formation was triggered by sideways shocks.  Although the analytic arguments in favor of this model are compelling, the complex nature of this problem, including nonlinear effects from higher-order coupling between hydrodynamics, cooling, and gravity, suggest that numerical studies may provide a more complete and detailed picture. Furthermore, numerical studies allow us to probe environments that may be observationally out of reach. For example, such studies may be able to parameterize the conditions of shock-induced star formation in the early universe, where direct observations of star-forming regions are not possible.  This could be important to understanding the role of jets in the feedback of active galactic nuclei (AGN) on their environment.  Numerical simulations also give us insight into the role of shock-induced star formation in other environments such as supernova shocks and cloud-cloud collisions.  The interaction of a strong shock with a single, non-radiative cloud has been the subject of many numerical studies.  A thorough analysis of this problem and a review of relevant literature is provided by \\citet{Klein94}.  A more recent study, considering the interaction of a strong shock with a system of clouds, is presented by \\citet{Poludnenko02}. For a non-radiative cloud the passing shock ultimately destroys the cloud within a few dynamical timescales.  Destruction results primarily from hydrodynamic instabilities at the interface between the cloud and the post-shock background gas.  Less well studied is the case of a strong shock interacting with a radiative cloud.  \\citet{Mellema02} showed that the effects of a shock on such a cloud are very different from the non-radiative case. Instead of re-expanding and quickly diffusing into the background gas, the compressed cloud instead breaks up into numerous dense, cold fragments.  These fragments survive for many dynamical timescales and are presumably the precursors to star formation.  In this work, we extend previous numerical studies by considering two primary cooling models.  The first model uses an equilibrium cooling curve with a tunable metallicity, set to solar in this work. As such, this model provides a reasonable approximation for cooling in metal-enriched clouds around low-redshift (FRI) radio galaxies. The second model solves the full non-equilibrium chemistry for a primordial gas mixture.  This model is very well suited to address jet-induced star formation in pristine clouds around high-redshift (FRII) galaxies.  With both models we explore a substantial subset of the parameter space over which shock-induced star formation may apply.  We proceed in \\S~\\ref{sec:shockcloud} by discussing the relevant physical processes, the timescales over which they act, and the limitations of the present study.  The numerical methods and models employed in the current work are discussed in \\S~\\ref{sec:method}, and the results of the models are presented in \\S~\\ref{sec:results}.  Further implications for star formation are discussed in \\S~\\ref{sec:implications}.  Our conclusions are briefly recapped in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have performed two-dimensional simulations of the evolution of radiatively cooling clouds subjected to strong shocks, as might arise from galactic jets.  The results of the models are summarized in Figures~\\ref{fig:equil} and \\ref{fig:chem}, which show the locations of the models in cloud density-shock velocity parameter space. Figure~\\ref{fig:equil} presents the results for models which include equilibrium cooling with solar abundances, while Figure~\\ref{fig:chem} shows the results for primordial clouds, which include non-equilibrium chemistry in which cooling at low temperatures is dominated by H$_2$ emission.  The numbers by each model indicate the percentages of the original gas that cools to below either 1000~K or 100~K. As can be seen from the figures, large fractions ($\\gtrsim 1/3$) of the clouds are able to cool below 1000~K.  Most of this cold gas ends up in very dense filaments, which have long dynamical lifetimes.  The final densities of these filaments are high enough in many cases that they are gravitationally unstable.  The subsequent gravitational collapse of these regions will likely result in a burst of star formation.  Our numerical results are consistent with previous work \\citep[e.g.][]{Mellema02} and have extended it by including new physics (chemistry and self-gravity) and exploring more of the parameter space of interest.  We applied our results to the specific example of Minkowski's Object. A number of important conclusions can be drawn. First, its peculiar morphology - bright star forming region orthogonal to the jet, and fainter filamentary features downstream from there - can be easily reproduced by our simulations (Fig \\ref{fig:jetcloud}). Second, the modest amount star formation required - $0.3 M_\\odot$ yr$^{-1}$ for the entire object, is also easily achieved for the plausible parameter space explored by our simulations.  Third, and most interestingly, we conclude that the star formation in Minkowski's Object could be induced by a moderate velocity jet ($9 \\times 10^4$ km s$^{-1}$) interacting with a collection of slightly overdense ($\\sim 10$ cm$^{-3}$), warm ($10^4$ K) clouds, i.e. it is NOT necessary to assume that this was an accidental collision between a jet and a preexisting gas-rich galaxy. This also suggests that the neutral hydrogen associated with Minkowski's Object ($3 \\times 10^8 M_\\odot$; W. van Breugel \\& J. van Gorkom 2003, private communication) may have cooled from the warm gas phase as a result of the radiative cooling triggered by the radio jet.  An update on the observations of Minkowski's Object will be presented in a forthcoming paper (S. D. Croft et al., in preparation).  Our results show that jets can drive radiative shocks in overdense regions in the intergalactic medium, resulting in star formation far away from the galaxies in which the jets originate.  This could be an important factor in understanding the feedback of AGN on their environment.  Our simulations may also be applied to radiative shocks and resulting star formation triggered by other mechanisms, such as in the wake of supernova blast waves \\citep{Preibisch02}, collisions of interstellar or intergalactic clouds \\citep{Smith80}, and ram pressure on gas-rich galaxies that fall into galaxy cluster atmospheres \\citep[e.g.][]{Gavazzi01}.  Our models can be used to infer the importance of shock-induced star formation in other regimes by simply adjusting the criteria used to develop Figure \\ref{fig:equil}.  Our models have confirmed the reality of the division between regions II and V (cooling and non-cooling) derived using approximate analytical methods, while the other region boundaries are set by physical limits for the clouds.  In future work, we shall extend our numerical simulations into regimes appropriate for other possible shock-cloud interaction scenarios."
    },
    "0311/astro-ph0311361_arXiv.txt": {
        "abstract": "I report recent results of numerical simulations designed to study the inner structure of Cold Dark Matter (CDM) halos. This work confirms the proposal of Navarro, Frenk \\& White (NFW) that the shape of $\\Lambda$CDM halo mass profiles differs strongly from a power law and is approximately independent of mass. The logarithmic slope of the spherically-averaged density profile, as measured by $\\beta=-d\\ln \\rho/d\\ln r$, decreases monotonically towards the center, becomes shallower than isothermal ($\\beta < 2$) inside a characteristic radius, $r_{-2}$, and shows no evidence for convergence to a well-defined asymptotic value ($\\beta_0$) at the center. The dark mass contained within the innermost radius resolved by the simulations places strong constraints on $\\beta_0$; cusps as steep as $r^{-1.5}$ are clearly ruled out in our highest resolution simulations. A profile where the radial dependence of the slope is a simple power law, $\\beta(r)\\propto r^{\\alpha}$, approximates the structure of halos better than the NFW profile and so may minimize errors when extrapolating our results inward to radii not yet reliably probed by numerical simulations. We compare the spherically-averaged circular velocity ($V_c$) profiles to the rotation curves of low surface brightness (LSB) galaxies in the samples of de Blok et al and Swaters et al. The $V_c$ profiles of simulated CDM halos are generally consistent with the rotation curves of LSB galaxies, but there are also some clearly discrepant cases. This disagreement has been interpreted as excluding the presence of cusps, but it may also just reflect the difference between circular velocity and rotation speed likely to arise in gaseous disks embedded within realistic triaxial halos. ",
        "introduction": "The pioneering efforts of Frenk et al (1988), Dubinski and Carlberg (1991), and Crone et al (1994), among others, led to the identification of a number of key features of the structure of dark matter halos assembled through hierarchical clustering. One important result of this early work concerns the remarkable similarity (``universality'') in the structure of dark matter halos of widely different mass. This was first proposed by Navarro, Frenk \\& White (1996, 1997; hereafter NFW), who suggested a simple fitting formula to describe the spherically-averaged density profiles of dark matter halos, \\begin{equation} \\rho(r)={\\rho_s \\over (r/r_s) (1+(r/r_s))^{2}}. \\label{eq:nfw} \\end{equation} A second result concerns the absence of a well defined central ``core'' of constant density in virialized CDM halos: the dark matter density grows apparently without bounds toward the center of the halo.  Subsequent work has generally confirmed these trends, but has also highlighted potentially important deviations from the NFW fitting formula. In particular, Fukushige \\& Makino (2001), as well as Moore and collaborators (Moore et al 1998, 1999), have reported that NFW fits to their simulated halos (of much higher mass and spatial resolution than the original NFW work) underestimate the dark matter density in the inner regions ($r<r_s$). These authors interpreted the disagreement as indicative of inner density ``cusps'' steeper than the NFW profile and advocated a simple modification to the NFW formula (referred to hereafter as the M99 profile) with asymptotic central slope $\\beta_0=1.5$.  The actual value of $\\beta_0$ is still being hotly debated in the literature (Jing \\& Suto 2000, Klypin et al 2001, Taylor \\& Navarro 2001, Power et al 2003, Fukushige, Kawai \\& Makino 2003, Hayashi et al 2003, Navarro et al 2003), but there is general consensus that CDM halos are indeed ``cuspy''. This has been recognized as an important result, since the rotation curves of many disk galaxies, and in particular of low surface brightness (LSB) systems, appear to indicate the presence of an extended region of constant dark matter density: a dark matter ``core'' (Flores \\& Primack 1994, Moore 1994).  Unfortunately, rotation curve constraints are strongest just where numerical simulations are least reliable.  Resolving CDM halos down to the kpc scales probed by the innermost points of rotation curves poses a significant computational challenge that has been met in very few of the simulations published to date.  Indeed, rotation curves are generally compared with extrapolations of the simulation data that rely heavily on the applicability and accuracy of fitting formulae such as the NFW profile to regions that may be compromised by numerical artifact.  The theoretical debate on the asymptotic central slope of the dark matter density profile, $\\beta_0$, has also led at times to unwarranted emphasis on the very inner region of the rotation curve datasets, rather than on an proper appraisal of the data over its full radial extent. De Block et al (2001), for example, attempt to derive constraints on $\\beta_0$ from the innermost few points of their rotation curves, and conclude that $\\beta_0\\sim 0$ for most galaxies in their sample. However, their analysis focuses on the regions most severely affected by non-circular motions, seeing, misalignments and slit offsets, and other effects that limit the accuracy of circular velocity estimates. It is perhaps not surprising, then, that other studies (van den Bosch et al 2000, Swaters et al 2003) have disputed the conclusiveness of these findings. The disagreement is compounded by the results of the latest cosmological N-body simulations (Power et al 2003, Hayashi et al 2003, Navarro et al 2003), which find scant evidence for a well defined value of $\\beta_0$ in simulated CDM halos. Given these difficulties, focusing the theoretical or observational analysis on $\\beta_0$ seems unwise. We shall try here to improve upon previous work by comparing directly the results of our simulations with the full radial extent of the rotation curves of LSB galaxies.   I report here results from a major computational effort pursued by an international collaboration bringing together computational cosmology groups in Victoria (Canada), Seattle (USA), Durham (UK), and Garching (Germany). We have now completed a series of about 20 simulations of $\\Lambda$CDM halos spanning roughly five decades in mass, from dwarf galaxies with circular velocities $V_c \\sim 50$ km s$^{-1}$ to galaxy cluster systems with $V_c\\sim 1500$ km s$^{-1}$. This series follows a thorough convergence study where the role of all relevant numerical parameters has been assessed in detail (Power et al 2003). As a result, we are able to identify in each simulated halo a minimum convergence radius ($r_{\\rm conv}$) beyond which the mass profile is reliably reproduced: circular velocity estimates are accurate to better than $10\\%$ for $r>r_{\\rm conv}$. Each of the simulated halos has several million particles within the virial radius, allowing for the mass profiles to be traced reliably down to within $1\\%$ of the virial radius. ",
        "conclusions": "Although it appears possible to reconcile dark matter cusps with LSBs by appealing to asphericity in CDM halos, it would be premature to argue that the problem has been fully solved. After all, it may come as no surprise that one is able to reproduce LSB rotation curves, given the number of extra ``free'' parameters afforded by relaxing the assumption of spherical symmetry. It is therefore important to build a more compelling case for this interpretation of rotation curve data, so as to render it falsifiable. We are in the process of identifying corroborating trends that may be used either to support or to rule out this interpretation, as the case may be. In particular, we would like to characterize better the $\\gamma>2$ systems: are such rotation curves in general asymmetric?  Is rotation around the minor axis expected? Can one verify the triaxial-halo interpretation in full velocity maps? In this sense, identifying a ``make or break'' prediction will be as important as the success of aspherical halos in reproducing the rich variety of shapes of LSB rotation curves. Only once this is accomplished shall we be able to conclude that LSB rotation curves do not preclude the presence of dark matter density cusps, freeing the CDM paradigm of one of its most vexing challenges on small scales."
    },
    "0311/astro-ph0311157_arXiv.txt": {
        "abstract": "We have performed systematic frequency analysis of the LMC Cepheids observed by OGLE project. Several new types of pulsation behaviour are identified, including triple-mode and amplitude-modulated double-mode pulsations. In $\\sim\\! 10$\\% of the first overtone Cepheids we find low amplitude secondary periodicities corresponding to nonradial modes. This is the first evidence for excitation of nonradial oscillations in Classical Cepheid variables. ",
        "introduction": "For decades Cepheids and RR~Lyrae stars have been considered primary examples of purely radial pulsators. This simple picture has changed in recent years, when nonradial modes have been detected in substantial fraction of RR~Lyrae variables in various stellar systems (Kov\\'acs 2002; Moskalik \\& Poretti 2003; Alcock et al. 2003). Motivated by these findings, we have undertaken a systematic search for nonradial Fpulsators among Pop.~I Cepheids of the Magellanic Clouds. In this paper we present our results for the LMC, with particular emphasis on the new types of pulsation behaviour detected during the survey. The analysis of the SMC Cepheids is currently underway and will be presented at a later date. ",
        "conclusions": ""
    },
    "0311/astro-ph0311227_arXiv.txt": {
        "abstract": "We summarize recent developments in modeling of mass outflows from accretion disks in AGN. We illustrate how the hydrodynamical and magnetohydrodynamical wind structures revealed in numerical simulations relate to observations. Important results from the numerical simulations include: no shielding from X-rays is needed to radiatively launch a wind from the inner disk; shielding is needed to accelerate the wind to high velocities. Disk winds launched magnetically seem to be too dense to be accelerated by line driving. It appears that we have an {\\it all or nothing} situation for radiation-driving: radiation driving has to be responsible for both launching and accelerating the wind or a wind launched by some other mechanism must have a very finely tuned density in order  to be  accelerated by radiation driving to velocities as high as measured in BAL QSOs. ",
        "introduction": "Mass outflows in active galactic nuclei (AGN) are fairly common. Broad absorption lines (BALs) in QSOs are perhaps the most compelling evidence for such outflows. The BALs are always blueshifted relative to the emission-line rest frame,\\footnote{A few SDSS BAL quasars have a redshifted and a blueshifted component; see Hall et al.\\ 2002.} indicating the presence of outflows from the active nucleus, with velocities in the ultraviolet as large as $66,000~\\rm km\\,s^{-1}$. Other evidence for AGN outflows include narrow absorption lines (NALs). NALs observed in ultraviolet (UV) observations of some QSOs can be blueshifted by as much as $\\sim~50,000~{\\rm km~s^{-1}}$ (Hamann et al.\\ 1997). NALs observed in UV observations of many Seyfert galaxies are blueshifted by several $100~{\\rm km~s^{-1}}$ (Crenshaw et al.\\ 1999). NALs can also be due to highly ionized species such as those observed in a high-resolution X-ray observation of the Seyfert galaxy NGC~5548 obtained by {\\em Chandra} (Kaastra et al.\\ 2000). We refer a reader to Arav, Shlosman, \\& Weymann (1999); Crenshaw, Kraemer, \\& George (2002); and these proceedings for more details  on observations of AGN outflows.  One of most plausible scenarios for AGN outflows is a wind from an accretion disk around a black hole. However, even if we agree that AGN outflows are indeed disk outflows, there are still many key questions  that await answers. Here we mention  just five: (i) From what radius on the disk do the outflows come? (ii) Do the outflows need shielding from X-rays to avoid overionization? (iii) What is the mechanism responsible for launching the outflows from the disk? (iv) What is the mechanism responsible for accelerating the outflows to their terminal velocities? (v) All AGN are believed to have accretion disks, why then do not all AGN exhibit outflows?  Our goal here is to review  various theoretical results and try to identify a physical disk wind model which is most promising in answering as many questions as possible concerning AGN flows with as few assumptions as possible.  Before we proceed, let us consider a specific question: is the terminal velocity of an outflow, $v_\\infty$, a good indicator of the radius, $r_l$, from which the outflow was launched? One can argue that the answer to this question is positive if the terminal velocity is correlated with the escape velocity, $v_{esc}$, from the location where the flow was launched. Such a correlation between $v_\\infty$ and $v_{esc}$ is one of the predictions of the radiation-driven wind theory (i.e., $v_\\infty \\approx~{\\rm a~few}~v_{esc}$, see Proga 1999 and references therein). ",
        "conclusions": ""
    },
    "0311/astro-ph0311011_arXiv.txt": {
        "abstract": "We consider effects of the MHD turbulence on the viscosity during the evolution of the thermal-viscous ionization instability in the standard $\\alpha$-accretion disks. We consider the possibility that the accretion onto a supermassive black hole proceeds through an outer standard accretion disk and inner, radiatively inefficient and advection dominated flow.  In this scenario we follow the time evolution of the accretion disk in which the viscosity parameter $\\alpha$ is constant throughout the whole instability cycle, as implied by the strength of MHD turbulence.  We conclude that the hydrogen ionization instability is a promising mechanism to explain the intermittent activity in AGN. ",
        "introduction": "The standard accretion disk (Shakura \\& Sunyaev 1973) is known to be subject to the thermal - viscous instability due to the partial hydrogen ionization (Meyer \\& Meyer-Hofmeister 1981; Smak 1984). As a result of this instability the disk cycles between the two states: a hot and mostly ionized state with a large local accretion rate and a cold, neutral state with a low accretion rate. This instability has originally been proposed to explain the large amplitude luminosity variations observed in cataclysmic variables (CV; Smak 1982; Meyer \\& Meyer-Hofmeister 1982).  It is also believed that the same mechanism is responsible for the eruptions in soft X-ray transients (SXT; see e.g. Cannizzo, Ghosh \\& Wheeler 1982; Dubus, Hameury \\& Lasota 2001; Lasota 2001 for review). The ionization instability was also shown to operate in the disk around a supermassive black hole in active galactic nuclei (AGN; Lin \\& Shields 1986; Clarke 1988; Mineshige \\& Shields 1990; Siemiginowska, Czerny \\& Kostyunin 1996).  The characteristic timescales of a cycle activity scale roughly with the mass of a compact object (Hatziminaoglou et al. 2001). Therefore, the observed timescale of cycle of order of years in binaries translates into thousands to millions of years in galaxies, which harbor a supermassive black hole.  The variability amplitudes and timescales observed in Galactic binary systems require the viscosity to be smaller in the quiescent disk than during the outburst.  Various {\\it ad-hoc} viscosity scaling laws have been used in the disk evolution models, including the simplest one with a constant viscosity, however usually $\\alpha$ is assumed to be 4-5 times larger during the outburst than in the quiescence (Cannizzo 1993). The value of $\\alpha_{\\rm hot}$ determines the timescale of an outburst, while $\\alpha_{\\rm cold}$ governs the separation between the subsequent outbursts. The difference between these two parameters determines the outburst amplitude.  The dependence of the viscosity parameter on the disk state is suggested to come out as a property of the magneto-rotational instability (MRI) mechanism (Balbus \\& Hawley 1991; 1998). Nonlinear development of this instability, which is assumed to be a primary source of the disk turbulence, is sensitive to the presence of resistive diffusion of magnetic field (Hawley, Gammie \\& Balbus 1996), measured by the magnetic Reynolds number $Re_{\\rm M}$. Gammie \\& Menou (1998) calculated the Reynolds number for the two states of the disk in CVs and show that it is low in the quiescence. It means that the magneto-hydrodynamical (MHD) turbulence dies away and the matter accumulates in the outer disk.  Menou \\& Quataert (2001) argued that the Reynolds number will not be low enough to suppress the MHD turbulence in the quiescent disk of AGN. Therefore the efficiency of angular momentum transport should be comparable in the hot and cold states ($\\alpha_{\\rm cold} \\approx \\alpha_{\\rm hot}$).  In such a case the outburst amplitude is dramatically reduced, and the ionization instability leads only to a small amplitude flickering (Siemiginowska et al. 1996).   The conclusion about the small amplitude of outbursts was obtained, however, for a geometrically thin, optically thick disk extending down to the marginally stable orbit. This assumption may not be correct. When the local accretion rate is low, the local radiative cooling may not be efficient (Rees et al. 1982; Begelman 1985). The innermost part of the disk may be replaced with an optically thin, hot and possibly two-temperature plasma. An example of such a solution, advection dominated accretion flow (ADAF) was introduced by Ichimaru (1977) and Narayan \\& Yi (1994). The development of the inner flow through the disk evaporation was recently discussed by several authors, starting with Meyer \\& Meyer-Hofmeister (1994).  In this paper we  first perform a self-consistency check of the $\\alpha$ description of the angular momentum transfer. Having justified $\\alpha_{\\rm cold}$ = $\\alpha_{\\rm hot}$ approach to AGN accretion disks, we consider the time evolution of the disk under the influence of the ionization instability. We also take into account the evaporation of the inner disk.  We define a parameter space where the instability zone (i.e. the partial hydrogen ionization zone) and the evaporation region overlap. The overlapping of these two regions strongly enhances the outburst amplitude due to the time evolution of the disk evaporation radius. We illustrate this behavior with exemplary lightcurves.  The structure of this article is as follows. In Section \\ref{sec:reynolds} we define the criterion for efficient development of the MHD turbulence in the accretion disk, taking into account its vertical structure. We confirm that indeed the bi-modal viscosity behavior is characteristic for Galactic binary systems. However, a single value of the viscosity parameter is appropriate in the case of AGN disks.  In Section \\ref{sec:evap} we check if the zone of partial hydrogen ionization overlaps with the disk evaporation region and we calculate the evolution of the disk luminosity in case of the evaporated inner disk. Finally, in Section \\ref{sec:diss} we discuss our results and give conclusions. ",
        "conclusions": "\\label{sec:diss}  We studied the effects of MHD turbulence and the efficiency of angular momentum transport in the quiescent accretion disks.  We analyzed the full disk vertical structure, which enabled us to determine accurately the coupling of the accreting gas to the magnetic field over the entire disk height.  Confirming the previous results of Menou \\& Quataert (2001), we find that only in the case of accretion disks in binary systems, e.g. around white dwarfs (CVs) or stellar mass black holes (SXTs), the timescale of the magnetic field decay is too short for the field to be frozen into the gas. Therefore, MHD turbulence does not operate efficiently and the large amplitude outbursts observed in these sources can be accounted for by the thermal ionization instability cycle in which $\\alpha_{\\rm cold} \\ll \\alpha_{\\rm hot}$.  On the other hand, in the case of accreting supermassive black holes in AGN the gas in the disk is well coupled to the magnetic field.  This effect ($Re_{\\rm M} \\gg 1$ and $Re_{\\rm A} \\gg 1$) can be only slightly reduced by encounters of free electrons with charged dust particles. In principle, the presence of dust in the disk can affect the number density of free electrons. However, the grains evaporate if the temperature exceeds $\\sim 1800-2000$ K, which is the case near the disk equatorial plane. As a result, the role of the dust in suppressing MHD turbulence is limited only to the uppermost disk layers and cannot have a global effect. We conclude, that the AGN accretion disks in quiescence have comparable angular momentum transport efficiency to the disks in outburst.  Therefore, we studied the ionization instability assuming a constant value of the viscosity parameter $\\alpha_{\\rm cold} = \\alpha_{\\rm hot}$. We took into account the evaporation of the innermost parts of the disk during the quiescence in AGN accretion disks. The transition to an ADAF below the evaporation radius during the evolution of the instability results in large luminosity variations ($\\Delta \\log L_{\\rm disk} \\sim 2$). Because only the outer parts of the disk ($R > R_{\\rm evap} \\sim 10^{3} R_{\\rm Schw}$) contribute to the total disk emission the minimum luminosity obtained within the instability cycle is lower than in the standard disk with no evaporation. We note, that also during the outburst the evaporation is important and ADAF remains in the innermost regions of the disk ($R_{\\rm evap} \\sim 25 R_{\\rm Schw}$), even at quite a high global accretion rate, $\\dot M = 0.1 \\dot M_{\\rm Edd}$.  The evaporation radius, however, strongly depends on the adopted value of the viscosity parameter $\\alpha$ (cf. Equation \\ref{eq:clasadaf}).  The $\\alpha$-disks in AGN are self-gravitating beyond a certain radius (Paczy\\'nski 1978; Hure 1998).  A typical distance of the self-gravitating zone is about a thousand Schwarzschild radii (Hure 2000). However, the greater the black hole mass and lower the parameter $\\alpha$, the closer to the center this zone is located. We can include a self-gravitation term equal to $-4\\pi G \\Sigma$ in Equation \\ref{eq:hydro} of the disk vertical structure. This additional term modifies the S-curve, so there is a single 'wiggle' on the middle branch of the curve, instead of two. The middle branch is also broader and reaches lower surface densities. The instability happens for a narrower range of effective temperatures, while the more strongly curved unstable middle branch covers broader range  of surface density values. Therefore, the amplitude of the outburst is not expected to be significantly modified in our model by the presence of self-gravity. However, whether the model actually can be extended into the self-gravity region, remains an open question. In particular, our description of convection, which influences the shape of the S-curve, may not be applicable in case of a self-gravitating disk. Also, MRI was not extensively studied in the self gravity region (but see Fromang et al. 2003) and the use of viscosity parameter $\\alpha$ may not be appropriate here. Instead, self-gravity may either provide a mechanism of angular momentum transfer (e.g. Paczy\\'nski 1978), or lead to the disk fragmentation and star formation (Collin \\& Zahn 1999). The physics of those processes is poorly known. Therefore the use of simple models like $\\alpha$-disks and comparison of their predictions with observed behavior of AGN seems to be justified.   The variability amplitudes predicted by our model are about two orders of magnitude and support the observational evidence of intermittent activity of AGN.  This kind of activity has been proposed to explain the formation of giant radio galaxies (Subrahmanyan et al. 1996) and the extended radio structures in the Giga-Hertz Peaked Spectrum (GPS) radio sources (Baum et al. 1990; Schoenmakers et al. 1998). The timescales of episodic activity in quasars were recently estimated by Martini \\& Schneider (2003). Recent observations of Compact Steep Spectrum (CSS) sources indicate that the evolutionary path of radio loud AGN is consistent with an on-off activity with duration governed by the mass of the central black hole (Marecki et al. 2003).  Also the black hole masses implied by the HST observations of normal galaxies and AGN suggest that every galaxy must undergo an active phase in its life (Ferrarese et al. 2001).   The estimated timescales of the activity cycle are within $10^{5}$ to 10$^{7}$~years (e.g. Reynolds \\& Begelman 1997 for GPS sources). Di Matteo et al. (2003) proposed a feedback mechanism to explain the cyclic activity in M87, which occurs on a timescale of 10$^{7}$ years. In our time-dependent calculations, we adopted a simplified approach to the disk evaporation process, and therefore underestimated the repetition timescale.  This is because we do not consider the time of the disk reconstruction after it has been evaporated, which increases the viscous timescale by a factor of $\\sim 2$. Assuming that the outburst starts at the outer edge of the instability zone we obtain the viscous timescale of $t_{\\rm visc} = 1/(\\alpha\\Omega)\\times (r/H)^{2} \\approx 0.9$ Myrs for $M_{\\rm BH} = 10^{8} M_\\odot$. Thus the outbursts will be less frequent than these presented in Figure \\ref{fig:ldisk_adaf}. However, because during the outburst there is no ADAF and the entire disk is present, the decay timescale is calculated properly and the duration of an outburst is correct. On the other hand, the outburst profiles will be less steep than the ones calculated here, since the disk density rises gradually at the beginning of the instability cycle. Although these two effects are important, the extension of the unstable region is not very large in comparison with the disk size and the uncertainty involved in our calculations is not crucial. The detailed model of the time evolution of the accretion disk with transition to ADAF is clearly worth investigating and we address it to the future work.   {\\sl Acknowledgments} AJ is grateful to F. Meyer and E. Meyer-Hofmeister for interesting discussion. This work was supported in part by grants 2P03D00322 and 2P03D00124 of the Polish State Committee for Scientific Research. Partial support for this work was provided by the National Aeronautics and Space Administration through Chandra Award Number GO-01164X issued by the Chandra X-Ray Observatory Center, which is operated by the Smithsonian Astrophysical Observatory for and on behalf of NASA under contract NAS8-39073."
    },
    "0311/hep-ph0311238_arXiv.txt": {
        "abstract": " ",
        "introduction": "It has recently been  pointed out that black holes could be  formed at future colliders if the Planck scale is of order a TeV, as is the case in some  extra-dimension  scenarios  \\cite{dimo,gidd}.  This  idea has driven a considerable amount of interest (see {\\it e.g.} \\cite{idea}). The same phenomenon could also occur due to  ultrahigh energy neutrino interactions in the atmosphere \\cite{ancho}. Most  works consider that those black holes could be described by the D-dimensional ($D  \\ge 5$) generalized Schwarzschild or Kerr metrics  \\cite{myers}. The aim of this paper is to study the experimental consequences of the existence of the Gauss - Bonnet term (as a step toward quantum gravity) if it is included in the D-dimensional action. This approach  should be more general  and  relies  on  a  real  expansion of supersymmetric  string theory. In Section 2, the basics of black hole formation at colliders and the  related cross sections are reminded. The details of the multi-dimensional Gauss - Bonnet black hole solutions and their thermodynamical properties are  given in Section 3. The flux  computation and the main analytical formulae  are explained in  Section 4. It is shown in Section  5 that the  Gauss - Bonnet (string) coupling constant can be measured in most cases, together with  the  dimensionality  of  space. Finally, some  possible  consequences  and  developments, especially with an additional cosmological constant, are discussed. ",
        "conclusions": "In  case  the Planck  scale  lies  in  the  TeV  range  due  to  extra dimensions, this study shows that, beyond the dimensionality of space, the  next  generation of  colliders  should  be  able  to  measure the coefficient  of  a  possible  Gauss - Bonnet term in  the  gravitational action. This would allow an important step forward in the construction of a full quantum theory of  gravity. It is also interesting to notice that  this  would  be  a  nice  example  of  the  convergence  between astrophysics and particle  physics in the final understanding of black holes and gravity in the Planckian region.  Nevertheless, those results could  be  improved and refined in several ways. First, the endpoint of the Hawking evaporation  process is still an unsolved  problem. In this paper, we have  considered that the time integral  of  the instantaneous  spectrum  is valid  up  to the  total disappearance of the black hole. Although usually a good approximation (as most  particles are emitted at masses close  to the initial mass), this can become a serious problem if the number of extra dimensions is high. In such cases, the mean  number of emitted particles can be very small and even smaller than one. The spectrum therefore {\\it  must} be truncated properly. A possibility could be to add a Heavyside function to  ensure  energy  conservation  while keeping the  same  probability distribution,  as   suggested   in   \\cite{stasastro},   but   a  full understanding of the phenomenon would be required as the analytical formulae derived in this work would not stand anymore.  Then, as studied in \\cite{cai,bir}, a cosmological constant could also be  included  in the action. On  the  theoretical side, this would  be strongly motivated by the great deal of attention paid to  the Anti-de Sitter and,  recently, de Sitter / Conformal Field  Theory (AdS and dS /CFT) correspondences. On the experimental  side,  this  would open an interesting window as there is no  unambiguous  relation  between  the D-dimensional and the 4-dimensional cosmological constants.  Finally, it would be  very interesting to extend this study to  Kerr - Gauss  -  Bonnet  black  holes \\cite{aurelnew} as the  holes  possibly produced   at   colliders  are  expected  to  be  spinning.   Although qualitatively   equivalent,   the   results   are   expected   to   be quantitatively quite different and probably more realistic."
    },
    "0311/astro-ph0311294_arXiv.txt": {
        "abstract": "We study the evolution of the global stellar mass density in a $\\Lam$ cold dark matter ($\\Lam$CDM) universe using two different types of hydrodynamical simulations (Eulerian TVD and SPH) and the analytical model of \\citet{Her03}.  We find that the theoretical calculations all predict both a higher stellar mass density at $z\\sim3$ than indicated by current observations, and that the peak of the cosmic star formation rate history should lie at $z\\gtsim 5$.  Such a star formation history implies that as much as (70\\%, 30\\%) of the total stellar mass density today must already have formed by $z=(1, 3)$.  Our results suggest that current observations at $z\\sim 3$ are missing as much as 50\\% of the total stellar mass density in the Universe, perhaps owing to an inadequate allowance for dust obscuration in star-forming galaxies, limited sample sizes, or cosmic variance. We also compare our results with some of the updated semi-analytic models of galaxy formation. ",
        "introduction": "\\label{section:intro}  Is the evidence for high redshift galaxy formation consistent with the concordance $\\Lam$CDM model?  Recent observational results include the discovery of Extremely Red Objects at $z\\ge 1$ \\citep[e.g.][]{Elston88, McCarthy92, Hu94, Cimatti03, Smail02}, Sub-millimeter galaxies at $z\\ge 2$ \\citep[e.g.][]{Smail97, Chapman}, Lyman-break galaxies (LBGs) at $z\\sim 3$ \\citep[e.g.][]{Steidel99}, and galaxies at $z\\gtsim 4$ either by their Lyman-$\\alpha$ emission \\citep[e.g.][]{Hu99, Rhoads, Taniguchi, Kodaira, Ouchi03a} or by their optical to near infrared (IR) colors \\citep[e.g.][]{Iwata, Ouchi03b, Dick03b}.  Multiband photometry including the near-IR band makes it possible to estimate the stellar mass of these high redshift galaxies by fitting the observed photometric results with artificial galaxy spectra generated by a population synthesis model.  Using this technique, several groups have now estimated the stellar mass density in the Universe in the redshift range of $0\\le z \\le 3$ \\citep[e.g.][]{Brinch, Cole, Cohen, Dick03a, Fontana03, Glazebrook04}. \\citet{Rudnick03} also estimated the stellar mass density at $z=0-3$ by combining the estimates of the rest-frame optical luminosity density and the mean cosmic mass-to-light ratio. These observational estimates constrain the evolution of the stellar mass density $\\Os$ as a function of redshift or cosmic time.  The observations indicate the presence of a significant stellar population at high redshift, and, by comparing with semi-analytic models of galaxy formation, some authors claim that $\\Lam$CDM models seriously {\\it underpredict} galaxy formation at $z\\sim 3$. For example, \\citet{Fontana03} compare their estimate of $\\Os$ from the Hubble Deep Field (HDF) South with the semi-analytic model of \\citet{Menci02}, and argue that the high-mass tail of the galaxy stellar mass function is not adequately described by CDM models.  Comparing to the same semi-analytic model, \\citet{Poli} argue that hierarchical models lack sufficient star formation at $z=2-4$, resulting in a failure to reproduce the pronounced brightening of the luminosity function at these redshifts.  Also, \\citet{Dick03a} find that their data from the HDF North suggest a steeper increase in $\\Os$ than some semi-analytic models \\citep[]{Kau99, Som, Cole00}, and some of the semi-analytic models predict {\\it higher} stellar mass density compared to their data points. Given these contradicting claims and the large parameter space available to the semi-analytic models as well as the limitations of the current observational samples, it is not clear at present if any of these discrepancies actually pose a serious problem to hierarchical evolution.  In this paper, we compare the observational data with the results from state-of-the-art cosmological simulations of the standard concordance $\\Lambda$CDM model and the theoretical model of \\citet{Her03} (hereafter H\\&S model), to show that, contrary to some of the claims, theory predicts a {\\it higher} $\\Os$ at $z=3$ than indicated by current observations and that the cosmic star formation rate (SFR) density peaks at $z\\geq 5$, earlier than suggested by most semi-analytic models. Also, we will explicitly compare our results with those from the updated semi-analytic models by \\citet{Som}, \\citet{Granato00} (GALFORM), and \\citet{Menci02}. ",
        "conclusions": "\\label{sec:conclusion}  We have shown that two independent different types of numerical hydrodynamic simulations both predict that the cosmic star formation rate density should peak at $z\\geq 5$, and that this relatively early peak in the SFR leads to a more rapid development of the stellar mass density than current observational estimates.  When all the results are scaled to the local value at $z=0$, the stellar mass density at $z=3$ in the simulations is larger than observed values by more than a factor of two. We also showed in Figure~1b that the semi-analytic models of \\citet{Som}, \\citet{Granato00}, and \\citet{Menci02} predict larger $\\Os$ than current observational estimates, but within the upper error bar of the \\citet{Fontana03} data point at $z=3$. In particular, the H\\&S model predicts higher $\\Os$ than those semi-analytic models by about a factor of two at $z=3$. This comes from the fact that both the simulations and H\\&S model has a peak of SFR at $z=5-7$, whereas the semi-analytic models have a peak at $z=2-4$.  The high $\\Os$ predicted by our simulations and the model of \\citet{Her03} suggest that current observations are missing nearly half of the total stellar mass density in the universe at high redshifts. This missing stellar mass could be hidden in a population of red galaxies that have not been detected previously in the optical ground-based data. In fact, \\citet{Franx} and \\citet{Daddi04} find such a population of red galaxies at $z\\geq 2$ in HDF-South that has a volume density half that of LBGs at $z=3$ and a stellar mass density comparable to that of LBGs. However, the data points by \\citet{Rudnick03} include this red galaxy population. Therefore the cosmic variance might be a stronger cause for the underestimate of $\\Os$ by the current observations. This is hinted by the absence of this red population in HDF-North. If these red populations are strongly clustered as suggested by \\citet{Daddi03}, observations with small fields-of-view could easily miss them. The stellar mass density in the simulations and observations agree reasonably well at $z<1$ where the rate of increase is much slower than at higher redshift. In addition, the dust extinction correction may seriously underestimate the fraction of early star formation that is heavily obscured. If we only had optical UV observations of nearby starburst galaxies -- rather than the full spectrum extending to the submillimeter -- we would greatly underestimate star formation in these systems \\citep[e.g.][]{Hughes98, Meurer99, Barger00, Calzetti00, Takagi03}. Future studies of starforming galaxies in the far-infrared wavelengths \\citep[e.g.][]{Kennicutt03} using the {\\it Spitzer Space Telescope}, in the submillimeter \\citep[e.g.][]{Tecza04}, in the millimeter \\citep[e.g.][]{Bekki00} using the {\\it Atacama Large Millimeter Array} (ALMA), and in X-rays \\citep[e.g.][]{Norman04} using the {\\it Chandra X-ray Observatory} would enable us to constrain the nature of starburst galaxies better, and estimate the intrinsic SFR more accurately.  It is reassuring that the two different sets of simulations, which use very different hydrodynamic methods (i.e. Eulerian TVD and SPH), both give a star formation history that peaks at $z\\geq 5$.  As argued by \\citet{Her03} this is to be expected, because the evolution of the cosmic SFR is driven mainly by a competition between gravity and the expansion of the Universe, with a weaker dependence on the details of star formation and feedback. The form of equation~(\\ref{eq:sfr}) given in Section~\\ref{sec:sfr} can be understood as follows. At high redshifts, when the cooling time is short, star formation is limited primarily by the gravitational growth of halos, which is independent of the dissipative gas dynamics. Thus, the parameter $\\beta$ and the factor $\\chi^{7/4}$ in the exponential of equation~(\\ref{eq:sfr}) are determined by the form of the matter power spectrum, and the description of star formation and feedback enter into $\\dot\\rhos(z)$ only logarithmically. At low redshifts, the supply of star forming gas is limited by the expansion rate of the Universe, fixing the dependence of $\\dot\\rhos(z)$ on $\\chi$ as $z\\rightarrow 0$. The explicit influence of the prescription for star formation and feedback is again subdominant and mainly affects the values of the normalization parameters $\\alpha$ and $\\dot\\rhos(0)$.  For these reasons, equation~(\\ref{eq:sfr}) can be generalized straightforwardly to other cosmologies and to include other physics \\citep[e.g.][]{Yoshida03}.  Moreover, as shown in e.g. figure 6 of \\citet{Her03}, the fact that the SFR density is regulated mainly by the competition between gravity and the expansion of the Universe means that the peak in $\\dot\\rhos(z)$ should lie at $z\\approx 5$, unless an implausible value is adopted for the parameters governing star formation. Hence, it is not surprising that the various sets of simulations should be consistent, with residual differences owing to details in the cosmology and, most important, cosmic variance and resolution limitations.  We plan to investigate these issues in the future using the algorithms described here, as well as adaptive mesh refinement codes.  For now, the agreement between our different numerical approaches supports the general arguments made by Hernquist \\& Springel (2003), that the SFR density should peak at $z\\geq 5$, mostly independent of the details of star formation and feedback.  As we have demonstrated here, such an early peak in the cosmic star formation history yields $\\Os$ that clearly exceeds current observational estimates and the results of semi-analytic models, suggesting that most of the stars in the universe at $z\\gtsim 3$ are ``missing.''"
    },
    "0311/astro-ph0311588_arXiv.txt": {
        "abstract": "}[2]{{\\footnotesize\\begin{center}ABSTRACT\\end{center} \\vspace{1mm}\\par#1\\par \\noindent {~}{\\it #2}}}  \\newcommand{\\TabCap}[2]{\\begin{center}\\parbox[t]{#1}{\\begin{center} \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable \\\\[2mm] \\footnotesize #2 \\end{center}}\\end{center}}  \\newcommand{\\TableSep}[2]{\\begin{table}[p]\\vspace{#1} \\TabCap{#2}\\end{table}}  \\newcommand{\\FigCap}[1]{\\footnotesize\\par\\noindent Fig.\\  % \\refstepcounter{figure}\\thefigure. #1\\par}  \\newcommand{\\TableFont}{\\footnotesize} \\newcommand{\\TableFontIt}{\\ttit} \\newcommand{\\SetTableFont}[1]{\\renewcommand{\\TableFont}{#1}}  \\newcommand{\\MakeTable}[4]{\\begin{table}[htb]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newcommand{\\MakeTableSep}[4]{\\begin{table}[p]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newenvironment{references}% { \\footnotesize \\frenchspacing \\renewcommand{\\thesection}{} \\renewcommand{\\in}{{\\rm in }} \\renewcommand{\\AA}{Astron.\\ Astrophys.} \\newcommand{\\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \\newcommand{\\ApJ}{Astrophys.\\ J.} \\newcommand{\\ApJS}{Astrophys.\\ J.~Suppl.~Ser.} \\newcommand{\\ApJL}{Astrophys.\\ J.~Letters} \\newcommand{\\AJ}{Astron.\\ J.} \\newcommand{\\IBVS}{IBVS} \\newcommand{\\PASP}{P.A.S.P.} \\newcommand{\\Acta}{Acta Astron.} \\newcommand{\\MNRAS}{MNRAS} \\renewcommand{\\and}{{\\rm and }} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311541_arXiv.txt": {
        "abstract": "{ We present stellar parameters and metallicities, obtained from a detailed spectroscopic analysis, for a large sample of 98 stars known to be orbited by planetary mass companions (almost all known targets), as well as for a volume-limited sample of 41 stars not known to host any planet. For most of the stars the stellar parameters are revised versions of the ones presented in our previous work. However, we also present parameters for 18 stars with planets not previously published, and a compilation of stellar parameters for the remaining 4 planet-hosts for which we could not obtain a spectrum. A comparison of our stellar parameters with values of T$_{\\mathrm{eff}}$, $\\log{g}$, and [Fe/H] available in the literature shows a remarkable agreement. In particular, our spectroscopic $\\log{g}$ values are now very close to trigonometric $\\log{g}$ estimates based on Hipparcos parallaxes. The derived [Fe/H] values are then used to confirm the previously known result that planets are more prevalent around metal-rich stars. Furthermore, we confirm that the frequency of planets is a strongly rising function of the stellar metallicity, at least for stars with [Fe/H]$>$0. While only about 3\\% of the solar metallicity stars in the CORALIE planet search sample were found to be orbited by a planet, this number increases to more than 25\\% for stars with [Fe/H] above $+$0.3. Curiously, our results also suggest that these percentages might remain relatively constant for values of [Fe/H] lower than about solar, increasing then linearly with the mass fraction of heavy elements. These results are discussed in the context of the theories of planetary formation. ",
        "introduction": "The discovery of now more than 115 giant planets orbiting solar-type stars\\footnote{See e.g. table at {\\tt http://obswww.unige.ch/Exoplanets} for a continuously updated version} has lead to a number of different studies on the formation and evolution of the newly found planetary systems (for a recent review see e.g. \\citet[][]{May03b} or \\citet{San03b}). With the numbers increasing, current analyses are giving us the first statistically significant results about the properties of the new systems \\citep[e.g.][]{Jor01,Zuc02,Udr03,San03,Egg03}. Amongst these, some deal with the planet-host stars themselves: they were found to be significantly metal-rich with respect to the average field dwarfs \\citep[e.g.][]{Gon97,Furetal98,Gon98,San00,San01,Gon01,Rei02,San03,Law03}.   Current studies seem to favor that this ``excess'' metallicity has a primordial origin, i.e., that the high metal content of the stars was common to the cloud of gas and dust that gave origin to the star-planet system \\citep[][]{Pin01,San01,San03}. Furthermore, it has been shown that the frequency of planetary companions is a strong function of the metal content of the star \\citep[][]{San01,San03,Rei02,San03}: it is much easier to find planets around metal-rich objects. Overall, the results suggest that the formation of giant planets (or at least of the kind we find now) is very dependent on the grain content of the disk, a result that has important consequences for theories of planetary formation \\citep[][]{Pol96,Bos02,Ric03}.  During the last few years we have gathered spectra for planet host stars, as well as of a sample of objects not known to harbor any planetary companion. The main results of our uniform study, concerning the metallicity of planet host stars, have been presented in \\citet[][]{San00}, \\citet[][]{San01}, and \\citet[][]{San03} (hereafter Papers\\,I, II, and III, respectively).  \\begin{table*} \\caption[]{Atomic parameters and measured solar equivalent widths for the \\ion{Fe}{i} and \\ion{Fe}{ii} lines used.} \\begin{tabular}{lccc|lccc} \\hline $\\lambda$ (\\AA) & $\\chi_{l}$ & $\\log{gf}$ & EW$_{\\sun}$ (m\\AA)& $\\lambda$ (\\AA) & $\\chi_{l}$ & $\\log{gf}$ & EW$_{\\sun}$ (m\\AA) \\\\ \\hline \\ion{Fe}{i} & & & & 6591.32 & 4.59 & $-$1.98 & 10.6  \\\\ 5044.22 & 2.85 & $-$2.04 & 73.4      & 6608.03 & 2.28 & $-$3.96 & 17.7  \\\\ 5247.06 & 0.09 & $-$4.93 & 66.8      & 6627.55 & 4.55 & $-$1.48 & 28.0  \\\\ 5322.05 & 2.28 & $-$2.90 & 60.4      & 6646.94 & 2.61 & $-$3.94 &  9.9  \\\\ 5806.73 & 4.61 & $-$0.89 & 53.7      & 6653.86 & 4.15 & $-$2.41 & 10.5  \\\\ 5852.22 & 4.55 & $-$1.19 & 40.6      & 6703.57 & 2.76 & $-$3.02 & 36.9  \\\\ 5855.08 & 4.61 & $-$1.53 & 22.4      & 6710.32 & 1.48 & $-$4.82 & 16.0  \\\\ 5856.09 & 4.29 & $-$1.56 & 33.8      & 6725.36 & 4.10 & $-$2.20 & 17.2  \\\\ 6027.06 & 4.08 & $-$1.18 & 64.3      & 6726.67 & 4.61 & $-$1.05 & 46.9  \\\\ 6056.01 & 4.73 & $-$0.50 & 72.4      & 6733.16 & 4.64 & $-$1.43 & 26.8  \\\\ 6079.01 & 4.65 & $-$1.01 & 45.7      & 6750.16 & 2.42 & $-$2.61 & 74.1  \\\\ 6089.57 & 5.02 & $-$0.88 & 35.0      & 6752.71 & 4.64 & $-$1.23 & 35.9  \\\\ 6151.62 & 2.18 & $-$3.30 & 49.8      & 6786.86 & 4.19 & $-$1.90 & 25.2  \\\\ 6157.73 & 4.07 & $-$1.24 & 61.9      &         &      &         &       \\\\ 6159.38 & 4.61 & $-$1.86 & 12.4      & \\ion{Fe}{ii} & &         &       \\\\ 6165.36 & 4.14 & $-$1.50 & 44.6      & 5234.63 & 3.22 & $-$2.23 & 83.7  \\\\ 6180.21 & 2.73 & $-$2.64 & 55.8      & 5991.38 & 3.15 & $-$3.53 & 31.5  \\\\ 6188.00 & 3.94 & $-$1.63 & 47.7      & 6084.11 & 3.20 & $-$3.78 & 20.8  \\\\ 6200.32 & 2.61 & $-$2.40 & 73.3      & 6149.25 & 3.89 & $-$2.72 & 36.2  \\\\ 6226.74 & 3.88 & $-$2.07 & 29.3      & 6247.56 & 3.89 & $-$2.35 & 52.2  \\\\ 6229.24 & 2.84 & $-$2.89 & 37.9      & 6369.46 & 2.89 & $-$4.13 & 19.2  \\\\ 6240.65 & 2.22 & $-$3.29 & 48.3      & 6416.93 & 3.89 & $-$2.64 & 40.1  \\\\ 6265.14 & 2.18 & $-$2.56 & 86.0      & 6432.69 & 2.89 & $-$3.56 & 41.5  \\\\ 6270.23 & 2.86 & $-$2.58 & 52.3      & 6446.40 & 6.22 & $-$1.91 &  4.2  \\\\ 6380.75 & 4.19 & $-$1.32 & 52.2      & 7479.70 & 3.89 & $-$3.59 & 10.0  \\\\ 6392.54 & 2.28 & $-$3.93 & 18.1      & 7515.84 & 3.90 & $-$3.43 & 13.4  \\\\ 6498.94 & 0.96 & $-$4.63 & 45.9      & 7711.73 & 3.90 & $-$2.55 & 46.0  \\\\ \\hline \\end{tabular} \\label{tab1} \\end{table*}   Most groups working on exoplanets are now convinced that planet host stars are really more metal-rich than average field dwarfs. This result is clearly independent of the kind of analysis done to obtain the stellar metallicity \\citep[e.g.][]{Gim00,Gon01,San01,Mur02,Martell02,Hei03}, and in Paper\\,III we showed that this result is not due to any observational bias. However, some authors have questioned the quality of the spectroscopic analyses we (and others) have been publishing. In particular, the relatively high surface gravities derived in our preceding papers led to some criticism regarding this matter.  In order to address this problem, in this paper we present a revised spectroscopic analysis for all the stars presented in Papers\\,II and III. The new derived surface gravities are now compatible with the ones obtained by other authors, and with trigonometric gravities derived using Hipparcos parallaxes \\citep[][]{ESA97}. Other stellar parameters (T$_{\\mathrm{eff}}$ and [Fe/H]) are also similar to the ones presented elsewhere in the literature, and not particularly different from the ones derived in Papers\\,II and III.  Furthermore, we have derived stellar parameters for 18 planet host stars not analyzed before, increasing to 98 the number of these objects for which we have precise spectral information. The new results unambiguously confirm the previously presented trends: stars with planetary companions are more metal-rich than average field dwarfs.  \\begin{table*} \\caption[]{Stars with planets and derived stellar parameters (HD number between 1 and 60\\,000). See text for more details.} \\begin{tabular}{lcccrccccc} \\hline HD     & T$_{\\mathrm{eff}}$ & $\\log{g}_{spec}$ & $\\xi_{\\mathrm{t}}$ & \\multicolumn{1}{c}{[Fe/H]} & N(\\ion{Fe}{i},\\ion{Fe}{ii}) & $\\sigma$(\\ion{Fe}{i},\\ion{Fe}{ii}) & Instr.$^a$ & Mass         & $\\log{g}_{hipp}$ \\\\ number & [K]                &  [cm\\,s$^{-2}$]  &  [km\\,s$^{-1}$]  &        &                              &                                    &            & [M$_{\\sun}$] & [cm\\,s$^{-2}$] \\\\ \\hline \\object{HD\\,142   }    &  6302$\\pm$56  &  4.34$\\pm$0.13 &  1.86$\\pm$0.17 &   0.14$\\pm$0.07 & 28,8 & 0.05,0.05 & [2]  & 1.28 &4.27\\\\ \\object{HD\\,1237  }    &  5536$\\pm$50  &  4.56$\\pm$0.12 &  1.33$\\pm$0.06 &   0.12$\\pm$0.06 & 37,7 & 0.05,0.06 & [1]  & 0.99 &4.56\\\\ \\object{HD\\,2039  }    &  5976$\\pm$51  &  4.45$\\pm$0.10 &  1.26$\\pm$0.07 &   0.32$\\pm$0.06 & 34,6 & 0.05,0.04 & [1]  & 1.18 &4.35\\\\ \\object{HD\\,3651  }    &  5173$\\pm$35  &  4.37$\\pm$0.12 &  0.74$\\pm$0.05 &   0.12$\\pm$0.04 & 31,5 & 0.04,0.05 & [4]  & 0.76 &4.41\\\\ \\object{HD\\,4203  }    &  5636$\\pm$40  &  4.23$\\pm$0.14 &  1.12$\\pm$0.05 &   0.40$\\pm$0.05 & 37,7 & 0.05,0.07 & [2]  & 1.06 &4.19\\\\ \\object{HD\\,4208  }    &  5626$\\pm$32  &  4.49$\\pm$0.10 &  0.95$\\pm$0.06 &$-$0.24$\\pm$0.04 & 37,7 & 0.04,0.05 & [2]  & 0.86 &4.48\\\\ \\object{HD\\,6434  }    &  5835$\\pm$59  &  4.60$\\pm$0.12 &  1.53$\\pm$0.27 &$-$0.52$\\pm$0.08 & 30,4 & 0.06,0.06 & [2]  & 0.82 &4.33\\\\ \\object{HD\\,8574  }    &  6151$\\pm$57  &  4.51$\\pm$0.10 &  1.45$\\pm$0.15 &   0.06$\\pm$0.07 & 30,7 & 0.06,0.04 & [4]  & 1.18 &4.28\\\\ \\object{HD\\,9826  }    &  6212$\\pm$64  &  4.26$\\pm$0.13 &  1.69$\\pm$0.16 &   0.13$\\pm$0.08 & 27,6 & 0.06,0.05 & [4]  & 1.30 &4.16\\\\ \\object{HD\\,10647 }    &  6143$\\pm$31  &  4.48$\\pm$0.08 &  1.40$\\pm$0.08 &$-$0.03$\\pm$0.04 & 34,6 & 0.03,0.03 & [1]  & 1.14 &4.43\\\\ \\object{HD\\,10697 }    &  5641$\\pm$28  &  4.05$\\pm$0.05 &  1.13$\\pm$0.03 &   0.14$\\pm$0.04 & 33,7 & 0.03,0.03 & [4]  & 1.22 &4.03\\\\ \\object{HD\\,12661 }    &  5702$\\pm$36  &  4.33$\\pm$0.08 &  1.05$\\pm$0.04 &   0.36$\\pm$0.05 & 34,8 & 0.04,0.03 & [3]  & 1.05 &4.34\\\\ \\object{HD\\,13445 }    &  5119$\\pm$43  &  4.48$\\pm$0.14 &  0.63$\\pm$0.07 &$-$0.25$\\pm$0.05 & 38,6 & 0.05,0.07 & [1]  & 0.67 &4.44\\\\ \\object{HD\\,13445 }    &  5207$\\pm$30  &  4.56$\\pm$0.11 &  0.82$\\pm$0.05 &$-$0.23$\\pm$0.04 & 38,5 & 0.03,0.05 & [2]  & 0.74 &4.52\\\\ \\object{HD\\,13445 }    &  5163         &  4.52     \t&  0.72          &\\multicolumn{1}{c}{$-$0.24}     \t   &      &           & avg. & 0.70 &4.48\\\\ \\object{HD\\,16141 }    &  5801$\\pm$30  &  4.22$\\pm$0.12 &  1.34$\\pm$0.04 &   0.15$\\pm$0.04 & 37,7 & 0.03,0.04 & [2]  & 1.05 &4.17\\\\ \\object{HD\\,17051 }    &  6252$\\pm$53  &  4.61$\\pm$0.16 &  1.18$\\pm$0.10 &   0.26$\\pm$0.06 & 34,6 & 0.05,0.07 & [2]  & 1.32 &4.49\\\\ \\object{HD\\,19994 }    &  6217$\\pm$67  &  4.29$\\pm$0.08 &  1.62$\\pm$0.12 &   0.25$\\pm$0.08 & 35,5 & 0.06,0.03 & [1]  & 1.37 &4.14\\\\ \\object{HD\\,19994 }    &  6290$\\pm$58  &  4.31$\\pm$0.13 &  1.63$\\pm$0.12 &   0.32$\\pm$0.07 & 33,6 & 0.06,0.05 & [2]  & 1.40 &4.17\\\\ \\object{HD\\,19994 }    &  6121$\\pm$33  &  4.06$\\pm$0.05 &  1.55$\\pm$0.06 &   0.19$\\pm$0.05 & 37,5 & 0.04,0.03 & [5]  & 1.34 &4.09\\\\ \\object{HD\\,19994 }    &  6132$\\pm$67  &  4.11$\\pm$0.23 &  1.37$\\pm$0.12 &   0.21$\\pm$0.08 & 35,6 & 0.06,0.09 & [3]  & 1.36 &4.10\\\\ \\object{HD\\,19994 }    &  6190         &  4.19     \t&  1.54          &   \\multicolumn{1}{c}{0.24}     \t   &      &           & avg. & 1.37 &4.12\\\\ \\object{HD\\,20367 }    &  6138$\\pm$79  &  4.53$\\pm$0.22 &  1.22$\\pm$0.16 &   0.17$\\pm$0.10 & 31,6 & 0.08,0.09 & [6]  & 1.21 &4.42\\\\ \\object{HD\\,22049 }    &  5073$\\pm$42  &  4.43$\\pm$0.08 &  1.05$\\pm$0.06 &$-$0.13$\\pm$0.04 & 37,6 & 0.05,0.04 & [1]  & 0.73 &4.55\\\\ \\object{HD\\,23079 }    &  5959$\\pm$46  &  4.35$\\pm$0.12 &  1.20$\\pm$0.10 &$-$0.11$\\pm$0.06 & 35,6 & 0.05,0.05 & [2]  & 1.01 &4.36\\\\ \\object{HD\\,23596 }    &  6108$\\pm$36  &  4.25$\\pm$0.10 &  1.30$\\pm$0.05 &   0.31$\\pm$0.05 & 36,6 & 0.04,0.04 & [3]  & 1.30 &4.22\\\\ \\object{HD\\,27442 }    &  4825$\\pm$107 &  3.55$\\pm$0.32 &  1.18$\\pm$0.12 &   0.39$\\pm$0.13 & 36,6 & 0.11,0.13 & [2]  &  --  &--  \\\\ \\object{HD\\,28185 }    &  5656$\\pm$44  &  4.45$\\pm$0.08 &  1.01$\\pm$0.06 &   0.22$\\pm$0.05 & 38,6 & 0.05,0.03 & [1]  & 0.98 &4.39\\\\ \\object{HD\\,30177 }    &  5591$\\pm$50  &  4.35$\\pm$0.12 &  1.03$\\pm$0.06 &   0.39$\\pm$0.06 & 37,4 & 0.06,0.05 & [1]  & 1.01 &4.34\\\\ \\object{HD\\,30177 }    &  5584$\\pm$65  &  4.23$\\pm$0.13 &  1.14$\\pm$0.07 &   0.38$\\pm$0.09 & 38,7 & 0.07,0.05 & [2]  & 1.01 &4.34\\\\ \\object{HD\\,30177 }    &  5588         &  4.29     \t&  1.08          &   \\multicolumn{1}{c}{0.39}     \t   &      &           & avg. & 1.01 &4.34\\\\ \\object{HD\\,33636$^b$} &  6046$\\pm$49  &  4.71$\\pm$0.09 &  1.79$\\pm$0.19 &$-$0.08$\\pm$0.06 & 37,6 & 0.05,0.04 & [2]  & 1.16 &4.56\\\\ \\object{HD\\,37124 }    &  5546$\\pm$30  &  4.50$\\pm$0.03 &  0.80$\\pm$0.07 &$-$0.38$\\pm$0.04 & 36,7 & 0.04,0.02 & [3]  & 0.75 &4.33\\\\ \\object{HD\\,38529 }    &  5674$\\pm$40  &  3.94$\\pm$0.12 &  1.38$\\pm$0.05 &   0.40$\\pm$0.06 & 34,7 & 0.05,0.06 & [2]  & 1.60 &3.81\\\\ \\object{HD\\,39091$^b$} &  5991$\\pm$27  &  4.42$\\pm$0.10 &  1.24$\\pm$0.04 &   0.10$\\pm$0.04 & 38,7 & 0.03,0.04 & [1]  & 1.10 &4.38\\\\ \\object{HD\\,40979 }    &  6145$\\pm$42  &  4.31$\\pm$0.15 &  1.29$\\pm$0.09 &   0.21$\\pm$0.05 & 24,9 & 0.04,0.07 & [4]  & 1.21 &4.38\\\\ \\object{HD\\,46375 }    &  5268$\\pm$55  &  4.41$\\pm$0.16 &  0.97$\\pm$0.06 &   0.20$\\pm$0.06 & 37,4 & 0.05,0.07 & [3]  & 0.82 &4.34\\\\ \\object{HD\\,47536 }    &  4554$\\pm$85  &  2.48$\\pm$0.23 &  1.82$\\pm$0.08 &$-$0.54$\\pm$0.12 & 37,6 & 0.11,0.09 & [2]  &   -- &--  \\\\ \\object{HD\\,49674 }    &  5644$\\pm$54  &  4.37$\\pm$0.07 &  0.89$\\pm$0.07 &   0.33$\\pm$0.06 & 33,5 & 0.06,0.04 & [4]  & 1.04 &4.50\\\\ \\object{HD\\,50554 }    &  6026$\\pm$30  &  4.41$\\pm$0.13 &  1.11$\\pm$0.06 &   0.01$\\pm$0.04 & 37,6 & 0.03,0.05 & [3]  & 1.09 &4.40\\\\ \\object{HD\\,52265 }    &  6076$\\pm$57  &  4.20$\\pm$0.17 &  1.38$\\pm$0.09 &   0.20$\\pm$0.07 & 39,7 & 0.06,0.07 & [1]  & 1.19 &4.32\\\\ \\object{HD\\,52265 }    &  6131$\\pm$47  &  4.35$\\pm$0.13 &  1.33$\\pm$0.08 &   0.25$\\pm$0.06 & 36,6 & 0.05,0.04 & [2]  & 1.21 &4.34\\\\ \\object{HD\\,52265 }    &  6103         &  4.28     \t&  1.36          &   \\multicolumn{1}{c}{0.23}     \t   &      &           & avg. & 1.20 &4.33\\\\ \\hline \\end{tabular} \\\\ $^a$ The instruments used to obtain the spectra were: [1] 1.2-m Swiss Telescope/CORALIE; [2] 1.5-m and 2.2-m ESO/FEROS; [3] WHT/UES; [4] TNG/SARG; [5] VLT-UT2/UVES; [6] 1.93-m OHP/ELODIE; [7] Keck/HIRES \\newline $^b$ The companions to these stars have minimum masses above 10\\,M$_{\\mathrm{Jup}}$, and are thus probably brown-dwarfs. \\label{tabplan1} \\end{table*} ",
        "conclusions": "In this paper we have derived stellar metallicities from a detailed spectroscopic analysis of a sample of 98 stars known to be orbited by planetary mass companions, as well as for a volume-limited sample of stars not known to host any planets. The main results are:  \\begin{itemize}  \\item The obtained stellar parameters (T$_{\\mathrm{eff}}$, $\\log{g}$, [Fe/H], and stellar masses) are compatible, within the errors, with the values derived by other authors using similar or different techniques. In particular, the derived surface gravities are only on average $\\sim$0.03\\,dex different to trigonometric estimates based on Hipparcos parallaxes.  \\item We confirm the previously known trends that stars with planets are more metal-rich than average field dwarfs. The average difference is of the order of 0.25\\,dex.  \\item We confirm previous results (e.g. Papers II and III) that have shown that the frequency of stars having planets is a strong rising function of the stellar metallicity. About 25-30\\% of dwarfs in the CORALIE planet search sample having [Fe/H]$>$0.3 harbor a planetary companion. This number falls to $\\sim$3\\% for stars of solar metallicity. The Sun is in the tail of this distribution, that seems to be rather flat for [Fe/H]$<$0.0 (i.e. for mass fractions of heavy elements Z$<$0.02), but increasing (maybe linearly) as a function of Z for higher values. Possible implications of these results are discussed.  \\end{itemize}  The main conclusions of this paper agree with previous results that have investigated the striking role that stellar metallicity seems to be playing in the formation of giant planets, or at least in the formation of the kind of systems ``planet-hunters'' are finding now. However, it is crucial that this kind of analysis is done on a continuous basis as new planets are added to the lists. In particular, the question of knowing whether the the Solar System is typical is particularly troubling, as the Sun falls in the tail of the [Fe/H] distributions of planet-host stars."
    },
    "0311/astro-ph0311631_arXiv.txt": {
        "abstract": "We describe our recent attempts to model substructure in dark matter halos down to very small masses, using a semi-analytic model of halo formation. The results suggest that numerical simulations of halo formation may still be missing substructure in the central regions of halos due to purely numerical effects. If confirmed, this central `overmerging' problem will have important consequences for the interpretation of lensing measurements of substructure. We also show that the spatial distribution of subhalos relative to the satellite companions of the Milky Way rules out at least one simple model of how dwarf galaxies form in low-mass halos. ",
        "introduction": "When numerical simulations of cold dark matter (CDM) halo formation reached sufficient resolution a few years ago, they revealed that a wealth of dense substructure, the undigested remains of ancient hierarchical merging, should still survive in these systems at the present day (Klypin et al.~1999; Okamoto \\& Habe 1999; Moore et al.~1999). The properties of halo substructure will be a sensitive test of the nature of dark matter if we can manage to quantify them observationally, for instance in strong gravitational lens systems (cf.\\ talks by P.~Schneider and N.~Dalal in these proceedings). Halo substructure must also be linked to satellite galaxies, such as the dozen dwarf satellites of the Milky Way, although the exact nature of the connection remains problematic, given the huge discrepancy between the number of subhalos predicted by simulations and the number of luminous satellites observed locally.  Before we can reach firm conclusions on the properties of halo substructure or the nature of the satellite-substructure connection, however, we should first confirm that current simulations of substructure have converged to a definite and accurate prediction of its properties. For many years, N-body simulations suffered from `overmerging', the artificial disruption of substructure due to numerical effects. It is still not clear that all the properties of simulated halo substructure have converged to level required for applications such as lens modeling. To test for the possibly of residual overmerging, we will compare purely numerical results with semi-analytic (SA) predictions, using a recently developed SA model of halo formation. ",
        "conclusions": ""
    },
    "0311/astro-ph0311407_arXiv.txt": {
        "abstract": "We introduce crowded field integral field (3D) spectrophotometry as a useful technique for the study of resolved stellar populations in nearby galaxies. The spectroscopy of individual extragalactic stars, which is now feasible with efficient instruments and large telescopes, is confronted with the observational challenge of accurately subtracting the bright, spatially and wavelength-dependent non-uniform background of the underlying galaxy. As a methodological test, we present a pilot study with selected extragalactic planetary nebulae (XPN) in the bulge of M31, demonstrating how 3D spectroscopy is able to improve the limited accuracy of background subtraction which one would normally obtain with classical slit spectroscopy. It is shown that due to the absence of slit effects, 3D is a most suitable technique for spectrophometry. We present spectra and line intensities for 5 XPN in M31, obtained  with the MPFS instrument at the Russian 6m BTA, INTEGRAL at the WHT, and with PMAS at the Calar Alto 3.5m Telescope. The results for two of our targets, for which data are available in the literature, are compared with previously published emission line intensities. The three remaining PN have been observed spectroscopically for the first time. One object is shown to be a previously misidentified supernova remnant. Our monochromatic H$_\\alpha$ maps are compared with direct Fabry-P\\'erot and narrowband filter images of the bulge of M31, verifying the presence of filamentary emission of the interstellar medium in the vicinity of our objects. We present an example of a flux calibrated and continuum-subtracted filament spectrum and demonstrate how the ISM component introduces systematic errors in the measurement of faint diagnostic PN emission lines when conventional observing techniques are employed. It is shown how these errors can be eliminated with 3D spectroscopy, using the full 2-dimensional spatial information and point spread function (PSF) fitting techniques. Using 3D spectra of bright standard stars, we demonstrate that the PSF is sampled with high accuracy, providing a centroiding precision at the milli-arcsec level. Crowded field 3D spectrophotometry and the use of PSF fitting techniques is suggested as the method of choice for a number of similar observational problems, including luminous stars in nearby galaxies, supernovae, QSO host galaxies, gravitationally lensed QSOs, and others. ",
        "introduction": "\\label{INTRODUCTION}  Understanding the formation and evolution of galaxies is one of the most prominent goals of modern cosmology. After a long-lasting discussion about the best evolutionary scenarios for elliptical galaxies, for example, either according to the classical picture of instantaneous dissipative collapse (Larson 1974), or rather in the hierarchical scenario through gradual mergers of spirals (White 1980), it seems today that the latter view is the more widely accepted one. The observational evidence, that nearby ellipticals show peculiarities like counter-rotating cores, rings, gas disks, and recent star formation, can be explained by merger and accretion processes over an extended period of time. On the theoretical side, numerical N-body simulations combined with the SPH technique in a $\\Lambda$CDM cosmology are capable of reproducing observed photometric and kinematical properties of elliptical galaxies fairly well, thus supporting the hierarchical picture, despite some remaining discrepancies (e.g. Meza et al.\\ 2003). On the other hand, Chiosi \\& Carraro 2002 presented SPH simulations producing elliptical galaxies, which are compatible with monolithic collapse, so that the problem is not yet finally solved.  Among the classical tests for the validity of any of the theoretical pictures, the integrated-light measurement of broad-band colors and of absorption line indices play an important role. The observations can be compared with integrated light stellar population models in order to distinguish between different star formation histories and chemical enrichment (e.g. Saglia et al.\\ 2000, Mehlert et al.\\ 2000).  However, the presence of dust and ionized gas in the interstellar medium of elliptical galaxies (Goudfrooij et al.\\ 1994, Bregman et al.\\ 1992) has motivated the question, whether or to which extent the observed gradients are real (dust effecting the surface photometry through extinction, gaseous emission filling in the absorption line profiles of H$_\\beta$, Mg{\\it b}, Mg$_2$). Moreover, systematic errors of long-slit spectroscopy due to the effects of a variable point-spread-function of the spectrograph optics along the slit have to be taken into account. For a discussion of these problems, see Mehlert et al.\\ 2000.  From a different point of view, Freeman \\& Bland-Hawthorn (2002) advocate in what they termed ``near-field cosmology'' to probe the 6-dimensional phase space and chemical compositions of very many individual stars in order to unravel details of the star formation history of the Milky Way and thus the signatures of its formation and evolution.  An alternative approach, now rapidly producing new abundance data and reaching out to significantly larger distances, is based on the spectroscopic analysis of individual resolved luminous stars in nearby galaxies, e.g.\\ M31 (Smartt et al.\\ 2001, Venn et al.\\ 2000), NGC6822 (Venn et al.\\ 2001), M33 (Monteverde et al.\\ 1997), or NGC300 (Urbaneja et al.\\ 2003). Using the ESO-VLT + FORS, the feasibility of high signal-to-noise stellar spectroscopy even beyond the local group was demonstrated by Bresolin et al.\\ (2001), who measured 7 supergiants of spectral types B, A, and F with V$\\approx$20.5 in NGC3621, a galaxy which is located at a distance of 6.7 Mpc. For a review on extragalactic stellar spectroscopy, see Kudritzki 1998.   Currently the only way to measure individual abundances from old or intermediate age stars in galaxies more distant than the Magellanic Clouds is through the emission line spectra of extragalactic planetary nebulae (XPN), see Walsh et al.\\ 2000.  This approach has some similarities with the standard method of measuring abundance gradients from individual H$\\;$II regions in the disk of spiral galaxies (Shaver et al.\\ 1983, Zaritsky et al.\\ 1994). As opposed to H$\\;$II regions, XPN metallicities can be derived in a homogeneous way for galaxies of {\\em any} Hubble type, and on all scales of galactocentric distances. The task of obtaining abundance gradients out to large radii, where a low surface-brightness precludes to measure reliable colors or absorption line indices, can be addressed with XPN, providing important constraints for galactic evolution models (Worthey 1999). Moreover, since radial velocities of XPN are measurable out to several effective radii, they are  potentially useful for probing the gravitational potential of galaxies (M{\\' e}ndez et al. 2001), and  for tracing merger events (Hui et al. 1995, Durrell et al.\\ 2003). Recently, XPN and an H$\\;$II region have been detected in the intracluster space of the Virgo cluster, giving an excellent opportunity to study the properties of this unique stellar population and, potentially, their star formation history and metallicity (Arnaboldi  et al.\\ 2002, Gerhard et al.\\ 2002, Feldmeier et al.\\ 2003).   When compared to integrated light studies, the analysis of single objects has the advantage that the abundance determinations are based on detailed physical models, and that an assessment of individual errors is possible. The study of resolved stellar populations in galaxies  out to the distance of the Virgo cluster has become a major science case for the proposed new generation of Extremely Large Telescopes (Wyse et al. 2000, Hawarden et al. 2003, Najita \\& Strom 2002).    \\begin{figure*}[htb] \\centering\\includegraphics[width=0.70\\linewidth,clip]{f1.eps} \\figcaption{H$_\\alpha$ narrow-band image of the central 6~arcmin of the Andromeda galaxy M31, obtained with Calar Alto 3.5m prime focus camera + Fabry-P\\'erot. The continuum has been subtracted using an off-band image to reveal the complex and fuzzy emission structure of the ISM. The tiny dark spots are XPN candidates, while spots with white substructures are indicative of residuals from imperfectly subtracted continuum sources (stars). The insert shows a broad-band DSS image with an outline of the H$_\\alpha$ frame, illustrating the rapid increase of continuum surface brightness towards the nucleus. Orientation: North to the top, East to the left. \\label{M31Halpha}} \\end{figure*}  However, the spectroscopy of individual point sources in nearby galaxies suffers generally from source confusion, either from the continuum light of unresolved stars with a small angular separation from the target, or from the emission line spectra of H$\\;$II regions and diffuse nebulosities of the interstellar medium (ISM), or from both components at the same time. A common feature of these two background components is that their spectral and spatial distribution is uncorrelated and inhomogeneous, making it difficult to provide accurate background estimates when one is using conventional slit spectrographs.  Roth et al.\\ (1997) proposed the development of {\\em crowded field integral field (3D) spectroscopy} as a superior method to disentangle the spectrum of a target point source from the background in a crowded field environment. 3D spectroscopy (for short: 3D) is an established, but not yet entirely common and therefore relatively new observing technique. It provides simultaneously spectra for many spatial elements over a 2-dimensional field-of-view (FOV), see review by Boulesteix (2002). 3D has been used very successfully for kinematic mapping and integrated light stellar population studies of galaxies (de Zeeuw et al.\\ 2002). Our reason for observing {\\em point sources} may not be immediately obvious, given the technical effort to generate and record up to thousands of spectra in any single exposure. The main argument, however, is that the knowledge of the point-spread-function (PSF) of a stellar object can be used to apply PSF-fitting techniques, thus discriminating the source against the background. Such a technique would be analogous to PSF-fitting CCD photometry, which has been so successful for the construction of globular clusters CMDs and the photometric study of resolved stellar populations in nearby galaxies (Mateo 1998).  To this end, the design of the new integral field spectrophotometer PMAS was put forward (Roth et al. 2000a), featuring a relatively small FOV (max. 32$\\times$32 element lenslet array), but a unique wavelength coverage (4096 pixels) and a good response down to the UV, compared to other 3D instruments (Boulesteix 2002). First PMAS observations of point sources in M31, which are relevant for the problems discussed below, were presented by Roth et al.\\ 2002c.  In this paper, we discuss a pilot study for the novel observing technique of {\\em crowded field 3D spectrophotometry} for the observation of XPN in the local group galaxy M31. At a distance of 770~kpc towards M31 (Madore \\& Freedman 1991), an angular separation of 1~arcsec corresponds to a projected linear scale of 3.7~pc, making each XPN appear as a point source for ground-based observations. In M31, more than 400 XPN candidates have been detected with on-band/off-band imaging techniques in the prominent [O$\\;$III] $\\lambda$5007 emission line (Ciardullo et al.\\ 1989, henceforth CJFN89), enabling us to readily select targets from a large range of apparent brightness (m$_{5007} \\approx 20.5 \\ldots 24.5$) and location over the host galaxy, facilitating telescope pointing and identification, and allowing us to investigate in detail faint diagnostic lines, which tend to suffer from poor background subtraction accuracy. In a future paper (Fabrika et al., in preparation) we shall present a similar study for 3D observations of luminous stars in M33.  The paper is organized as follows: in \\S~\\ref{OBSPROBL}, we review the peculiar observational problems for the spectroscopy of XPN and explain, how 3D spectrophotometry can be used to overcome these problems. \\S~\\ref{INSTRUMENT} presents a brief overview of 3D techniques and the instruments that we have used for this work. We describe our observations in \\S~\\ref{OBSERV}, and data reduction/analysis in \\S\\S~\\ref{DATARED},\\ref{RESULTS}. We discuss our results in \\S~\\ref{DISCUSSION}, and conclude with a summary in \\S~\\ref{SUMMARY}. ",
        "conclusions": "\\label{SUMMARY}  Our pilot study has demonstrated the usefulness of integral field (3D) spectroscopy for spectrophotometry of faint planetary nebulae in M31. We have observed 5 objects, two of which (PN27, PN29) could be compared with previously published data. Of the three new spectra, one reveals that the object (PN276) had been misclassified as an XPN. Spatial extension, characteristic properties of the spectrum and positional coincidence with a previously detected X-ray source indicate with a high level of confidence that the source is a SNR.  Thanks to the full 2-dimensional spatial information of the observing technique, we make explicitely use of the a priori knowledge of the PSF, which is conveniently determined from the bright XPN [O$\\;$III] $\\lambda$5007 emission line. The PSF is used as a constraint to disentangle the point source flux from contaminating contributions that are due to diffuse or filamentary gaseous emission in the bulge of M31. We find that the low excitation XPN line intensities [O$\\;$II] I($\\lambda$3727,3729), [N$\\;$II] $\\lambda$6548,6583, and [S$\\;$II]$\\lambda$$\\lambda$6717,6731 tend to be overestimated in previous measurements with classical spectrographs. We have measured the emission line intensity of the ISM component and find that the fluxes corresponding to a typical spectrograph aperture are compatible with a level of background contamination required to explain the systematic discrepancies in line intensities, which have been reported in the literature. Simulated slit spectroscopy from a datacube for PN29 allowed us to produce arbitrary line ratios for I([O$\\;$III] 5007{\\AA})/I(H$_\\beta$), suggesting an explanation for the discrepancy reported by JC99 and RSM99. We have employed successfully the {\\em cplucy} deconvolution algorithm to achieve the best separation of point sources from a highly complex background, both in terms of spectral and of spatial variation.  A fundamental conclusion is that the limited geometry of conventional slit spectroscopy is ill-suited to solve the problem of background subtraction of faint sources embedded in complex, high surface brightness crowded fields. Our results support and quantify the previously suspected notion that such measurements suffer significantly from systematic errors. 3D spectroscopy offers the unique opportunity to make full use of the 2-dimensional spatial and 1-dimensional spectral information of a datacube to construct the best possible background model and eliminate these problems. At the other extreme, single fiber multi-object spectroscopy in high surface brightness regions is obviously a very limited technique and most vulnerable to systematic errors since there is no practical way to provide reliable background estimates in the immediate vicinity of the point source (but see Durrell et al.\\ 2003 for excellent results from a {\\em radial velocity} survey based on  [O$\\;$III] of XPN in M51). It must be stressed, however, that 3D does not necessarily provide dramatic improvements in S/N for photon-starved observations, but rather allows one to {\\em eliminate the systematic errors} discussed in this paper.  As a side-remark, we note that, in contrast to the conclusions of Allington-Smith \\& Content (1998a), residual cross-talk between adjacent spectra on the CCD is potentially dangerous for PSF-fitting algorithms, since the net effect is a distortion of the PSF (cf. Fig.~\\ref{PSFANALYSIS}), which must be expected to vary with FOV, wavelength, and time. Analogous to the reasoning of Mateo (1998) concerned with CCD photometry, the distortion is not important for the fitting of faint stars, however very much so for applications requiring a good definition of the PSF (cf. Wisotzki et al.\\ 2003).   In continuation of our successful pilot study, we are pursuing a guaranteed observing time programme targeting XPN in M31 and other local group galaxies. The analysis of our present and the expected new data by comparison with ionization models following the receipe described by JC99 will be published in a forthcoming paper.  We have begun to use 3D spectrophotometry for a number of other applications which benefit from the evaluation of spatial information and the PSF, e.g.\\ spectrophotometry of SN2002er (Christensen et al.\\ 2003; see also Aldering et al.\\ 2003 for a description of the SNAP experiment), the quadruple lensed QSO HE~0435-1223 (Wisotzki et al.\\ 2003), optical counterparts of superluminous X-ray sources (Lehmann et al., in prep.), LBV candidates in M33 (Fabrika et al., in prep.). Although our XPN study was facilitated by the availability of the bright emission in [O$\\;$III] $\\lambda$5007{\\AA}, which is a rather fortunate case to determine a well-defined PSF, we stress that the PSF fitting technique, in particular when using the two-channel {\\em cplucy} algorithm in combination with high-resolution HST images, opens a new window for solving the difficult observational problems of resolved stellar populations in nearby galaxies.  Although an upgrade of the present IFU with a 32$\\times$32 element lens array is underway, increasing the diameter of the FOV by a factor of 2, the multiplex advantage of PMAS for the rather sparsely distributed XPNe in nearby galaxies is very limited. This fact is a direct consequence of the baseline design, which favours wavelength coverage over the maximum number of spectra in terms of detector space, making PMAS a rather unique instrument. Normally, PMAS can observe only one target at a time. With a typical rate of one object per night, the efficiency may appear prohibitively low. On the other hand, we note that in a recent paper by Magrini et al.\\ (2003), out of 42 XPN candidates in M33, observed during 2 nights with the WYFFOS multi-object fiber spectrograph at WHT, only 3 spectra of were found to be of sufficient quality for an abundance analysis.  Clearly, the spectrophotometry of our XPN targets would benefit substantially from the light-collecting power and image quality of a modern 8-10m class telescope, pushing the source confusion problem to fainter limits and providing better S/N for the weak lines which are so crucial for the plasmadiagnostic analysis ([O$\\;$III]$\\lambda$4363, [S$\\;$II] $\\lambda$6717,6731, etc.). Multi-object spectroscopy of supergiant stars has already been demonstrated to be feasible out to distances beyond the local group with reasonable integration times on the order of $\\approx$10~hours (Bresolin et al.\\ 2001). The single-target approach of our pilot study would be unrealistically expensive. However, the use of a multi-object deployable IFU instrument in combination with the techniques described in this paper would eventually become an extremely powerful tool for crowded-field spectroscopy and the study of resolved stellar populations in local galaxies, one of the key science cases of the proposed new generation of Extremely Large Telescopes with diameters of 30-100m. As a promising alternative strategy, complex 3D instruments with arcmin FOV and of order 10$^5$-10$^6$ spatial elements are currently under investigation for the 2nd Generation Instrumentation Plan at the ESO-VLT (Bacon et al.\\ 2002, Morris et al.\\ 2002)."
    },
    "0311/astro-ph0311139_arXiv.txt": {
        "abstract": "In this paper we present results from our exploratory mid-IR study of Centaurus A circumnuclear environment using high-angular resolution imaging at the Magellan 6.5m telescope with the MIRAC/BLINC camera. We detected emission from a compact region surrounding the nuclear source, and obtained photometry at 8.8~$\\mu$m and in the N band. Our analysis suggests that the nuclear region is resolved with a size of $\\approx$3~pc. The mid-IR emission from this region is likely associated with cool dust with an estimated temperature of $\\sim$160~K, surrounding the central ``hidden'' AGN. We discuss the characteristics of this emission in relation to other mid-IR observations and the implications on models of dust formation in AGNs. ",
        "introduction": "The nearest radio-bright AGN is located at a distance of ``only'' $\\sim 3.5$~Mpc in the giant elliptical galaxy NGC5128 (Cen A) ({\\it vis} review by Israel 1998). The central engine powering this galaxy is thought to be a supermassive black hole enshrouded in a region containing dust and gas heated by the central source (Marconi {\\it et al.} 2001). Multi-wavelength observations of Cen A reveal many complex multi-scale structures, at scales ranging from a sub-pc to hundreds of kpc (eg. Karovska {\\it et al.} 2002). A radio and X-ray jet powered by the central source  extends across the galaxy to a distance of several hundreds of kpc from the nucleus. Optical and near-IR images of Cen A show large spatial scale (several kpc in size) dark bands stretching across the middle of the galaxy and obscuring the central region, probably due to absorption by dust and other cool material. These dust lanes are thought to be a remnant of a merger with a smaller spiral galaxy (Schiminovich {\\it et al.} 1994). Using mid-infrared and sub-millimeter wavelength observations of the  central region of Cen A, Mirabel {\\it et al.} (1999) detected dust emission from a large scale bisymmetric structure (5kpc in size) resembling a barred spiral galaxy.   Because of its proximity, Cen A provides a unique possibility for high-spatial resolution studies of the circumnuclear environment of AGNs using ground- and space-based large-aperture telescopes (0.1'' corresponds to $\\sim 1.5$~pc at the Cen A distance of $\\sim$3.5~Mpc). However, exploring the nuclear region of Cen A at optical wavelengths presents a major challenge because of the high extinction (eg. $A_V \\sim 10$~mags; Alexander et al. 1999) due to heavy  obscuration by the dust in the dust lanes.  Although the diffraction limited telescope resolution is significantly lower in the mid-IR when compared to the optical ($\\sim$ 10 times), the spectral range from 8 to 20~$\\mu$m is, nevertheless, particularly suitable for studying the emission from the region surrounding the nucleus because of at least ten times lower opacity. The lower opacity of astronomical dust at these wavelengths (Draine \\& Lee 1984) permits penetration of the optically opaque dust lanes. Most of the thermal radiation from dust is in fact emitted at these wavelengths, as seen in the mid-IR ISOCAM spectra of Cen A which show the presence of the characteristic silicate broad absorption feature at 9.8~$\\mu$m, plus a possible PAH emission line at 11.3~$\\mu$m. Alexander et al. 1999 suggested that the presence of a very deep 10~$\\mu$m silicate absorption feature in the ISOCAM spectrum could be due to self-absorption in a dusty circumnuclear region, in addition to the extinction by the dust in the dust lanes (optical depth $\\tau_\\nu \\simeq 1$ at 10~$\\mu$m).   We report here results from our mid-IR imaging of Cen A with the MIRAC/BLINC mid-IR camera at Magellan, showing that the nuclear region is resolved at a sub-arcsecond scale. The observations and the techniques for data acquisition and reduction are described in the next section. Results and analysis are presented in section 3, and discussed in relation to other available observations and models in section 4. ",
        "conclusions": "The results of our mid-IR imaging and photometry of the Cen A nuclear region suggest that the emission originates from a resolved source $\\approx$3~pc in diameter. This emission is likely to be thermal radiation from dust at $\\approx$150~K heated by  UV/optical radiation from a ``hidden'' central AGN source (e.g. Peterson 1997; Whysong \\& Antonucci 2002). It is possible that we are detecting the external cooler region, possibly surrounding hotter inner dust ($T \\sim 1000$~K), as suggested by the deep silicate absorption feature in the ISO spectra around 10~$\\mu$m.  Resolving the source of mid-IR emission in the nuclear region of Cen A provides an important key for understanding dust formation processes  and the structure of the nuclear region of AGNs and quasars. For example, the current ``unified'' model of AGNs predicts a dusty torus with a size of $\\sim 10$~Ly surrounding the accretion disk and a black hole central engine (Antonucci and Miller 1985; Krolik and Begelman 1988; Peterson 1997). The predicted size of the torus  is of the order of the measured size of the Cen A circumnuclear region using our `mid-IR imaging. We note that the estimated size of the Cen A circumnuclear region  is also similar to the size of the dusty region ($\\sim 3$~pc) surrounding the central source of NGC~1068, recently resolved AGN using the VLTI observations (ESA Press release 17/03).  The measured size of the mid-IR emitting region is also comparable to the distance from the central source at which dust will condense in an outflowing AGN wind.  As recently demonstrated by Elvis {\\it et al.} (2002), dust can be created from cooling Broad Emission Line clouds in AGN winds, and can survive in the quasar environment due to self shielding. This implies that Cen A could effectively be recycling dust from the host galaxy interstellar medium into IGM while changing the dust size and distribution and the dust-to-gas ratio. The possibility that AGN or quasars can provide an additional path for dust formation has important cosmological implications, since it could explain in a natural way the observed heavy obscuration of distant quasars (Omont et al. 2001; Elvis {\\it et al.} 2002),  and provide a new means of forming dust at early cosmological times.  The current observations do not have enough accuracy to determine the true geometry of the emitting region or to distinguish between different geometries suggested by AGN models. The observations were challenging because of the limited telescope resolution. Further multi-wavelength observations of Cen A with higher spatial resolution are crucial to explore the geometry of the extended nuclear region and derive the physical characteristics and the origin of the emission. Multi band mid-IR observations combined with spectroscopy will allow determination of the characteristics of the circumnuclear dust, including its temperature, and an estimate of the mass and mass loss rate. The results may further constrain the unified AGN model, and models of dust formation in quasar outflows."
    },
    "0311/astro-ph0311413_arXiv.txt": {
        "abstract": "We describe an optical simulation program that models a complete, coupled-cavity interferometer like those used by the Laser Interferometer Gravitational-Wave Observatory (LIGO) Project. A wide variety of interferometer deformations can be modeled, including general surface roughness and substrate inhomogeneities, with no \\textit{a priori} symmetry assumptions about the nature of interferometer imperfections. Several important interferometer parameters are optimized automatically to achieve the best possible sensitivity for each new set of perturbed mirrors. The simulation output data set includes the circulating powers and electric fields at various points in the interferometer, both for the main carrier beam and for its signal-sideband auxiliary beams, allowing an explicit calculation of the shot-noise-limited gravitational-wave sensitivity of the interferometric detector to be performed. Here we present an overview of the physics simulated by the program, and demonstrate its use with a series of runs showing the degradation of LIGO performance caused by realistically-deformed mirror profiles. We then estimate the effect of this performance degradation upon the detectability of astrophysical sources of gravitational waves. We conclude by describing applications of the simulation program to LIGO research and development efforts. ",
        "introduction": "Introduction}  The LIGO Project \\cite{rf1} is part of the current initiative to detect gravitational radiation via its perturbing effects on resonant laser interferometers. This initiative involves several large collaborations around the world, including VIRGO \\cite{rf2}, GEO \\cite{rf3}, TAMA \\cite{rf4}, and ACIGA \\cite{rf5}.  A paramount issue in all such projects is the maximization of interferometer sensitivity to detect the extremely weak signals that are expected from even the most powerful astrophysical sources \\cite{rf6}. To that end, coupled-cavity systems with multiple resonant stages will be used to maximize the shot-noise-limited signal-to-noise ratio of the gravitational wave (GW) signal readout. For these complex interferometric detectors, intensive modeling is necessary to estimate their performance in the presence of general optical imperfections.  A hierarchy of approaches has been used to estimate detector performance, each method negotiating the trade-off between accuracy and computational complexity. Analytical methods may suffice for the consideration of optical defects that can be treated as pure losses, or for some cases involving geometric \\cite{rf7} or randomized \\cite{rf8} mirror deformations. A matrix model that evaluates the coupling between the first few lowest-order TEM laser modes has been shown to be useful for the study of mirror tilts and beam displacements \\cite{rf9}; and another matrix model using discrete Hankel transforms exists for problems with axial symmetry \\cite{rf10}. Such models, which consider the exchange of power between a limited set of pre-specified modes, allow one to obtain fast results that can be used for predicting certain important interferometer behaviors (e.g., time-dependent detector responses \\cite{rf11,rf12,rf13}), at the expense of some sophistication in modeling the detailed interferometer steady-state power buildup. For the consideration of highly general optical imperfections, however, the most comprehensive method is the complete modeling of the transverse structure of the laser field wavefronts; this method simulates the electric fields on large grids, and uses (Fourier transform-based) numerical computations for the propagations of these laser beams through long cavities \\cite{rf14,rf15}. This technique is useful in the case of mirrors with complex deformations, and for mirrors with significant losses (due to absorption, scattering, and/or diffractive loss from finite-sized apertures), that violate the assumption of unitarity \\cite{rf9} for mirror operators in the matrix models, thus introducing complications into that approach \\cite{rf12}.  Grid-based simulations of intra-cavity laser fields have a long history (e.g., \\cite{rf16}), and have been applied previously (e.g., \\cite{rf17,rf18,rf19}) to the study of the interferometric GW detectors now being implemented. But simplifications are generally imposed, such as restricting the optical deformations that are studied (e.g., considering only geometric imperfections, like tilt and curvature mismatch \\cite{rf17,rf18}), and/or by modeling a relatively simple cavity system \\cite{rf19}. In this paper, we describe a simulation program (originally based upon the work of Vinet \\textit{et al.} \\cite{rf15}) that has been extended to efficiently model the complex fields that build up between realistically imperfect optical components, while resonating within complete, coupled-cavity interferometers like those used for the LIGO detectors. Versions of our program have been used for a variety of applications by the gravitational wave community, including numerous design and performance estimation tasks conducted by the LIGO group itself (see Section \\ref{s6}), as well as for collaborative investigations between LIGO scientists and other groups \\textemdash ~ such as ACIGA-LIGO efforts to explore alternative interferometer length control schemes \\cite{rf20}, and TAMA-LIGO efforts to estimate the effects of mirror imperfections \\cite{rf21} and thermal lensing \\cite{rf22} upon the performances of their future, large-scale interferometers.  A wide variety of interferometer imperfections can be modeled with our program, including mirror tilts and shifts, beam mismatch and/or misalignment, diffractive loss from finite mirror apertures, mirror surface figure and substrate inhomogeneity profiles \\textemdash ~ in particular, using deformation phase maps that are adapted from measurements of real mirror surfaces and substrates \\textemdash ~ as well as fluctuations of reflection and transmission intensity across the mirror profiles. In addition to the main (carrier frequency) laser field, auxiliary fields (i.e., radio frequency sidebands) for LIGO's heterodyne signal detection scheme are also modeled, allowing us to give absolute numbers for the shot-noise-limited GW-sensitivity of a single LIGO interferometer. Furthermore, a number of active optimization procedures are performed continuously during code execution, guaranteeing that several key parameters in the LIGO configuration will be brought to their optimum values upon program completion; the GW-sensitivity is thus maximized for each specific ``realistic'' interferometer, given the particular imperfections being simulated in that run. Considering the time needed for program execution, we note that though our principal usage of the program has been on various supercomputing platforms, a full run of the simulation code (with all options included) can be performed in a non-prohibitive amount of time on a modest Sun SPARCstation, even for significantly (and non-symmetrically) deformed mirrors. This has been achieved through a combination of fast iteration techniques, efficient parameter optimization routines, and procedures designed to carefully choose the initial guesses for laser fields that are computed via relaxation. Lastly, we note that versions of the code are available for both the first-generation LIGO and advanced-LIGO (Dual Recycling \\cite{rf23}) configurations \\cite{rf28,newref1}, though we will focus primarily upon the former in this paper.  The discussion is organized as follows: in Section \\ref{s2}, we give a description of the physical system that is modeled here (i.e., a first-generation LIGO interferometer), and the computed electric fields that are necessary for the calculation of its shot-noise-limited GW-sensitivity function. In Section \\ref{s3}, we provide an overview of the technical details of the program's operations, including the modeling of the optics, the iterative method used for computing the interferometer's electric fields, and the various optimization procedures that are performed to maximize its sensitivity. In Section \\ref{s4}, we present an array of runs with representative sets of mirror substrate and surface deformation maps adapted from real mirror measurements, in order to demonstrate how realistically-deformed mirrors will reduce the circulating power stored in a LIGO interferometer (as well as altering its modal structure), thus degrading the interferometer sensitivity and interfering with its control systems. In Section \\ref{s5}, we discuss the impact of optical imperfections upon LIGO science capabilities, by estimating how deformed-mirror effects will reduce LIGO's ability to detect astrophysical sources of gravitational waves, such as non-axisymmetric pulsars and coalescing black hole binaries \\cite{rf6}. In Section \\ref{s6}, we conclude with a discussion of various LIGO research and development initiatives to which this grid-based simulation work has contributed. ",
        "conclusions": ""
    },
    "0311/astro-ph0311625_arXiv.txt": {
        "abstract": "{ We detect near-infrared H$_2$ emission in DG Tau using the Infrared Camera and Spectrograph (IRCS) on the 8.2-m SUBARU telescope. The spectra obtained along the jet axis show that the centroidal position of the 1-0 S(1) emission is offset by 0.2'' from the star towards the jet, while those obtained perpendicular to the jet axis show a marginal extension, indicating that the emission line region has a typical width of $\\sim$0.6''. Their line profiles show a peak velocity of $\\sim$15 km s$^{-1}$ blueshifted from the systemic velocity. These results indicate that the emission originates from a warm molecular wind with a flow length and width of $\\sim$40 and $\\sim$80 AU, respectively. The line flux ratios ($I_{1-0 S(0)}/I_{1-0 S(1)}$ and an upper limit for $I_{2-1 S(1)}/I_{1-0 S(1)}$) suggest that the flow is thermalized at a temperature of $\\sim$2000 K, and is likely heated by shocks or ambipolar diffusion. The observed velocity and spatial extension suggest that the H$_2$ and forbidden line emission originate from different components of the same flow, i.e., a fast and partially ionised component near the axis and a slow molecular component surrounding it. Such a flow geometry agrees with model predictions of magneto-centrifugal driven winds. ",
        "introduction": "Understanding the mechanism of mass accretion and driving of jets/winds is one of the most important key issues for star formation theories. There is growing evidence that such outflows are powered by disk accretion. Disk accretion rates estimated through UV excess measurements are correlated with outflow signatures such as forbidden line luminosities (Hartigan, Edwards, \\& Ghandour 1995). Recent high-resolution observations have suggested that the jet appears to rotate around the flow axis (Davis et al. 2000; Bacciotti et al. 2002), supporting the scenario that the jet removes the excess angular momentum from the circumstellar disk, allowing the accretion to occur.  Among the models proposed to date, magneto-centrifugal wind models seem the most promising as the mechanism for driving the jet/wind. According to these models, magnetic and centrifugal forces act together to launch the jet/wind along magnetic field lines, from a region of the disk between $\\le$0.1 (``X-wind'' --- see Shu et al. 2000) and a few tens of AU (``disk wind'' --- see, e.g., K\\\"onigl \\& Pudritz 2000; Ferreira 1997). The magnetic field lines act as solid wires up to the ``Alfv\\'en surface'', located at less than 20 AU from the disk. These accelerate the flow particles outwards and upwards in a 'bead-on-a-wire' fashion, simultaneously removing the angular momentum from the accretion disk. Although present high-resolution facilities cannot resolve such kinematics in the innermost region, the models seem to explain some key observed properties, including (1) the observed mass ejection per mass accretion rates ($\\sim$0.1, see e.g., Calvet 1997; Richer et al. 2000), (2) motion in the jet consistent with rotation around the axis (Davis et al. 2000; Bacciotti et al. 2002), and in the same sense as the circumstellar disk or molecular envelope (Wiseman et al. 2001; Testi et al. 2002).  These models also predict the presence of slow and poorly collimated streamlines surrounding a well-collimated jet. Although atomic lines at UV-to-IR wavelengths have been powerful tools to study the jet/wind over decades (see e.g., Eisl\\\"offel et al. 2000a and Hartigan et al. 2000 for reviews), they may originate from only a limited fraction of the entire flow close to the jet axis (e.g., Li \\& Shu 1996; Cabrit, Ferreira, \\& Raga 1999; Shang et al. 2002). The presence of an ``unseen'' wide-angled component has been suggested by the presence of a cavity in the molecular core/envelope (see Davis et al. 2002, and references therein), and a broad range of morphology of molecular bipolar outflows (see e.g., Lee et al. 2000, 2001). In addition, Davis et al. (2002) and Saucedo et al. (2003) revealed the presence of shock-heated H$_2$ along the cavity wall in some young stellar objects (YSOs), presumably due to interaction with a wide-angled wind. Direct detection of such a flow component is desired to test the validity of the magneto-centrifugal wind scenario, to investigate the angular transfer problem in detail, and to understand the nature of molecular outflows.  We report the detection of near-infrared H$_2$ emission in DG Tau, which originates from a warm molecular wind close to the star. The object is one of the most active T Tauri stars known, and exhibits an energetic jet like Class I protostars (e.g., Eisl\\\"offel \\& Mundt 1998; Davis et al. 2003). Since the object is not heavily embedded like Class I protostars ($A_V$=2.2 --- Muzerolle, Calvet, \\& Hartmann 1998), it has been observed as one of the best targets for studying outflowing activities close to the star (see, e.g., Solf \\& B\\\"ohm 1993; Kepner et al. 1993; Bacciotti et al. 2000, 2002; Pyo et al. 2003). In particular, Bacciotti et al. (2000, 2002) observed the forbidden line outflow using the {\\it Hubble Space Telescope}, and revealed the presence of an \"onion-like\" kinematic structure, i.e., a continuous bracketing of the higher speed central flow within the lower speed, less collimated, broader flow.  In this paper, we show that the warm molecular wind appears as a part of the ``extensive'' onion-like kinematic structure observed in the forbidden line outflow. The rest of the paper is organised as follows: we describe the details of the observations and results in Sect. 2; we compare our results with previous observations of near-infrared H$_2$ emission in other YSOs, and investigate the driving, location of the launching region, and excitation/heating of the warm molecular wind in Sect. 3; we then give our conclusions in Sect. 4. Throughout the paper, we adopt a distance to the object of 140 pc based on {\\it HIPPARCOS} observations (Wichmann et al. 1998). ",
        "conclusions": "We detected near-infrared H$_2$ emission in DG Tau, one of the most active T Tauri stars known. The spatial extension along and perpendicular to the jet, and measured blueshift of 15 km s$^{-1}$, indicate that the emission originates from a warm molecular wind with a flow length and width of $\\sim$40 and $\\sim$80 AU, respectively. Via detailed comparison of the spatial scale and velocity field, we conclude that H$_2$ and forbidden line emission originate from different regions of the entire flow, which includes (1) fast, well-collimated and partially ionised streamlines, and (2) slow, poorly-collimated, and molecular streamlines. Such a flow geometry agrees with predictions of magneto-centrifugal disk wind models. Although it could also be explained by the entrainment of the surrounding material by the collimated jet, we consider it unlikely.  The measured line flux ratios ($I_{1-0 S(0)}/I_{1-0 S(1)}$ and the upper limit of $I_{2-1 S(1)}/I_{1-0 S(1)}$) indicate that the flow is thermalized at a temperature of $\\sim$2000 K. The flow is likely to be heated by shocks or ambipolar diffusion."
    },
    "0311/astro-ph0311555_arXiv.txt": {
        "abstract": "{We present our Disk Mass project as the main science case for building a new fiber IFU-module for the PMAS spectrograph, currently mounted at the Cassegrain focus of the 3.5m telescope on Calar Alto. Compared to traditional long-slit observations, the large light collecting power of 2-dimensional Integral Field Units dramatically improves the prospects for performing spectroscopy on extended low surface brightness objects with high spectral resolution. This enables us to measure stellar velocity dispersions in the outer disk of normal spiral galaxies. We describe some results from a PMAS pilot study using the existing lenslet array, and provide a basic description of the new fiber IFU-module for PMAS. ",
        "introduction": "A major roadblock in testing galaxy formation models is the disk-halo degeneracy; density profiles of dark matter haloes as inferred from rotation curve decompositions depend critically on the adopted M/L of the disk component (Figure 1). An often used refuge to circumvent this degeneracy is the adoption of the maximum-disk hypothesis (van Albada \\& Sancisi 1986). However, this hypothesis remains unproven. Recently, Bell \\& de Jong (2001) have shown that stellar population synthesis models yield plausible {\\it relative} measurements of stellar M/L in old disks, but uncertainties in the IMF prevent an {\\it absolute} measurement of total disk M/L from photometry. Another tool to determine the M/L, and specifically whether disks are maximal, is the Tully-Fisher relation, e.g. by looking for offsets between barred vs. un-barred galaxies, but this too is only a {\\it relative} measurement. Evidently, none of these methods are suited to break the degeneracy, and without an independent measurement of the M/L of the stellar disk, it is not possible to derive the structural properties of dark matter haloes from rotation curve decompositions.  An absolute measurement of the M/L can be derived from the vertical component $\\sigma_{\\rm z}$ of the stellar velocity dispersion. For a locally isothermal disk, $\\sigma_{\\rm z} = \\sqrt{\\pi G(M/L)\\mu z_o}$, with $\\mu$ the surface brightness, and $z_0$ the disk scale height. The latter is statistically well-determined from studies of edge-on galaxies (de Grijs \\& van der Kruit 1996; Kregel et al. 2002). Thus, $\\sigma_{\\rm z}$ provides a direct, kinematic estimate of the M/L of a galaxy disk and can break the disk-halo degeneracy.  This kinematic approach has been attempted before (e.g. Bottema 1997), but with long-slit spectroscopy on significantly inclined galaxies. These observations barely reached out to 1.5 disk scale-lengths, required significant radial binning, and because of the high inclinations of these galaxies, the measured velocity dispersions required large and uncertain corrections for the tangential ($\\sigma_\\phi$) and radial ($\\sigma_{\\rm r}$) components of the stellar velocity dispersion ellipsoid.  With the advent of 2-dimensional spectroscopy using Integral Field Units (IFU), the observational prospects for measuring $\\sigma_{\\rm z}$ have improved dramatically. This is because a wide-field (1$^\\prime$$\\times$1$^\\prime$) IFU can collect and pipe much more light to a spectrograph than a single 2$^{\\prime\\prime}$ wide long-slit.  The power of IFUs with face-on galaxies lies in the ability to azimuthally average many fibers. This yields clean $\\sigma_{\\rm z}$ measurements well beyond 2 disk scale-lengths where contamination from bulge stars is negligible and the rotation curve of the stellar disk has reached its flat part. Here, we describe our Disk Mass project which employs two custom built IFUs to measure $\\sigma_{\\rm z}$. ",
        "conclusions": ""
    },
    "0311/astro-ph0311249_arXiv.txt": {
        "abstract": "We present CCD $UBVI$ photometry of the nearby, intermediate age open cluster NGC 6939. Using the synthetic Colour - Magnitude Diagrams technique we estimate the following parameters: age between 1.3 and 1.0 Gyr (depending on whether or not overshooting is considered), reddening 0.34 $\\leq$ \\ebv $\\leq$ 0.38 and distance modulus 11.3 $\\leq$ \\mmm $\\leq$ 11.4. ",
        "introduction": "A vast amount of  information on the formation and evolution of our Galaxy can be gathered from the study of open clusters (OCs). In particular, old OCs may be useful to derive information not only on the present day situation, but also on the time evolution of the disc, since their ages cover the whole existence of the disc, reaching to about 10 Gyr (Friel 1995). In order to fully exploit the information provided by OCs we must derive for them in a homogeneous way: accurate absolute ages (and a consistent age ranking), distances, reddenings, and metal content. With this aim we are studying a number of old OCs (see Bragaglia \\& Tosi 2003 and Sandrelli et al. 1999  for references), and we now add another cluster to our sample.  Given its relative proximity, NGC 6939 (C2030+604) has been the target of several studies in the past: the first bibliographic entry is 80 years ago (Kustner 1923) but, surprisingly, the first CCD data appeared only in 2002 (Rosvick \\& Balam 2002, hereafter RB02). As usual, the cluster parameters found in literature do not agree with each other, and we present in this paper new and improved determinations for this intermediate age open cluster located at RA(2000) = 20:31:32, DEC(2000) = +60:39:00, or l = 95.88,  b = 12.30.  Photometry has been previously presented by several authors, but old photographic photometry only reached about one magnitude below the main sequence Turn-Off. Mermilliod, Huestamendia, \\& del Rio (1994) took UBV photoelectric photometry of 37 member stars  all in the red clump phase, with the intent of discriminating between different evolutionary models (with or without overshooting) by comparison with theoretical  isochrones.  It turned out that most of the bright stars lying in the cluster direction  are indeed cluster members (their  Table 1 lists only 4 non-members among 45 objects). The recent work by Rosvick \\& Balam (2002) has presented the first deep BVI CCD data.  They used the 1.85m Dominion Astrophysical Observatory telescope, covering the central part of NGC 6939 (more or less our field A, see later). Their CMDs show considerable scatter, which they do not attribute to contamination from field stars (even if they could not prune their diagrams on the main sequence, since no proper motion survey on this cluster has reached so deep), but to differential reddening [\\ebv = 0.29 to 0.41], which also influences the distance derivation [$(m-M)_V$ = 12.21 to 12.39]. Using the Girardi et al. (2000) solar metallicity isochrones, they obtain a cluster age of 1.6 $\\pm$ 0.3 Gyr  Canterna et al. (1986)  give a metal abundance of [Fe/H] = --0.1 based on Washington photometry, a value confirmed by the revision of this method done by Geisler, Clari\\'a \\& Minniti (1991), who give  --0.13. Thogersen, Friel \\& Fallon (1993) determine a cluster mean radial velocity of $-42 \\pm 10$ kms$^{-1}$, and [Fe/H]=--0.14$\\pm$0.13, from low resolution spectra of four cluster giants. The metallicity has recently been revised to [Fe/H]=--0.19$\\pm$0.09 by Friel et al. (2002). Other measures of radial velocities were given by Geisler (1988), who found $-24.1 \\pm 3.1$ based on medium resolution spectroscopy of 8 stars; by  Milone (1994) who, on the basis of higher resolution spectra of 26 giants, found a mean cluster radial velocity of $-19.89 \\pm 0.19$ kms$^{-1}$; and by  Glushkova \\& Rastorguev (1991), who derived $-19.3$ as average for 29 cluster members. These last also determined proper motions (hence membership) for 136 stars in the cluster vicinity. Zhao et al. (1985) determined membership probabilities based on proper motion for more than 200 stars in the cluster field, and we used their information to confirm the turn-off and red clump positions in our data.  Dutra and Bica (2000) cite a  reddening  \\ebv=0.39 based on the Schlegel, Finkbeiner \\& Davis (1998) maps.  At least two groups used existing literature data to re-derive relevant properties. Twarog, Ashman \\& Anthony-Twarog (1997), in their work trying to derive on homogeneous grounds the properties of Galactic open clusters, found \\ebv = 0.45 to 0.49, [Fe/H] = 0.03, distance modulus, $(m-M)_V$ = 12.45.\\footnote{The value listed in their Table 2 is 13.05, but Bruce Twarog kindly informed us that their correct modulus is 12.45, because, due to a typo, 13.05 is actually not the modulus but the mean visual magnitude of the red clump.} Carraro, Ng \\& Portinari (1998) derived, using the synthetic CMD techniqe with the Padova tracks and the Thogersen et al. (1993) metallicity, an age of 1.4 Gyr, and  R$_{\\rm GC}$ = 9.7 kpc.  Finally, Robb \\& Cardinal (1998) tried to detect variables in this cluster and actually found six, of which at least two are eclipsing variables.  We describe our data in Section 2, and the resulting colour - magnitude diagrams in Section 3. Section 4 is devoted to the derivation of the cluster parameters, while summary and conclusions are presented in Section 5.  \\begin{figure} \\vspace{14cm} \\special{psfile=andreuzzi_fig1.ps vscale=90 hscale=90 voffset=0 hoffset=40} \\caption{Position of our three fields (A, B and external), each one of 9.4 $\\times$9.4 arcmin; the map is 15 x 50 arcmin, and is oriented with north up and East left.} \\label{fig-map} \\end{figure}  \\begin{figure} \\vspace{12.5cm} \\special{psfile=andreuzzi_fig2.ps vscale=70 hscale=80 voffset=-120 hoffset=-20} \\caption{Comparison of our U,B,V photometry with the photoelectric values given in Mermilliod et al. (1994).} \\label{fig-photo} \\end{figure}  \\begin{figure} \\vspace{12cm} \\special{psfile=andreuzzi_fig3.eps vscale=75 hscale=75 voffset=-135 hoffset=-115} \\caption{B, U-B CMD for NGC 6939.The mean errors per interval of magnitude B are also plotted.} \\label{fig-cmd1} \\end{figure} ",
        "conclusions": "We have examined UBVI CCD data of the open cluster NGC 6939, and determined its distance, reddening, age and (approximate) metallicity making use of the synthetic CMDs technique and three different sets of theoretical tracks. We found that the best combination of the cluster parameters is: solar abundance, 0.34 $\\leq$ \\ebv $\\leq$ 0.38, 11.3 $\\leq$ \\mmm $\\leq$ 11.4, 1.0 Gyr $\\leq$ age $\\leq$ 1.3 Gyr. We also found indication of the presence of binary systems.  Our findings do not agree completely with literature data, although quite similar in some cases (see Introduction). For instance, when comparing with RB02, we do not find a substantial differential reddening and our age is  slightly lower than theirs. The age difference is likely due to the use of different versions of the Padova tracks: they adopt the Girardi et al. (2000) solar metallicity isochrones, while we prefer the older BBC tracks .  This seems to be confirmed by our best concordance with Carraro et al. (1998) results, based on our same BBC tracks. Our intrinsic distance modulus is slightly larger than that (10.95) derivable from the corrected $(m-M)_V$ = 12.45 of Twarog et al. (1997) assuming their mean reddening \\ebv = 0.47. However, if our mean reddening \\ebv = 0.36 is applied to their $(m-M)_V$, the resulting intrinsic modulus is 11.3, identical to ours.  Regarding the differential reddening claimed by RB02, we cannot substantiate their values. When looking at the CMDs we derive from our data (Sect. 3.3) we do not see any clear indication of its presence. Furthermore, the synthetic CMDs are able to reproduce the cluster sequences without invoking differential reddening; we think that a $\\Delta$ E(B-V) as large as suggested by RB02 would have broadened the MS much more than observed. We agree, however, that a small variation of $\\Delta$\\ebv $\\leq$0.03 over the field may allow us to well reproduce the TO morphology. We suggest that the presence of a fraction around 30\\% of binary stars in the cluster would further improve the CMD reproduction. Indeed, the broadening effect on the MS of binaries and of differential reddening differ from each other. Binarism affects both magnitudes and colours and leads to an increasing widening of the MS towards fainter magnitudes, with the binary MS converging on the single MS at the TO region. Vice versa, differential reddening shifts the colours by the same amount at all magnitudes, thus making the effect much more evident in the brighter, tighter TO regions.   The rather poor fit of the lower MS obtained for all the models employed is not unheard of, and cannot be taken as a general failure to measure cluster properties. Very recently Grocholski \\& Sarajedini (2003) presented a comparative study of isochrone fits to CMDs in the visual and infrared filters. They took five sets of theoretical models, compared them in the luminosity - temperature plane, and proceeded to fit the observed CMDs of six open clusters of assorted ages and metallicities with the isochrones transformed into the observational planes. The soundness of the comparison of different evolutionary codes is  probably undermined by the fact that Grocholski \\& Sarajedini (2003) use the isochrones as transformed from the theoretical plane (luminosity and temperature) to the observational one (magnitudes and colours) by each group, i.e. following quite different prescriptions. This introduces additional effects,  which we prefer to avoid using a single transformation for all tracks. Unfortunately, we have no theoretical set in common with them (they use yet another version of the Padova tracks, by Girardi et al. 2002), but their results are quite interesting: there is no single set that is able to match all CMDs in all colours, at all ages, and for all metallicities. All of them behave rather well in some cases, and do not fit well in others. However, when the observed CMDs are deep enough, the fit to the lower main sequence is often (but curiously not always) bad; this happens in general around M$_{\\rm V} \\simeq $ 8 (the same we find for NGC 6939) for the visual colours, and is attributed by Grocholski \\& Sarajedini (2003) to some still missing ingredient in the model atmospheres for low mass stars.  Finally, we wish to stress that the roughly solar abundance we find  for NGC 6939 is more an indication than a well founded measurement. It is encouraging that it's in  reasonable agreement with what other methods find, but high resolution spectroscopy and detailed abundance analysis should be performed to determine the cluster elemental composition.    \\bigskip\\noindent ACKNOWLEDGEMENTS  This work is based on observations made with the Italian Telescopio Nazionale Galileo (TNG) operated on the island of La Palma by the Centro Galileo Galilei of the INAF (Istituto Nazionale di Astrofisica) at the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias. We thank M. Bellazzini for the data acquired on August 2001 and P. Ventura for the FST models. We are grateful to B. Twarog for having informed us of the correct value of their distance modulus and for having pointed out a problem of ours. The bulk of the simulation code was originally provided by L.Greggio. Financial support to this project has come from the MURST-MIUR through Cofin98 ''Stellar Evolution'', and Cofin00 ''Stellar Observables of Cosmological Relevance''. This research has made use of the Simbad database, operated at CDS, Strasbourg, France. Finally we acknowledge the use of the valuable BDA database, maintained by J.-C. Mermilliod, Geneva.  \\begin{figure} \\vspace{5cm} \\special{psfile=andreuzzi_fig10.ps vscale=75 hscale=75 voffset=-130 hoffset=-95} \\caption{V luminosity function of the reference observed stars of Fig.\\ref{empcmd} (dots) and of a typical synthetic case in good agreement with the data (line). Errorbars reflect the counting uncertainty as derived from a Poisson distribution.} \\label{lf} \\end{figure}"
    },
    "0311/astro-ph0311280_arXiv.txt": {
        "abstract": "{% We assess the cosmological variability of the fine-structure constant $\\alpha$ from the analysis of an ensemble of \\ion{Fe}{ii}~$\\lambda1608$, $\\lambda2344$, $\\lambda2374$, $\\lambda2383$, $\\lambda2587$, and $\\lambda2600$ absorption lines at the redshift $z=1.15$ toward the QSO HE~0515--4414 by means of the standard many-multiplet (MM) technique and its revision based on linear regression (RMM). This is the first time the MM technique is applied to exceptional high-resolution and high signal-to-noise QSO spectra recorded with the UV-Visual Echelle Spectrograph (UVES) at the ESO Very Large Telescope (VLT). Our analysis results in $\\langle\\Delta\\alpha/\\alpha\\rangle_\\mathrm{MM}=(0.1\\pm1.7) \\cdot10^{-6}$ and $\\langle\\Delta\\alpha/\\alpha\\rangle_\\mathrm{RMM}=(-0.4 \\pm1.9\\pm2.7_\\mathrm{sys})\\cdot10^{-6}$, which are the most stringent bounds hitherto infered from an individual QSO absorption system. Our results support the null hypothesis $\\langle\\Delta\\alpha/\\alpha\\rangle=0$ at a significance level of 91~percent, whereas the support for the result $\\langle\\Delta\\alpha/\\alpha\\rangle=-5.7\\cdot10^{-6}$ presented in former MM studies is 12~percent. ",
        "introduction": "% \\label{sc:intro} \\enlargethispage{2pt}  Modern 10~m class telescopes equipped with instruments like the High-Resolution Echelle Spectrograph (HIRES) at the Keck Observatory or the UV-Visual Echelle Spectrograph (UVES) at the ESO Very Large Telescope (VLT) facilitate the accurate observation of QSO absorption (or emission) lines in order to study the hypothetical variability of fundamental physical constants like the fine-structure constant $\\alpha=e^2/(4\\pi\\epsilon_0\\hbar c)$ or the proton-electron mass ratio $\\mu=m_\\mathrm{p}/m_\\mathrm{e}$. The interest in these studies is motivated by the unification theories incorporating varying fundamental constants \\citep[for a review see, e.g.,][]{UzanJ_2003}.  From the astronomical point of view the cosmological variability of the fine-structure constant is assessed as \\begin{equation} \\Delta\\alpha/\\alpha = (\\alpha_z-\\alpha_0)/\\alpha_0, \\label{eq:dalpha} \\end{equation} where $\\alpha_0$ and $\\alpha_z$ denote the values of the fine-structure constant in the laboratory and the specific absorption (or emission) line system at redshift $z$, respectively. While observational studies based on the fine-structure splitting of intergalactic alkali-doublet (AD) absorption lines \\citep[e.g.,][]{LevshakovS_1994,MurphyWFPW_2001} and intrinsic QSO emission line doublets \\citep{BahcallSS_2003} have provided robust upper bounds on $|\\Delta\\alpha/\\alpha|$, \\citet[hereafter MFW]{MurphyFWDP_2003} have recently detected a non-zero expectation value $\\langle\\Delta\\alpha/\\alpha\\rangle$ for a sample of 143 complex metal absorption systems identified in QSO spectra recorded with the Keck/HIRES: $\\langle\\Delta\\alpha/\\alpha\\rangle=(-5.7\\pm1.1) \\cdot10^{-6}$ in the redshift range $0.2<z<4.2$. This remarkable statistical evidence for a cosmological variation of the fine-structure constant is achieved by means of the many-multiplet (MM) technique, which is a generalization of the AD method incorporating the multi-component profile decomposition of many transitions from different multiplets of different ionic species \\citep[hereafter MWF, and references therein]{DzubaFW_1999,MurphyWF_2003}. While the MM technique considerably improves the formal accuracy, it also shows some immanent deficiencies. \\citet{LevshakovS_2003} illustrates in detail that the prerequisite assumption that the spatial distribution is the same for all ionic species is not valid in typical intergalactic absorption systems. Consequently, it is more reliable to apply the MM technique to samples of absorption lines arising from only one ionic species. Furthermore, in comparison to the AD calculations the incorporation of more transitions over a wider wavelength range results in a stronger susceptibility to systematic effects. Nevertheless, sources of error like wavelength miscalibration, spectrograph temperature variations, atmospheric dispersion, and isotopic or hyperfine-structure effects do demonstrably not explain the detected non-zero expectation value \\citep[MWF, MFW]{MurphyWFCP_2001}. The observational discrepancy grows since the radioactive decay rates of certain long-lived nuclei deduced from geophysical and meteoritic data provide a stringent bound, $\\langle\\Delta\\alpha/\\alpha\\rangle=(8\\pm8)\\cdot10^{-7}$, back to the epoch of Solar system formation, $z\\le0.45$ \\citep{OlivePQMV_2003}. \\enlargethispage{2pt}  In the optical spectroscopy, the most striking deficiency inherent to all decomposition techniques are non-linear inter-parameter correlations preventing the accurate optimization of the model parameters and possibly causing ambigous results. In fact, the spectral resolution attained in QSO observations at the 10~m class telescopes is still not sufficient to resolve the metal lines with an expected minimum thermal width of about 1~km\\,s$^{-1}$. Clearly, in order to solve the profile decomposition problem spectral observations with the highest possible resolution and the highest possible signal-to-noise ratio are desirable.  \\begin{table} \\centering \\caption{Atomic data of the \\ion{Fe}{ii} transitions between $\\lambda1608$ and $\\lambda2600$. The laboratory wavelengths $\\lambda$, oscillator strengths $f$, and relativistic correction coefficients $q$ are excerpted from \\citetalias[Table~2]{MurphyWF_2003} and \\citet[Table~1]{DzubaFKM_2002}. The sensitivity coefficient $Q$ is defined in Sect.~\\ref{sc:rmm}. Estimiated errors are indicated in parentheses} \\begin{tabular}{@{}lllll@{}} \\hline \\hline Tr.\\rule[-5pt]{0pt}{15pt} & $\\lambda$ (\\AA) & $f$ & $\\phantom{-}q$ (cm$^{-1}$) & $\\phantom{-}Q$\\\\ \\hline $1608$\\rule{0pt}{10pt} & $1608.45080\\,(8)$ & $0.058$ & $-1300\\,(300)$ & $-0.021\\,(5)$\\\\ $2344$ & $2344.2130\\,(1)$ & $0.114$ & $\\phantom{-}1210\\,(150)$ & $\\phantom{-}0.028\\,(4)$\\\\ $2374$ & $2374.4603\\,(1)$ & $0.0313$ & $\\phantom{-}1590\\,(150)$ & $\\phantom{-}0.038\\,(4)$\\\\ $2383$ & $2382.7642\\,(1)$ & $0.320$ & $\\phantom{-}1460\\,(150)$ & $\\phantom{-}0.035\\,(4)$\\\\ $2587$ & $2586.6496\\,(1)$ & $0.06918$ & $\\phantom{-}1490\\,(150)$ & $\\phantom{-}0.039\\,(4)$\\\\ $2600$ & $2600.1725\\,(1)$ & $0.23878$ & $\\phantom{-}1330\\,(150)$ & $\\phantom{-}0.035\\,(4)$\\\\ \\hline \\end{tabular} \\label{tb:adata} \\end{table}  In this study, we present exceptional high-resolution and high signal-to-noise spectra of the notably bright intermediate redshift QSO HE~0515--4414 \\citep[$z=1.73$, $B=15.0$]{ReimersHRW_1998} recorded with the VLT/UVES. The spectra reveal a multi-component complex of metal absorption lines associated with a sub-damped Lyman-$\\alpha$ (sub-DLA) system at the redshift $z=1.15$ \\citep{VargaRTBB_2000,QuastBR_2002,ReimersBQL_2003}. We analyze a homogenous subsample of \\ion{Fe}{ii} $\\lambda1608$, $\\lambda2344$, $\\lambda2374$, $\\lambda2383$, $\\lambda2587$, and $\\lambda2600$ absorption lines by means of the standard and a revised MM technique in order to assess the cosmological variability of the fine-structure constant. ",
        "conclusions": "%  Our results strongly support the null hypothesis of a non-varying fine-structure constant, but do not contradict the MFW result $\\langle\\Delta\\alpha/\\alpha\\rangle=-5.7\\cdot10^{-6}$ at a significance level higher than $88$ percent. Nevertheless, we conclude: (i)~The MM technique has the capability to provide stringent bounds on $\\langle\\Delta\\alpha/\\alpha\\rangle$ even if the analysis is restricted to the \\ion{Fe}{ii} lines only. (ii)~The RMM technique is illustrative and methodically more transparent than the standard MM technique. In addition, the regression analysis facilitates the consideration of systematic errors inherent to the wavelength calibration of QSO spectra. (iii)~The accuracy of the individual $\\Delta\\alpha/\\alpha$ assessment is principally limited by systematic errors inherent to the wavelength calibration. The accuracy attainable with high-quality QSO spectra recorded with the VLT/UVES is limited to~$10^{-5}$. (iv)~Optically thick profiles are susceptible to systematic effects biasing the expectation value $\\langle\\Delta\\alpha/\\alpha\\rangle$. (v)~The analysis of an extensive homogenous sample of \\ion{Fe}{ii} absorption lines is inevitable and will provide an independent and crucial test of the MFW result.  In particular, HE~0515--4414 is the brightest known intermediate redshift QSO in the sky and is therefore predestinated for spectroscopy with the new High Accuracy Radial velocity Planet Searcher (HARPS) operated at the ESO La Silla 3.6~m telescope. This very high-resolution spectrograph is specified to provide an efficient wavelength calibration facilitating the performance of radial velocity measurements with an accuracy of better than 1~m\\,s$^{-1}$ \\citep{PepeMR_2002}. Presumably, HARPS will improve the accuracy of the individual $\\Delta\\alpha/\\alpha$ assessment by more than an order of magnitude."
    },
    "0311/astro-ph0311005_arXiv.txt": {
        "abstract": "We explore the types of slow-roll inflationary potentials that result in scalar perturbations with a {\\em constant} spectral index, {\\em i.e.,} perturbations that may be described by a single power-law spectrum over all observable scales.  We devote particular attention to the type of potentials that result in the Harrison--Zel'dovich spectrum. ",
        "introduction": "\\label{sec:intro}  Inflation, a cornerstone of the modern framework for understanding the  early universe \\cite{guth81,lrreview}, predicts the initial conditions for the formation of structure and the cosmic microwave background (CMB) anisotropies.  During inflation, the primordial scalar (density) and tensor (gravitational wave) perturbations generated by quantum fluctuations are redshifted beyond the Hubble radius, becoming frozen as perturbations in the background metric \\cite{muk81,hawking82,starobinsky82,guth82,bardeen83}.  However, even when there is only one scalar field --- the {\\it inflaton} --- the number of inflation models proposed in the literature is large \\cite{lrreview}. Determination of the properties of the scalar perturbations and tensor perturbations from CMB and large-scale structure observations allows one to constrain the space of possible inflation models \\cite{dodelson97,kinney98a,probes,wmapinf,barger,kkmr,liddleleach}.  It is often adequate to characterize inflationary perturbations in terms of four quantities: the scalar and tensor power spectra, $\\mathcal{P}_{\\mathcal{R}}$ and $\\mathcal{P}_{g}$, and the scalar and tensor spectral indices $n$ and $n_T$.  In this paper we focus on the {\\em scalar} spectral index which, unless explicitly indicated otherwise, we refer to simply as the `spectral index'. Successful inflation models predict $n$ close to $1$ (the so-called Harrison--Zel'dovich spectrum), and $n$ typically has a small scale dependence. The best data available to date, combining the Wilkinson Microwave Anisotropy Probe \\cite{wmap} and Sloan Digital Sky Survey \\cite{SDSS} data sets, indicate that the evidence for anything other than a scale-invariant spectra is marginal at best, with no evidence for significant running of the scalar spectral index \\cite{tegmarksloan}. Moreover, one of us has recently argued that when information criteria are used to carry out cosmological model selection based on the current data sets available, then the best present description of cosmological data uses a scale-invariant ($n=1$) spectrum \\cite{liddleparameters}. It therefore makes sense to be considering the inflationary potentials associated with that spectrum.  It is known that inflaton potentials $V(\\phi)= \\exp(-\\alpha\\phi)$ for constant $\\alpha^2<2$ lead to perturbation spectra that are exact power laws, {\\it i.e}.~$n$ is a constant \\cite{lucchin}. However, there has not yet been a systematic analysis of the types of inflaton potentials that yield constant $n$. Here we take a first step in that direction, classifying those potentials within the framework of the slow-roll approximation \\cite{SRref}.  In the next section the basic results employed to calculate the properties of the perturbation spectrum using the slow-roll parameterization of the inflaton potential are reviewed. In Sec.~\\ref{sec:method} two exact differential equations connecting the potential and the field to the slow-roll parameters are derived and the general method used to calculate all the relevant cosmological quantities is outlined. In Sec.~\\ref{sec:app} this method is applied to the determination of the inflationary potential yielding a $k$-independent density spectral index: both the Harrison--Zel'dovich $(n=1)$ and the general $(n=1-2n_0^2)$ case are considered to lowest order and to next order in the slow-roll parameter approximation. In Sec.~\\ref{Sect:flow} the flow of $\\epsilon$ is examined to understand the number of solutions that arise. The conclusions are contained in Sec.~\\ref{sec:dis}. ",
        "conclusions": "\\label{sec:dis}  The analysis that has been carried out shows that inflaton potentials yielding the Harrison--Zel'dovich flat spectrum can be determined to \\textit{lowest-order} and \\textit{next-order} approximation in the slow-roll parameters. Similarly, potentials producing a $k$-independent spectral index slightly different from unity have been derived to \\textit{lowest-order} and to \\textit{next-order}.  It is also possible to speculate that the same procedure can be carried out to any order of expansion in the slow-roll parameters. This is because the implications of the spectral index $k$-independence are not as trivial as they may seem at first glance. Notice in fact that every time a higher approximation order is assumed, new slow-roll parameters will appear in the expression for the spectral index: going from lowest-order to next-order, for example, $\\xi^2$ was introduced. This is hardly surprising, though, because these new parameters just correspond to higher derivatives of $V(\\phi)$ or $H(\\phi)$ (whatever is the degree of freedom chosen to express the slow-roll parameters) and a higher order treatment necessarily needs to take into account more derivative terms of the potential. However, the requirement of the spectral index to be $k$-independent implies not only a particular value for $n(k)$ but also that all its derivatives are equal to zero: \\begin{equation}\\label{nn=0bis} \\frac{d^i n(k)}{d (\\ln k)^i}=0, \\textrm{ with } i=1,2,... \\end{equation}  Furthermore, the expression for the ${(l_0-1)}^{th}$ derivative of the spectral index contains slow-roll parameters up to the $l_0^{th}$ one. So once the approximation order $l_0$ is chosen, the problem is characterized by $l_0+1$ parameters and $l_0$ equations of constraint relating them. This allows the expression of \\textit{all} the slow-roll parameters $\\lambda_H^l$ as functions of $\\epsilon$. The choice of $\\epsilon$ is not arbitrary, because once the expression for $(\\eta-\\epsilon)$ appropriate for the approximation level assumed is derived, the \\textit{exact} expressions for $d\\epsilon/d\\phi$ and for $d\\ln V /d\\epsilon$, Eqs.~(\\ref{depsdphi2}) and (\\ref{dlnVdeps}), can be exploited to compute $\\phi$ and $V$ as functions of $\\epsilon$ thus yielding the map $\\phi \\rightarrow V(\\phi)$."
    },
    "0311/astro-ph0311143_arXiv.txt": {
        "abstract": "The maximum value of the ratio of the tangential force to the mean background radial force is a useful quantitative measure of the strength of nonaxisymmetric perturbations in disk galaxies. Here we consider the distribution of this ratio, called $Q_g$, for a statistically well-defined sample of 180 spiral galaxies from the {\\it Ohio State University Bright Galaxy Survey} and the {\\it Two Micron All-Sky Survey}. $Q_g$ can be interpreted as the maximum gravitational torque per unit mass per unit square of the circular speed, and is derived from gravitational potentials inferred from near-infrared images under the assumptions of a constant mass-to-light ratio and an exponential vertical density law.  In order to derive the most reliable maximum relative torques, orientation parameters based on blue-light isophotes are used to deproject the galaxies, and the more spherical shapes of bulges are taken into account using two-dimensional decompositions which allow for analytical fits to bulges, disks, and bars. Also, vertical scaleheights $h_z$ are derived by scaling the radial scalelengths $h_R$ from the two-dimensional decompositions allowing for the type dependence of $h_R/h_z$ indicated by optical and near-infrared studies of edge-on spiral galaxies. The impact of dark matter is assessed using a ``universal rotation curve\" parametrization, and is found to be relatively insignificant for our sample. In agreement with a previous study by Block et al. (2002), the distribution of maximum relative gravitational torques is asymmetric towards large values and shows a deficiency of low $Q_g$ galaxies. However, due to the above refinements, our distribution shows more low $Q_g$ galaxies than Block et al. We also find a significant type-dependence in maximum relative gravitational torques, in the sense that $Q_g$ is lower on average in early-type spirals compared to late-type spirals. The effect persists even when the sample is separated into bar-dominated and spiral-dominated subsamples, and also when near-infrared types are used as opposed to optical types. ",
        "introduction": "Nonaxisymmetric features are a pervasive and complex aspect of disk galaxies. In normal, relatively non-interacting galaxies, these features are in the forms of bars or spirals. It is well-known that the presence of nonaxisymmetric structures in galaxy disks can impact the evolution of morphology. For example, bars may drive spiral density waves (Kormendy and Norman 1979), generate resonance rings of gas (Schwarz 1981; Buta \\& Combes 1996), impact abundance gradients (Martin \\& Roy 1994), or induce gas inflow that may lead to bar destruction and bulge growth (Norman, Sellwood, \\& Hasan 1996). A spiral may trigger shocks, inducing star formation (Roberts, Roberts, \\& Shu 1975), or may rearrange stochastically-induced star-forming regions into a more organized pattern (McCall 1986). It is clear that nonaxisymmetric features, with their associated pattern speeds and resonances, are extremely important to galactic evolution, and understanding how these features develop is one of the principal problems in galaxy formation and dynamics.  The source of much of the evolution caused by bars and spirals is gravity torques due to tangential forces. Combes \\& Sanders (1981; see also Sanders \\& Tubbs 1980) suggested that these forces could provide a useful measure of the strengths of nonaxisymmetric features such as bars, if the potential could be determined. The idea is to derive the maximum value of the ratio of the tangential force to the mean background (or axisymmetric) radial force, which would give a single dimensionless number indicating the relative importance of nonaxisymmetry in the potential of a galaxy. This ratio, which is physically the same as the maximum gravitational torque per unit mass per unit square of the circular speed, will be referred to in this paper as $Q_g$, while the method for deriving $Q_g$ will be referred to as the gravitational torque method (or GTM).  The advent of routine near-infrared imaging of galaxies has made application of the GTM more practical than ever. Near-infrared images trace the stellar mass distribution of galaxies, due to their emphasis on the older, dominant stellar populations. Potentials can be derived from such images using fast Fourier transform techniques in conjunction with assumptions concerning the mass-to-light ratio and the vertical density distribution (e.g., Quillen, Frogel, and Gonz\\'alez 1994, hereafter QFG). From this potential, the radial and tangential components of the forces in the plane of the galaxy can be derived, and the Combes \\& Sanders ratio can be estimated. Recent studies by Buta \\& Block (2001), Block et al. (2001), Laurikainen, Salo, \\& Rautiainen (2002), Laurikainen \\& Salo (2002), and Block et al. (2002) have provided the first attempts to derive the maximum force ratios for significant samples of galaxies. However, in these cases, the samples were either ill-defined statistically, based entirely on relatively short exposure Two Micron All-Sky Survey (2MASS, Skrutskie et al. 1997) near-infrared images, or used deprojected images that did not allow for the typically rounder shapes of bulges or the most reliable estimates of vertical scaleheights.  There are good reasons for trying to derive the maximum force ratio for a large, statistically well-defined sample of galaxies using a refined version of the GTM. Firstly, Sellwood (2000) has argued that we could evaluate scenarios of bar formation in disk galaxies if we knew the observed distribution of bar strengths. Various bar formation scenarios, such as the natural ``bar instability\" (Miller, Prendergast, \\& Quirk 1970; Hohl 1971; Sellwood \\& Wilkinson 1993 and other references therein) or tidal bar formation (e.g., Noguchi 1996; Miwa \\& Noguchi 1998), may predict different distributions of maximum relative bar torques, and an observed distribution may distinguish which mechanism is most important. Secondly, recurrent bar formation due to accretion of external gas would impact the distribution of maximum force ratios (Bournaud \\& Combes 2002). The idea is that bars can be the engines of their own destruction in the presence of gas (see, for example, Das et al. 2003), but may reform or regenerate later if a galaxy accretes significant quantities of external gas during a Hubble time that may cool the disk sufficiently (see also Sellwood \\& Moore 1999). Thus, accretion can impact the ``duty cycle\" of bars. This idea was evaluated by Block et al. (2002) using an application of the GTM to the Ohio State University Bright Galaxy Survey (OSUBGS, Eskridge et al. 2002). Block et al. concluded that the distribution of maximum relative torques favored the idea that galaxies accrete enough gas to double their mass in 10$^{10}$ years.  In this paper, we re-examine the distribution of maximum relative torques in spiral galaxies based on application of a much refined version of the GTM to basically the same OSUBGS sample as used by Block et al., supplemented by a few larger galaxies with images from the 2MASS database. Our goal is to derive a reliable distribution of maximum relative bar and spiral torques in disk galaxies that can be compared with model predictions. The refinements we use account for the shapes of bulges, improved estimates of the galaxy orientation parameters, vertical scaleheights inferred from type-dependent scalings of the radial scalelength, and a statistical evaluation of the impact of dark matter. The $Q_g$ values we use are from Laurikainen et al. (2003). Only a few of the technical details connected with these values will be provided here, and we refer the reader to Laurikainen et al. (2003) for a full accounting of our application of the GTM. Our approach allows us to derive the most reliable maximum relative torques, and therefore the most accurate distribution of these torques. ",
        "conclusions": ""
    },
    "0311/astro-ph0311469_arXiv.txt": {
        "abstract": "In this paper, we compare the results of the modeling of a protoplanetary disk chemical evolution obtained with the UMIST\\,95 and ``New Standard Model'' (NSM) chemical databases. Assuming the same initial conditions, it is found that the substitution of one chemical network by another causes almost no difference for the disk ionization degree. In contrast, the NSM and UMIST\\,95 abundances of CO can differ by a factor of a hundred at some regions of the disk surface. However, relevant CO vertical column densities differ much less, at most by a factor of a few. In addition, we synthesize the single-dish CO(J=3-2) line by means of the 2D line radiative transfer for both considered chemical networks. It is shown that the intensity of this line in the case of the UMIST\\,95 abundances is lower compared to the NSM case by $\\sim 15\\%$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311375_arXiv.txt": {
        "abstract": "We apply a uniform procedure to select very soft sources from point sources observed by {\\it Chandra} in $4$ galaxies. This sample includes one elliptical galaxy (NGC 4967), $2$ face-on spirals (M101 and M83), and an interacting galaxy (M51). We report on some intriguing results, including the following.  \\noindent (1) We have found very soft X-ray sources (VSSs) in every galaxy. Some of these fit the criteria for canonical supersoft sources (SSSs), while others are somewhat harder. These latter have characteristic values of $kT \\laeq 300$~eV; we refer to them as quasisoft sources (QSSs). We found a combined total of $149$ VSSs in the $4$ galaxies we considered; 77 were SSSs and 72 were QSSs.  \\noindent (2) The data are consistent with the existence of a large VSS population, most of whose members we cannot observe due to the effects of distance and obscuration. The total VSS population of sources with $L > 10^{37}$ erg s$^{-1}$ in each galaxy could be on the order of $1000$.  \\noindent (3) Whereas in M31 only $\\sim 10\\%$ of all X-ray sources detected by {\\it Chandra} are VSSs, more than $35\\%$ of all detectable X-ray sources in the face-on galaxy M101 fit the phenomenological definition of VSSs. This difference may be due to differences in $N_H$ between typical lines of sight to sources in each galaxy.  \\noindent (4) SSSs can be super-Eddington for Chandrasekhar mass objects.  \\noindent (5) We find evidence for SSSs and QSSs with luminosities $10^{36}$ erg s$^{-1} < L < 10^{37}$ erg s$^{-1}$. These sources have luminosities lower than those of the $\\sim 30$ soft sources used to establish the class of SSSs.   \\noindent (6) In the spiral galaxies M101, M83 and M51, a large fraction of the SSSs and QSSs appear to be associated with the spiral arms. This may indicate that some SSSs are young systems, possibly younger than $10^8$ years.  \\noindent (7) In addition to finding hot white dwarfs and soft X-ray binaries, our method should also be efficient at selecting supernova remnants (SNRs). A small fraction of the VSSs in the spiral arms of M101 appear to be associated with SNRs. ",
        "introduction": "\\chandra\\ and \\xmm\\ observations make it possible to systematically discover and study very soft X-ray sources (VSSs) in external galaxies. Until now, the canonical examples of VSSs were luminous supersoft X-ray sources (SSS). The SSSs that established the class lie in the Magellanic Clouds and in the Milky Way; they have $kT$ in the range of tens of eV and luminosities between roughly $10^{37}$ erg s$^{-1}$ and a few times $10^{38}$ erg s$^{-1}$. \\chandra\\ and \\xmm\\ surveys in other galaxies are important, since gas obscures more than 99\\% of the Milky Way's SSSs, so that direct studies of the size and characterics of galactic populations of SSSs can be carried out only in other galaxies.  This paper is the third is a series on the selection of VSSs in external galaxies. In paper 1 (Di Stefano \\& Kong 2003a), we presented a set of strict hardness ratio conditions (the ``HR conditions) to select only the softest sources. Sources selected by the HR conditions are likely to have spectra very similar to those of the SSSs observed in the Galaxy and Magellanic Clouds. Also in paper 1 we tested the HR conditions by applying them to {\\it Chandra} data from $4$ galaxies. In paper 2 (Di Stefano \\& Kong 2003b), we presented an algorithm that starts with the HR conditions, but which then proceeds to progressively relax them. Each set of relaxed conditions chooses sources that deviate from the expected broadband spectra of SSSs in some way. For example, the $3\\sigma$ conditions select sources that may be as soft as SSS-HRs, but which provide too few photons to satisfy the stronger HR conditions. Any source satisfying either the HR or $3\\sigma$ conditions is called a classical supersoft source, or simply an SSS (SSS-HR or SSS-$3\\sigma$). In addition, there is a set of $7$ conditions which can select either highly absorbed SSSs, or sources that are genuinely harder than the SSSs studied in the Galaxy and Magellanic Clouds. We refer to any source which satisfies one of these weaker conditions as a quasisoft source (QSS). For example, the ``NOH'' conditions can select quasisoft sources (QSS-NOH) located behind large gas columns. Other conditions (e.g., the ``$\\sigma$'' conditions) select sources (QSS-$\\sigma$) which may have a dominant component as soft as SSSs, but which also include a small hard component. All sources selected by the algorithm, whether they are supersoft or quasisoft, are referred to as very soft sources (VSSs).  In this paper, we apply the full algorithm to data from the same $4$ galaxies considered in paper 1: M101, M83, M51, and NGC 4472. Because the HR conditions have already been applied, we could have left SSS-HR sources out of this analysis. We decided however, to include them, because it is useful to explicitly compare their properties with those of the other soft sources identified by higher steps in the algorithm. ",
        "conclusions": "The work described in this paper has served three purposes. First, it tests the selection algorithm presented in the companion paper. By fitting spectra to those sources providing adequate numbers of counts, we find that the criteria uniformly select soft sources. As expected from the tests conducted on simulated data (Di\\,Stefano \\& Kong 2003a), those VSSs fit by blackbody spectra had $kT < 250$ eV, while all VSSs emitted less than 10\\% of their flux in photons with energies greater than 1.5 keV. Second, we establish the existence of QSSs. Sources with kT on the order hundreds of eV and $L > 10^{36}$ erg s$^{-1}$ appear not to have been observed before. By now, observations in several nearby galaxies, including M31 (\\rd\\ et al. 2003a) and M104 (\\rd\\ et al. 2003b) verify the result. Finally, investigations have also yielded interesting information about extragalactic SSSs, providing some first steps toward answering the $7$ questions posed in Di\\,Stefano \\& Kong (2003a). We summarize this progress below.  (1)  {\\sl What are typical galactic populations of very soft sources?} It is too soon to claim an answer to this question. We can say, however, the SSSs and QSSs seem to exist in virtually every galaxy. (Ongoing analyses of additional galaxies lend further support to this conclusion.) We have so far carried out only a crude population analysis of the $4$ galaxies studied here. In each case though, the data is consistent with a population of at least $\\sim 500$ VSSs with $L > 10^{37}$ erg s$^{-1}$ active at the time of the observation. The true underlying population of such luminous sources could be several times larger. Observations so far provide evidence for galactic populations of low-luminosity VSSs ($L < 10^{37}$ erg s$^{-1}$), but do not constrain its size. With regard to the potential of VSSs to ionize galactic ISMs, we note that even the most conservative estimate of VSS population for these galaxies yields more than $10^{40}$ erg s$^{-1}$ of highly ionizing radiation released by SSSs in typical galaxies. Although the luminosity is smaller than that associated with young stars, photons from VSSs may dominate at high energies.  (2) {\\sl Are any spiral galaxy parameters related to the relative sizes of SSS and QSS populations?} Among the spiral galaxies in this study, there is not enough variation among the values of $M_B$ to allow a search for correlations between $M_B$ and the true size of the SSS/QSS population. The ratio of the number of detected SSSs and QSSs to the total number of detected X-ray sources does vary systematically with inclination ($39\\%$, $37\\%$, and $30\\%$ for M101, M83, and M51, respectively). If we limit consideration to sources selected by the HR conditions, which seem to be the softest sources selected, the trend appears even stronger ($17\\%$, $11\\%$, and $7\\%$ for M101, M83, and M51, respectively). This  might suggest similar sized populations in each of these galaxies.  (3) {\\sl Do elliptical galaxies house large SSS/QSS populations?} The existence of SSSs and QSSs in NGC 4697 demonstrates that there are SSSs in ellipticals. (See also Sarazin, Irwin, \\& Bregman 2001.) It is difficult to estimate the size of the underlying SSS/QSS populations in galaxies like NGC 4697, largely because we do not yet have good models for the characteristics of the SSS/QSS population in these galaxies. We have demonstrated, however, that, if the sources we detect in NGC 4697 are similar to the SSSs in the Milky Way and Magellanic Clouds that were used to define the class, its population of SSSs is large--about 1250 ``classical'' SSSs per detected example of a ``classical'' SSS in NGC 4697. It is possible, however,  that many of the $14$ detected VSSs are examples of new types of sources.  Finally, we note that the data so far are consistent with the hypothesis that a significant portion of the diffuse soft emission  found in elliptical galaxies may be due to unresolved SSSs (Fabbiano, Kim, \\& Trinchieri 1994).   (4) {\\sl Within spiral galaxies, what are the relative populations of SSSs/QSSs in the galaxy bulges and disks?} In the $3$ spirals we have studied, $f_{SSS+QSS}$, the fraction of SSSs and QSSs located within $1$ kpc of the galaxy center is the same as $f_{all}$, the fraction of all X-ray sources located within $1$ kpc of the galaxy center. The values of $f_{SSS+QSS}$ vary significantly among the spiral galaxies, however.   (5) {\\sl Do galaxies with massive central black holes have more SSSs/QSSs located within 1 kpc of the nucleus than comparable galaxies without massive central black holes?} To answer this question, we have to rely on the dynamical measurements of the central supermassive BH. M51 is a Seyfert 2 galaxy with a low-luminosity AGN (see, e.g., Ho, Filippenko, \\& Sargent 1997) and it is very likely that there is a supermassive, $\\sim 10^7 M_{\\odot}$ (Hagiwara et al. 2001) BH in the center.  M83 might also harbor a $\\sim 10^7 M_{\\odot}$ BH (Thatte et al. 2000). The concentration factors for SSSs/QSSs in both M83 and M51 are high, approximately $50$ for each galaxy. The concentration factors for all sources in each galaxy is comparable, but somewhat smaller. \\\\ \\noindent M101 might have a $< 10^6 M_{\\odot}$ supermassive BH (Moody et al. 1995). The relative values of the concentration factors among the spiral galaxies suggests that the concentration of SSSs/QSSs, and in fact the concentration of all X-ray sources, scales with the mass of the central supermassive BH. It is interesting, but not at this point statistically significant, that the concentration factors are somewhat higher for SSSs/QSSs. \\\\ \\noindent NGC 4697 appears to have the highest-mass central BH of any of the galaxies in our sample ($1.2 \\times 10^8 M_\\odot$). It is also the galaxy that has the highest concentration factor for all X-ray sources (${\\cal F}_{all} = 140),$ the highest concentration factor for SSSs (${\\cal F}_{SSS+QSS} = 360$), and the largest difference between the central concentration of SSSs and other X-ray sources. From this  it seems that interactions near the galaxy center may be playing an important role in creating X-ray sources, and that this is especially true for SSSs/QSSs. (An alternative hypothesis is that the characteristics of the stellar population near the center differ from those elsewhere in the galaxy.) A variety of interactions could enhance the formation of X-ray binaries near the center of the galaxy; it may be that the additional enhancement for SSSs is due to the presence of the cores of stars that have been tidally disrupted by the central BH.   (6) {\\sl For all galaxies, are the positions of SSSs and QSSs correlated to the positions of other objects, such as HII regions, planetary nebulae, supernova remnants, or globular clusters?} In the $3$ spiral galaxies we have studied, the SSSs seem to be concentrated in the spiral arms. Ongoing work on SSSs and QSSs in M101 indicates that they may be preferentially located near HII regions and SNRs, suggesting that some SSSs/QSSs in M101 may be related to young populations.  (7) Are SSSs significant contributors to the rates of Type Ia SNe? It is too soon to say. To use the data to answer this question, we need to estimate the fraction of the SSSs in each galaxy that are accreting WDs. This is difficult to do, even in the Magellanic Clouds. There are, however, some signatures which can help. First, the WD models developed to date invoke donor stars that too old to be concentrated in the spiral arms; the fraction of sources located away from the spiral arms and from regions with recent star formation can provide clues about the size of the NBWD population. Second, some of the accreting WDs that may be Type Ia progenitors should be surrounded  by large, distinctive ionization nebulae (\\rd\\ 1996).\\footnote{The reasons are the following. (1) SSSs emit enough photons to ionize large ($> 10$ pc) regions of the surrounding ISM, if the number density in their surroundings is moderate, on the order of $1$ atom cm$^{-3}$  (Rappaport, Kallman, \\& Malina 1994). (2) These nebulae have distinctive optical emission signatures that can be significantly different from other nebulae, such as PNe and SNRs (Rappaport \\& Chiang 1996; \\rd\\, Paerels, \\& Rappaport 1996). (3) These nebulae could be detectable for much of the lifetime of the SSS state ($\\sim 10^7$ yrs), as long as the duty cycle is $> \\sim 1\\%$ (Chiang \\& Rappaport 1996).} In nearby galaxies where optical counterparts can be identified, we may be able to detect some of the distinctive disk features of close binary SSSs ($P_{orb}< 3$ days) or to detect the donor stars in wide binary SSSs ($P_{orb}$ generally is more than 100 days).  \\subsection {New Questions}  \\noindent (8) {\\sl What are the ultraluminous SSSs?} A natural suggestion is that they are BHs. In fact, if there are BHs with masses filling the gap between the $\\sim 10 M_\\odot$ BHs thought to be stellar remnants, and the $> 10^6 M_\\odot$ BHs found   at the centers of galaxies, then SSS behavior is predicted for a wide range of masses, source luminosities, and efficiency factors. If BHs with masses in this intermediate range do not exist, or are not found in binaries, then the explanation of ultraluminous  SSSs may be more complicated.  \\noindent (9) {\\sl What are the low-luminosity high-temperature VSSs?} The data provide a hint of a population of sources with temperatures near or over $100$ eV, but with luminosities between $10^{36}$ erg s$^{-1}$ and $10^{37}$ erg s$^{-1}$. NBWDs at this temperature are generally expected to be more luminous. If evidence for this high-T/low-L SSS population is found in other investigations, a new model for some SSSs will be required. At present, there is no ``canonical'' model for QSSs. Intermediate-mass BHs provide one possible explanation.  Consider, e.g., a 200 eV source with $L_X\\sim 5\\times10^{37}$ erg s$^{-1}$. This could correspond to a BH of roughly 50 M$_\\odot$ accreting at a rate of roughly 1\\% $\\dot{m}_{Edd}$. (Note that if the bolometric luminosities were significantly smaller, the source might not be expected to be in the soft state.)  \\noindent (10) {\\sl Is there a significant population of quasisoft sources and, if so, what is the nature of the sources?} Some of the SSSs and QSSs may be SNRs. If so, we expect the flux we measure in observations taken during the next few years to be comparable to the values already measured; some may be associated in optical nebulae. It will be important to search for quasisoft sources that are not SNRs.  \\noindent (11) {\\sl If a subset of VSSs are young ( $< 10^8$ yrs), what are they?} Some of them may be SNRs, but we have argued (\\S 5.4.2) that the majority of them are not SNRs. Some may be BHs, some may be neutron stars, perhaps analogs of RX J0059.2-7138. It is even possible that some may be NBWDs. One scenario we propose (R. Di\\,Stefano, in preparation) is one in which a primordial binary with $2$ relatively massive stars ($> 3\\, M_\\odot$) evolves into a binary with a massive WD in orbit with a massive star. Even if the primary never fills its Roche lobe, the system could become a symbiotic which experiences an SSS phase. If the primary {\\it will} fill its Roche lobe, an epoch of irradiation-driven winds prior to contact could also lead to SSS behavior. Finally, in some of these wind-driven systems, the primary mass may be decreased enough that mass transfer is not dynamically unstable when contact is established, and the SSS behavior can continue.  \\subsection{Prospects}  {\\it Chandra} observations are beginning to help us answer key questions about SSSs and to demonstrate that there are new horizons in SSS research. An unanticipated result of this research is the discovery that galaxies contain a significant number of X-ray sources that appear to be hotter than the canonical SSSs that established the class. Discovering the nature of these QSSs presents a fresh challenge. Continuing observations with both {\\it Chandra} and {\\sl XMM} should lead to important results. With the publication of this paper we will make our code available to other researchers, and hope that it can facilitate comparative studies of the VSS populations of different galaxies."
    },
    "0311/astro-ph0311302.txt": {
        "abstract": "We have examined X-ray and optical observations of two ultra-luminous X-ray sources, X7 and X10 in NGC4559, using \\xmm, \\chandra\\/ and \\hst. The UV/X-ray luminosity of X7 exceeds $2.1\\times10^{40}$ erg s$^{-1}$ in the \\xmm\\/ observation, and that of X10 is $>1.3\\times10^{40}$ erg s$^{-1}$. X7 has both thermal and power-law spectral components, The characteristic temperature of the thermal component is $0.12$ keV. The power-law components in the two sources both have slopes with photon index $\\simeq 2.1$. A timing analysis of X7 indicates a break frequency at 28 mHz in the power spectrum, while that for X10 is consistent with an unbroken power law. The luminosity of the blackbody component in the X-ray spectrum of X7 and the nature of its time variability provides evidence that this object is an intermediate mass black hole accreting at sub-Eddington rates, but other scenarios which require high advection efficiencies from a hollowed-out disk might be possible. The emission from X10 can be characterised by a single power-law. This source can be interpreted either as an intermediate mass black hole, or as a stellar-mass black hole with relativistically-beamed Comptonised emission. There are four optical counterparts in the error circle of X7. No counterparts are evident in the error circle for X10. ",
        "introduction": "\\label{sec:intro}  It has been known for some time that there are point-like, non-nuclear sources in some nearby galaxies, with luminosities exceeding $10^{39}$ erg s$^{-1}$ (eg. Long, 1982, Fabbiano 1989, Colbert \\& Mushotzky 1999, Roberts \\& Warwick 2000, Makishima et al. 2000). The nature of these ultra-luminous compact X-ray sources (ULX) is still unclear.  Fabbiano (1989) showed statistically that they are not foreground objects, or background AGNs.  Some of the objects have been identified with young supernova remnants (Fabbiano \\& Trinchieri 1987), but many exhibit significant variability on long time scales. Others may be white dwarfs emitting via nuclear burning on their surfaces (Fabbiano et al. 2003). Low signal-to-noise spectra obtained by {\\it Chandra} and {\\it ASCA} show at least two classes of spectra: a pure power-law and a thermal blackbody-like spectrum (Strickland et al.\\ 2001, Makishima et al.\\ 2000). They are most likely to be accreting compact objects (Colbert \\& Mushotzky 1999), but because their luminosity exceeds the Eddington luminosity of a 20 \\Msun\\/ compact object ($3\\times10^{39}$ ergs s$^{-1}$), sometimes by a factor of $\\sim15$, this scenario suggests that the compact  object may be an intermediate mass black hole (IMBH) of mass \\Mbh$\\sim100$\\Msun.  King et al. (2001) argued that the production of IMBH is difficult, especially when taking into account the constraints on the mass-donor companion stars. However, they can be formed directly from the collapse of high mass stars in low metallicity environments (Heger et al 2003). ULX may also be numerous in elliptical galaxies (Angelini, Loewenstein \\& Mushotzky 2001) although this is still unclear (Irwin, Athey \\& Bregman 2003).  An extreme case is their creation as a primordial population in the first burst of star formation (Population III stars, Madau \\& Rees 2001).  Alternatively, indirect channels may exist for the formation of IMBH, either through binary evolution (Fryer \\& Kalogera  2001) or through the mergers of massive stars in Superstar Clusters  (Portegies Zwart \\& McMillan 2002). IMBH may also be formed through the mergers of lower mass black hole (BH) in cluster environments (Lee 1995, Miller \\& Hamilton 2002).  For higher metallicities, BH masses tend to be \\ltae$20$ \\Msun. Direct-collapse black holes (the most likely origin of $\\gamma$-Ray Bursts) have probable progenitors which yield BH remnant masses with \\Mbh$\\simeq 14-23$ \\Msun\\/ (Fryer 1999). If other effects such as disk winds and magnetic fields are important, these masses could be lower. (For stars with insufficient angular momentum, weak $\\gamma$-Ray Burst explosions or no explosions are produced, and the remnant can be much more massive, up to the mass of the star.) So, if ULX are powered by a stellar mass BH (which we define here as \\Mbh$\\leq20$\\Msun), beaming or super-Eddington luminosities are required to explain the  observed luminosities.  Two different types of beaming can in principle explain the luminosity of ULX in external galaxies. The beaming could be due to anisotropic emission of the radiation (King et al.\\ 2001), or the relativistic motion of the radiation sources (for example Fabrika \\& Mescheryakov 2001).  King et al.\\ (2001) pointed out that if the accretion disk around the BH has a much lower scattering optical depth over a restricted range of solid angle ($\\Delta \\Omega$) than in other directions, almost all the emitted X-rays will emerge in these directions. For a beaming factor $b (=\\Delta \\Omega / 4 \\pi) $, the accreting BH could have a luminosity $L = b L_{\\rm sph}$, where $L_{\\rm sph}$ is the luminosity inferred in the assumption of spherical accretion/isotropic emission. Such beaming is due to the special geometry of the accretion disk, thus it does not require the emitting plasma to be in relativistic motion relative to the observer. The spectrum of the emission could be thermal blackbody-like or a power-law. The chance of discovering such sources depends on the opening angle of the beaming cone, the duty cycle and the life span of the phases in which the mass transfer rate from the donor star is high enough to account for the observed X-ray luminosity.  The other type of beaming is caused by the relativistic motion of the emitting plasmas. This would be the case, for example, if we were viewing the emission from BL Lac objects or the jets of SS~433-type systems at an angle almost parallel to the jet orientation (Fabrika \\& Mescheryakov 2001, K\\\"{o}rding, Falcke, \\& Markoff 2002, Georganopoulos, Aharonian \\& Kirk 2002). The emission is Doppler-boosted by the relativistic motion, and it appears to be ultra bright, depending on the jet speed. As the emission is likely to be non-thermal direct synchrotron emission or (optically thin) Comptonised emission from relativistic beamed electrons, we would expect a power-law spectrum instead of a blackbody-like spectrum. The radiation should have strong linear polarization (even in the X-ray bands). The discovery probability of these type of sources is restricted by the opening angle of the jets. Fabrika \\& Mescheryakov (2001) estimated that about one such relativistic beamed source would be observed in each irregular or spiral galaxy, which is similar to the number of ULX per galaxy in recent observations.  It has proved difficult to distinguish whether ULX are stellar-mass Galactic BH binaries with beamed X-ray emission or the more exotic IMBH systems.  SS~433- and BL Lac-type relativistic beaming models can be excluded for sources with a thermal blackbody-like spectrum. Non-relativistic beaming (anisotropic emission) models will have difficulties if the X-ray observed luminosities are too high. Moreover, these kind of models would impose some signatures on the spectra either by Compton scattering or photoelectric absorption. As the geometrical beaming model requires that the accretion disk is viewed almost face-on, any spectral features are expected to indicate a low inclination geometry. If the geometrical beaming model is correct, the power spectrum of variability would be shifted to higher frequencies than in the IMBH model because the characteristic frequencies at the inner edge of the disk scale as 1/\\Mbh. For the relativistic beaming model, large amplitude variability should occur if the emission is direct synchrotron radiation.  \\begin{figure*} \\begin{center} \\includegraphics[width=\\columnwidth]{cropper_fig1a.ps} \\hspace*{0.75\\columnsep} \\includegraphics[width=\\columnwidth]{cropper_fig1b.ps} \\end{center} \\caption{(a) The  \\xmm\\/ EPIC-pn image of NGC4559. North is up and East is left. X7 is the bright source near the centre with the slightly fainter X10 located above and to the left. The inset shows the region around X10 and nucleus, with the \\chandra\\/ positions superimposed. (b) The \\xmm\\/ OM UVW1 image of NGC4559. The locations of X7 and X10 are  marked with arrows and small red circles and the centre of the galaxy is marked with a magenta cross.} \\label{fig:xmm_image} \\end{figure*}  Further progress on ULX therefore depends on a detailed spectral study, a temporal study extending over as wide a range of frequencies as possible and a study of the location of the ULX in the host galaxies, in particular with respect to star forming regions. Ultimately, kinematic studies of confirmed optical counterparts will prove to be definitive.  It almost goes without saying that ULX are of significant interest.  If they are IMBH, they are the link between stellar mass BH and supermassive BH. Their origin may be primordial, or the result of stellar or BH mergers. They will have a major impact on their surroundings. If they are beamed stellar mass BH, the beaming mechanism and the underlying structure of the disk are also astrophysically important, and again, the scaling to AGN-generated jets is of interest. ",
        "conclusions": "the luminosity of the blackbody-type component, and the temporal properties of the source. In the first line of argument, the fluxes are derived directly from the spectral fits. The UV/X-ray luminosities of this component can be inferred to be $0.8-1.8\\times10^{40}$ erg s$^{-1}$. It is difficult to construct a scenario where significant beaming $(b<0.1)$ of this component can be achieved even if the emitting material itself is collimated. An IMBH accreting at $L\\sim0.1L_{Edd}$ provides a simple explanation for the observed luminosities and blackbody temperatures, and at the same time a consistent explanation for the power-law component in terms of Comptonised emission. The masses derived from the disk fitting, or simple blackbody area measurements indicate that \\Mbh$\\sim2\\times10^3$\\Msun. This could be modified, however, by the (modest) beaming and inclination factors.  On the other hand, a model for a lower mass BH (\\Mbh$<20$\\Msun) fed from an accretion disk truncated at a large radius cannot be ruled out, if most of the energy within this radius is advected with high efficiency ($\\sim 99$ percent).  In the second line of argument, the power spectrum exhibits a break at 0.028 mHz. If the relationship between break frequency and compact object mass can be interpolated between stellar and supermassive BH, then this indicates X7 contains an IMBH. The masses determined from the temporal characteristics are \\Mbh$\\sim40$\\Msun\\/ or \\Mbh$\\sim1.3\\times10^3$\\Msun, depending on the calibration used.  In the case of X10, the absence of the blackbody-type component prevents any inference of the mass from disk parameters. The BH mass can be determined only from the luminosity: if this is sub-Eddington then we infer an IMBH for the compact object, while if we allow super-Eddington accretion or relativistic beaming it can be $<20$\\Msun. The power-law component has the same slope as that for X7, and we interpret this to arise by a similar mechanism, Compton scattering. Among the beaming models, Doppler-boosted Compton emission from a relativistic jet can explain the high observed fluxes without invoking an IMBH. The power-law is unlikely to be produced by direct  beamed synchrotron emission as there is no evidence of sufficiently bright optical and radio counterparts, and the large amplitude temporal signatures expected are absent."
    },
    "0311/astro-ph0311583_arXiv.txt": {
        "abstract": "We report the detection of dust emission at sub-millimeter wavelengths from HD~107146, a G2V star with an age estimated to lie between 80 and 200~Myr. The emission is resolved at 450~\\micron\\ with a size $300~{\\rm AU}\\times 210~{\\rm AU}$. A fit to the spectral energy distribution gives a dust temperature of 51~K and dust mass of $0.10~M_\\oplus$. No excess emission above the photosphere was detected at 18~\\micron\\ showing that there is very little warm dust and implying the presence of a large inner hole, at least 31~AU ($\\sim 1''$) in radius, around the star. The properties of this star-disk system are compared with similar observations of other systems.  We also discuss prospects for future observations that may be able to determine whether the inner hole is maintained by the dynamical effect of an unseen orbiting companion. ",
        "introduction": "Protostellar disks dissipate over a period of several Myr as their constituent dust and gas either accretes onto the star, is dispersed by processes such as stellar winds and photoevaporation, or aggregates into planetesimals (Hollenbach, Yorke, \\& Johnstone 2000). Possible observational examples of disk systems in the process of dissipating are the ``transitional disks'', a thus far small class of objects which includes sources such as V819 Tau (Skrutskie et al.\\ 1990), HR4796A (Jura et al.\\ 1993; Jayawardhana et al.\\ 1998; Koerner et al.\\ 1998; Schneider et al.\\ 1999; Telesco et al.\\ 2000) and HD141569A (Weinberger et al.\\ 1999; Augereau et al.\\ 1999; Fisher et al.\\ 2000).  Dramatic evidence for disk dispersal through planetesimal formation is the detection of numerous planets, and some planetary systems, around nearby stars (Marcy, Cochran, \\& Mayor 2000). Planetesimals are believed to grow into planets through collisional agglomeration. But collisions between planetesimals are also expected to create a cascade of smaller particles, and in this way a relatively old circumstellar disk may regenerate its own dust.  Such second generation dust was first detected around main-sequence stars by IRAS (Aumann et al. 1984).  The IRAS results indicated that $\\sim 15$\\% of main-sequence stars possess dusty disks (Lagrange, Backman, \\& Artymowicz 2000), the brightest of which have been detected at sub-millimeter wavelengths (Zuckerman \\& Becklin 1993; Greaves et al.\\ 1998; Sylvester, Dunkin, \\& Barlow 2001; see also Wyatt, Dent, \\& Greaves 2003). Sub-millimeter observations are useful because they place strong constraints on the dust mass of the disk, due to the low optical depth of dust grains in the sub-millimeter compared to their opacities at shorter wavelengths. Because of their small angular size and weak emission, only four debris disk systems have been spatially resolved at these wavelengths to date (Holland et al. 1998; Greaves et al. 1998). The maps of these systems show interesting asymmetries that have been interpreted as the dynamical signature of a planetary companion (Holland et al. 2003; Wilner et al. 2002; Greaves et al. 1998).  Although the majority of the well-studied debris (and transitional) disk systems surround early-type stars (Vega, Fomalhaut, $\\beta$ Pic, HR4796A, and HD141569A are all A stars), the expectation is that the process of disk dissipation and the regeneration of dust through planetesimal collisions is also a part of the evolutionary history of lower mass stars and, indeed, our own solar system.  Thus, it is of considerable interest to identify debris disks associated with young solar-mass stars, since the detailed study of these objects can provide direct insight into the evolutionary history of our solar system.  The $\\epsilon$ Eri system ($d=3.2$ pc) is one exciting example (Greaves et al.\\ 1998; Hatzes et al.\\ 2000). In this {\\it Letter}, we present sub-millimeter observations of a dusty disk around a close solar analogue, the nearby young G2V star HD~107146.  HD~107146 was first identified as an ``excess dwarf'' on the basis of IRAS colors by Silverstone (2000). We observed it as part of a program to provide ground-based support for the ``Formation and Evolution of Planetary Systems'' SIRTF Legacy Project\\footnotemark[7]\\footnotetext[7]{Information concerning this program can be found at http://feps.as.arizona.edu}. HD~107146 was selected for inclusion in the Legacy project based on its distance (28.5~pc; Perryman et al. 1997) and its high Ca II H and K index $\\log R^\\prime_{\\rm HK}=-4.28$, which is indicative of youth. An age range can be estimated from its lithium equivalent width, 125~m\\AA\\ (Wichmann, Schmitt, \\& Hubrig 2003), which places it at the lower envelope of the 125~Myr Pleiades distribution (Soderblom et al. 1993) but above the distribution for the 250~Myr old M34 (Jones et al. 1997) at the effective temperature of a G2V star. Wichmann et al. also show that its space motions are similar to the Pleiades moving group. Its x-ray luminosity, $L_X=2\\times 10^{29}$~erg~s\\e\\ (Voges et al. 1999), is again similar to solar-type Pleiads (Micela et al. 1999) but weaker than that of the average solar-type stars in the 80~Myr old Alpha Per cluster (Prosser et al. 1996). Together, these indicators suggest an age for HD~107146 between 80 and 200~Myr. Its location in the HR diagram is also consistent with this age range but allows the possibility that it may be as young as 30~Myr if the star has just reached the main sequence.  Our observations, made at sub-millimeter and mid-infrared wavelengths, are detailed in \\S 2. The results from the imaging and fits to the spectral energy distribution (SED) are presented in \\S 3. The SED shows clear evidence for a large inner hole, although the hole is too small to be resolved in the sub-millimeter maps presented here. In \\S4, we compare the properties of the HD107146 star-disk system to other debris and transitional disk systems. We also discuss prospects for future observations that may be able to determine whether the inner hole is maintained by the dynamical effect of an unseen orbiting companion. ",
        "conclusions": "We have detected strong sub-millimeter excess emission from HD~107146, a G2V star with an age estimated to lie in the range $80-200$~Myr. Based on a fit to the SED of the system, the mass of the emitting dust is estimated to be $0.10~M_\\oplus$ or larger. We also find that the disk is marginally resolved at 450~\\micron. The constraints placed by these observations on the mass, temperature, and physical extent of the dust in the system are compared in Table 2 with the properties of other well-studied debris disk systems.  Four of the six stars in Table~2 are A stars, a probable bias due to the relatively high luminosity of these systems at far-infrared and sub-millimeter wavelengths. Nevertheless, despite the range in central star masses and luminosities, the disks have similar properties and follow several trends. The HD~107146 disk is quite massive, comparable in mass to the $\\beta$~Pic disk.  Although it lies noticeably above the trend of decreasing mass with age that is defined by HR4796A, $\\beta$ Pic, Fomalhaut, and Vega (Holland et al.\\ 1998), it is within the scatter in the mass-age relation found for larger samples of dust disks (Wyatt et al.\\ 2003). The fractional dust luminosity ($L_d/L_*$) of HD107146 is also large; it is half that of $\\beta$ Pic, $>10$ times that of Fomalhaut, and is consistent with the trend of decreasing $L_d/L_*$ with age (Spangler et al.\\ 2001). Table 2 also indicates an apparent trend of decreasing outer disk radius with age.  Perhaps this trend, admittedly of low statistical significance at present, will be verified when a larger number of debris disk systems have been spatially resolved.  We also find that the local minimum in the SED at 25~\\micron\\ places a strong limit on the amount of warm dust in the HD~107146 disk, $M_{\\rm d}({\\rm T=51 K})/M_{\\rm d}({\\rm T=100 K})\\simgt 140$. This lack of warm dust implies that the disk does not extend all the way to the star. The dust temperature, and therefore the size of this inner hole, depends on the grain size distribution and optical properties. For grains that emit as blackbodies at 850~\\micron, 51~K dust would lie at 31~AU (Wyatt et al. 1999). This is a lower limit to the inner radius as smaller grains could achieve this temperature at greater distances from the star. The other disks listed in Table~2 have similar size inner holes, as determined from spatially resolved images in thermal emission or scattered light, and are comparable in size to the Kuiper belt in our solar system.  Inner holes are not, by themselves, long lasting because Poynting-Roberston drag will cause dust to spiral in from the outer disk onto the central star on $\\sim 10$~Myr timescales. Such inward migration would be evident in the SED by the presence of warm dust (Jura et al. 1998). Thus, the existence of inner holes has been explained as a consequence of either the dynamical sweeping of an orbiting companion (e.g., Vega; Wilner et al. 2002) or the sublimation of icy grains (e.g., HR~4796A; Jura et al. 1998). Since water ice sublimates at temperatures greater than 100~K (Pollack et al. 1994), this explanation can not apply to the HD~107146 disk where the shape of the SED places a strong limit on the mass of such warm dust.  The alternative explanation, that the inner hole is dynamically maintained by a closely orbiting companion, would require a fairly low-mass companion. The best current limit on the existence of close companions is from Palomar adaptive optics imaging by Metchev \\& Hillenbrand (2002) who found no companions at detection limits of 11.2, 11.7, and 15.2 in absolute K band magnitude at angular separations of $0\\farcs 5, 1''$, and $2''$ respectively. The corresponding mass limits for an age of 100 Myr are approximately 30, 25, and $10~M_{\\rm J}$ (Burrows et al. 1997).  While it has been well-recognized that the dynamical sculpting of disks by orbiting companions can produce inner holes and ring-like structures, recent studies have shown that the migration of dust in the presence of a residual gas disk can also induce ring-like structures in the dust distribution. For example, Takeuchi \\& Artymowicz (2001) have shown that the dust structures seen in the HR4796A and HD141569A systems are qualitatively similar to those expected to result from the coupling between gas and dust in disks. However, orbiting companions may also induce significant departures from axisymmetry in the dust distribution (e.g., Liou \\& Zook 1999; Ozernoy et al.\\ 2000; Quillen \\& Thorndike 2002; Moro-Martin \\& Malhotra 2002), whereas such non-axisymmetric structures cannot be produced by dust migration. Thus, demonstrating the existence of asymmetries in the dust distribution, either in thermal emission or scattered light, as well as measuring the gas content of the HD107146 disk, is needed to distinguish between these two possibilities.  At a distance of 28.5~pc, HD~107146 is relatively nearby. Compared to the other sources listed in Table~2, it is more distant than the debris disk systems that have been spatially resolved at submillimeter wavelengths ($\\beta$ Pic, Vega, Fomalhaut, and $\\epsilon$ Eri), but it is at half the distance of HR4796A.  Since the HD~107146 disk is just resolved at the shortest operating wavelength of SCUBA, and the submillimeter excess is relatively bright, the Submillimeter Array should be able to map the morphology of the emitting dust in greater detail.  Phase referencing will provide a more accurate absolute position, and such observations could determine whether the dust is distributed axisymmetrically (e.g., in a ring) or in a more asymmetric distribution. For example, these observations may confirm the marginal evidence for an offset between the stellar position and the peak of the $450\\micron$ emission found in the observations presented here (Figure~1). The offset, if real, may result from an asymmetric dust distribution, as has been found for several of the other debris disk systems in Table 2.  The dust disk asymmetry can also be addressed by imaging the disk in scattered light. The possibility of detecting scattered light from Vega-like stars is of great interest given the few such systems that have been detected thus far, despite extensive deep surveys (e.g., Kalas \\& Jewitt 1996). Using the quantities listed in Table~2, we can make a rough estimate of the relative strengths of the scattered light from each disk. If we ignore the disk inclination and assume a similar dust grain size distribution and albedo, the scattering area will be proprtional to $M_d$ and the flux of scattered light will be roughly proportional to $M_d L_*/R_{\\rm out}^2d^2$, where $d$ is the distance to the star and $R_{\\rm out}$ is the physical extent of the disk. The factor of $R_{\\rm out}^2$ arises in this expression because grains of a given size located farther from the star intercept less starlight. Hence the surface brightness of the disk will be proportional to $M_d L_*/R_{\\rm out}^4$. A potentially more relevant comparison is of the contrast between the expected surface brightness from the disk and the flux from the star, which is proportional to $M_d d^2/R_{\\rm out}^4$ (cf. Jura et al. 1998). For the disk around HD107146, the expected surface brightness of the scattered light is one half that of the disk around $\\beta$ Pic and the expected contrast is nine times larger, suggesting that the scattered light from HD107146 may not only be quite bright but may also stand out against the glare of the star. Given the low dust temperature deduced for the HD107146 disk ($<100$K), the grains are expected to be icy with a high albedo, which favors the detection of scattered light.  Nevertheless, there are significant uncertainties associated with this estimate. Detecting reflected light will be more challenging if the disk is face-on rather than edge-on. In addition, the sub-millimeter measurements from which $M_d$ is derived are primarily sensitive to large grains ($\\sim 100~\\micron$), whereas the scattered light observations will be particularly sensitive to much smaller grain sizes. Thus, scattered light from the HD107146 disk will be weaker if the grain size distribution is significantly skewed to large grain sizes. Conversely, scattered light measurements can help to constrain the grain size distribution of the disk (e.g., Artymowicz, Burrows, \\& Paresce 1989).  Future spectroscopic observations of this disk with SIRTF, particularly in the $30-40$~\\micron\\ range, will place additional constraints on the grain composition and size distribution (e.g. Wolf \\& Hillenbrand 2003). SIRTF is also expected to discover many more disks over a large range of ages and stellar masses with consequent improvements for our understanding of the formation and evolution of planetary systems."
    },
    "0311/astro-ph0311060_arXiv.txt": {
        "abstract": "We present a comprehensive study of the physical properties of $\\sim 10^5$ galaxies with measurable star formation in the SDSS. By comparing physical information extracted from the emission lines with continuum properties, we build up a picture of the nature of star-forming galaxies at $z<0.2$. We develop a method for aperture correction using resolved imaging and show that our method takes out essentially all aperture bias in the star formation rate (SFR) estimates, allowing an accurate estimate of the total SFRs in galaxies.  We determine the SFR density to be $1.915^{+0.02}_{-0.01} \\mathrm{(rand.)}^{+0.14}_{-0.42} \\mathrm{(sys.)}$ h$_{70} 10^{-2}$ \\Msun/yr/Mpc$^{3}$ at $z=0.1$ (for a Kroupa IMF) and we study the distribution of star formation as a function of various physical parameters.  The majority of the star formation in the low redshift universe takes place in moderately massive galaxies ($10^{10}$--$10^{11}\\Msun$), typically in high surface brightness disk galaxies.  Roughly 15\\% of all star formation takes place in galaxies that show some sign of an active nucleus. About 20\\% occurs in starburst galaxies.  By focusing on the star formation rate per unit mass we show that the present to past-average star formation rate, the Scalo $b$-parameter; is almost constant over almost three orders of magnitude in mass, declining only at $\\Mstar>10^{10}\\Msun$. The volume averaged $b$ parameter is $0.408^{+0.005}_{-0.002} \\mathrm{(rand.)}^{+0.029}_{-0.090} \\mathrm{(sys.)} \\mathrm{h}_{70}^{-1}$. We use this value constrain the star formation history of the universe.  For the concordance cosmology the present day universe is forming stars at least 1/3 of its past average rate. For an exponentially declining cosmic star formation history this corresponds to a time-scale of $7_{-1.5}^{+0.7}$Gyr. In agreement with other work we find a correlation between $b$ and morphological type, as well as a tight correlation between the 4000\\AA\\ break (D4000) and $b$. We discuss how D4000 can be used to estimate $b$ parameters for high redshift galaxies. ",
        "introduction": "\\label{sec:intro}  One of the most active areas of research during the last decade has been the study of the present and past star formation histories of galaxies. This includes detailed studies of the stellar population of our own Galaxy \\citep[e.g.][]{2000A&A...358..869R} and of Local Group dwarfs, both individually and as relic evidence of the global star formation history of the Universe \\citep{2000MNRAS.317..831H,2001ApJ...558L..31H}. While potentially very powerful, the techniques used in these studies are difficult to extend much beyond the Local Group.  At somewhat larger distances, efforts have gone towards understanding the global properties of more massive galaxies and the dependence of star formation on mass and/or morphological classification \\citep[e.g.][]{1983ApJ...272...54K,1994ApJ...435...22K,1996A&A...312L..29G,2001AJ....121..753B}. This has been accompanied by careful mapping of both the star formation and the gas distribution in galaxies \\citep[e.g.][]{2001ApJ...555..301M,2002ApJ...569..157W,1995ApJS...98..219Y}. Based on these studies, a picture has emerged in which more massive galaxies undergo a larger fraction of their star formation at early times than less massive ones. This has often been said to present a challenge for existing models of galaxy formation.  These studies have traditionally been carried out with relatively small samples of local galaxies. For example, the comparatively large sample of \\citet{2002A&A...396..449G} contains 369 spiral galaxies. With the advent of large fiber-fed surveys such as the 2dF \\citep{2001MNRAS.328.1039C} and the SDSS \\citep{2000AJ....120.1579Y}, it is now possible to extend such studies dramatically in size.  Large galaxy surveys have most recently been exploited by \\citet{Hopkins-et-al-2003}, who carried out a comprehensive study of different SFR indicators, and by \\citet{Nakamura-Fukugita-2003} who studied  the local \\ha\\ luminosity density.  Alongside these careful studies of individual local galaxies, a global picture of the star formation history of the universe has been assembled through careful studies of star forming galaxies as a function of look-back time \\citep{1988prun.proc....1C,1990ApJ...348..371S,1996ApJ...460L...1L,1996MNRAS.283.1388M,Hogg-SFR,1998ApJ...498..106M}. The main emphasis of this work has been the star formation history of the universe as a whole. Some workers have also tried to study the evolution of different morphological types \\citep{1998ApJ...499..112B,1997ApJ...489..559G}.  These studies have adopted a wide variety of star formation indicators \\citep{1999MNRAS.306..843G,2001ApJ...558...72S,2001AJ....122..288H,Glazebrook-cosmspec,2000ApJ...536...68A,1999ApJ...517..148F}. Despite differences between different indicators it is clear that there has been a decline in the rate of star formation since $z\\sim 1$, although there is still no consensus on how rapid this decline was. There is also no clear consensus on the star formation history before $z=2$. Studies in the literature give answers ranging from a clear increase back to $z=8$ \\citep{2002ApJ...570..492L}, through a constant rate of star formation back to $z=5$ \\citep[e.g.][]{2002ApJ...569..582B,2003astro.ph..9065G}, to a decline at high $z$ \\citep[e.g.][]{1998ApJ...498..106M}.  One of the key unanswered questions is what physical parameter(s) drive changes in the star formation rate in individual galaxies. Although an increasing body of evidence points to turbulence being a key influence on star formation on small scales \\citep[e.g.][ and references therein]{2003astro.ph..1093L}, the origin of this turbulence and its relevance for star formation on galaxy-sized scales is unclear. It is therefore of major interest to investigate empirically what global quantities correlate strongly with star formation activity. It is logical, however, to expect that the star formation rate is a function of many variables. To investigate this dependence it is necessary to study the star formation rate density function in various projections \\citep[see also the approach by][]{2002ApJ...570..492L}. This is also becoming a favoured approach to analyse theoretical models \\citep[e.g.][]{2003MNRAS.339..312S,2003MNRAS.339.1057Y}.  We will pursue this goal here by trying to tie together detailed information on individual galaxies with their contribution to the overall cosmic star formation activity. This is possible both through studying a snapshot in time and by comparing current and past average star formation distributions. On a global scale, this provides a nearly model-independent complement to the studies of the cosmic spectrum by \\citet{2002ApJ...569..582B} and \\citet{Glazebrook-cosmspec}.  This kind of study requires a survey that spans a large range in galaxy properties and with accurate photometric and spectroscopic information.  The Sloan Digital Sky Survey \\citep[SDSS ][]{2000AJ....120.1579Y} is an ideal candidate for this kind of work. Its accurate photometry allows good characterisation of the selection function and the structural parameters of galaxies. The large database of accurately calibrated spectra allows a much more careful study of the spectral properties of galaxies than has hitherto been possible.  With such unprecedented quality of data it is pertinent to improve modelling over what has generally been done before. One aspect in particular that is generally missing in earlier work is a rigorous treatment of line formation and dust. This is especially important for the SDSS as it spans a large range in luminosity so the physical parameters of the galaxy can be reasonably expected to vary. In this paper we have therefore built on the work of \\citet{2002MNRAS.330..876C} and fit all strong emission lines simultaneously to simultaneously constrain these physical properties.   Our paper is structured as follows: We start in section~\\ref{sec:data} by briefly describing the data used.  Section~\\ref{sec:models} gives an overview of the models and the modelling needed for our analysis. This includes a discussion of the relationship between \\ha-luminosity and SFR.  Section~\\ref{sec:agn_contrib} discusses in some detail the influence of AGN activity on our estimates of star formation rates (SFRs). Section~\\ref{sec:aperture_effects} discusses the very important question of aperture effects in a fibre survey such as the SDSS and sets out our method for removing these.  With these preliminaries out of the way, we discuss the local SFR density in Section~\\ref{sec:sfr_density} and give an inventory of star formation in Section~\\ref{sec:overall_prop}. The specific star formation rate and its variation with mass is discussed in section~\\ref{sec:spec_sfr} and the variation with other physical parameters is covered in Section~\\ref{sec:b_vs_phys_param}.  The paper is concluded in Section~\\ref{sec:discussion}.  We use a cosmology with $H_0=70$km/s/Mpc, $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$.  Where appropriate we define $h=H_0/100$km/s/Mpc.  We use the \\citet{2001MNRAS.322..231K} universal IMF. To convert to the popular choice of a \\citet{1955ApJ...121..161S} IMF between 0.1 and 100\\Msun, one should multiply our SFR estimates by 1.5, the ratio of mass in the two IMFs for the same amount of ionising radiation. We will also make use of stellar mass estimates below and to convert these to a Salpeter IMF it is again necessary to multiply the masses by 1.5. ",
        "conclusions": "\\label{sec:discussion}  The preceding sections have provided a detailed inventory of the star formation activity in the local universe. One possibly surprising result is that the dominant contribution to the SFR density comes from galaxies with on-going star formation in their inner few kiloparsecs (the SF and low S/N SF classes). These galaxies contribute $\\sim 72$\\% of the total SFR density, whereas galaxies with signs of AGN (Composites and AGN) contribute about 15\\% of the total. The remaining 12\\% of $\\rhoSFR$ is located in galaxies we are unable to classify based on their fibre spectra. These are almost certainly galaxies with substantial bulges, in which stars are forming in the outer disk. A thorough examination of this question has to be postponed to a different article, but it is clear that the relative importance of star formation in the outer disk is a function of the mass or luminosity of the galaxy. When we look at narrow redshift ranges we find that the median aperture corrections increase by a factor of 2 from stellar masses around $10^9\\Msun$ to $\\Mstar \\sim 10^{11}\\Msun$. The natural question is how this relates to star formation in bulges as compared to disks, but this will have to await a study of the resolved properties of our galaxies.  We found that galaxies in the SF class have SFR profiles that follow the \\sdssr-band light. For these galaxies, it is possible, and accurate, to remove aperture bias by simple scaling with the \\sdssr-band flux \\citep[cf.][]{Hopkins-et-al-2003}.  We have, for the first time, been able to calibrate the conversion factor from \\ha\\ to SFR as a function of mass and we show that the standard \\citet{1998ARA&A..36..189K} conversion factor is a very good average correction for most galaxies in the local Universe. By  lucky coincidence,  trends in the conversion factor and in the intrinsic Case B ratio of \\ha\\ to \\hb\\ tend to cancel each other out. As a result,  it will be sufficient to take a fixed Case B ratio and a fixed $\\eta^0_{H\\alpha}$ for most studies, even though neither of the two assumptions is correct.  One of the main questions we want to address is what excites large-scale star formation in galaxies. Seen from this angle the different parameters we have studied fall into two categories. First there is the incidental set -- those parameters which are controlled by star formation activity rather than the other way around. Among these are D4000 and  colour. Of more interest are those parameters that may have some direct effect on star formation such as \\Mstar, $\\mu_*$, Hubble type and size.  We have found that the stellar mass correlates well with the SFR, a result found previously for a much smaller sample by \\citet{1996A&A...312L..29G}. This is quite likely because even though the gas {\\em fractions} of galaxies decline with increasing mass, the total gas masses still increase as a function of mass \\citep{2001AJ....121..753B}.  Because of this overall correlation, the SFR density distribution functions are rather similar to the mass density distribution functions discussed by \\citet{2003MNRAS.341...33K}. We find that the majority of star formation takes place in high surface brightness disk-dominated galaxies spanning a range of concentrations. These span a broad range of masses, peaking for galaxies slightly less massive than $10^{11}\\Msun$ and with half-light radii around 3 kpc. The star forming gas has typically solar or slightly super-solar metallicity.  Our most striking result is contained in Figure~\\ref{fig:sfr_dens_fig1}, which shows that the majority of star formation in the local universe takes place in massive galaxies. This is a robust prediction of hierarchical models and the shape of our distribution matches that of \\citet{2003MNRAS.339..312S} (although these authors show $\\rhoSFR$ versus the halo mass). This clearly warrants a more in-depth comparison and we will follow this up in a forthcoming paper.  Another interesting question is whether star formation in the local universe is predominantly bursty or relatively quiescent. As we saw in Figure~\\ref{fig:mstar_vs_b}, the answer depends strongly on the mass of the galaxy.  It is safe to say, however, that the vast majority of star formation in the local universe takes place in galaxies that are not undergoing extreme starbursts.  Quantitatively we find that 28\\% of $\\rhoSFR$ comes from galaxies with SFR$/\\langle\\mathrm{SFR}\\rangle>1$ and 50\\% at SFR$/\\langle\\mathrm{SFR}\\rangle>0.54$, but only 3\\% of $\\rhoSFR$ comes from galaxies with SFR$/\\langle\\mathrm{SFR}\\rangle>10$.  The traditional definition of star bursts is however $b\\ge 2$--$3$ and with that definition around 20\\% of the local SFR density takes place in star bursts. We should caution that our method of inferring the SFR from optical emission lines is likely to miss part of the SFR occurring in extremely dusty FIR-dominated galaxies. Given that such systems are very rare at low redshifts \\citep[e.g.][]{1996ARA&A..34..749S}, this does not significantly affect our results.  Although IR-luminous galaxies are rare in the local universe, most local star formation occurs in galaxies with substantial optical attenuation in their central parts. It is important to realise that this quantity might suffer from aperture biases so the galaxy averaged attenuation could be lower. The relationship between dust attenuation and SFR activity has been discussed by many authors. The most recent discussion is given by \\citet{Hopkins-et-al-2003} for a subsample of our SDSS data. Our results are in good agreement with theirs. Our main result is that the SFR-weighted dust attenuation at \\ha\\ is $A_{H\\alpha} \\sim 1.0\\pm 0.5$ magnitudes, with a good correlation between dust attenuation and SFR. The exact relationship between dust and star formation remains unclear. We see evidence for a strong correlation between stellar mass and dust attenuation, presumably reflecting a correlation between metallicity and dust content. The correlation between SFR and dust may simply reflect a more fundamental correlation between dust and metallicity. This question clearly merits further investigation. Another consequence of the moderately strong dust attenuation present in most galaxies is that the colours of the galaxies that dominate the SFR density have intermediate values --- this is why we needed a second colour to get a good constraint on star formation activity in section~\\ref{sec:aperture_effects}.  To isolate the influence of other factors that may regulate the star formation, it is clearly necessary to take out the dominant trend with mass. One should probably normalise by the mass of the gas reservoir, but as this information is not  available we have normalised by stellar mass. Previous studies by \\citet{2002A&A...396..449G} and \\citet{2001AJ....121..753B} have shown that SFR$/\\Mstar$ correlates with mass. We confirm this trend. The median value of $\\sSFR$ is fairly constant over a wide range at low masses and then decreases smoothly at masses greater than $\\Mstar\\sim10^{10}\\Msun$. The simplest interpretation is that low mass galaxies form stars at an approximately constant rate, whereas the star formation history is increasingly peaked towards earlier times for more massive galaxies.  These trends have arguably been seen before. They agree with the ``downsizing'' scenario of \\citet{1996AJ....112..839C}. What has not been apparent up to now is the fact that the extremes of the distribution change noticeably. As we go up in mass, there are fewer and fewer galaxies with very high $\\sSFR$ and a corresponding tail towards low $\\sSFR$ starts to appear.  It is clear from the previous section that no single factor stands out as controlling $\\sSFR$. This is not entirely surprising since the properties we have used are simple and somewhat limited. It might turn out  that more detailed properties correlate better with $b$, for example lopsidedness \\citep{2000ApJ...538..569R},  local density \\citep{2003ApJ...584..210G} or  merging state \\citep[e.g.][]{2000MNRAS.311..565L}. We will study these relations further in future work.   We mentioned previously that the $\\sSFR^V$ value for the sample as a whole provides a strong constraint on the star formation history of the universe.  This has been discussed in some detail by \\citet{2001AJ....121..753B}, but that study did not use a SFH that is compatible with look-back studies -- it assumed an exponential star formation history for all galaxies. Our $b^V$ value with our cosmology corresponds to an effective $\\tau = 7_{-1.5}^{+0.7}$Gyr (including systematic errors) if one considers the whole universe to have an exponential SFH.   We will however adopt the parameterisation of the SFH of the universe used by \\citet{2002ApJ...569..582B}. In general form this reads \\begin{equation} \\label{eq:SFH_formula} \\mathrm{SFR} \\propto \\left\\{ \\begin{array}{ll} (1+z)^\\alpha & \\mbox{for $z\\in [z_d, z_c]$} \\\\ (1+z)^\\beta & \\mbox{for $z<z_d$} \\\\ 0 & \\mbox{for $z>z_c$}\\\\ \\end{array} \\right., \\end{equation} where we will adopt $z_d=1$, $z_c=5$ to be consistent with \\citet{2002ApJ...569..582B} and \\citet{Glazebrook-cosmspec}. Those authors showed that the local ``cosmic spectrum'' is well fit by a model with $\\beta=3$ and $\\beta=0$.  To calculate $b^V$ for the SFH defined by Equation~(\\ref{eq:SFH_formula}) we simply integrate over time and define \\begin{equation} \\label{eq:mstar_from_model} \\langle \\mathrm{SFR} \\rangle = \\frac{1}{t_0} \\int_{0}^{t_0} \\mathrm{SFR}(t) dt, \\end{equation} with $t_0$ being the age of the universe at $z=0.1$. This, together with equation~(\\ref{eq:SFH_formula}), then gives the present to past-average SFR for the model \\begin{equation} \\label{eq:b_model} b^V_{\\mathrm{model}} = t_0 \\frac{\\mathrm{SFR}(t_0)}{\\int_{0}^{t_0} \\mathrm{SFR}(t) dt}. \\end{equation}  It is interesting to note that because of the way we measure $\\Mstar$ and the SFR, there is a dependence on the Hubble parameter in our $b$ value calculations. That dependence cancels out in equation~(\\ref{eq:SFH_formula}). To constrain $\\alpha$ and $\\beta$ we thus need a prior on $H_0$. We will take the simple approach and choose $h=0.7$. We use the constraint on $b^V$ given in equation~(\\ref{eq:present_to_past_total}).  \\begin{figure} \\centering \\includegraphics[width=84mm]{fig30.ps} \\caption{The constraint on $\\alpha$ and $\\beta$ (see Equation~\\ref{eq:SFH_formula}) from the $b^V$ value calculated for our sample. The solid line shows the locus of $b^V = 0.400$ whereas the dashed line shows the 68\\% confidence limits from the likelihood distributions and the dotted the locus of our conservative systematic error estimates.} \\label{fig:alpha_beta_constrains} \\end{figure}   Figure~\\ref{fig:alpha_beta_constrains} shows the resulting constraints on $\\alpha$ and $\\beta$. Since the $b^V$ value only constrains the $z=0.1$ value of the SFR and the integral under the curve, there is a strong degeneracy even with a fixed functional form. This degeneracy cannot be lifted without independent constraints, but it is still of considerable interest. Most noticeably it is clear that any really strong evolution, $\\beta > 4$, at $z<1$ is excluded by the current data.  The dotted lines in Figures~\\ref{fig:alpha_beta_constrains} show the effect of systematic uncertainties on these constraints (for simplicity we have symmetrised these). The dashed lines show the purely statistical errors (1$\\sigma$). It is clear that if we could remove all systematic uncertainties, the method would give a very powerful constraint on the star formation history of the universe. It is therefore worth asking what is required to further reduce the systematic uncertainties.  A thorough test of our aperture corrections using resolved spectroscopy would substantially reduce the systematic uncertainty in our SFR estimates. To reduce uncertainties due to evolution, one could limit oneself to a narrow range in redshift, but that would mean a corresponding increase in the uncertainty in extrapolating to faint galaxies. Alternatively, it might be possible to directly constrain the evolution inside the sample, but that is a rather delicate operation due to the complex aperture effects.  We have postponed this to a later paper.  Systematic uncertainties notwithstanding, the results  presented here provide remarkably good constraints on the cosmic star formation history using only data at $z<0.22$. It is clear that a similar study carried out at higher redshift will be of great interest and provide a very interesting complement to recent studies of the mass-assembly of the universe \\citep{2000ApJ...536L..77B,2001ApJ...562L.111D,2003ApJ...587...25D,2003ApJ...594L...9F,2003astro.ph..7149R}.  In general it is our firm belief that future studies of high redshift data should try as far as possible to study the evolution of the full distribution functions in question rather than summarising the distributions to a single value. As we have seen, important changes might occur in the tails of the distributions with little influence on the median value. These physically interesting details can only be extracted through the use of distribution functions. It has been standard to study luminosity functions. With the techniques introduced in the present paper it is now possible to extend the analysis to a wide variety of distribution functions. This has been possible due to the unprecedented dataset provided by the SDSS survey."
    },
    "0311/astro-ph0311256_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "} In this introduction to $\\nu$-astronomy, we focused mostly on supernova neutrinos. We aimed at helping the orientation of a reader in this field, so we did not attempt to give a comprehensive study (i.e.\\ we did not consider all theoretical possibilities or scenarios, or reactions to detect neutrinos). Rather, we offered a selection of the background information, provided some few formulae, reference numbers, and showed illustrative calculations. Let us conclude by recalling some of the important points we touched: \\vskip2mm  \\noindent $\\star$ Neutrino astronomy is theoretically appealing and rich of promises. Supernova neutrinos are a very well defined and interesting possibility.  \\noindent $\\star$ Neutrino observations from SN1987A are not in contradiction with the general theoretical picture. However, supernova explosions are still mysterious, and this warrants more discussion and stimulates more efforts.  \\noindent $\\star$  Next galactic supernova will permit us much more precise observations and this will be certainly very helpful to progress. In particular, the response from new generation neutrino detector(s), sensitive to various types of neutrinos and reactions could be of major importance (we discussed in some detail the case of a detector like ICARUS, that combines a large mass with a high resolution and detection efficiency).  \\noindent $\\star$ The effects of oscillations are important and have to be included. Conversely, one could combine experiments and use theoretical information in order to attempt to make inferences on oscillations, but astrophysical uncertainties should be thought as an essential systematics for this purpose. (In other words, there are chances to learn something on neutrinos, but, in our view, the primary aim of these observations is just supernova astrophysics.)  \\noindent $\\star$ All this is fine; the most important task left is an exercise of patience, $0-100$ years for next galactic supernova.  \\vskip1cm"
    },
    "0311/astro-ph0311126_arXiv.txt": {
        "abstract": "It has been suggested in the literature that the (Fe/$\\alpha$) abundance ratio may be used as a chronometer, due to a delay in this ratio reaching its solar value as predicted by galactic chemical evolution models. Using grids of photoionization models along a sequence of the (Fe/Mg) abundance ratio vs.\\ metallicity with time in a giant elliptical starburst scenario, we investigate the relationship between the (Fe/Mg) abundance ratio and the FeII/MgII emission line flux ratio under the assumption that these lines originate in photoionized clouds within the broad emission line regions of quasars. ",
        "introduction": "Nearly 20 years ago, in their paper investigating the physics of FeII emission in quasars, Will, Netzer, \\& Wills (1985; WNW) suggested the importance of measuring the (Fe/Mg) abundance ratio in quasars, in that an overabundance of iron in quasars  and active galaxies ``may have important implications for their evolution.'' Hamann \\& Ferland (1993; 1999) discussed the possible use of the FeII/MgII $\\lambda$2800 emission line flux ratio as a chronometer in quasars. With Type~II SNe producing most of the $\\alpha$ elements (O, Mg, etc), and Type~Ia SNe accounting for most of the iron enrichment, a characteristic delay of $\\sim$~1 Gyr is expected for (Fe/$\\alpha$) to reach solar values. However, the value of this delay has been a matter of recent controversy, with some models of Matteucci \\& Recchi (2001) suggesting much shorter delay time scales, depending upon the environment.  Recent quasar surveys, such as the Sloan Digital Sky Survey, have discovered large numbers of high redshift quasars. Infrared spectroscopy of these quasars have revealed their emission in MgII and UV-optical FeII. Elston et~al.\\ (1994) was the first to report strong FeII emission in $z > 3$ quasars. Since then several investigators have measured the UV FeII/MgII emission ratio in quasars with redshifts up to $z \\approx 6$, with hopes of gauging early chemical evolution in the centers of host galaxies and pinning lower limits to the age of the universe at high redshift (Kawara et~al.\\ 1996; Thompson et~al.\\ 1999; Aoki et~al\\ 2002; Dietrich et~al.\\ 2002; Iwamuro et~al.\\ 2002; Pentericci et~al.\\ 2002; Freudling et~al.\\ 2003; Dietrich et~al.\\ 2003a; Maiolino et~al.\\ 2003; Barth et~al.\\ 2003). Yoshii et~al.\\ (1998) used galactic chemical evolution models and FeII/MgII measurements of a $z = 3.62$ quasar to constrain cosmological models, based upon the delay in the iron enrichment and two important assumptions. First, as suggested by WNW, they assumed that a flux ratio FeII(UV $+$ opt)/MgII $\\approx 3$ represents solar (Fe/Mg). This has never been established theoretically. Second, they assumed that the flux ratio scales directly with the abundance ratio. It is this latter relationship that we investigate, as an important first step toward possibly establishing the FeII/MgII flux ratio in quasars as a meter stick of time and/or (Fe/$\\alpha$) enrichment history (see Di~Matteo et~al.\\ 2003). A linear relation (logarithmic slope $\\approx 1$) between the flux and abundance ratios is not expected, due to the thermostatic effects of significant coolants. The logarithmic slope in the relation may even change since MgII is only a doublet, while FeII can redistribute its cooling to a vast array of UV and optical multiplets (where one may not be measuring). ",
        "conclusions": ""
    },
    "0311/astro-ph0311310_arXiv.txt": {
        "abstract": "The total area coverage by starspots is of interest for a variety of reasons, but direct techniques only provide estimates of this important quantity. Sunspot areas exhibit a lognormal size distribution irrespective of the phase of the activity cycle, implying that most sunspots are small. Here we explore the consequences if starspot areas were similarly distributed. The solar data allow for an increase in the fraction of larger sunspots with increasing activity. Taking this difference between the size distribution at sunspot maximum and minimum, we extrapolate to higher activity levels, assuming different dependencies of the parameters of the lognormal distribution on total spot coverage. We find that even for very heavily spotted (hypothetical) stars a large fraction of the spots are smaller than the current resolution limit of Doppler images and might hence be missed on traditional Doppler maps. ",
        "introduction": "\\label{sec:intro} In recent years, an ever increasing number of spotted stars have been mapped using Doppler imaging. The maps reveal the surface distribution of starspots, which in general are large compared to even the largest sunspots. While Doppler images do a good job of catching starspots that modulate the line profile, it is extremely difficult to detect a background of small starspots more or less homogeneously distributed over the stellar surface.  While TiO-band mapping still suffers from notable uncertainties (see Sec.~2), it should in principle be able to pick up non-modulating and homogeneous spot distributions. Typically, techniques using TiO bands to determine spot temperatures and surface areas tend to find larger covering fractions than Doppler imaging techniques (though see also \\scite{berdyugina2002potsdam}). Such differences in apparent spot coverage only provide a hint for unresolved starspots. One star for which the distribution of spot sizes is known in great detail is the Sun. \\scite{bogdan88} found the size distribution to be well represented by a lognormal function. This implies that the number of small sunspots is much larger than that of large spots. This supports the idea that there could be additional small, i.e. unresolved, starspots on more active stars as well. The total starspot coverage is of interest as a measure of stellar magnetic activity, in  order to establish the proper ratio of starspot to stellar plage \\cite{radick98} and to obtain improved estimates of the total magnetic flux carried by the star (with the possible exception of Zeeman Doppler imaging where most of the magnetic signal appears to come from penumbral-type structures, techniques of stellar magnetic field measurement mainly sample plage fields \\cite{saar88,solanki92cs7}).  Here we explore hypothetical scenarios for extrapolating the solar spot-size distribution to activity levels typical for much more active stars. The basic assumption is that the size distribution of star spots can be described by a lognormal function, as in the case of sunspots. This assumption is not unreasonable since the magnetic fields on both the Sun and on more active cool stars are thought to be produced by a dynamo residing at the base of the convection zone \\cite{petrovay2001,schussler2002}. From there flux tubes carry the field to the solar surface. The fragmentation of these tubes during their passage through the convection zone is thought to give rise to the observed lognormal distribution \\cite{bogdan88}. Lognormal distributions can, however, differ significantly from each other in their parameters.  In order to constrain these parameters for active stars we investigate the possible range of behaviour between solar activity minimum and maximum and use these to extrapolate to larger levels of activity. Hence we assume that the processes which lead to the flux-tube size distribution do not change qualitatively with increasing activity. Such an assumption has in the past helped to reproduce, e.g., the high latitudes of starspots \\cite{schuessler92,schussler96buoy,schrijver2001}, or the presence of active longitudes on the Sun and Sun-like stars \\scite{berdyugina2003longitude}.  \\begin{table} \\label{tab:starlist} \\caption[]{A selection of RS\\, CVn stars whose spot coverages have been determined using different techniques. The first column gives the name of the object, the second and third columns the surface and spot effective temperatures according to \\scite{oneal98tio}. The fourth column gives the inclination angle, usually taken from the Doppler imaging papers listed in Sec.~2.1. Note that not all groups agree on the spot temperatures and inclination angles. } \\centerline{ \\begin{tabular}{lccc} \\hline Object       & T$_{\\rm star}$ & T$_{\\rm spot}$       & inclination \\\\ \\hline II Peg       &       4750    & 3530                  & 60\\deg        \\\\ EI Eri       &       5600    & 3700                  & 46\\deg        \\\\ $\\sigma$ Gem &       4600    & 3850                  & 60\\deg        \\\\ DM UMa       &       4600    & 3570                  & 55\\deg        \\\\ HD 199178    &       5350    & 3800                  & 40\\deg        \\\\ \\hline \\end{tabular} } \\end{table} ",
        "conclusions": "\\label{sec:conclusion}  We consider some of the consequences if spot sizes follow a lognormal distribution as on the sun. \\scite{bogdan88} proposed that the passage of magnetic flux tubes from the dynamo to the stellar surface through the turbulent convection zone leads to a certain fragmentation of the flux, producing a lognormal distribution of umbral areas (flux-tube cross-sections). In this picture a significant fraction of the total starspot area of any star is in the form of starspots below the resolution limit of Doppler images. The sunspot-size distribution is consistent with a shift towards larger spots at solar activity maximum. We have used the possible range in the solar parameters to extrapolate to higher activity levels in different ways and have compared the resulting fractional spot areas that can be resolved by Doppler imaging. This has been done assuming  that the starspots are randomly distributed on the stellar surface, which implies that most starspots resolved by Doppler imaging are single spots. Hence we assume that starspots do not clump. If they clump together as in solar active regions then Doppler images detect a larger fraction of the starspots than suggested by our analysis.  Recent calculations of \\scite{berdyugina2002potsdam} suggest that at least for the RS\\,CVn star II~Peg, spectral synthesis of the TiO band reproduces the observations of these molecular lines without requiring starspots in addition to those present on the corresponding Doppler image.  If starspot distributions on RSCVn stars are indeed log-normal and follow a sun-like pattern, this result would suggest that the starspots are tightly clumped and that the starspots resolved by Doppler imaging are actually conglomerates of smaller spots. The low photospheric temperature of II~Peg (Table~\\ref{tab:starlist}), however, means that CN lines blending the analysed TiO band are also present in the spectrum of the immaculate star (Berdyugina, priv.~comm.). Due to uncertainties in the exact temperature there are also some uncertainties in the above result and further such calculations for hotter stars would be of great interest.  If larger spots increase in number more rapidly than smaller spots as stars become more active, we also expect starspots as a whole to be more numerous relative to smaller magnetic flux tubes, i.e., bright magnetic elements.  The ``switch-over'' between activity-bright and activity-dark stars seen with increasing activity level \\nocite{radick87,radick90,radick98,lockwood92} (Radick et al.~1987; 1990; 1998; Lockwood et al.~1992) fits well into this, as does the strong increase of spot area relative to facular area from solar activity minimum to maximum \\cite{chapman97}. The fact that model calculations based on extrapolations from solar values do reproduce the switch-over at about the correct activity level \\cite{knaack98diplom}, supports our general approach of extrapolating the size distribution of spots from the Sun to more active stars."
    },
    "0311/hep-ph0311233_arXiv.txt": {
        "abstract": "For a universe containing cosmological constant together with uniform arrangements of magnetic fields, strings, or domain walls, exact solutions to the Einstein equations are shown to lead to a universe with ellipsoidal expansion. The magnetic field case is the easiest to motivate and has the highest possibility of finding application in observational cosmology. ",
        "introduction": "The first year WMAP results \\cite{Bennett:2003bz,Spergel:2003cb,Hinshaw:2003ex} contain interesting large-scale features which warrant further attention \\cite{Tegmark:2003ve,deOliveira-Costa:2003pu}. First, there is the suppression of power at large angular scales ($\\theta \\gtrsim 60^{\\circ}$), which is reflected most distinctly in the reduction of the quadrupole $C_2$. This effect was also seen in the COBE results~\\cite{Smoot:1992td,Kogut:1996us}.  However, WMAP was able to make more precise measurements of $C_2$. Next, the quadrupole $C_2$ and octupole $C_3$ are found to be aligned.  In particular, the $\\ell=2$ and $3$ powers are found to be concentrated in a plane inclined about $30^{\\circ}$ to the Galactic plane.  In a coordinate system in which the equator is in this plane, the $\\ell= 2$ and $3$ powers are primarily in the $m=\\pm \\ell$ modes. The axis of this system defines a unique direction and supports an idea of power in one direction being suppressed compared to the power in the orthogonal plane. These effects seem to suggest one (longitudinal) direction may have expanded differently from the other two (transverse) directions, where the transverse directions describe the equatorial plane mentioned above.  After the COBE experiment, Monte Carlo studies were used~\\cite{Berera:1997wz} to cast doubt on quadrupole suppression and analyses of the WMAP data \\cite{Bennett:2003bz,Spergel:2003cb,Hinshaw:2003ex,Tegmark:2003ve,deOliveira-Costa:2003pu} have arrived at similar conclusions. Nevertheless, interesting physical effects are not ruled out and a realistic physical model that could explain some or all effects in the WMAP data would be of interest.  It would be most satisfying to explain global anisotropy in the universe by a simple modification of the conventional Friedman-Robertson-Walker (FRW) model. As a first step in this direction, we will consider an energy-momentum tensor which is spatially non-spherical or spontaneously becomes non-spherical. Such situations could occur when defects or magnetic fields are present. Magnetic fields \\cite{Kronberg:1993vk} and cosmic defects \\cite{Hindmarsh:1994re} can arise in various ways. Moreover, large scale magnetic fields exist in the universe, perhaps up to cosmological scales \\cite{Kronberg:1993vk,Wick:2000yc}. These considerations motivate us to examine the effect magnetic fields and defects have on the expansion of the universe~\\cite{RZNB}.  To begin to understand the form, significance and implications of an asymmetric universe, we here modify the standard spherically symmetric FRW cosmology to a form with only planar symmetry \\cite{Taub:1950ez}. Our choice of energy-momentum will result in non-spherical expansion from a spherical symmetric initial configuration: an initial co-moving sphere will evolve into an ellipsoid that can be either prolate or oblate depending on the choice of matter content. We then explore the effects of this type of asymmetric expansion on modes of a multipole expansion. For the sake of clarity, we first give some general properties of cosmologies with planar symmetry (The universe looks the same from all points but they all have a preferred direction.) and then investigate one case in detail. Our exemplar will be a universe filled with uniform magnetic fields and cosmological constant.  This is perhaps the most easily motivated, analytically solvable case to consider and it will give us a context in which to couch the discussion of other examples with planar symmetry and cases where planar symmetry is broken.  To set the stage, consider an early epoch in the universe at the onset of cosmic inflation, where strong magnetic fields have been produced in a phase transition \\cite{Savvidy:1977as,Matinian:1976mp,Vachaspati:nm,Enqvist:1994rm,Berera:1998hv}. The magnitude of the magnetic field and vacuum energy ($\\Lambda$) densities initially are about the same, but eventually $\\Lambda$ dominates. It was estimated~\\cite{Vachaspati:nm} that the initial magnetic field energy produced in the electroweak phase transition was within an order of magnitude of the critical density. Other phase transitions may have even higher initial field values \\cite{Enqvist:1994rm,Berera:1998hv}. Hence, it is not unphysical to consider a universe with B-fields and $\\Lambda$ of comparable magnitudes. If the B-fields are aligned in domains, then some degree of inflation is sufficient to push all but one domain outside the horizon. (Below we also discuss cases where there are still a few domains within the horizon.) ",
        "conclusions": "Since the time of COBE, various speculations have been put forward to explain the origin of the suppressed quadrupole. One idea has been to explore topological features with specific emphasis on the ``small universe'' picture with a toroidal $T^3$ topology~\\cite{ellis}. Another idea has been to study the possibility of a small inflation, which somehow gets tuned to about the scale of the present-day Hubble radius~\\cite{Berera:1997wz,Cline:2003ve,Bastero-Gil:2003bv}.  However, very little has been suggested for physical models that could explain the new features suggested by the WMAP data. An obvious modification of the $T^3$-model would be to make one dimension smaller than the other two and get a $T_a^2 \\times T_b^1$ asymmetric toroidal topology \\cite{deOliveira-Costa:2003pu}. Other topologies could also work. There are many choices \\cite{Gibbons:1996aq,Ratcliffe:rb} and the results could be profound.  In this paper we have not departed radically from conventional cosmology, but rather considered a simple modification of the FRW model. We determined how magnetic fields, strings, and walls affect the expansion rate of the universe. Before solving the Einstein equations we found the asymptotic values and bounds for the ratio of longitudinal and transverse sizes for several interesting cases. Next, we presented exact analytical solutions for a universe filled with magnetic field and gave asymptotic results for the case of defects. Finally, we mentioned particle physics motivated models that could yield the sorts of magnetic fields and/or defects of importance for asymmetric expansion. We have found that a universe filled with aligned B-fields and $\\Lambda$ expands non-spherically: a co-moving sphere at $t_\\i$ evolves into an ellipsoid. The eccentricity is calculable and fixed by the initial conditions. Aside from being potentially applicable to explaining some features in the WMAP data, the models and the results presented in this paper differ from those in the already mentioned Ref.~\\cite{RZNB}. The first of the papers in Ref.~\\cite{RZNB} examined the situation where there is a uniform magnetic field and no cosmological constant. By adding cosmological constant, we have found the exact solution for the model which is physically relevant. The second paper in Ref.~\\cite{RZNB} treats only small anisotropies and determines only asymptotic large time behavior. Our calculations apply to anisotropy of arbitrary strength and we have obtained the exact solutions to the Einstein equations which are valid at all times.  The effects of non-spherical expansion and small inflation are potentially measurable depending on the degree of eccentricity and expansion attained and the quality of data. In realistic particle physics motivated models, where the B-field in question is associated with particular gauge fields and the cosmological constant $\\Lambda$ arises from some specific inflaton model, additional predictions may emerge. For example, a strongly aligned B-field at the beginning of inflation may affect density perturbations in a measurable way~\\cite{Tsagas:1999ft}, perhaps through polarization effects, and a small inflation that follows may mean these perturbations are observable at the largest scales today. We have provided only homogeneous solutions. To apply our results to structure formation and CMB data, the next step will be to begin an analysis of linear density perturbations in such a Universe.  The model developed here can be extended to cases of less symmetric evolution. Some examples are models with an aligned B-field plus cosmic strings neither aligned parallel nor perpendicular to this field, or models with various kinds of misaligned cosmic defects. Another possibility with more physical motivation is a model with B-fields and $\\Lambda$ where there are still a few misaligned magnetic domains within the horizon at the end of inflation. While this model is more complicated, it may also be more realistic as the initial domain size is more in line with the models of cosmological phase transitions \\cite{Savvidy:1977as,Matinian:1976mp,Vachaspati:nm,Enqvist:1994rm,Berera:1998hv}. Finally, while there could be other combinations of cosmological constant, magnetic fields, defects, and geometry that arise in brane world physics and produce asymmetric expansion, we believe it is important to have considered the simplest well motivated case first."
    },
    "0311/astro-ph0311195_arXiv.txt": {
        "abstract": "{ As part of our efforts to unveil extragalactic large-scale structures behind the southern Milky Way, we here present redshifts for 764 galaxies in the Hydra/Antlia, Crux and Great Attractor region ($266^\\circ \\le \\ell \\le 338^\\circ$, $|\\,b\\,| \\la 10^\\circ$), obtained with the Meudon-ESO Fibre Object Spectrograph (MEFOS) at the 3.6-m telescope of ESO. The observations are part of a redshift survey of partially obscured galaxies recorded in the course of a deep optical galaxy search behind the southern Milky Way (Kraan-Korteweg 2000; Woudt \\& Kraan-Korteweg 2001). A total of 947 galaxies have been observed, a small percentage of the spectra ($N = 109$, 11.5\\%) were contaminated by foreground stars, and 74 galaxies (7.8\\%) were too faint to allow a reliable redshift determination. With MEFOS we obtained spectra down to the faintest galaxies of our optical galaxy survey, and hence probe large-scale structures out to larger distances ($v \\la 30\\,000$ km s$^{-1}$) than our other redshift follow-ups using the 1.9-m telescope at the South African Astronomical Observatory (Kraan-Korteweg et al.~1995; Fairall et al.~1998; Woudt et al.~1999) and the 64-m Parkes radio telescope (Kraan-Korteweg et al.~2002). The most distinct large-scale structures revealed in the southern Zone of Avoidance are discussed in context to known structures adjacent to the Milky Way. \\keywords { Catalogs -- Surveys -- ISM: dust, extinction -- Galaxies: distances \\& redshifts -- clusters: individual: ACO 3627 -- large-scale structure of Universe} } ",
        "introduction": "At the same time when significant advances are made in characterising large-scale structures of galaxies in the Universe through dedicated surveys such as the Sloan Digital Sky Survey (Zehavi et al.~2002) and the 2dF Galaxy Redshift Survey (Colless et al.~2001), the study of large-scale structures {\\sl behind the Milky Way} has also progressed enormously. Dedicated deep optical searches, near infrared all-sky surveys (2MASS and DENIS), HI all-sky surveys and X-ray surveys, have all resulted in the detection of voids, clusters and superclusters at low Galactic latitude (for a review, see Kraan-Korteweg \\& Lahav (2000) and references therein).  We have focussed our efforts on an extended region behind the southern Milky Way (Kraan-Korteweg 2000; Woudt \\& Kraan-Korteweg 2001). This region of space harbours a large local overdensity of galaxies, the Great Attractor (Dressler et al.~1987; Lynden-Bell et al.~1988; Kolatt et al.~1995; Tonry et al.~2000). Local and complex large-scale structures appear to extend across the Galactic Plane (GP) in this part of the sky (Kraan-Korteweg et al.~1995 (Paper I); Fairall et al.~1998 (Paper II); Woudt et al.~1999 (Paper III)).   Our deep optical galaxy search behind the southern Milky Way is divided into three separate areas: the Hydra/Antlia region ($266^{\\circ} \\le \\ell \\le 296^{\\circ}$, $-10^{\\circ} \\le b \\le +8^{\\circ}$) (Kraan-Korteweg 2000), the Crux region ($289^{\\circ} \\le \\ell \\le 318^{\\circ}$, $|\\, b\\, | \\le 10^{\\circ}$), and the Great Attractor region ($316^{\\circ} \\le \\ell \\le 338^{\\circ}, |\\, b\\, | \\le 10^{\\circ}$) (Woudt \\& Kraan-Korteweg 2001). The optical survey in these regions has revealed over 11\\,000 previously unrecorded galaxies with observed major diameters $\\ga 0\\farcm2$.  Recently, the survey has been extended towards the Puppis region (at lower Galactic longitudes), and the Scorpius region (towards the Galactic bulge).  \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics{mefos_fig1.ps}} \\caption{The distribution of galaxies in the Hydra/Antlia, Crux and GA region (upper panel, from right to left). The dashed line marks these three different survey regions. The contours are lines of equal Galactic foreground extinction, taken from the Galactic reddening maps of Schlegel \\etal (1998). The contours correspond to $A_{\\rm B}$ = 1$^{\\rm m}$, 3$^{\\rm m}$ (thick line) and 5$^{\\rm m}$. In the upper panel, the big dots show the location of the 48 observed MEFOS fields. The lower panel shows the galaxy density contours in the Hydra/Antlia, Crux and GA region. The contours mark 0.5, 5 (dotted line), 10, 25 (thick line) and 50 galaxies per square degree, respectively. The position of the peak of the reconstructed mass density field associated with the GA (Kolatt et al.~1995) is clearly marked.} \\label{meffig1} \\end{figure*}   Following our deep optical galaxy search, we have been engaged in a redshift survey of these newly found partially obscured galaxies. This redshift survey consists of three different, yet complementary approaches in tracing the extragalactic large-scale structures. They are:  \\begin{itemize} \\item{Optical spectroscopy of individual early-type galaxies with the 1.9-m telescope of the South African Astronomical Observatory (SAAO) (Papers I -- III).} \\item{H\\,I line spectroscopy of low surface brightness spiral and irregular galaxies with the Parkes 64-m radio telescope (Kraan-Korteweg \\etal 2002).} \\item{Low resolution, multi-fibre spectroscopy of galaxies in regions of high galaxy-density with the 3.6-m telescope of the European Southern Observatory (ESO), initially in combination with Optopus, later with MEFOS (Meudon-ESO Fibre Object Spectrograph).} \\end{itemize}  MEFOS, with its 29 object fibres and large field of view (1-degree diameter), is optimal for covering the highest density regions in our optical survey. These multi-fibre observations are an important part of our redshift survey, specifically in mapping the centre of the GA region and the extended overdensity around ACO 3627.  In this paper -- paper IV in the series of papers on extragalactic large-scale structures behind the southern Milky Way -- the results from the MEFOS multi-fibre spectroscopy in the Hydra/Antlia, Crux and Great Attractor regions are presented. ",
        "conclusions": "Multi-fibre spectroscopy with MEFOS on the 3.6-m telescope of ESO has resulted in 764 redshifts in the southern Zone of Avoidance. Contrary to the SAAO and Parkes components of our redshift survey in the Zone of Avoidance, MEFOS has focussed on the faint end of the galaxies in our optical galaxy catalogue in the high galaxy-density regions. As a result, MEFOS has contributed significantly to the knowledge of low Galactic latitude clusters at $v \\ga 12\\,500$ km s$^{-1}$."
    },
    "0311/astro-ph0311476_arXiv.txt": {
        "abstract": "Since pulsating subdwarf B (sdBV or EC14026) stars were first discovered (Kilkenny et al, 1997), observational efforts have tried to realize their potential for constraining the interior physics of extreme horizontal branch (EHB) stars. Difficulties encountered along the way include uncertain mode identifications and a lack of stable pulsation mode properties. Here we report on Feige~48, an sdBV star for which follow-up observations have been obtained spanning more than four years, which shows some stable pulsation modes.  We resolve the temporal spectrum into five stable pulsation periods in the range 340 to 380 seconds with amplitudes less than 1\\%, and two additional periods that appear in one dataset each. The three largest amplitude periodicities are nearly equally spaced, and we explore the consequences of identifying them as a rotationally split $\\ell=1$ triplet by consulting with a representative stellar model.  The general stability of the pulsation amplitudes and phases allows us to use the pulsation phases to constrain the timescale of evolution for this sdBV star.   Additionally, we are able to place interesting limits on any stellar or planetary companion to Feige~48. ",
        "introduction": "To date, over 30 pulsating subdwarf B (EC~14026 or sdBV) stars have been identified, with pulsation periods ranging from 68 to 528 seconds and with amplitudes generally less than 50 millimagnitudes (mmag). For recent reviews of this class of stars, see Kilkenny (2001), and Reed, Kawaler \\& Kleinman (2000), for observational properties; Charpinet, Fontaine \\& Brassard (2001 and references therein) describe in detail some important aspects of pulsation theory in sdB stars. Most sdBV stars show periods at the short end of the range, and probably represent stars close to the Zero Age Horizontal Branch (ZAHB).  PG~1605+072 is the longest period, and lowest gravity, sdBV star, with Feige~48 being an intermediate object.  In general, the longer-period sdBV stars represent more highly evolved objects.  Feige~48 was identified as a ``faint blue star'' as part of the Feige survey (Feige, 1958). It was re-categorized as an sdB star when it was observed as part of the Palomar-Green survey (Green, Schmidt, \\& Liebert, 1984).   Koen et al. (1998; hereafter K98) identified five pulsation periods in Feige~48 in six observing runs from 1997 May to 1998 February. The periods detected by K98 range from 342 to 379 seconds with the largest amplitude being 6.4 mmag.  Amplitude variability led K98 to conclude that mode beating was probably present implying that other unresolved modes were present in their data. This provided the motivation for our follow-up observations. Heber, Reid, \\& Werner (2000; hereafter HRW) obtained a high resolution (0.09\\AA) spectrum of Feige~48, from which they determined log $g$=5.50$\\pm$0.05 and T$_{\\rm eff}$=29500$\\pm$300 K. This places Feige~48 among the coolest sdBV stars known with a surface gravity intermediate between PG~1605+072 and the rest of the class.  Here we report on our multi-year campaign of high-speed photometry of Feige~48. In Section~2, we outline our observations.  Section~3 describes the time series analysis and period identifications.  We report on a stellar model fit to Feige~48 in Section~4.  The phase stability of pulsations is described in Section~5, where we use this stability to place interesting limits on any possible planetary companion. Section 6 gives our conclusions and outlines future observations for Feige~48. ",
        "conclusions": "From our multi-season photometry of Feige~48, we have consistently detected five pulsation periods. Of these five, three ($f$1, $f$3 and $f$4) are consistent with K98. One frequency ($f$5) differs by a daily alias, while K98's fifth frequency (2874~$\\mu$Hz) is not detected in our data. In data sets V and IX, we also detect new pulsation frequencies at 2890 and 2843~$\\mu$Hz respectively, indicating that there may be some stochastically excited pulsations in Feige~48. This is consistent with pulsation behavior seen in another sdBV star, PG~1605+072 (Reed 2001).  Our attempted model fit follows the strategy that has been successfully applied to other classes of pulsating stars, but has rarely worked for sdBV stars; namely using observed frequency splittings to impose $\\ell$ constraints on models.  Using standard evolutionary models, our preliminary model grid includes a model that is able to explain all but the lowest frequency stable pulsations. Most exciting is the fact that such standard models fail to explain all five frequencies.  Thus, despite its relative simplicity and the richness of the parameters available, the failure of this model suggests that standard stellar evolution theory does not fully explain the evolution of sdB stars and\\/ or the nature of pulsations within them. We have something new to learn.  Our modeling example shows that Feige~48 should also serve as an interesting test for other methods of mode identification. Though optical multicolour photometry was not useful for identifying pulsation modes in KPD2109+4401(Koen 1998), we expect that UV multicolour photometry as developed by Robinson, et al. (1995), will be useful to determine \\emph{if} high--order $\\ell$ modes are present in sdBV stars (as indicated by Brassard et al. 2001 and Bill\\`eres et al. 2000). The argument for high--order $\\ell$ values is particularly interesting in light of the frequencies detected in Group V's data. If the lowest frequency mode is disregarded, the remaining modes have frequency spacings of 13.3, 26.4, 12.8, and 16.3~$\\mu$Hz respectively. If these were all parts of a single, rotationally split mode, it would require $\\ell\\geq 3$. Such an $\\ell$ value should be apparent in UV multicolour photometry (see, for example, Kepler et al 2000). As such, we look forward to the analysis of HST data obtained by Heber (2002). Should Heber's (2002) HST data agree with our $\\ell$=1 interpretation, Feige~48 would make an excellent star to calibrate other mode identification methods in sdBV stars such as optical time-series spectroscopy (O'Toole et al 2002; Woolf et al 2002).  The results of our $O-C$ analysis are consistent with a non-binary nature for the star within the  data limits. It also indicates that using the $O-C$ diagram to detect planets around evolved stars is possible, though in this case we did not detect any. We plan to continue to monitor Feige~48 over the next several years to tighten the constraints on planetary companions.  Our limit on a stellar companion also addresses the origin of sdB stars. Binary evolution is a candidate for producing sdB stars, either through common-envelope evolution (Sandquist, Taam \\& Burkert 2000; Green, Liebert \\& Saffer 2000) or via Roche-lobe overflow near the tip of the red giant branch (Green, Liebert \\& Saffer 2000). Though observations indicate that a great many sdB stars \\emph{are} in binaries (Reed \\& Steining 2003; Green, Liebert \\& Saffer 2000; Han et al 2002), either the evolutionary sequence that produces sdB stars is independent of binary evolution (D'Cruz et al 1996), is bimodal, or has several paths that can result in the production of an sdB star (perhaps including the merger of two low mass white dwarfs as described by Iben \\& Tutukov 1986). For the case of Feige~48, it would appear that it is either in a short period binary (whose orbital period is commensurate with the $\\sim$10 hour rotation period predicted with our model), in a long period binary with an orbital period substantially longer than our data (which would rule out Roche-lobe overflow, so the companion would have no effect on the evolution of the pre-sdB star), in a binary with an extremely low inclination or a single star."
    },
    "0311/astro-ph0311238_arXiv.txt": {
        "abstract": "As the presumed remnant of SN1181, 3C~58 houses one of the youngest known neutron stars in the Galaxy. The properties of this young pulsar and its associated wind nebula differ considerably from those of the Crab, and may well offer a more typical example of the endpoint of massive star collapse. High resolution X-ray studies reveal structures in the inner nebula that may be associated with the pulsar wind termination shock, a jet that may be aligned with the rotation axis, and other regions of enhanced emission. Spectral variations in the PWN are consistent with the expected evolution of the postshock flow, and complex loops of emission are seen in the nebula interior. Limits on the NS surface temperature fall below standard cooling models, indicating that some more rapid neutrino cooling process is required. The outer regions of 3C~58 show thermal emission with enhanced levels of Ne, indicative of shocked ejecta bounding the PWN. ",
        "introduction": "3C~58 (G130.7+3.1) is a Crab-like nebula believed to be associated with the historical supernova SN~1181. X-ray observations reveal reveal a compact source at the center, and a power law spectrum whose index steepens with radius, characteristic of particle injection from a central pulsar. 3C~58 appears to be quite different from the Crab Nebula in many respects, however. The nebula is larger, and the radio luminosity is ten times lower than that of the Crab. The X-ray luminosity is nearly 2000 times lower, perhaps suggesting that any central pulsar has undergone a significant reduction in its injection power. However, recent \\chandra\\ HRC observations (Murray et al.\\ 2002) have revealed PSR~J0205+6449, the second most energetic pulsar known in the Galaxy, powering the nebula; the characteristic age is $\\tau \\equiv P/2\\dot P \\approx 5 \\times 10^3$~yr which, assuming a true age of $\\sim 820$~yr based on the historical association, implies an initial spin period of $P_0 \\approx 60$~ms --- much longer than that inferred for the Crab pulsar.  The emission from the center-most region of 3C~58 is extended, with elongation in the N-S direction, perpendicular to the long axis of the main nebula. The pulsar resides at the center of this compact core nebula, and a jet-like feature protrudes westward. At the western edge resides a narrow radio wisp discovered by Frail \\& Moffett (1993), who suggested that this lies along the pulsar wind termination shock. Using the measured value of $\\dot E = 10^{37.6}{\\rm\\ erg\\ s}^{-1}$ and the pressure determined for the nebula interior, the expected radius of the termination shock is $r_w \\approx 12^{\\prime\\prime}$, in good agreement with the north-south extent of the core nebula in 3C~58 (Slane et al. 2002). This suggests a toroidal structure with an inclination angle of $\\sim 70^\\circ,$ and with a projected pulsar spin axis in the E-W direction. The elongated shape of the entire PWN is consistent with a pulsar axis oriented in the E-W direction (Begelman \\& Li 1992, van der Swaluw 2003).  While unresolved from the compact surrounding nebula, limits on the X-ray flux from PSR~J0205+6449 provide important constraints on models for NS cooling. The upper limit to the thermal emission from the NS surface falls well below predictions from standard cooling models and appears to require the presence of some rapid cooling mechanism (Slane, Helfand, \\& Murray 2002). Such rapid cooling may imply an above-average NS mass, for which an increased proton fraction in the interior allows the direct Urca process to proceed (Yakovlev et al. 2002), or may indicate the presence of pion condensates in the NS interior, which also lead to rapid cooling (Tsuruta et al. 2003).  \\begin{figure}[t] \\centerline{\\psfig{file=slanep1_1.ps,width=5in}} \\caption{ 350 ks X-ray image of 3C~58 taken with the \\chandra\\ ACIS detector. The compact region in the center is extended, and reveals a jet-like feature protruding to the west. On larger scales, complex loop structures are observed throughout the PWN. } \\end{figure}  \\begin{figure}[t] \\centerline{\\psfig{file=slanep1_2.ps,width=2.5in} \\psfig{file=slanep1_3.ps,width=2.5in}} \\caption{ Left: Variation of the spectral index with radius in 3C 58. Right: Spectrum of the outermost regions showing the best-fit power law model (1); a power law + thermal model (2) that requires an overabundance of Ne; and the contribution from only the power law (3) from the best-fit two component model. } \\end{figure} ",
        "conclusions": "3C~58 is a young PWN powered by an energetic pulsar. The compact interior shows evidence of structure from the pulsar wind termination shock. Limits on thermal emission from the neutron star itself require rapid cooling mechanisms such as the direct Urca process or the presence of exotic particles in the NS interior. The overall structure of the PWN itself is dominated by a significant east-west elongation, consistent with a pulsar rotation axis in this same direction, along with complex loop-like structures that may represent magnetic structures in the nebula. Thermal emission from the outer regions of the nebula appear to correspond to ejecta-rich material that is bounding the PWN."
    },
    "0311/astro-ph0311524_arXiv.txt": {
        "abstract": "Motivated by recent observations of deuterated molecules in the interstellar medium, we examine the production of HD and D$_2$ molecules on dust grain surfaces. A mechanism for the enhanced production of these deuterated molecules is studied. This mechanism applies on grain surfaces on which D atoms stick more strongly than H atoms, under conditions of low flux and within a suitable range of temperatures. It is shown that under these conditions the production rates of HD and D$_2$ are greatly enhanced (vs. the H$_2$ production rate) compared with the expected rates based on the adsorption of gas-phase atomic abundances of D and H. The enhancement in the formation rate of HD is comparable with the enhancement due to gas-phase ion-molecule reactions in diffuse clouds. ",
        "introduction": "\\label{intro.}  The formation of molecular hydrogen on the surfaces of dust grains in the interstellar medium is a process of fundamental importance because the formation rate in the gas phase cannot account for the observed abundance \\citep{Gould1963}. Once produced on grain surfaces and ejected or desorbed into the gas, hydrogen molecules are a necessary component in the initiation of gas-phase chemical reaction networks that give rise to the chemical complexity observed in interstellar clouds. The formation of molecular hydrogen has been studied both theoretically \\citep{Hollenbach1970, Hollenbach1971a, Hollenbach1971b, Williams1968, Smoluchowski1981, Smoluchowski1983, Aronowitz1985, Duley1986, Pirronello1988, Sandford1993,Takahashi1999} and experimentally \\citep{Brackmann1961, Schutte1976, Pirronello1997a, Pirronello1997b, Pirronello1999, Manico2001}. Although formation of H$_{2}$ is clearly the most studied surface reaction pertaining to interstellar clouds, there are many other reactions that are thought to take place, for the most part involving H atoms and other species.  Reactions are typically considered to occur via the Langmuir-Hinshelwood mechanism, in which weakly-bound surface species diffuse towards one another.  Some authors have considered different mechanisms, where chemisorption plays a role \\citep{Cazaux02}.  The computational modeling of diffusive chemical reactions on dust grains in the interstellar medium is typically done using rate equation models \\citep{Pickles1977, d'Hendecourt1985,  Brown1990, Brown1990b, Hasegawa1992,  Hasegawa1993a, Hasegawa1993b,  Caselli1993, Willacy1993,  Shalabiea1994}. As long as a sizable number of reactive species exists on a grain surface, the rate equation results are valid. However, in the limit of small average surface populations ($\\le 1$) of reactive species, rate equations are not always suitable,  because they take into account only average densities and ignore the fluctuations as well as the discrete nature of the populations of the atomic and molecular species \\citep{Tielens1995,Tielens1982,Charnley1997,Caselli1998,Shalabiea1998,Stantcheva2001, Stantcheva2002}. Depending on the parameters chosen, this limit can pertain mainly to small grains or can even to ``classical'' grains, namely grains of 0.1$\\mu$ radius. For example, as the number of H atoms, $N_{\\rm H}$, on a grain fluctuates in the range of 0, 1 or 2, the H$_2$ formation rate cannot be obtained from the average number alone. This can be easily understood, since  recombination  requires at least two H atoms simultaneously on the surface.  Recently, a master equation approach was proposed that is suitable for the simulation of diffusive chemical reactions on interstellar grains in the accretion limit \\citep{Biham2001,Green2001,Biham2002}. The master equation takes into account both the discrete nature of the reactive species as well as the fluctuations in their populations. \\citet{Green2001} considered a simple system in which H and O atoms land on classical grains to produce H$_2$, OH,  and O$_2$. Biham and co-workers focused on the formation of H$_2$. With a slow diffusion rate for H based on the analysis by \\citet{Katz1999} of experimental work by Pirronello, Vidali, \\& co-workers \\citep{Pirronello1997a,Pirronello1997b}, they found the master equation approach to be needed mainly for small grains. For the formation of H$_2$, the dynamical variables are the probabilities $P_{\\rm H}(N_{\\rm H})$ that there are $N_{\\rm H}$ atoms on the grain at time $t$. The time derivatives $\\dot{P}_{\\rm H}(N_{\\rm H})$, $N_{\\rm H}=0, 1, 2, \\dots,$ are expressed in terms of the adsorption, reaction and desorption terms. The master equation provides the time evolution of $P_{\\rm H}(N_{\\rm H})$, $N_{\\rm H}=0, 1, 2, \\dots$, from which the recombination rate can be calculated. The master equation has also been used for the study of reaction networks involving multiple species \\citep{Stantcheva2002,Stantcheva2003}.  In addition to the master equation method, a Monte Carlo method \\citep{Tielens1982} was proposed \\citep{Charnley2001}, and applied to surface reaction networks involving multiple species \\citep{Caselli2002,Stantcheva2003}. The weakness of the Monte Carlo approach to surface chemistry is that it cannot be used easily for models in which gas-phase reactions take place simultaneously. Finally, a semi-empirical approach, known as the modified rate method, was proposed by \\citet{Caselli1998}.  Although artificial, the modified rate method is easy to employ and has been used with mixed success \\citep{ Stantcheva2001, Caselli2002}.  In addition to H, which is the most abundant atom in the universe, there is a density of D atoms, which is smaller by a factor of about $10^{-5}$ compared with the density of H atoms in regions where the gas is primarily atomic. If deuterated hydrogen molecules, namely HD and D$_2$, are formed on grain surfaces in the same way as H$_2$, these molecules are expected to be rare due to the low abundance of D atoms. Of course, some differences between the formation rates of HD and D$_2$ and that of H$_2$ are expected because of the possibly different diffusion and desorption rates of H and D atoms on the surface. It was realized early in the study of interstellar chemistry, however, that the formation of HD in diffuse interstellar clouds occurs more efficiently in the gas phase via the sequence of reactions  \\begin{eqnarray} {\\rm H}^{+}  &+&  {\\rm D}  \\longrightarrow {\\rm D}^{+}  +  {\\rm H} \\nonumber \\\\ {\\rm D}^{+}  &+&  {\\rm H}_{2} \\longrightarrow {\\rm HD}  +  {\\rm H}^{+}, \\end{eqnarray}  \\noindent which can lead to enhancements of $\\approx 10^{2}$ over surface formation \\citep{Watson1973,Dalgarno1973} and explain the observed HD/H$_{2}$ abundance ratio in diffuse clouds.  By the stage in which dense interstellar clouds have been formed, HD has become the major repository of deuterium, and is the starting point for a rich chemistry of deuterium fractionation that occurs in the gas phase via ion-molecule reactions  \\citep{Watson1973, Millar1989, Charnley1997,Roberts2000a, Roberts2000b,Rodgers2001, Roberts2002}. Since gas-phase fractionation produces a large D/H atomic abundance ratio in dense clouds, it then becomes possible that surface processes can produce multiply deuterated species in high abundance following the adsorption of D atoms \\citep{Charnley1997,Rodgers2001, Caselli2002,Stantcheva2003}.  Although it has been thought for some time that the formation of HD in diffuse clouds occurs primarily in the gas, the surface processes for the formation of HD and D$_{2}$ have not been looked at carefully. In this paper we propose a mechanism taking place on the surfaces of dust grains that, under certain conditions, can provide a significantly enhanced production of deuterated molecules. We demonstrate the mechanism for the simple case that involves only H and D atoms in diffuse clouds, and examine the formation rates of HD and D$_2$ vs. that of H$_2$. The mechanism is based on the assumption that D atoms stick to the surface more strongly than H atoms so that their desorption rate is lower. An isotope effect of this type has been observed in various experimental situations \\citep{Koehler1988,Hoogers1995}.  Due to the very low cosmic abundance ratio of D vs. H ($10^{-5}$) most of the atoms that hit the grain surfaces in diffuse clouds are H atoms. However, the residence time of D atoms on grains is expected to be longer than that of H atoms, due to their lower desorption rate. Thus, the D/H atomic abundance ratio on the grains is enhanced compared with the gas-phase ratio. As a result, newly adsorbed H (or D) atoms are more likely to find D atoms already residing on the grains, giving rise to an enhanced production of HD molecules. This enhancement turns out to be even larger on very small grains on which only a handful of H and D atoms typically reside. This effect is expected to be significant in diffuse clouds, where the density of the gas phase is low and the grain size distribution is broad and dominated by very small grains \\citep{Li2001}.  The paper is organized as follows. The interaction between hydrogen atoms and dust grains in interstellar clouds in described in Sec. 2. The formation of molecular hydrogen on large grains is described in Sec. 3 using the rate equation model. The formation of molecules on small grains is presented in Sec. 4 using the master equation approach, followed by an asymptotic analysis for very small grains in Sec. 5. The simulations and results are presented in Sec. 6. The results are discussed in Sec. 7 and summarised in Sec. 8. In Appendix A we consider the relevance of the isotope effect studied here to related experiments. ",
        "conclusions": "In this paper we have used a simple description of the grains as spherical objects. In practice the grains are expected to be amorphous with rough surfaces and may include pores. These features may increase the surface area of the grain and introduce a distribution of diffusion and desorption barriers. However, no solid data are available to quantify such effects. Moreover, the analysis of experimental results on molecular hydrogen production on olivine and amorphous carbon surfaces \\citep{Pirronello1997a,Pirronello1997b,Pirronello1999} did not require more than one energy barrier for each process \\citep{Katz1999}.  The analysis presented in this paper is restricted to the case in which the mobility of adsorbed atoms on the surface is due to thermal hopping. Tunneling of H and D atoms on the surface is not included since experimental results indicate that tunneling does not play a significant role in the formation of molecular hydrogen under astrophysically relevant conditions \\citep{Pirronello1997a,Pirronello1997b,Pirronello1999}. In a recent analysis of the results of these experiments, chemisorption processes were taken into account \\citep{Cazaux02}. The study of chemisorption processes is motivated by observations of molecular hydrogen in photon-dominated regions in which the typical grain temperature is around 40 K, making the production of H$_2$ molecules by physisorbed H atoms hopelessly inefficient.  The mechanism presented here for an enhanced production of HD and D$_2$ molecules, vs. H$_2$, takes place in a window of grain temperatures. For the surface parameters used here this temperature window is between 16 - 20 K. In this temperature range the production efficiency of H$_2$ is reduced, while the efficiency of HD and D$_2$ production remains steady. The effect is observed for grains of all sizes, but the relative enhancement is larger on small grains.  It will be interesting to see if the enhancement plays a significant role in chemical models of diffuse and translucent clouds; relevant model calculations are planned by Le Petit and Roueff (private communication)."
    },
    "0311/astro-ph0311462_arXiv.txt": {
        "abstract": "This document presents a ray-optics analysis of the response of a two-element radio interferometer to the presence of a plane gravitational wave.  A general expression for the differential phase observed between the two receiving stations as a result of an arbitrary gravitational wave is determined, as well as the specific responses to monochromatic and black hole ring-down waveforms.  Finally, the possibility of gravitational wave detection via this mechanism is discussed in the context of interferometer noise. ",
        "introduction": "A gravitational wave is a perturbation to the geometry of space that propagates through space at the speed of light.  It's effect is to stretch and shrink the distances between objects in the directions transverse to its direction of travel.  The distortion is volume-preserving, so while it is shrinking distances in one direction, it stretches them in another.  The stretching and shrinking effect is not a co-ordinate artifact.  A measuring stick used to monitor the distance between two objects freely-falling in the presence of a gravitational wave will measure one separation distance at one time and a different separation later.  The stretching and shrinking is also linear:  given the change in distance between two objects separated by a distance \\(L\\), the change in distance between two objects separated by \\(2L\\) will be doubled.  The amplitude of a gravitational wave is expressed as a dimensionless quantity giving the change in distance per unit distance --- a shear strain in space.  Any apparatus capable of accurately measuring the distance between two objects, or the change in the distance between two objects, as a function of time is, in principle, capable of detecting a gravitational wave.  The question of the effectiveness of the apparatus in accomplishing this task is one of sensitivity and signal-to-noise ratio.  One design being actively pursued by several research groups today is based on a Michelson-type laser interferometer.  In this design, a laser beam is split and sent down the two orthogonal arms of the apparatus, and then at the end of each is reflected off a suspended mass back toward the splitter where the beams are recombined \\cite{abramovici:1992, caron1997, danzmann1997, luck1997, kawabe1997}.  If the length of one or the other arms changes, the result will be a change in the differential phase between the returning laser beams and so a passing gravitational wave will appear as a shift in the interference pattern at the output of the interferometer.  A variation on this technique that is also being pursued by several research groups is a detector consisting of an Earth-based transmitter/receiver and a spacecraft carrying a radio or optical transponder.  There are currently proposals to launch specially-designed vehicles dedicated to this task using optical links from the ground to the vehicle and back \\cite{ni2002}.  It is also possible to use existing missions.  Deep-space vehicles generally carry radio transponders that are used as aids to navigation.  By maintaining a coherent up-link/down-link channel at high frequency stability with appropriate modulations one can measure the range and range-rate to the spacecraft with high precision and thereby search for the passage of a gravitational wave through the solar system \\cite{bertotti1998}.  In this document, another type of gravitational wave detection apparatus similar to this last technique is introduced and investigated.  Figure \\ref{fig1} shows a radio interferometer arrangement with a distant radio source and two receiving elements. \\begin{figure} \\begin{center} \\begin{picture}(0,0)% \\includegraphics{concept.ps}% \\end{picture}% \\setlength{\\unitlength}{3947sp}% \\begingroup\\makeatletter\\ifx\\SetFigFont\\undefined% \\gdef\\SetFigFont#1#2#3#4#5{% \\reset@font\\fontsize{#1}{#2pt}% \\fontfamily{#3}\\fontseries{#4}\\fontshape{#5}% \\selectfont}% \\fi\\endgroup% \\begin{picture}(2446,3990)(-22,-3068) \\put(151,-2086){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{10}{12.0}{\\familydefault}{\\mddefault}{\\updefault}2}}} \\put(2176,-2086){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{10}{12.0}{\\familydefault}{\\mddefault}{\\updefault}1}}} \\put(901,-1636){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{10}{12.0}{\\familydefault}{\\mddefault}{\\updefault}polarization}}} \\put(151,-1111){\\makebox(0,0)[lb]{\\smash{\\SetFigFont{10}{12.0}{\\familydefault}{\\mddefault}{\\updefault}g.w.}}} \\end{picture} \\end{center} \\caption{A gravitational wave moving across the lines of sight in the plane of a two-station interferometer.  One of the gravitational wave's polarization axes is aligned (ignoring parallax) with the lines of sight.} \\label{fig1} \\end{figure} It also shows a gravitational wave travelling across the lines of sight to the radio source in the plane defined by the source and the two interferometer elements (to be referred to, hereinafter, as the plane of the interferometer).  For the purpose of illustration, one of the gravitational wave's polarization axes is aligned (ignoring parallax) with the lines of sight.  Because the gravitational wave travels at finite speed, any perturbation to distances transverse to its direction of travel will influence the signal arriving at receiver 1 in advance of its influence on the signal arriving at receiver 2.  The presence of a gravitational wave will, therefore, manifest itself as a time-varying difference between the arrival times of radio signals at the receivers --- a time-varying differential phase between the stations.  Such a detector is similar in operation to the single arm, spacecraft-based, coherent up-link/down-link type of detector.  In this case, however, a second arm is introduced to act as a phase reference for the first which permits the use of a broad-band radio signal.  What follows is an analysis of the magnitude of the differential phase between the stations when a gravitational wave is present, and a preliminary assessment of the signal-to-noise ratio in this type of detector.  We will consider a radio signal emanating from the radio source toward each of the receiving stations.  Each radio signal is treated as a null geodesic in space-time, originating on the radio source's world line.  The intersections of the geodesics with the world lines of the receivers define events in space-time giving the co-ordinate times of arrival of the signal at each receiving element.  By determining the difference in the times of emission of the signals being received simultaneously at the two receiving elements, and multiplying this time difference by the angular frequency of the radio emission one obtains the differential phase between the two receiving elements.  The calculation will first be done for arbitrary receiver locations and arbitrary gravitational waveforms, then a simplified interferometer arrangement will be introduced in order to more easily discuss the response to specific gravitational waves.  Throughout the calculations, the units for length will be chosen so as to make the numerical value of the speed of light, \\(c\\), equal to 1. ",
        "conclusions": "The technique, introduced here, of using the time-dependent phase difference between two radio receiving stations to watch for the passage of a gravitational wave would seem to be a promising method of gravitational wave observation worthy of further investigation.  In the future, a more detailed analysis of realistic interferometer configurations and data analysis techniques will shed more light on true sensitivity of this type of detector."
    },
    "0311/astro-ph0311181_arXiv.txt": {
        "abstract": "{We consider generation and evolution of small-scale magnetic fields in neutron stars. These fields can be generated by small-scale turbulent dynamo action soon after the collapse when the proto-neutron star is subject to convective and neutron finger instabilities. After instabilities stop, small-scale fields should be frozen into the crust that forms initially at high density $\\sim 10^{14}$ g/cm$^{3}$ and then spreads to the surface. Because of high crustal conductivity, magnetic fields with the lengthscale $\\sim 1-3$ km can survive in the crust as long as 10-100 Myr and form a sunspot-like structure at the surface of radiopulsars. ",
        "introduction": "Our knowledge of the magnetic field strength in neutron stars comes mainly from radiopulsars with measured spin-down rates. With the assumption that the spin-down torque on the pulsar is determined by its magnetodipole radiation (Ostriker \\& Gunn 1969), the spin period, $P$, and spin-down rate, $\\dot{P}$, are related to the field strength at the magnetic pole, $B_{d}$, by \\begin{equation} P \\dot{P} = \\frac{8 \\pi^{2} B_{d}^{2} R^{6}}{3 c^{2} I} \\end{equation} where $I$ is the moment of inertia and $R$ is the radius (we assume that the magnetic and rotation axes are perpendicular). The magnetic fields inferred from the spin-down data range from $\\sim 5 \\times 10^{13}$ to $\\sim 10^{8}$G but, most likely, these fields characterize the global magnetic configuration of neutron stars rather than a fine magnetic structure near the stellar surface.  Measurements of the spin-down rate are not, however, the only way to obtain information about the neutron star magnetic fields. Recent observations of the X-ray spectra features of some pulsars provide one more opportunity to look into the magnetic field near the surface. Absorbtion features in the spectrum of the isolated pulsar 1E 1207.4-5209 were associated by Sanwal et al. (2002) with atomic transitions of once-ionized helium in the neutron star atmosphere with a strong ``surface'' magnetic field, $B_{s} \\sim 1.5 \\times 10^{14}$G. An estimate of the spin-down rate of this pulsar provides the strength of the ``dipole'' magnetic field, $B_{d} \\sim (2-4) \\times 10^{12}$G (Pavlov et al. 2002) that is typical for radio pulsars of that age ($\\sim 0.2-1.6$ Myr). Becker et al. (2002) found an emission line in the X-ray spectrum of PSR B1821-24 that could be interpreted as cyclotron emission from the corona above the pulsar's polar cap. This emission line is likely formed in a magnetic field $B_{s} \\sim 3 \\times 10^{11}$G which is approximately two orders of magnitude stronger than the ``dipole'' field inferred from $P$ and $\\dot{P}$. Haberl et al. (2003) interpreted a broad absorption feature in the spectrum of the isolated neutron star RBS1223 as a cyclotron line produced by protons in the magnetic field $\\sim (2-6) \\times 10^{13}$ G. These measurements provide strong evidence that the local magnetic fields on the neutron star surface can exceed the conventional ``dipole'' field.  Another piece of evidence on the distinction between the ``dipole'' and ``surface'' magnetic fields comes from the data on emission of radio pulsars. Recently, Gil \\& Mitra (2001) and Gil \\& Melikidze (2002) argued that the formation of a vacuum gap in radio pulsars is possible if the actual surface magnetic field near the polar cap is very strong, $B_{s} \\sim 10^{13}$G, irrespective of the field measured from the $P$-$\\dot{P}$ data. Also, radio emission from the recently discovered pulsar PSR J2144-3933 with the longest period 8.5 s, which lies extremely far beyond the conventional death line, can be understood if this pulsar has a strong surface magnetic field of complex geometry (Gil \\& Mitra 2001). Radio emission of many other pulsars which lie in the pulsar graveyard and should be radio silent can be explained if one adopts the model with a strong and complex surface field with a small curvature of the field lines ($< 10^{6}$ cm). This model is consistent with the conclusion of Arons \\& Scharlemann (1979) and Arons (1993) that pulsars with very long periods ($\\geq 5$ s) require a more complex field configuration than a dipole if pair creation is essential for the mechanism of radio emission. Analysing the phenomenon of drifting subpulses observed in many pulsars, Gil \\& Sendyk (2000) found that their behaviour is consistent with the vacuum gap maintained by a strong sunspot-like magnetic field. Following this idea Gil, Melikidze \\& Mitra (2002 b) found that in the famous case of PSR 0943+10 (Deshpande \\& Rankin 1999, 2001) the ``surface'' field should be approximately 7 times stronger than the ``dipole'' component inferred from the spin-down rate. A sunspot-like configuration of the surface magnetic field is also suggested by the spin-down index in some pulsars (Cheng \\& Ruderman 1993, Ruderman, Zhu \\& Cheng 1998). Also, Cheng \\& Zhang (1999), analysing the X-ray emission from the polar regions of the rotation-powered pulsars, argued that $B_{s} \\sim 10^{13}$G and the characteristic curvature of the field lines is around $10^{5}$cm resembling the sunspot-like structure.  Rapidly growing amount of evidence on the distinction between the local field strength at the stellar surface and the global ``dipole'' field suggests that this can be a general phenomenon in neutron stars. In this paper, we propose a scenario of the formation of such complex magnetic configurations with the strength of a small scale ``surface'' field in excess of the ``dipole'' component. It is generally accepted that neutron stars are subject to hydrodynamic instabilities soon after their birth in the core collapse (see, e.g., Epstein 1979, Livio, Buchler \\& Colgate 1980, Burrows \\& Lattimer 1986). The convective stage lasts about 30-40 s (Miralles, Pons \\& Urpin 2000, 2002) and, under certain conditions, turbulent motions can amplify the magnetic field via dynamo action (Thompson \\& Duncan 1993, Xu \\& Busse 2001). The generated magnetic field must be frozen into the crust that is formed in the course of neutron star cooling. Because of high crustal conductivity, magnetic fields with a relatively small lengthscale, $\\sim 10^{5}$ cm, can survive during the active lifetime of radio pulsars.  The paper is organized as follows. In Section 2, we consider convection and associated dynamo action in protoneutron stars, and estimate the field that can be generated during this stage. We discuss the crust formation and small scale crustal magnetic structures in Section 3. Our results are briefly summarized in Section 4. ",
        "conclusions": "We have considered the formation and evolution of small-scale magnetic structures in the neutron star surface layers. A wide spectrum of these structures can be generated by the small-scale turbulent dynamo action during the unstable phase that lasts $\\sim 30-40$ s after the neutron star birth. There are two substantially different unstable regions in the proto-neutron star, with the convective instability active in the inner region and the neutron-finger instability more efficient in the outer region. Generally, the small-scale dynamo can generate small-scale magnetic structures in both unstable zones. Due to high conductivity of nuclear matter, the lengthscale of generated magnetic structures spreads from the main scale of turbulence, $L \\sim 1-3$ km (comparable to the pressure scale height), to extremely short lengthscales determined by ohmic dissipation. After instabilities stop, structures with short lengthscales decay on a short timescale $\\propto \\ell^{2}$ due to finite electrical conductivity whereas fields with larger lengthscales can survive for a longer time. The crust formation that starts almost immediately after instabilities stop may have the decisive influence on the evolution of such magnetic fields. These magnetic structures can be frozen into the crystallized matter and then evolve in the crustal layers. The neutron star cooling increases the conductivity of the crust and, as a result, small-scale magnetic structures decay extremely slowly. Our calculations show that structures with the lengthscale $L \\sim 1-3$ km can survive as long as 10-100 Myr that is basically comparable to the active life-time of radiopulsars.  The strength of the magnetic field generated by the small-scale dynamo action is approximately determined by equipartition and decreases with the decreasing lengthscale. For magnetic structures with $L \\sim 1-3$ km, the magnetic field can be as strong as $\\sim 5 \\times 10^{12}-2 \\times 10^{13}$ G even for radiopulsars as old as $\\sim 10-100$ Myr. Note that the decay of small-scale magnetic fields is qualitatively similar to that of the large-scale crustal field considered by Urpin \\& Konenkov (1997). The origin of a large-scale field in neutron stars is still debatable, but it is possible that this field was generated by some mechanism in the layer with $\\rho < 2 \\times 10^{14}$ g/cm$^{3}$ that corresponds to the crust (see, e.g., Bonanno, Rezzolla \\& Urpin 2003). Then, the thickness of a layer occupied by the large-scale field can generally differ from $L$. If the initial depth of a large-scale field is smaller than $L$, then the small-scale field decreases slower, and we observe a radiopulsar with magnetic \"spots\" where the field is stronger than the dipole field inferred from the spin-down rate. On the contrary, if the initial depth of a large-scale field is larger than $L$, then small-scale structures decay faster than the dipole field, and the resulting magnetic field becomes more regular with the age.  Small-scale magnetic structures with $L \\sim 1-3$ km and $B \\sim 5 \\times 10^{12}-2 \\times 10^{13}$ G at the surface may have an important influence on many properties of radiopulsars. As already mentioned in the Introduction, this range of the surface magnetic field is favourable for the inner vacuum gap formation in pulsars (Gil \\& Melikidze 2002). The sparking discharge of this gap produces filaments of electron-positron plasma, whose presence seems absolutely necessary for generation of coherent pulsar radio emission. Moreover, the ${\\bf E} \\times {\\bf B}$ drift of spark plasma should be manifested as the observed subpulse drift, provided that the actual surface field has some degree of axial symmetry, with a tendency to converge at the local pole  (see Gil, Melikidze \\& Geppert, 2003 for review). These authors assumed that one small-scale structure (a \"spot\") coincides to some extent with the canonical (dipolar) polar cap in the sence that the magnetic field lines form a complex magnetic configuration near the pole but connect smoothly with a subset of open dipolar field lines at a larger radius (Gil, Melikidze \\& Mitra 2002a). This means that the actual polar cap is defined by those non-dipolar field lines which penetrate the light cylinder.  A study of small-scale magnetic structures can also provide information regarding the very early evolutionary stage of neutron stars since the observed fields were frozen into the crust soon after the neutron star birth. For example, the presence of small-scale magnetic structures in radiopulsars can be a good evidence that proto-neutron stars pass the turbulent stage in their evolution."
    },
    "0311/astro-ph0311304_arXiv.txt": {
        "abstract": "The discrete-dipole approximation (DDA) is a powerful method for calculating absorption and scattering by targets that have sizes smaller than or comparable to the wavelength of the incident radiation. We present a new prescription -- the Surface-Corrected Lattice Dispersion Relation (SCLDR) -- for assigning the dipole polarizabilities that takes into account both target geometry and finite wavelength. We test the SCLDR in DDA calculations using spherical and ellipsoidal targets and show that for a fixed number of dipoles, the SCLDR prescription results in increased accuracy in the calculated cross sections for absorption and scattering. We discuss extension of the SCLDR prescription to irregular targets. ",
        "introduction": "The discrete-dipole approximation (DDA) is a numerical technique for calculating scattering and absorption of electromagnetic radiation by targets with sizes smaller than or comparable to the incident wavelength. The method consists of approximating the target by an array of polarizable points (dipoles), assigning polarizabilities at these locations based on the physical properties of the target, and solving self-consistently for the polarization at each location in the presence of an incident radiation field. This procedure can yield arbitrarily accurate results as the number of dipoles used to approximate the target is increased. However, computational considerations limit the number of dipoles that can be used. Hence, methods for increasing the accuracy for a fixed number of dipoles are desirable.  A key factor in determining the level of accuracy that can be reached for a given number of dipoles is the prescription for assigning dipole polarizabilities. In this work, we present a new polarizability prescription that takes into account both target geometry and the finite wavelength of incident radiation. We test this technique in calculations of absorption and scattering by spherical and ellipsoidal targets and show that for a fixed number of dipoles, it generally provides increased accuracy over previous methods. In Section~\\ref{pol} we discuss previous polarizability prescriptions and develop the new method. In Section~\\ref{calc} we present calculations testing the new prescription, and in Section~\\ref{conclusion} we discuss our results. ",
        "conclusions": "Conclusion}  We introduce a new DDA polarizability prescription -- the Surface-Corrected Lattice Dispersion Relation (SCLDR). This technique builds on previous work, principally by Draine \\& Goodman \\cite{drai94} and Rahmani, Chaumet \\& Bryant \\cite{rahm02}, to account properly for both finite wavelength and target geometry. We have tested the new polarizability prescription in calculations of absorption and scattering by spherical and ellipsoidal targets. These tests show that the SCLDR performs generally better than previous prescriptions which took account either of finite wavelength or of target geometry but not both. The SCLDR technique is most easily applicable to target shapes for which there exists an analytical solution to the electrostatic applied field problem, but it can be applied to any dielectric target (homogeneous or inhomogeneous, isotropic or anisotropic) provided that the electrostatic problem can at least be solved numerically to obtain the tensors ${\\bf C}_i$ (see eq.~\\ref{eq:Em_and_E0}). In such cases, it generally provides a significant increase in accuracy over previous methods, especially for highly absorptive materials."
    },
    "0311/astro-ph0311418_arXiv.txt": {
        "abstract": "In this paper, we discuss the low frequency spectrum of SN 1993J with GMRT. We observed SN 1993J at several epochs in 20cm, 50cm, 90cm and 125cm wavelengths and achieved near simultaneous spectra. We fit synchrotron self absorption (SSA) and free-free models to the data. We compare the size of SN obtained using SSA fits to that of size extrapolated from VLBI measurements at various epochs using public data at earlier epochs. We find  that the synchrotron self absorption process is insufficient to reproduce the observed size of the supernova under the assumption of equipartition between magnetic fields and relativistic electrons. We also derive the evolution of spectral index and magnetic field at several epochs. ",
        "introduction": "\\label{sec:1} SN 1993J, a typical type IIb supernova,  exploded on March 28, 1993. It was the nearest extragalactic supernova observed (3.63 Mpc away) and due to its high positive declination, it was easily accessible by most of the telescopes for observations for most part of the year. For this reason SN 1993J is the most detailed studied supernova after SN 1987A, in all wavebands. The radio emission in supernovae is due to the synchrotron emission of relativistic electrons in the presence of magnetic field. The radio emission is absorbed in its early phases due to the presence of dense circumstellar medium. It can also be absorbed because of its own highly dense ejecta (synchrotron self absorption).   In section 2 we describe the GMRT observations. In section 3, we discuss the spectra, the derived physical parameters from the various fits to the spectra  and their physical interpretations. ",
        "conclusions": ""
    },
    "0311/astro-ph0311132_arXiv.txt": {
        "abstract": "{This paper presents a detailed study of the chemistry in the outflow associated with the low-mass protostar NGC~1333-IRAS2A down to 3\\arcsec\\ (650 AU) scales. Millimeter-wavelength aperture-synthesis observations from the Owens Valley and Berkeley-Illinois-Maryland-Association interferometers and (sub)millimeter single-dish observations from the Onsala Space Observatory 20~m telescope and Caltech Submillimeter Observatory are presented. The interaction of the highly collimated protostellar outflow with a molecular condensation $\\sim$15000~AU from the central protostar is clearly traced by molecular species such as HCN, SiO, SO, CS, and CH$_3$OH. Especially SiO traces a narrow high velocity component at the interface between the outflow and the molecular condensation.  Multi-transition single-dish observations are used to distinguish the chemistry of the shock from that of the molecular condensation and to address the physical conditions therein. Statistical equilibrium calculations reveal temperatures of 20 and 70~K for the quiescent and shocked components, respectively, and densities near $10^6$~${\\rm cm}^{-3}$. The line-profiles of low- and high-excitation lines are remarkably similar, indicating that the physical properties are quite homogeneous within each component. Significant abundance enhancements of two to four orders of magnitude are found in the shocked region for molecules such as CH$_3$OH, SiO and the sulfur-bearing molecules. HCO$^+$ is seen only in the aftermath of the shock consistent with models where it is destroyed through release of H$_2$O from grain mantles in the shock. N$_2$H$^+$ shows narrow lines, not affected by the outflow but rather probing the ambient cloud. The overall molecular inventory is compared to other outflow regions and protostellar environments. Differences in abundances of HCN, H$_2$CO and CS are seen between different outflow regions and are suggested to be related to differences in the atomic carbon abundance. Compared to the warm inner parts of protostellar envelopes, higher abundances of in particular CH$_{3}$OH and SiO are found in the outflows, which may be related to density differences between the regions. ",
        "introduction": "One of the manifestations of a newly formed low-mass protostar is the presence of a highly collimated and energetic outflow or jet. A natural consequence of the propagation of such high velocity outflows through the protostellar envelope and the ambient molecular medium are shockfronts \\citep{reipurth99}. Shocks both heat and compress the gas and also trigger chemical reactions in the gas-phase, leading to a different chemistry than observed otherwise. Shock processing of dust grains may lead to the injection of atoms and molecules back into the gas, which further distinguishes the chemistry in the shocked region from that of a quiescent protostellar environment \\citep[e.g.,][]{tielens99}. This paper presents a study of the physics and chemistry of the outflow associated with a well-known class 0 young stellar object \\object{NGC~1333-IRAS2} using high-resolution millimeter wavelength interferometer and multi-transition single-dish observations. The \\object{NGC~1333-IRAS2} outflow provides a unique opportunity to study the effects of outflows on ambient molecular clouds, as the main shock is well separated from the central protostar and shows a relatively simple morphology. The combination of single-dish and interferometry observations makes it possible to discuss the physical and chemical properties of the outflowing gas and to address the spatial differentiation of the chemistry in the outflow region resolved by the interferometer observations.  Studies of molecular abundances in regions of high outflow activity provide insight into the dependence of the chemical reaction networks on temperature and density. Furthermore it is important to recognize the effect of outflow-triggered chemistry in the inner protostellar envelope, to disentangle it from emission from a circumstellar disk or to address the effect of passive heating by the central protostar. In the central part of the protostellar envelope, thermal evaporation of dust grain mantles can lead to a distinct chemistry as is seen in the case of low-mass protostars \\citep[e.g IRAS~16293-2422,][]{ceccarelli00a,ceccarelli00b,schoeier02,cazaux03}.  \\object{NGC~1333-IRAS2} (also known as \\object{IRAS~03258+3104}; hereafter simply \\object{IRAS2}) is located in the \\object{NGC~1333} molecular cloud, harboring several class 0 and I objects, first identified through IRAS maps by \\cite{jennings87}. Continuum observations reveal that \\object{IRAS2} is a binary source with two components, \\object{IRAS2A} and 2B, separated by 6500~AU (30\\arcsec) \\citep{sandell94knee,blake96,looney00}. \\object{IRAS2A} is responsible for a highly collimated east-west outflow giving rise to a strong shock $\\sim$15000~AU from the central continuum source (Fig.~\\ref{scuba_overview}). A strong CO outflow in the north-south direction has also been observed \\citep{liseau88,engargiola99}, which originates within a few arcseconds from \\object{IRAS2A} \\citep{n1333i2art}. \\begin{figure*} \\resizebox{\\hsize}{!}{\\rotatebox{90}{\\includegraphics{0216_1.eps}}} \\caption{Overview of the \\object{IRAS2A} outflow region. The colored image shows the SCUBA 850~$\\mu$m emission tracing the cold dust. N$_2$H$^+$ BIMA line observations and SiO OVRO line observations (this paper) are indicated by blue and green contours, respectively. CO 2-1 emission from the \\cite{engargiola99} is indicated by the black/white dashed contours. The red stars indicate the position of IRAS2A and IRAS2B from \\cite{n1333i2art}.}\\label{scuba_overview} \\end{figure*}  \\cite{langer96} mapped the entire \\object{NGC~1333} region in CS and identified two peaks in CS emission toward the \\object{IRAS2} outflow. They suggested that these are associated with red-shifted (eastern) and blue-shifted (western) bow shock components of the outflow. \\cite{sandell94knee} reported bright CH$_{3}$OH emission toward the red-shifted (eastern) outflow lobe and correspondingly a high CH$_{3}$OH abundance, possibly enhanced by the shock. The more detailed structure of the CH$_{3}$OH emission from the \\object{IRAS2A} east-west outflow was discussed by \\cite{bachiller98}, who mapped the outflow positions at $\\approx$~3\\arcsec\\ using the IRAM interferometer and 30~m single-dish telescope. \\citeauthor{bachiller98} derived the physical conditions in the shock interaction zone from LVG calculations and obtained a density of $\\sim 10^6$~cm$^{-3}$ and temperature of $\\sim 100$~K.  \\citeauthor{bachiller98} also found that the observed methanol emission translates to a large enhancement of CH$_3$OH by a factor $\\sim 300$ in the \\object{IRAS2A} outflow. CH$_{3}$OH is thought to be released directly from the dust grain mantles and is often seen to be associated with protostellar outflows \\citep[e.g.,][]{bachiller95}. Other often-used tracers of shocks associated with protostellar outflows are Si-bearing species, in particular SiO \\citep[e.g.,][]{martinpintado92,codella99,garay02}. High abundances of these species may mark a clear distinction of shocked gas from unprocessed gas in the envelopes around low-mass protostars \\citep{bachiller97,garay98,bachiller01}.  In this paper we present a study of the detailed chemistry of the shock associated with the \\object{IRAS2A} outflow based on observations of a wide range of molecular lines at $\\sim 3-6$\\arcsec\\ resolution from the Owens Valley Radio Observatory (OVRO) and Berkeley Maryland Illinois Association (BIMA) millimeter interferometers, together with millimeter and submillimeter single dish observations from the Onsala 20~m telescope (OSO) and the Caltech Submillimeter Observatory 10.4~m telescope (CSO). Parts of the OVRO observations have previously been presented by \\cite{blake96}. Sect.~\\ref{observations} describes the observations and reductions. The maps from the interferometry observations are presented and discussed in Sect.~\\ref{interferometry}, while the single dish observations are treated in Sect.~\\ref{singledish}. The physical and chemical properties of the shock region are analyzed using statistical equilibrium calculations as described in Sect.~\\ref{analysis} and molecular abundances are derived. Sect.~\\ref{discussion} discusses the inferred chemistry and compares it to other well-studied outflow regions, to other types of star-forming environments and to available models for the chemistry in outflow regions. The main findings are summarized in Sect.~\\ref{conclusion}. A companion paper \\citep{n1333i2art} presents details of a millimeter-wavelength interferometer study of the environment surrounding the central protostellar system. ",
        "conclusions": "} A (sub)millimeter study of the shock associated with the \\object{NGC1333-IRAS2A} outflow has been presented. Both single-dish and interferometry line observations are presented, which allows for a detailed discussion of both the physical and chemical properties in the shocked region and the spatial distribution of emitting species. The main findings are as follows:  \\begin{enumerate} \\item Interferometer observations of the outflow region reveal a distinct morphology with a narrow high velocity feature in CS, SiO and HCN, while the low velocity part traces the more spatially extended material. HCO$^+$ is seen only in the aftermath of the shock, whereas N$_2$H$^+$ shows very narrow lines and does not seem to be present in the outflowing gas.  \\item Statistical equilibrium calculations show that the region can be divided into two distinct components: low velocity material in the ambient, quiescent cloud and high velocity material associated with the outflow. It is found that the conditions in the quiescent material are consistent with a temperature of $\\sim$20~K and density of $1\\times 10^6\\, {\\rm cm}^{-3}$, while the shocked gas is warmer with a temperature of $\\sim$70~K and density of $2\\times 10^6\\, {\\rm cm}^{-3}$. Within these components, however, the physical conditions are remarkably homogeneous as indicated by the similarities between the lineshapes for different transitions of various molecules and between the wing emission in spectra from interferometry and single-dish observations.  \\item The chemistry in the outflow and quiescent regions are significantly different. CH$_3$OH, SiO and the sulfur-bearing species are significantly enhanced in the outflowing gas, whereas HCN, H$_2$CO, HCO$^+$ and N$_2$H$^+$ show abundances more similar to those found in molecular clouds and protostellar envelopes. Compared to the well-studied \\object{L1157} outflow the CS, HCN and H$_2$CO abundances are markedly lower. This could be due to differences in the amount of atomic carbon in the shocked or pre-shocked gas. Higher abundances of the shock-tracing molecules (in particular SiO and CH$_{3}$OH) in outflows compared to the inner warm envelopes may be related to more rapid destruction of these molecules in the envelopes where the densities are higher.  \\item A scenario is suggested where the highly collimated protostellar outflow is progressing into a region with a steep density gradient, focusing the shock and giving rise to the narrow morphology observed. This leads to a shock-induced chemistry, which can explain both the morphologies and qualitatively the abundances of the different molecules. In particular, CH$_{3}$OH is seen to be greatly enhanced reaching an abundance of about 5\\% of the observed CO abundance in the shocked gas. CH$_{3}$OH is, however, not observed to as high velocities as seen in SiO, possibly as a consequence of CH$_{3}$OH being more volatile. Thus together with, e.g., HCN and H$_2$CO, CH$_{3}$OH is not able to survive at higher velocities. \\end{enumerate}  This work illustrates the large impact of protostellar outflows in shaping the physical and chemical properties of their parental environment. The combination of high-resolution interferometer observations and single-dish spectra makes it possible to address the physical and chemical conditions in the shocked and ambient gas and to investigate the spatial variation and time-scales characteristic for the shock induced chemistry. So-far only a few shocks have been studied in great chemical detail. Similar systematic studies of a large number of different outflows will allow for a more detailed comparison between outflows and shocks of different velocities and energetics and in different environments. Future observations with facilities such as the SMA, CARMA, and ALMA will allow further studies of the variation of physical and chemical conditions in shocks through high resolution, high sensitivity multi-transition molecular line observations. Also high spatial resolution observations of H$_2$O lines with Herschel-HIFI can confirm the anticorrelation between HCO$^+$ and H$_2$O. All such more detailed observational studies will serve as important starting points for more detailed physical and chemical models for shocks in protostellar environments."
    },
    "0311/astro-ph0311510.txt": {
        "abstract": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  We report new chemical abundances of 23 bright red giant members of the globular cluster M3, based on high-resolution (R~$\\sim$~45000) spectra obtained with the Keck~I telescope. The observations, which involve the use of multislits in the HIRES Keck~I spectrograph, are described in detail.  Combining these data with a previously-reported small sample of M3 giants obtained with the Lick 3m telescope, we compare metallicities and [X/Fe] ratios for 28 M3 giants with a 35-star sample in the similar-metallicity cluster M13, and with Galactic halo field stars having [Fe/H]~$<$~--1. For elements having atomic number A~$\\geq$~A(Si), we derive little difference in [X/Fe] ratios in the M3, M13 or halo field samples.  All three groups exhibit C depletion with advancing evolutionary state beginning at the level of the red giant branch ``bump'', but the overall depletion of about 0.7 to 0.9~dex seen in the clusters is larger than that associated with the field stars. The behaviors of O, Na, Mg and Al are distinctively different among the three stellar samples. Field halo giants and subdwarfs have a positive correlation of Na with Mg, as predicted from explosive or hydrostatic carbon burning in Type~II supernova sites. Both M3 and M13 show evidence of high-temperature proton capture synthesis from the ON, NeNa, and MgAl cycles, while there is no evidence for such synthesis among halo field stars. But the degree of such extreme proton-capture synthesis in M3 is smaller than it is in M13: the M3 giants exhibit only modest deficiencies of O and corresponding enhancements of Na, less extreme overabundances of Al, fewer stars with low Mg and correspondingly high Na, and no indication that O depletions are a function of advancing evolutionary state as has been claimed for M13.  We have also considered NGC~6752, for which Mg isotopic abundances have been reported by Yong \\etal\\ (2003). Giants in NGC~6752 and M13 satisfy the same anticorrelation of O abundances with the ratio (\\iso{25}{Mg}+\\iso{26}{Mg})/\\iso{24}{Mg}, which measures the relative contribution of rare to abundant isotopes of Mg. This points to a scenario in which these abundance ratios arose in the ejected material of 3-6 solar mass cluster stars, material that was then used to form the atmospheres of the presently evolving low-mass cluster stars. It also suggests that the low oxygen abundance seen among the most evolved M13 giants arose in hot-bottom O to N processing in these same intermediate-mass cluster stars. Thus mixing is required by the dependence of some abundance ratios on luminosity, but an earlier nucleosynthesis process in a hotter environment than giants or main-sequence stars is required by the variations previously seen in stars near the main sequence. The nature and the site of the earlier process is constrained but not pinpointed by the observed Mg isotopic ratio. ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Thirty years have elapsed since the initial discovery (Osborn 1971\\nocite{osb71}) of CN variations among giants in an otherwise mono-metallic globular cluster. A flood of evidence has now accumulated indicating that several light elements, not only carbon and nitrogen, but also oxygen, sodium and aluminum have large and often correlated abundance variations among the giants of individual clusters (see reviews by Kraft 1994\\nocite{kra94}, Briley \\etal\\ 1994\\nocite{bri94}, Wallerstein \\etal\\ 1997\\nocite{wal97}, DaCosta 1998\\nocite{dac98}). The range of variation in some of these elements often exceeds an order of magnitude, far more than expected from conventional ``first dredge-up'' predictions (see Iben \\& Renzini 1984\\nocite{ibe84}). It is clear that the surface chemical compositions of most cluster giants are reflections of the products of multiple proton-capture chains (Denissenkov \\& Denissenkova 1990\\nocite{den90}; Cavallo \\etal\\ 1996\\nocite{cav96}; Langer, Hoffman \\& Zaidins 1997\\nocite{lan97}; Denissenkov \\etal\\ 1998\\nocite{den98}) which involve conversion not only of C to N, but also O to N, Ne to Na and Mg to Al.  Less well understood is the site (or indeed sites) in which proton-capture syntheses actually take place. Two possible scenarios have been invoked. In the first, the ``primordial'' picture, the variations are a result of pre-existing syntheses which made contributions to the gas out of which the presently observable low-mass stars were assembled. This approach has two subcases: {\\it (1)} the stars formed from gas already possessing these variations at the epoch of cluster formation, a truly ``primordial'' picture, or {\\it (2)} the presently observed low-mass stars were ``polluted'' by the accretion of processed material ejected by intermediate-mass (3-6~\\Msun) stars in their late evolutionary stages (Cottrell \\& DaCosta 198l\\nocite{cot81}, Cannon \\etal\\ 1998\\nocite{can98}, Ventura \\etal\\ 2001\\nocite{ven01}, Yong \\etal\\ 2003\\nocite{yon03}). In either subcase, variations in C, N, O, Na and Al abundances should exist among main-sequence, or near main-sequence turnoff, stars in clusters.  Observational examples supporting this picture are readily forthcoming. For instance, CN bimodality persists among giants well down onto the subgiant branch in several clusters (\\eg, Hesser, Hartwick, \\& McClure 1977\\nocite{hes77}, Smith \\& Norris 1982\\nocite{smi82}, Norris \\& Smith 1984\\nocite{nor84}, Cannon \\etal\\ 1998\\nocite{can98}, Smith 2002a,b\\nocite{smi02a,smi02b}). C and N abundance inhomogeneities have been observed among subgiant and even main-sequence stars in M5 and NGC 6752 (Suntzeff \\& Smith 1991\\nocite{sun91}, Cohen \\etal\\ 2002\\nocite{coh02}) and M71 (Cohen 1999\\nocite{coh99}). Moreover, variations involving other light elements among such relatively unevolved stars have recently come to light, \\eg, the correlation of CN band strengths and Na in main sequence stars of 47 Tuc (Briley \\etal\\ 1996\\nocite{bri96}), and the anticorrelation of O and Na as well as Al and Mg among main-sequence stars of NGC 6752 (Gratton \\etal\\ 2001\\nocite{gra01}). O and Na also appear to be anticorrelated among subgiants in M71 (Ramirez \\& Cohen 2002\\nocite{ram02}). Evidence of proton-capture synthesis of Ne into Na and Mg into Al has been found among a small sample of turnoff stars in M92 (King \\etal\\ 1998\\nocite{kin98}). Large ranges in Na in the case of M5 have been found to persist all the way from the giants to the main sequence (Ramirez \\& Cohen 2003\\nocite{ram03}), and a substantial range in C abundance has been found among main-sequence stars in M13 (Briley \\etal\\ 2003\\nocite{bri03}). Evidence favoring a primordial source for the variations among these light elements therefore seems well established.  But there is also evidence that the stars we presently observe reshuffle the abundances of many light elements as a result of deep mixing during the ascent of the red giant branch. Material of the stellar envelope is exposed to CN, ON, and in some cases NeNa and possibly even MgAl cycling as it is circulated through regions near or even in the CNO-burning shell of an evolving low-mass, metal-poor star. Among field metal-poor ([Fe/H]~$\\lesssim$~--1.0) halo stars, Gratton \\etal\\ (2000)\\nocite{gra00} found a progressive decrease in total C and \\carbiso\\ ratios, with corresponding enhancement of N, as stars evolve up the first giant branch. These observed abundance changes far exceed those predicted from conventional evolutionary models. Among globular cluster giants having [Fe/H]~$<$~--1.0, these changes in C and N are even more marked (Langer \\etal\\ 1986\\nocite{lan86}, Bellman \\etal\\ 2001\\nocite{bel01}, Trefzger \\etal\\ 1983\\nocite{tre83}, Briley \\etal\\ 1990\\nocite{bri90}, Briley \\etal\\ 2003\\nocite{bri03}) and contrary to the case of field stars, often involve also the ON cycle, so that many cluster giants exhibit depletion of O (Pilachowski 1988\\nocite{pil88}, Kraft \\etal\\ 1997\\nocite{kra97}, Sneden \\etal\\ 1997\\nocite{sne97} and references therein) and corresponding enhancement of N. At the same time,  in many clusters (e.g., NGC 288, NGC 362, M5, M10, M13, M15) there is an O versus Na anticorrelation, explicable as evidence of NeNa cycling in a region experiencing ON cycling (Langer \\& Hoffman 1995\\nocite{lan95}, Cavallo \\etal\\ 1996\\nocite{cav96}).  It is not known if the O versus Na anticorrelation is a result of primordial events or a result of deep mixing in the giants we observe. Even more controversial is the origin of the correlation between Al and Na observed among giants in several clusters (Shetrone 1996\\nocite{she96a}), since CNO-burning shell temperatures in low-mass, metal-poor red giants do not reach temperatures high enough to permit significant operation of the \\iso{24}{Mg}/\\iso{27}{Al} cycle (Cavallo \\etal\\ 1998\\nocite{cav98}, Powell 1999\\nocite{pow99}, Langer \\etal\\ 1997\\nocite{lan97}). Shetrone also discovered that the M13 giants of lowest oxygen abundance had anomalously high ratios of (\\iso{25}{Mg}+\\iso{26}{Mg})/\\iso{24}{Mg}, and a similar result has recently been found in a large sample of giants in NGC 6752 (Yong \\etal\\ 2003\\nocite{yon03}). Such overabundances of the lesser isotopes of Mg can readily be produced in the evolution of low-metallicity 3 to 6 solar mass AGB stars (Siess \\etal\\ 2002\\nocite{sie02}), and this strongly suggests that a pollution scenario is in operation. However, the possibility of enhancement of \\iso{27}{Al} through proton-captures on \\iso{25,26}{Mg}, operating in a region favorable to NeNa cycling, seems possible in the case of giants in M4 (Ivans \\etal\\ 1999\\nocite{iva99}). Could such a process also be in operation in M13 and in NGC 6752? Finally, as Cavallo \\& Nagar (2000)\\nocite{cav00} have most recently emphasized (see also Pilachowski \\etal\\ 1996\\nocite{pil96b}, Kraft \\etal\\ 1997\\nocite{kra97}), the mean Al and Na abundances of M13 giants sharply increase in the last 0.4 mag before the red giant tip. This is hard to understand unless it is the result of deep mixing which brings up the ashes of \\iso{22}{Ne} to \\iso{23}{Na} and \\iso{25}{Mg} or \\iso{26}{Mg} to \\iso{27}{Al} processing in the low-mass giants we presently see. Finally, in indirect support of at least a modest degree of deep mixing is the finding that although main-sequence stars in M13 exhibit a wide range in [C/Fe] ratios, many of which have [C/Fe] well below --0.3 (Briley \\etal\\ 2003\\nocite{bri03}), the mean [C/Fe] ratio is much lower among giants than dwarfs. It is also true that seven out of 30 M13 giants have [O/Fe]~$<$~--0.4 (Kraft \\etal\\ 1997\\nocite{kra97}), an additional result that is hard to understand on a primordial ``pollution'' scenario.  A growing consensus emerges that most clusters have stars whose abundance ratios reflect primordial variations, but that these variations are further modified by deep mixing (see, \\eg, Briley \\etal\\ 1994\\nocite{bri94}, Cavallo \\& Nagar 2000\\nocite{cav00}, Briley \\etal\\ 2003\\nocite{bri03}) as stars evolve up the giant branch. Of special interest is the contention that deep mixing may also play a role as a driver of the so-called ``second parameter'' effect (Sweigart 1997\\nocite{swe97}) in which clusters of closely the same metallicity may nevertheless have quite different horizontal branch morphologies. Critical pairs of clusters exhibiting these differences, for example, are M3 versus M13 and NGC 362 versus NGC 288. Many investigators have argued, based on the study of cluster color-magnitude diagrams, that age is the most likely driver of the second parameter (\\eg, Buonanno \\etal\\ 1994\\nocite{buo94}, Sarajedini \\etal\\ 1997\\nocite{sar97}, Rey \\etal\\ 2001\\nocite{rey01}). Others have argued that in at least some of these pairs, particularly M3/M13, the age difference inferred from the relative location of the cluster turnoff luminosity is too small to account for the difference in HB morphology (Stetson \\etal\\ 1996\\nocite{ste96}, Johnson \\& Bolte 1998\\nocite{joh98}).  Mixing, if it is deep enough to penetrate the hydrogen burning shell, will no doubt increase the He/H ratio in the convective envelopes of the red giant progenitors, so the resulting HB branch descendants are likely to be bluer than would be the case in the absence of deep mixing (Sweigart 1997\\nocite{swe97}, Kraft \\etal\\ 1998\\nocite{kra98}). Since deep mixing appears to be present to a significant degree among M13 giants close to the tip of the first giant branch (Kraft \\etal\\ 1997\\nocite{kra97}, Cavallo \\& Nagar 2000\\nocite{cav00}), we would expect M13 to possess a very blue HB, and indeed it does in comparison with M3. At the same time, evolved M3 giants show a much smaller degree of deep mixing than do comparable giants in M13 (Kraft \\etal\\ 1992\\nocite{kra92}, Cavallo \\& Nagar\\nocite{cav00}). However, owing to the fact that M3 is considerably farther from the Sun than is M13, the stars are fainter and less amenable to abundance analysis at high spectral resolution. Thus the published M3 high resolution sample is much smaller than is the case for M13 (10 versus 35). Would more evidence of deep mixing be found in M3, if a larger sample of giant star spectra could be obtained? In addition, Na abundances based on lower resolution spectra are available for 95 additional M13 giants, and of these Mg abundances have been derived for 74 (Pilachowski \\etal\\ 1996\\nocite{pil96b}).  In this paper, we examine more fully the [X/Fe] ratios in M3 based on high resolution analysis of 28 M3 giants. We confirm the tentative earlier result that the Na versus O, Mg versus Na, and Mg versus Al diagrams of M3 differ significantly from those of M13. We compare these diagrams for M3 and M13 giants with those of metal-poor field halo giants. We address also the abundances of $r$- and $s$-process elements Eu, La and Ba. We finally return to a discussion of these results in connection with the second parameter problem.     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  The major concern of this paper centers on the remarkable contrast of certain [X/Fe] ratios among giants in M3 in comparison with comparably-evolved giants in M13, on the one hand, and the halo field giants on the other. We divide the discussion into the familiar groups by nucleosynthetic origin, and discuss M3 stars along with their M13 and halo field surrogates within each group. We do not, in any one of these cases, have a ``complete'' sample of stars selected by \\MV0. Brighter than \\MV0~=~--1.5, we have high resolution analyses for 19 giants in M3 and 22 giants in M13. Total sample sizes in M3 and M13 are 28 and 32, respectively; in each cluster, our sample includes a few stars having \\MV0~$>$~--0.5. We therefore regard our two samples, comparable in size, as at least representative, and therefore anticipate that larger samples will not lead to unexpected surprises.      %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsection{Fe-peak and Heavier Elements} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Existing research on cluster abundances indicates that there is little star-to-star intracluster variation in [X/Fe] ratios for elements having atomic weights $A$~$\\geq$~$A$(Si). There are, of course, exceptions, most notably among $r$-process species in the very metal-poor cluster M15 (Sneden \\etal\\ 1997\\nocite{sne97}, Sneden \\etal\\ 1999\\nocite{sne00}). There are also cases in which $\\rm\\langle[X/Fe]\\rangle$ differs significantly from one cluster to another at a similar value of [Fe/H], as for example, the substantial differences in Si, Ba, La and Eu that exist when one compares M4 with M5 (Ivans \\etal\\ 2001\\nocite{iva01}). These variations emerge when stellar sample sizes are moderately large (N~$\\gtrsim$~20), and since our M3 and M13 samples exceed this number, we begin by exhibiting in Figure~\\ref{f3} the [X/Fe] ratios for Si and heavier elements versus \\teff\\ for the M3 stars of Table~4 and 6, and M13 stars of Table~7. Horizontal dashed lines refer to the mean values of [X/Fe] for M3 stars only.  Inspection of Figure~\\ref{f3} and examination of the accidental errors tabulated at the feet of Tables~4, 6, and 7 suggests that the spread in [X/Fe] is for the most part small and attributable to a paucity of lines (in the cases of, \\eg, \\ion{Sc}{2}, \\ion{Ba}{2}, \\ion{La}{2}, \\ion{Eu}{2}) and/or line weakness (EW~$<$~10~m\\AA, \\eg, for \\ion{V}{1}, \\ion{La}{2}, \\ion{Eu}{2}). Exceptions appear to be the very high abundances of La and Ba abundances in two different M13 giants, and the very low value of Eu in a single M3 giant (B4.4). None of these stars exhibit unusual abundances of other elements. The data of Figure~\\ref{f3} also show that the mean values of [X/Fe] in M3 and M13 differ no more than $\\sim$0.1~dex. In addition, the pattern also follows that of field halo stars with metallicities [Fe/H]~$\\sim$~--1.5 (Fulbright 2002\\nocite{ful00}), with the possible exception of Ba, which on average appears to be about 0.15~dex more Ba-rich in both clusters compared with the field. On the other hand, [Eu/Fe] in both clusters is quite close to the mean value for field stars.     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsection{Interrelated C, N, and O Abundances} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Carbon and nitrogen abundances were obtained for a substantial sample of CN-weak and CN-strong giants in both M3 and M13 in a pioneering study by Suntzeff (1981).\\nocite{sun81} A number of these stars were reanalyzed by Smith \\etal\\ (1996)\\nocite{smi96} who, coupling these results with what were then freshly determined values of [O/Fe] (Kraft \\etal\\ 1992,1993\\nocite{kra92,kra93}), examined whether the relative abundance of these species was consistent with CNO tricycle processing. Their cyanogen-band index S(3839) was shown to be anticorrelated with [C/Fe] and [O/Fe], and correlated with [N/Fe]; they concluded that the strength of the cyanogen band was driven largely by the abundance of N. Further modeling led them to conclude that [C/Fe] and [O/Fe] were correlated, whereas [N/Fe] remained anticorrelated with both [C/Fe] and [O/Fe] (their Figure~5). They also found that the total abundance C+N+O was, within the errors, essentially the same for all stars in both M3 and M13. These results support the view that the behavior of C, N, and O is compatible with CNO tricycle processing.  The new M3 oxygen abundances reported here add only a small amount of additional information, largely because the newly analyzed stars, lying as they do mostly near the center of M3, are not the ones studied previously for C and N abundances. Included among stars with newly derived [O/Fe]$_{\\rm II}$ ratios\\footnote{ Hereafter in the text we employ the notation [O/Fe]$_{\\rm II}$ to emphasize that this ratio has been formed between abundances derived from [\\ion{O}{1}] and \\ion{Fe}{2} lines. The notation [Na, Mg, or Al/Fe]$_{\\rm I}$ is used because these ratios have been formed between neutral-species abundances of these elements and \\ion{Fe}{1} abundances. Carbon abundances, taken without alteration from the literature, are written simply as [C/Fe]. The ``I'' and ``II'' subscripts are not written in the figures.} are I-21, IV-77, IV-101, and A. Although they have known S(3839) indices, only I-21, IV-77, and IV-101 have freshly derived carbon abundances (Pilachowski \\etal\\ (2003).\\nocite{pil03} In addition, oxygen abundances for M3 giants previously analyzed by Kraft \\etal\\ (1992,1993\\nocite{kra92,kra93}) have been somewhat revised in this paper. In Figure~\\ref{f4} we plot both the revised and newly derived abundances of [O/Fe]$_{\\rm II}$ against S(3839). The anticorrelation is morphologically the same as that found previously by Smith \\etal\\ (1996)\\nocite{smi96}. The reader should note that M3 and M13 stars lie along the same locus, but in contrast with M13, there are no M3 stars with very low oxygen abundance.  In Figure~\\ref{f5}, we plot also the oxygen versus carbon correlation, following Smith \\etal\\ (1996)\\nocite{smi96}, but with the revised oxygen abundances tabulated here, plus the new carbon abundances for the three M3 stars noted above. The results are in the expected direction: whenever oxygen is depleted via CNO processing, carbon also is depleted, and the abundances are correlated. There is a slight hint that at a given [O/Fe]$_{\\rm II}$, [C/Fe] is slightly smaller in M13 than in M3.  The conclusion that the relationship between C, N, and O abundances is consistent with CNO tricycle processing does not, reveal the site of the corresponding nucleosynthetic reshuffling. However, Suntzeff's (1981)\\nocite{sun81} early study showed that $\\rm\\langle[C/Fe]\\rangle$ declined with evolutionary state along the first giant branches of both M3 and M13. In the case of M3, this result was strengthened owing to a recent study by Smith \\etal\\ (2002a).\\nocite{smi02a} These authors found that $\\rm\\langle[C/Fe]\\rangle$ declines from $\\sim$--0.2 at $M_V$~$\\sim$~+0.6 to $\\sim$--1.1 at $M_V$~$\\sim$~--2.3 (the red giant tip). They also detected a slight offset in carbon abundances between CN-weak and CN-strong stars: $\\rm\\langle[C/Fe]\\rangle$ was --0.14 lower in the CN-weak group.  Plotting Suntzeff's M13 data on top of the M3 data (Figure~\\ref{f6}), we find a similar decline in $\\rm\\langle[C/Fe]\\rangle$ (albeit with larger scatter) over the same luminosity range. Additionally, the M13 giants achieve even lower carbon abundances than do their M3 counterparts (by $\\sim$0.2~dex). Even more striking is the result of Briley \\etal\\ (2003).\\nocite{bri03} They found a range of [C/Fe] values from --0.8 to 0.0 among main sequence turnoff stars in M13, but at the same time found that $\\rm\\langle[C/Fe]\\rangle$ declined from $\\sim$--0.4 at the main sequence to the value of $\\sim$--1.3 obtained earlier by Suntzeff for stars at the red giant tip.  Owing to stellar flux limitations, nothing is known of oxygen and nitrogen abundances among main sequence members of either M3 or M13. However, the spread in [C/Fe] among these stars in M13 suggests that CNO processing must have occurred in material that was incorporated in these stars before they achieved their present evolutionary state. Nevertheless, there is no escaping the fact that, on the average, $\\rm\\langle[C/Fe]\\rangle$ declines as M13 and M3 stars evolve above $M_V$~$\\sim$~+0.6, the luminosity level at which the hydrogen shell burns through the molecular weight discontinuity left behind from the deepest penetration of the outer convection zone (the so-called red giant ``bump'').  Among halo field giants, $\\rm\\langle[C/Fe]\\rangle$ also begins to decline at the same red giant ``bump'' luminosity (Gratton \\etal\\ 2000\\nocite{gra00}), but contrary to the situation in M3 and M13, the degree of carbon depletion at its most severe does not exceed --0.7, and averages about --0.5. This limit is strikingly similar to the highest [C/Fe] ratios seen among giants in M3 and M13 (see Figure~\\ref{f5}). It is therefore perhaps not surprising to find that, among halo field giants, Gratton \\etal\\ (2000) found no change in [O/Fe] with advancing evolutionary state, fixed at a value of $\\rm\\langle[O/Fe]\\rangle$~=~+0.34.  In Figure~\\ref{f7} we illustrate [O/Fe]$_{\\rm II}$ as a function of \\MV0\\ in M3 (Tables~4 and 6) and M13 (Table~7). Our samples contain no stars with \\MV0~$\\gtrsim$~+0.1, so we cannot test the behavior of [O/Fe]$_{\\rm II}$ at luminosities below that of the red giant bump. We can describe the distribution of the stars brighter than \\MV0~=~+0.1, by dividing them into two groups near log~$g$~=~0.8 (\\MV0~$\\sim$~--1.7), and by introducing three levels of oxygen ``deficiency''. Typically among halo field giants near [Fe/H]~=~--1.5, [O/Fe]$_{\\rm II}$ lies in the range 0.0 to +0.5 (average near +0.3), and we speak of these as ``oxygen-normal''. Those in the clusters with 0.0~$>$~[O/Fe]$_{\\rm II}$~$>$~--0.4, we call ``oxygen-poor''. Finally we refer to those in the range --0.4~$>$~[O/Fe]$_{\\rm II}$~$>$~--1.2 as ``super oxygen-poor''.  Of the 11 M3 stars having \\MV0~$>$~--1.7, nine (82\\%) are O-normal and two (18\\%) are O-poor; there are no super-O-poor stars in this range. The absence of super-O-poor stars is also a characteristic of the 16 brighter M3 stars. The split among normal and O-poor stars in this group is 72\\% to 28\\%, respectively, roughly similar to the lower luminosity group. (We arbitrarily assigned one star, B.4.2, having [O/Fe]$_{\\rm II}$~=~0.00 to the ``normal'' group.) The situation in M13 is not so different for the 11 stars in the fainter group: 73\\% are O-normal, 18\\% are O-poor, and one star (I-12) is super-O-poor. But among the stars of the brighter group, the situation is dramatically different. Of the 23 brighter stars, only seven (30\\%) are O-normal, nine (39\\%) are O-poor, and seven (30\\%) are super-O-poor. (Again we arbitrarily assign one star, L261, having [O/Fe]$_{\\rm II}$~=~0.00 to the ``normal'' group). Six of the nine super-O-poor stars are brighter than \\MV0~=~--2.3, and are therefore essentially at the tip of the giant branch.  We summarize the results of this section as follows:  \\begin{enumerate}  \\item Metal-poor halo stars, whether in the field or in such clusters as M3 and M13, all show a decline in $\\rm\\langle[C/Fe]\\rangle$ starting when stars evolve through the red giant bump. This is superimposed on top of a pre-existing spread in [C/Fe] exhibited among main sequence and turnoff stars at least in M13, and one would not be surprised if this were true also in M3 (although this is not yet observationally demonstrated). A spread in [C/Fe] is found also among field halo subdwarfs (Gratton \\etal\\ 2000\\nocite{gra00}), but it appears to be smaller than that found among main sequence and turnoff stars in M13.  \\item The decline is driven by stellar evolution and presumably results from the deep mixing of envelope material through the region just ahead of the H-burning shell where conversion of C to N takes place. The deficiency of $\\rm\\langle[C/Fe]\\rangle$ among halo giants is typically --0.5 dex, but is larger in clusters, reaching --0.7 to --0.9~dex in M3 and M13.  \\item We divide (somewhat arbitrarily) our sample of giants into three groups depending on the degree of oxygen deficiency: O-normal stars have [O/Fe]$_{\\rm II}$~$>$~0.0, O-poor stars lie in the range 0.0~$>$~[O/Fe]$_{\\rm II}$~$>$~--0.4, and super-O-poor stars have [O/Fe]$_{\\rm II}$~$<$--0.4. Among halo field stars, there are no O-poor or super-O-poor objects, and giants appear to have on the average the same oxygen abundances as main sequence subdwarfs (see also Gratton \\etal\\ 2000\\nocite{gra00}). In M3 there are no super-O-poor stars, and of the 28 giants in the sample, about one-third are O-poor. If this is attributable to O to N conversion in the material making up the envelope of some giants in M3 (either pre-existing processing or processing within the stars themselves), we conclude that both the degree and the frequency of the effect is rather mild. The situation in M13 is dramatically different: fewer than half the giants are O-normal, and of the remaining, 30\\% are super-O-poor. Most of these lie near the red giant tip. Indeed of the giants brighter than \\MV0~=~--2.3, 60\\% are super O-poor. It is hard to escape the conclusion that the majority of the brightest M13 giants mix envelope material deep enough into the H-burning shell that significant processing of O to N occurs within the stars themselves.  \\end{enumerate}     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsection{Interrelationships Among the Light Elements: O, Na, Mg, Al} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsubsection{The Na versus O Anticorrelation} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Giants in most low metallicity globular clusters exhibit an anticorrelation of [Na/Fe] and [O/Fe], and M3 is no exception. In Figure~\\ref{f8}, we compare our samples of giants in M3 (Tables~4 and 6) with those of M13 (Table~7), as well as giants of the halo field (Shetrone 1996\\nocite{she96a}, Gratton \\etal\\ 2000), the last-named group consisting of stars with [Fe/H]~$<$~--1.0. Small vertical and horizontal lines drawn through the points denote stars known to be CN-strong or CN-weak, respectively (Smith 2002a,b\\nocite{smi02a,smi02b}); most field halo giants are CN-weak (Hesser \\etal\\ 1977\\nocite{hes77}; Langer \\etal\\ 1992\\nocite{lan92}). In contrast with earlier plots of this kind (\\eg, Kraft 1994\\nocite{kra94}, his Figure~8), we have included the nLTE Na corrections recommended by Gratton \\etal\\ (2002), which have the effect of enhancing values of [Na/Fe] with increasing luminosity. We also normalize [O/Fe] to [Fe/H] based on \\ion{Fe}{2} and [Na/Fe] to [Fe/H] based on \\ion{Fe}{1}, in accordance with precepts argued by Kraft \\& Ivans (2003)\\nocite{kra03}.  Inspection of the three panels leads us to the following conclusions.  \\begin{enumerate}  \\item Halo field giants show no sign of a Na versus O anticorrelation. There is a considerable range of values of [Na/Fe]$_{\\rm I}$ (--0.5 to +0.3), and a smaller range in [O/Fe]$_{\\rm II}$ (0.0 to +0.6), a result noted earlier by Kraft \\etal\\ (1993), Hanson \\etal\\ (1998), and Gratton \\etal\\ (2000). Mean values of [O/Fe]$_{\\rm II}$~= +0.3 and [Na/Fe]$_{\\rm I}$~=~--0.1 are generally regarded as numbers characteristic of Type~II supernovae ejecta.  \\item M13 giants exhibit a range of [Na/Fe]$_{\\rm I}$ (a factor of $\\sim$8) and a range of [O/Fe]$_{\\rm II}$ (a factor of $\\sim$25) that is larger than that seen in any other cluster (although the range in $\\omega$~Cen may be comparable, see Norris \\& DaCosta 1995\\nocite{nor95}), and the anticorrelation is decidedly marked. Na and O abundances in the most Na-poor and O-rich giants of M13 are comparable to what is seen in halo field giants. As expected, these stars are CN-weak; CN-strong stars are those with deficiencies of O. In the present sample, there is a tendency for giants that are less advanced in evolutionary state (\\logg~$>$ 0.8) to favor higher [O/Fe]$_{\\rm II}$ and lower [Na/Fe]$_{\\rm I}$ than are found in stars nearer the red giant tip (\\logg~$<$ 0.8).  \\item M3 giants exhibit an anticorrelation that is ``intermediate'' between the halo field and M13 stars. M3 stars follow quite closely the anticorrelation defined by the M13 stars, but ``cutoff'' near [O/Fe]$_{\\rm II}$~= --0.2, [Na/Fe]$_{\\rm I}$~= +0.5. There are no super-O-poor stars in M3, as already noted. It is also the case that the distribution of stars in the diagram is essentially independent of evolutionary state, in contrast to the situation in M13. The shape and boundaries of the M3 Na versus O distribution are, in fact, quite similar to those found recently in NGC~6752 (Yong \\etal\\ 2003\\nocite{yon03}), a point to which we will return in Section \\S 4.1.1.  \\end{enumerate}      %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsubsection{[Mg/Fe]$_{\\rm I}$ versus [O/Fe]$_{\\rm II}$} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Previous work (Figure~9 of Kraft \\etal\\ 1997\\nocite{kra97}) indicated that M13 giants with lower than average [O/Fe] also had lower than average [Mg/Fe]. In the present Figure~\\ref{f9}, we plot [Mg/Fe]$_{\\rm I}$ versus [O/Fe]$_{\\rm II}$ for the sample of M13 giants of that study observed with either the Keck~I HIRES spectrograph or the post-1995 version of the Lick Hamilton spectrograph. As before, we find $\\rm\\langle[Mg/Fe]_I\\rangle$~=~+0.22 ($\\sigma$~=~$\\pm$~0.09) for giants with [O/Fe]$_{\\rm II}$~$>$~--0.2 whereas $\\rm\\langle[Mg/Fe]_I\\rangle$~=~--0.01 ($\\sigma$~=~0.07) for giants with [O/Fe]$_{\\rm II}$~$>$~--0.2. In Figure~\\ref{f9}, we superimpose the O and Mg abundances for 22 of the 23 giants in our M3 sample (Table~4). One star has [O/Fe]$_{\\rm II}$~$>$~--0.2, and it does not have a low Mg abundance; on the other hand, the one M3 giant (B4.5) having a very low, indeed negative, value of [Mg/Fe]$_{\\rm I}$ does not have a derivable oxygen abundance (the S/N of the spectrum is too low). However, it is clearly the case that our M3 sample contains few, if any, stars with the low Mg abundances characteristic of many of the brightest giants in M13. We address later implications of the change in the abundance ratio of Mg isotopes with [O/Fe]$_{\\rm II}$ among giants in M13 (Shetrone 1996b).     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsubsection{[Al/Fe]$_{\\rm I}$ versus [Mg/Fe]$_{\\rm I}$} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  The expectation that Al may arise owing to proton captures on the various isotopes of Mg (\\eg, Kraft \\etal\\ 1997\\nocite{kra97}, Shetrone 1996b\\nocite{she96b}, Ivans \\etal\\ 1999\\nocite{iva99}) prompts an investigation of [Al/Fe]$_{\\rm I}$ versus [Mg/Fe]$_{\\rm I}$ which is shown in Figure~\\ref{f10}. All data are taken from Tables~4 and~7. We divide the stars at \\logg~=~+0.8 into two groups by evolutionary state, as we did previously in the discussion of the Na versus O anticorrelation (\\S 4.2). Although the sample sizes are small, the distribution suggests that in each cluster there are two groups with distinctly differing Al abundances. Whereas the higher and lower luminosity giants are distributed among the two groups about equally in the case of M3, the distribution appears to be quite different in the case of M13, with most of the brighter giants lying in the group with high Al abundances. The highest Al abundances in M13 also appear to be about a factor of two higher than the highest Al abundances in M3.  Generally speaking, the observational accuracy of our Al abundances, which are based on two lines of moderate EW and excitation potential, is likely higher than that of our Mg abundances, which are based on one or two rather strong lines of rather high excitation potential. This may in part account for the rather displeasing result that the M3 giants of higher luminosity seems to have [Mg/Fe]$_{\\rm I}$ ratios that are on average about 0.1~dex higher than the stars of lower luminosity (omitting the anomalous star (B4.5) with a negative [Mg/Fe]$_{\\rm I}$ ratio). The result suggests the presence of a systematic error in our adopted \\teff\\ or \\logg\\ scales, as a function of luminosity. If an arbitrary adjustment to correct the effect in M3 were also applied to the sample of stars in M13, it would strengthen the conclusion that the more luminous M13 giants have a lower Mg abundance than the less luminous stars. However, the same conclusion would be reached even if such a correction were not applied to the M13 sample.     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsubsection{[Al/Fe]$_{\\rm I}$ versus [Na/Fe]$_{\\rm I}$} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   As expected, these two ``odd-elements'' are correlated in both M3 and M13, but in each case, there appear to be two well-separated groups (Figure~\\ref{f11}). M3 and M13 giants satisfy essentially the same correlation, but the M13 giants extend toward higher values of both Na and Al by a factor of about two. The population of M13 stars having the highest values of Na and Al abundances is dominated by stars in the most advanced evolutionary state (\\logg~$<$~0.8). All halo field giants considered here lie among the group of stars with lowest Na and Al abundances.     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsubsection{``Odd'' versus ``Even'': [Na/Fe]$_{\\rm I}$ versus [Mg/Fe]$_{\\rm I}$} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Hanson \\etal\\ (1998)\\nocite{han98}, using the the observational results of Pilachowski \\etal\\ (1996)\\nocite{pil96a}, noted that there exists a Na versus Mg correlation among halo field giants having [Fe/H]~$<$~--1.0. In Figure~\\ref{f12} (left-hand panel), we plot [Na/Fe]$_{\\rm I}$ versus [Mg/Fe]$_{\\rm I}$ in these giants, the data having been ``corrected'' for the nLTE effect in Na. In the right hand panel of Figure~\\ref{f12}, we plot the results of Fulbright (2002)\\nocite{ful02} for halo field subdwarfs, the straight line taken from a fit (by eye) of the correlation for the halo giants. The Mg abundances of Fulbright have been modified, first to eliminate all Mg lines except $\\lambda$5528 and/or $\\lambda$5711 (the lines used exclusively for studies of the giants), and second to allow for differences in the adopted log~$gf$ values assumed by Fulbright for the 5528/5711 pair. The two panels should therefore be on the same system of Mg and Na abundances.  Inspection of the two panels of Figure~\\ref{f12} suggests that, (1) there is no change in the correlation whether the samples are drawn from the main sequence or from the giants, and (2) there is no change in the distribution whether one deals with giants having gravities greater or less than \\logg~=~0.8. In other words, there is no dependence of the correlation on evolutionary state among metal-poor halo field stars.  Sodium and magnesium are near neighbors in the periodic table and are believed to arise during carbon and neon burning in Type~II SNe, giving rise to the so-called ``odd-even'' effect. The correlation suggests that there is a range in the [Na/Mg] ratio at a given metallicity, since all stars regardless of metallicity below [Fe/H]~=~--1.0 seem to satisfy the same correlation. Thus along the correlation, the [Na/Mg] ratio varies from --0.6 to -0.1, with a mean near --0.4; this is in quite good agreement with Arnett's (1971) hydrostatic and explosive carbon-burning models, in which [Na/Mg] ranges from --0.8 to --0.2 for stars having [Fe/H]~$<$~--1.0. The more recent models reported by Timmes \\etal\\ (1995)\\nocite{tim95} suggest a lower mean [Na/Mg] ratio, near --0.7, for such low-metallicity stars. We note that the stars with [Na/Mg] ratios closest to the theoretical value of Timmes \\etal\\ are the one shown by Fulbright (2002)\\nocite{ful02} to favor Galactic orbits having relatively high values of orbital energy and (absolute) angular momentum.  Corresponding plots of [Na/Fe]$_{\\rm I}$ versus [Mg/Fe]$_{\\rm I}$ for M3 and M13 are shown respectively in the left and right hand panels of Figure~\\ref{f13}, the data being taken from Tables~4, 6, and 7. The straight line is the same as that in the two panels of Figure~\\ref{f12}. Again we see the somewhat unpleasing result that in M3 the stars of low luminosity have systematically lower Mg abundances than the stars of higher luminosity (\\logg~$<$~0.8), but otherwise the two kinds of giants are essentially intermixed. The giants of M13 present a rather different picture. Here the more luminous giants prefer a location well off the halo-defined line, in the direction of low Mg, high Na, whereas the lower luminosity giants, except for three exceptions (out of a sample of 10 stars) prefer a location near the normal halo line.  The effect is illustrated more dramatically in Figure~\\ref{f14}, in which [Na/Fe]$_{\\rm I}$ is plotted against [Mg/Fe]$_{\\rm I}$ for M13, this time using the very much larger sample of M13 giants studied by Pilachowski \\etal\\ (1996b)\\nocite{pil96b}. The Mg abundances were adjusted slightly to conform to the system of Mg abundances adopted here (by plotting [Mg/Fe]$_{\\rm I}$ values of Table~7 against the [Mg/Fe] values of Pilachowski \\etal\\ for common M13 giants), and the Na abundances adjusted for nLTE as explained earlier. Again we see that M13 giants of high luminosity are confined to the upper left-hand part of the diagram, although they do not have that domain to the exclusion of stars having \\logg~$>$~0.8.      %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsection{Implications of the Interrelationships among O, Na, Mg, and Al} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsubsection{Deep Mixing versus Primordial (or Pollution) Scenarios} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  In all of the interrelationships illustrated in \\S 4.3, M13 giants inhabit a more ``extreme'' domain than do M3 giants, when the stars are compared at essentially the same evolutionary state. As discussed in \\S 4.3.1, most M13 giants near the red giant tip have, in comparison with corresponding M3 giants, very low O and very high Na abundances, suggesting that first ascent M13 giants immediately prior to the He core flash undergo a period of deep mixing not experienced by similar giants in M3. But a recent study of giants in NGC 6752 (Yong \\etal\\ 2003\\nocite{yon03}), a cluster having [Fe/H] essentially the same as that of M3 and M13, casts doubt on such an interpretation. Thus, in what follows, we describe similarities and differences among giants in NGC 6752, M13 and M3.  From a sample of 20 (mostly) first ascent NGC~6752 giants with --1.0~$\\leq$~$M_{\\rm V}$~$\\leq$~--2.5, Yong \\etal\\ (2003)\\nocite{yon03} discovered that the abundances of the minor isotopes \\iso{25}{Mg} and \\iso{26}{Mg}, relative to the dominant \\iso{24}{Mg}, were ofter higher than in the solar system (\\iso{24}{Mg}:\\iso{25}{Mg}:\\iso{26}{Mg}~= 79:10:11, Lodders 2003\\nocite{lod03} and references therein) and normal stars. The minor Mg isotopes were also very much more abundant than the values predicted for the ejecta of Type~II SNe (Timmes \\etal\\ (1995)\\nocite{tim95}. For NGC~6752, Yong \\etal\\ found a range of abundance ratios running from 63:08:30 to 84:08:08, with one NGC~6752 giant (star 702) having an unusually low ratio of 53:09:39. Hydrogen-burning shell temperatures in first ascent NGC~6752 giants are too low to allow efficient operation of the reaction \\iso{24}{Mg}(p,$\\gamma$)\\iso{25}{Al} (Powell \\etal\\ 1999\\nocite{pow99}), and thus the chain of proton captures leading to the destruction of \\iso{24}{Mg} and the production of \\iso{25}{Mg}, \\iso{26}{Mg} and \\iso{27}{Al}. Therefore Yong \\etal\\ concluded that deep mixing was ruled out as a source of the anomalous Mg abundance ratios as well as the enhancement of \\iso{27}{Al}. They proposed instead a scenario in which the abundance ratios in NGC~6752 giants actually reflect those found in the ejecta of 3-6~\\Msun\\ cluster giants that had already evolved: either the low mass giants we presently see formed from such secularly ejected material, or their envelopes were seriously polluted or even ablated and replaced by this material. A scenario like this is supported both observationally and theoretically. On the one hand, most of the abundance anomalies in Na, Mg and Al found among NGC~6752 giants, that had earlier been attributed to deep mixing, were found to exist already in their main sequence progenitors (Gratton \\etal\\ 2001\\nocite{gra01}), which points to a pre-existing origin. In addition, models which predict the abundance ratios in the ejected envelope material of intermediate-mass low-metallicity AGB stars (Siess \\etal\\ 2002\\nocite{sie02}, Karakas \\& Lattanzio 2003\\nocite{kar03}) show that \\iso{25}{Mg} and \\iso{26}{Mg} are overproduced relative to \\iso{24}{Mg}. For example, Karakas \\& Lattanzio found that in the envelope of a 4~\\Msun\\ AGB star having Z~=~0.004, there can develop a situation in which the abundances of \\iso{24}{Mg}, \\iso{25}{Mg} and \\iso{26}{Mg} are approximately equal; in a 6~\\Msun, Z~=~0.004 AGB star, \\iso{24}{Mg} virtually disappears relative to the other isotopes. At the same time, hot bottom CNO cycling in the same stars can lead to an overabundance of N at the expense of O and excess production of \\iso{23}{Na}. Qualitatively these are just the sorts of anomalies found in the spectra of present-day NGC 6752, M3 and M13 giants.  So far, we have discussed recent results found in giants of NGC 6752. What bearing does this have on M3 and/or M13? Our M3 spectra do not have sufficient S/N near the MgH bands at $\\lambda$5135 to permit discussion of the Mg isotopes. But a previous study of the MgH bands in six M13 giants (Shetrone 1996b\\nocite{she96b}) permits us to compare ratios of (\\iso{25}{Mg} + \\iso{26}{Mg})/\\iso{24}{Mg} in M13 with NGC 6752. Shetrone's spectra, of somewhat lower resolution than those of Yong \\etal (2003)\\nocite{yon03}, did not permit the separation of lines of \\iso{25}{Mg}H from \\iso{26}{Mg}H, so we resort to discussing the quantity (\\iso{25}{Mg}+\\iso{26}{Mg})/\\iso{24}{Mg}). In Figure~\\ref{f15} we plot [O/Fe]$_{\\rm II}$ versus the quantity (\\iso{25}{Mg}+\\iso{26}{Mg})/\\iso{24}{Mg} for 19 of the 20 giants studied by Yong \\etal, plus the six M13 giants of Shetrone. The missing NGC 6752 star is No. 702, the one with the lowest abundance of \\iso{24}{Mg}. Yong \\etal\\ were unable to measure [O/Fe]$_{\\rm II}$ for this star, owing to lower than average S/N for this spectrum.  Inspection of Figure~\\ref{f15} shows that there is an anticorrelation between oxygen abundance and the overabundance of the rare isotopes of Mg relative to \\iso{24}{Mg}. Further, the plot indicates that whatever is the nucleosynthetic process that depletes oxygen and enhances the rarer isotopes of Mg, NGC~6752 and M13 share in that process. Moreover, if a primordial ablation or pollution scenario is responsible for what is seen in NGC~6752, the same scenario must have operated to any even greater extent in M13. Finally, the existence of this anticorrelation means that the extreme oxygen depletion seen in M13 giants near the red giant tip is probably not a result of deep mixing.  Parallel support for this conclusion is found Figure~\\ref{f16}, which is a replot of Figure~\\ref{f9}, this time excluding the giants of M3, but noting the Mg abundance ratios of the six stars studied by Shetrone (1996a,b)\\nocite{she96a,she96b}. (Because Shetrone could not resolve \\iso{25}{Mg} from \\iso{26}{Mg}, he simply assigned them equal weight in the blended spectra.) The stars of lowest oxygen abundance also have the lowest abundance of \\iso{24}{Mg} as well as the highest abundances of the rarer isotopes, as already noted by Shetrone. If the current nuclear reaction rates are correct, deep mixing again cannot be responsible for the depletion of \\iso{24}{Mg} found among most of the M13 giants of lowest oxygen abundance and highest luminosity.  What about M3? Even though we cannot report on the Mg isotopic ratios in this cluster, we can compare some key abundance-ratio plots in M3 with those of NGC 6752. The comparison is somewhat clouded by the possibility of systematic differences between investigators. For example, Yong \\etal\\ (2003)\\nocite{yon03} employ a completely different set of \\ion{Fe}{1} and \\ion{Fe}{2} lines compared with the set adopted in our analyses of M3 and M13. NGC~6752 stars have not been investigated by our group; as Yong \\etal\\ note, previous investigations yielded [Fe/H] ranging from --1.42 to --1.62. It seems best then, in comparing abundance ratio diagrams of NGC 6752 with those of M3 and M13, to concentrate on the morphology and recognize the possibility of zero-point differences of 0.1 or 0.2~dex in values of [X/Fe].  In the two panels of Figure~\\ref{f17} we compare the [Na/Fe]$_{\\rm I,nLTE}$ versus [O/Fe]$_{\\rm II}$ diagrams of M3 and NGC~6752. The two diagrams are morphologically quite similar, and both are remarkably different from the corresponding M13 diagram. In Figure~\\ref{f18}, we compare the [Na/Fe]$_{\\rm I,nLTE}$ versus [Mg/Fe]$_{\\rm I}$ diagrams of M3 and NGC~6752, again noting that both are morphologically dissimilar to the corresponding diagram for M13 (Figure~\\ref{f14}). M3 and NGC 6752 do not match perfectly, as one can easily see: M3 giants seem to have a larger scatter of Mg abundances. The main point is that neither cluster has the large population of giants lying far from the field giant relationship in the low Mg versus high Na regime, such as one finds in M13.     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\subsubsection{HB Morphology} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  The main difference in the morphology of the color-magnitude diagrams for these three clusters is the distribution of stars along the HB, a phenomenon usually referred to as ``the second parameter problem''. The Lee \\etal\\ (1994)\\nocite{lee94} second parameter index is nearly the same for M13 vs NGC 6752 (0.97 versus 1.00), which indicates that the HB is very blue, whereas the value of the index for M3 is 0.08, which corresponds to a uniformly populated HB. Although age is rather widely believed to be the factor responsible for the second parameter (\\eg, Sarajedini \\etal\\ 1997\\nocite{sar97}, Rey \\etal\\ 2001\\nocite{rey01}), other possibilities have been invoked. For example, if deep mixing were to significantly penetrate into the hydrogen shell burning region, evolution to the red giant tip would be slowed and the He/H ratio in the envelopes of evolving red giants would be increased (Sweigart 1997\\nocite{swe97}). The HB descendants would then be bluer than their unmixed counterparts. Another possibility would involve an ``ab initio'', or ``primordial'' difference in He abundance between clusters (Johnson \\& Bolte 1998\\nocite{joh98}), perhaps brought about if the low mass giants we presently see formed from the He-enriched ejecta of 3-6~\\Msun\\ cluster giants (Ventura \\etal\\ 2002\\nocite{ven02}). Qualitatively at least, an increase in the He/H ratio would favor the development of a blue HB, all other factors being equal.  But if deep mixing were responsible for the blue HB's of M13 and NGC 6752, and if seriously depleted O were the most significant ``symptom'' of the deep mixing phenomenon, then one might expect NGC~6752 to exhibit a Na vs O relationship as ``advanced'' as that of M13. This is clearly not the case. As we have seen, in both the Na versus O and Na versus Mg plots, NGC~6752 is much more like M3, a cluster with an ``intermediate'' HB morphology. However, these remarks apply only to the case in which an increase in the He/H ratio is a result of deep mixing. It does not rule out a picture in which NGC~6752 and M13 share the property of a ``primordial'' overabundance of He, relative to M3. It would clearly be of interest to explore the Mg isotopic ratios in M3 as a test (Ventura \\etal\\ 2002\\nocite{ven02}) of the He/H ratio among M3 giants close to the red giant tip.     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0311/astro-ph0311242_arXiv.txt": {
        "abstract": "We study the possibility of magnetic mass detection by the gravitational microlensing technique. Recently the theoretical effect of magnetic mass in the NUT space on the microlensing light curve has been studied. It was shown that in the low photometric signal to noise and sampling rate of MACHO experiment light curves, no signature of NUT factor has been found. In order to increase the sensitivity of magnetic mass detection, we propose a systematic search for microlensing events, using the currently running alert systems and complementary telescopes for monitoring the Large Magellanic Clouds stars. In this strategy of observation, we obtain the magnetic mass detection efficiency and also the lowest observable limit of the NUT factor. This method of survey for gravitational microlensing detection can also be used as a tool for searching other exotic space-times. ",
        "introduction": "The gravitational microlensing method for detecting MAssive Compact Halo Objects (MACHOs) in the Milky Way halo has been proposed by Paczy\\'nski (1986). Many groups have contributed to this experiment and have detected hundreds of microlensing candidates in the direction of the galactic bulge, spiral arms and Large and Small Magellanic Clouds (LMC \\& SMC). Due to the low probability of the microlensing detection, less than 20 events have been observed by the EROS and MACHO groups in the direction of the Magellanic clouds (Lasserre et al. 2000; Alcock et al. 2000). The low statistics not only causes ambiguities in identifying the galactic model of the Milky Way, but also in some cases the microlensing results are at variance with the results of other observations (Gates and Gyuk 2001).\\\\ Comparing LMC microlensing events with the theoretical galactic models can give us the mean mass of MACHOs and the fraction of halo mass in the form of MACHOs. In the case that we use a Dirac Delta mass function for the MACHOs, the mass of MACHOs in standard halo model is obtained about $\\sim 0.5 M_{\\sun}$. This means that the initial mass function of MACHO progenitors in the galactic halo should be different from that of the disk, because we neither see the low mass stars which should still exist nor heavier stars that would have exploded in the form of supernova (Adams \\& Laughlin 1996; Chabrier, Segretain \\& Mera 1996). Another contradiction is that if there were as many white dwarfs in the halo, as suggested by the microlensing experiments, they would increase the abundance of heavy metals via Type I Supernova explosions (Canal., Isern $\\&$ Ruiz-Lapuente 1997). Also, recently Green \\& Jedamzik (2002) and Rahvar (2004) showed that the observed distribution of duration of microlensing events is not compatible with what is expected from the standard and non-standard halo models. They showed that the observed distribution is significantly narrower compared to what is expected from the galactic models.\\\\ The mentioned problems can be a motivation for establishing the next generation of the microlensing experiments. The new surveys will have the potential to increase the number of microlensing candidates and reduce the ambiguities due to Poisson statistics. The other improvements of the new surveys can be the higher sampling rates and the higher precision photometry of the light curves. More precise light curves will enable us to distinguish the deviations between the standard and non-standard light curves due to parallax or source finite-size effects (Rahvar et al. 2003). In the so-called non-standard microlensing candidates the degeneracy partially can be broken between the lens parameters, such as the distance and the mass of a lens. A better determination of the distance and the mass distributions of the lenses can help us to better identify the Milky Way halo model (Evans 1994).\\\\ Although the mentioned effects are in the background of a Schwarzschild space, it is also be possible that a MACHO which plays the role of the lens, resides in an exotic space-time such as the Kerr or the NUT space. Deviation of the space-time from the Schwarzschild metric causes deviation of the microlensing light curve from the standard one. Thus, studying the microlensing light curves not only can be used to determine the dark matter in the form of MACHOs but also as a unique tool to explore the other exotic space-times as well.\\\\ In the paper by Nouri-Zonoz and Lynden-Bell (1997) the gravitational lensing effect on the light rays passing by a NUT hole has been considered, using the fact that all the geodesics in the NUT space, including the null ones, lie on cones. The extension of this work to the microlensing light curve in the NUT space has been studied by Rahvar and Nouri-Zonoz (2003) and the possible existence of magnetic mass on the light curves of the MACHO group microlensing candidates has been tested. According to the analysis of the light curves, no magnetic mass effect has been found. Although the result showed that the effect of the NUT factor is almost negligible, one can not rule out the existence of NUT charge on that basis. The next generation of microlensing experiments may prove the (non-) existence of magnetic mass through a more careful study of the microlensing light curves.\\\\ Here in this work we simulate the microlensing light curves in the NUT metric according to a strategy for the next generation of microlensing surveys. The aim of this work is to obtain the observational efficiency for the magnetic mass detection and to find the lowest limit for the NUT charge that can be observed. Large Magellanic Cloud (LMC) stars are chosen as the target stars for monitoring. The advantage of using LMC stars as compared to the spiral arms and the galactic bulge stars is the lower contamination by blending and source finite-size effects, which can affect the NUT light curves. The other advantage of LMC monitoring is that it enables us to increase the microlensing statistics to put a better limit on the mass of the lenses and the mass fraction of the galactic halo in form of MACHOs. \\\\ The outline of this paper is as follows. In Section 2, we give a brief account on the microlensing light curve in the NUT metric and compare it with the Schwarzschild case. In Section 3, we introduce the observational strategy and perform a Monte-Carlo simulation to generate the microlensing light curves. Section 4 contains the fitting process to the simulated light curves to obtain the observational efficiency of the magnetic mass detection. The results are discussed in Section 5. ",
        "conclusions": "In this work we proposed a new strategy for microlensing observation that not only can be used for searching MACHOs of the galactic halo through observing the LMC stars, but also it can be a useful tool to explore exotic space-times around compact objects such as the NUT metric. As a result of our Monte Carlo simulation, we obtained the detection efficiency for magnetic mass. The minimum value for $R$ that can be observed by this method is about $0.1$. In order to evaluate the amount of detectable magnetic mass {\\it l}, we use the relation between the magnetic mass and $R$ (Rahvar and Nouri-Zonoz 2003): \\begin{equation} R = c\\sqrt{\\frac{{\\it l}}{2GM}}. \\label{rl} \\end{equation} EROS and MACHO experiments results propose that the mean value of the mass of MACHOs is about $0.5M_{\\odot}$ (Alcock et al. 2000; Lasserre et al. 2000). It is worth to mention that this result is obtained in the standard halo model where the mean mass of MACHOs depends on the model that is used for the Milky Way. Assuming standard model for the Milky Way halo, according to equation (\\ref{rl}) the minimum observable magnetic mass ${\\it l}$ is evaluated to be about $14$ m. Non-existence of magnetic mass signal in the microlensing light curves can also put an upper limit for the value {\\it l} $<$ 14 m in the MACHOs of the Milky Way."
    },
    "0311/astro-ph0311074_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1} A key goal in understanding the connection between nuclear activity (accretion onto a super massive black hole) and host galaxy structure (e.g. morphology, stellar populations, etc.) is to find out about the physical conditions of quasars and the galaxy hosts they reside in.  Because of its large angular size and its brightness at the redshift of $z\\approx 0.37$ \\cite{boroson84}, 3C 48 is an ideal target for high-resolution imaging and spectroscopy with large telescopes. The near-infrared (NIR) is sensitive for the mass dominating stellar populations and at the same time it is less affected by extinction than the optical. The radio-loud galaxy shows an excess far-infrared (FIR) emission, $L_\\mathrm{FIR} =5\\times10^{12} L_\\odot$ \\cite{neug85}, and a one-sided radio jet extending about $0.5''$ to the north, fanning out to the east to about $1''$ \\cite{wilk91}. The quasar is rich in molecular gas, $M_{\\mathrm{H}_2}\\sim  2.7\\times 10^{10} M_\\odot$, as CO(1-0) observations show \\cite{scoville93, wink97}. There is also morphological evidence for a recent merger event. Such are a possible double nucleus and a tidal tail extending several arcseconds to the northwest \\cite{stockton91, chatzi99, canalizo2000}. Yet it is still unclear, whether the brightness peak about $1''$ northeast (NE) of the QSO is really the nucleus of the companion galaxy, or a region of interaction of the radio jet with the interstellar medium of the host galaxy inducing star formation \\cite{canalizo2000}. ",
        "conclusions": ""
    },
    "0311/astro-ph0311505_arXiv.txt": {
        "abstract": "We present a new approach to gauge-invariant cosmological perturbations at second order, which is also covariant.  We examine two cases in particular for a dust Friedman-Lema\\^\\i tre-Robertson-Walker model of any curvature: we investigate gravity waves generated from clustering matter, that is, induced tensor modes from scalar modes; and we discuss the generation of density fluctuations induced by gravity waves~-- scalar modes from tensor perturbations.  We derive a linear system of evolution equations for second-order gauge-invariant variables which characterise fully the induced modes of interest, with a source formed from variables quadratic in first-order quantities; these we transform into fully-fledged second-order gauge-invariant variables. Both the invariantly defined variables and the key evolution equations are considerably simpler than similar gauge-invariant results derived by other methods. By finding analytical solutions, we demonstrate that non-linear effects can significantly amplify or dampen modes present in standard linearised cosmological perturbation theory, thereby providing an important source of potential error in, and refinement of, the standard model. Moreover, these effects can dominate at late times, and on super-Hubble scales. ",
        "introduction": "}}}  The standard model of cosmology is built using perturbations of the homogeneous and isotropic Friedman-Lema\\^\\i tre-Robertson-Walker (FLRW) models. These perturbations have, with a few exceptions, been limited to first-order thus far; at first-order, the generic large-scale features of the universe can be accounted for. However, there are good reasons for investigating higher-order perturbations as a refinement of the accuracy of this model. The unprecedented precision of present and upcoming cosmic microwave background (CMB) measurements, such as WMAP and Planck, are reaching levels where non-linear effects may play an important role, especially in any non-Gaussianity present. Indeed, it is rapidly becoming an important area of research to try to predict possible imprints of quantum gravity effects in the CMB; but this can only properly succeed if we also fully understand the range of classical non-linear effects which may produce very similar fluctuations. Non-linear effects may also play an important role in inflation as in the stochastic approach to inflation, which partly includes second-order back-reaction effects~\\cite{GLMM}, and also from a fully second-order approach~\\cite{ABMR2}.  Another reason to consider second-order perturbations is that there is no way, within linear theory, to decide when the perturbations have become too large for the theory to handle~-- i.e., when non-linearities should be taken into account. The reason is that, while it would seem sensible to compare the magnitudes of linearised objects to some invariant background variable, such as the energy density, say, this does not always work, a simple example being gravitational waves in flat space. The study of cosmology is not yet at an advanced enough stage for fully non-linear numerical predictions, as is the case for many astrophysical situations, so second-order perturbation theory is the tool which must be developed for this.  For these reasons, second-order cosmological perturbation theory is beginning to be investigated, following work in the sixties by Tomita~\\cite{tomita} who first investigated the second-order terms from scalar modes. For example, other authors have considered second-order perturbations of a flat dust FLRW models, both with~\\cite{MTB} and without~\\cite{MMB} a cosmological constant. Inflation has also been considered at second order~\\cite{ABMR2}, providing the prediction of the bispectrum of perturbations from inflation. Meanwhile, the evolution of the curvature perturbation on super-Hubble scales after inflation has been shown to be constant at second order~\\cite{MW,BCLM}. More recently, it has been shown that second-order effects can lead to detectable non-Gaussianity in the CMB~\\cite{BMR1,BMR2}, as well as polarization~\\cite{MHM}. In this paper we complement these results by demonstrating that second-order effects can dominate at late times and on super-Hubble scales over first-order modes.   A crucial aspect of relativistic perturbation theory is the mapping, or gauge choice, between the (fictitious) background and the perturbed spacetime; many perturbation approaches are not invariant with respect to this mapping (see~\\cite{SW,BS} for a comprehensive discussion of this). In terms of coordinates, objects are typically not invariant under infinitesimal coordinate changes. This means that objects and modes specified in one gauge do not always correspond to objects and modes in another, making the interpretation of variables, and the physical situation, difficult. In fact, a gauge-dependent object is not observable~\\cite{BS}, which follows from the principle of general covariance. It is therefore important to have a method of describing perturbations using variables which are invariant under gauge transformations. A gauge-invariant object at \\emph{first-order} is a quantity which vanishes (or is constant) in the background; this is the Stewart-Walker Lemma~\\cite{SW}. This gauge problem was first solved in cosmology by Bardeen~\\cite{bardeen} by providing a full set of gauge-invariant quantities with which to describe the perturbed spacetime. Bardeen's variables are formed from gauge-invariant (GI) linear combinations of gauge-dependent variables, given in a particular coordinate system; but this prescription can make their interpretation difficult.   This situation has been clarified more recently by the covariant and gauge-invariant perturbation approach of Ellis, Bruni, and others~\\cite{EB,BDE}, using the covariant 1+3 approach~\\cite{EvE}. Its strength in cosmological applications lies in the fact that it is well adapted to the system it is describing: all essential information can be captured in a set of (1+3) covariant variables (defined with respect to a preferred timelike observer congruence $u^a$), that have an immediate physical and geometrical significance. These variables satisfy a set of {evolution} and {constraint} equations, derived from Einstein's field equations, and the Bianchi and Ricci identities, which form a closed system of equations when an equation of state for the matter is chosen. The covariant and gauge-invariant linearisation procedure is easy and transparent: it consists of deciding which variables are `first order' (or `of order $\\epsilon$') and those which are `zeroth order'~-- i.e., those which do not vanish in the background, which is usually a FLRW model. Products of first-order quantities can then be ignored in the equations. The key point of the approach is that it deals with physically or geometrically relevant quantities, such as the fractional density gradient, ${\\cal D}_a=a\\sdel_a\\mu/\\mu$ (where $a$ is the scale factor, $\\mu$ is the gauge-dependent energy density in the perturbed spacetime, and $\\sdel_a$ is the derivative operator in the observers' rest space), and the comoving spatial expansion gradient ${\\cal Z}_a=a\\sdel_a\\theta$ (where $\\theta$ is the gauge-dependent expansion), being the most important for scalar perturbations, while the electric and magnetic parts of the Weyl tensor, $E_{ab}$ and $H_{ab}$, respectively, represent the non-local parts of the gravitational field, and describe, amongst other things, the propagation of gravitational waves.  All 3-vectors and projected, symmetric, trace-free tensors in this approach, are automatically GI because of the symmetry of the background.     The gauge problem at second-order is quite a challenge: gauge transformations at second-order for second-order variables are \\emph{huge} in general~-- often as complicated as the evolution equations themselves~(see e.g., \\cite{MMB,BS,ABMR2}).  This is a problem far greater than in linear theory in terms of extracting the physics of non-linear effects, simply because it is far harder to intuit the physics by considering a selection of different gauges; the equations are just too complicated, and the range of possible gauges vast. So even choosing a specific gauge at first and second order, one has no guarantee that one has fully understood the problem. The key problem is that for an object to be gauge-invariant at \\emph{second-order} it must vanish in the background \\emph{and} at first-order (or be constant)~\\cite{BS}. This problem has hints of a solution in the recent work of Nakamura~\\cite{nakamura}, using the methods given in~\\cite{BS}, in which a general formalism for deriving GI quantities at second-order, assuming the existence of a method at first-order, is given. Thus it is possible that a fully gauge-invariant second-order theory may be presented in the near future using Bardeen's method for constructing first-order GI quantities together with Nakamura's approach. There has been recent progress along these lines;~\\cite{ABMR2} have given a gauge-invariant definition of the comoving curvature at second-order for scalar perturbations. Nevertheless, it is likely that when such a theory is found it will be very difficult to understand exactly what everything means, especially as the construction is not unique, although it will likely be crucial for numerical prediction of observable quantities such as the CMB power spectrum.  In this paper we discuss an alternative route to second-order gauge-invariant perturbations in cosmology following the covariant route initiated in~\\cite{EB}. The covariant 1+3 approach is a good method for perturbation theory simply because it is manifestly covariant; many of the usual gauge-problems arise as a result of coordinate transformations at each order. Once we have a complete set of GI variables, then we know that they are observable~\\cite{BS}, and correspond to physical or geometrical quantities.\\footnote{It should be noted that there is a freedom inherent in the 1+3 approach, even after a complete set of GI variables have been specified, and this freedom lies in ones choice of observers at the relevant perturbative order: but this is {not} a gauge choice in perturbation theory, as it has nothing to do with ones mapping between the background and perturbed spacetime; it is simply a choice of observers in the perturbed spacetime, a choice one must make in any non-vacuum spacetime. For example, the gauge-invariant density gradient at first-order, ${\\cal D}_a$, can be set to zero by a first-order Lorentz boost, at the expense of introducing acceleration, say. Under a zeroth-order boost, however, ${\\cal D}_a$ would no longer be gauge-invariant as it would not vanish in the background, and would correspond to unusual observers in the background. In metric-based approaches (where the components of the metric are explicitly solved for), one must specify a perturbed velocity \\emph{relative to} the background, as well as having this frame freedom. This is because the Stewart-Walker lemma and its generalisations refer to the principle of general covariance, and not the principle of relativity: different observers measure different values for physical quantities such as the energy density, the magnetic field or the electric part of the Weyl tensor, for example. The only objects which are observer (or tetrad) independent are the metric, the Riemann tensor (and its trace and trace-free parts) and the Maxwell tensor.}  Using a zeroth-order background of a dust FLRW model with zero cosmological constant, we discuss `mode-crossing' perturbations; that is, perturbations at second order generated from an `orthogonal mode' at first order~-- second-order  tensor modes generated by first-order scalar modes (Section~\\ref{sec1}), and viceversa (Section~\\ref{sec2}). \\footnote{We ignore rotational modes throughout. Only if they are included at first-order will rotational modes occur which do not occur at first-order, in the absence of acceleration.} We define a complete set of second-order gauge-invariant quantities which fully describe the induced modes, and convert all products of first-order quantities appearing in the second-order equations into new second-order variables, thus making the equations explicitly GI at second order, in a manner very similar to~\\cite{CMBD}. This simplifies the presentation of the equations by converting them into a \\emph{linear} system of DE's, a procedure which also simplifies their solution. In the case of first-order tensor modes, this system is infinite dimensional reflecting the coupling of modes of each wavelength feeding into the induced scalar modes. The standard zeroth-order harmonic functions can be used on second-order variables to remove the tensorial nature of the equations, and remove spatial gradients so converting the system into odes which may then be easily integrated. We explicitly integrate both cases under investigation to show that scalar-mode coupling effects may play an important role at both early and late times, and on small and large scales. ",
        "conclusions": "}}}    Understanding non-linear effects in cosmology will likely play a pivotal role in cosmology in coming years, as observations pin down the overall properties of the universe ever more tightly. Most of the problems preventing a systematic study of non-linear relativistic effects are mathematical: fully non-linear relativistic simulations are beyond the horizon in cosmology at present, leaving perturbation theory as our main probe. But non-linear perturbation theory is not particulary simple as gauge problems and large equations make it quite untidy. In this paper we have considered for the first time the 1+3~covariant approach to gauge-invariant non-linear perturbation theory.     The main problem  the 1+3 covariant approach has at second order is intricately related to the gauge-problem: while second-order equations, such as Eq~(\\ref{sigwavenl}), can be written down~-- by crossing off all the third order terms~-- they can't in general be integrated. This is because the derivative operators, through the commutation relations, contain lots of first order stuff, which the covariant approach does not solve for at first-order (it only solves for physical variables, not operators, in contrast with the metric approach). In order to integrate the second-order equations therefore, derivative operators (dot and $\\sdel$) must operate only on variables which vanish at first and zeroth order~-- SOGI variables, in other words. We created these SOGI variables in two ways: by appropriate derivatives of first-order quantities which vanish at first-order~-- exactly the same trick used to create FOGI variables from zeroth-order scalars; and from explicit products of FOGI quantities.     The GI approach we develop here has the advantage of producing a relatively simple linear system of odes, which is easy to integrate. By integrating the SOGI equations, we have found that non-linear effects may well play an important role at both early and late times. Indeed in both cases we examined we found that second-order modes will dominate the perturbation spectrum at late times, which is a consequence of the fact that second-order modes arise partly as an integrated effect, and can't always be neglected by assumption. In addition, tensor modes generated by non-linear effects may play an key role on super-Hubble scales. This analysis suggests that linear perturbation theory may not be a sufficiently accurate tool to understand some subtle dynamical aspects of the universe, but further investigation is required to find out exactly what these are."
    },
    "0311/astro-ph0311219_arXiv.txt": {
        "abstract": "We investigate numerically and semi-analytically the collapse of low-mass, rotating prestellar cores. Initially, the cores are in approximate equilibrium with low rotation (the initial ratio of thermal to gravitational energy is $\\alpha_0 \\simeq 0.5$, and the initial ratio of rotational to gravitational energy is $\\beta_0 = 0.02\\;{\\rm to}\\;0.05$). They are then subjected to a steady increase in external pressure. Fragmentation is promoted -- in the sense that more protostars are formed -- both by more rapid compression, and by higher rotation (larger $\\beta_0$).  In general, the large-scale collapse is non-homologous, and follows the pattern described in Paper I for non-rotating clouds, viz. a compression wave is driven into the cloud, thereby increasing the density and the inflow velocity. The effects of rotation become important at the centre, where the material with low angular momentum forms a central primary protostar (CPP), whilst the material with higher angular momentum forms an accretion disc around the CPP. More rapid compression drives a stronger compression wave and delivers material more rapidly into the outer parts of the disc. Consequently, (i) there is more mass in the outer parts of the disc; (ii) the outer parts of the disc are denser (because the density of the material running into the accretion shock at the edge of the disc is higher); and (iii) there is less time for the gravitational torques associated with symmetry breaking to redistribute angular momentum and thereby facilitate accretion onto the CPP. The combination of a massive, dense outer disc and a relatively low-mass CPP renders the disc unstable against fragmentation, and leads to the formation of one or more secondary protostars. At their inception, these secondary protostars are typically four or five times less massive than the CPP.  For very rapid compression there is no CPP and the disc becomes more like a ring, which then fragments into two or three protostars of comparable mass.  For more rapid rotation (larger $\\beta_0$), the outer disc is even more massive in comparison to the CPP, even more extended, and therefore even more prone to fragment. ",
        "introduction": "It is now well established that stars form in  dense cores embedded in  molecular clouds (see e.g. Andr\\'e et al. 2000). However the conditions under which these cores form, become unstable, collapse and fragment are still a matter of debate.  In a previous study, Hennebelle et al. (2003, Paper I) have investigated the possibility that the collapse of a prestellar core is driven from the outside by an increase in the external pressure. This model seems to reproduce many of the key features observed in nearby star forming cores. In particular, (i) during the pre-stellar phase, the density profile is approximately flat in the centre (Abergel et al. 1996, Bacmann et al. 2000, Motte \\& Andr\\'e 2001). (ii) For slow to moderate compression rates, subsonic infall velocities develop in the outer parts of the core during the prestellar phase (Tafalla et al. 1998, Williams et al. 1999, Lee et al. 1999). (iii) During the Class 0 phase, subsonic velocities persist in the outer parts of the core (Belloche et al. 2002), and transsonic velocities develop in the inner parts ($r \\la 2000\\,{\\rm AU}$). (iv) There is an initial short phase of rapid accretion onto the central protostar (the Class 0 phase), followed by a longer phase of slower accretion (the Class I phase), as inferred from the studies of Greene et al. (1994), Kenyon \\& Hartmann (1995), Bontemps et al. (1996) and Motte \\& Andr\\'e (2001).  A question that was not addressed in Paper I is the formation of multiple protostars by fragmentation of a collapsing core. This question is critical, because it is now well established (Duquennoy \\& Mayor 1991, Fischer \\& Marcy 1992, Ghez et al. 1997) that a large fraction of stars is in binary -- or higher multiple -- systems (i.e. at least 50\\% for mature solar-type stars in the field, and a higher percentage for Pre-Main Sequence stars in young associations).      \\subsection{Previous Work}  The gravitational fragmentation of a collapsing cloud has been intensively investigated for over two decades (see Bodenheimer et al. 2000 for a review), mostly with numerical simulations, but also by semi-analytic means. The enduring hope has been to derive a robust general theorem defining the conditions required for fragmentation, for instance, in terms of the initial ratio of thermal to gravitational energy, $\\alpha_0$, and the initial ratio of rotational to gravitational energy, $\\beta_0$. However, robust theorems have proved elusive. This is in part because the parameter space of initial conditions and constitutive physics is very large. Therefore the generality of the results obtained is hard to establish. In addition, fragmentation appears to depend sensitively on the complicated radiative-transport and thermal-inertia effects which come into play when star-forming gas switches from being approximately isothermal to being approximately adiabatic. The computational resources required to simulate this radiative transport properly are not yet available. Furthermore, all such theorems inevitably beg the question: How were the unstable initial conditions created in the first place? Here one probably has to appeal loosely to the chaotic effects of supersonic turbulence in star forming molecular clouds (e.g. Elmegreen 2000; Padoan \\& Nordlund 2002)  The simplest case to treat is a spherical cloud with uniform density, solid-body rotation, and isothermal equation of state; this is sometimes called the standard model. Using semi-analytic arguments, based on the classical stability analysis of Maclaurin spheroids (Lyttleton 1953; Chandrasekhar 1969), Tohline (1981) concluded that all such clouds should fragment. Subsequently, Miyama, Hayashi \\& Narita (1984) used numerical simulations to derive a fragmentation criterion of the form $\\alpha_0 \\beta_0 \\la 0.12$. A similar criterion was obtained semi-analytically by Hachisu \\& Eriguchi (1984, 1985), and Miyama (1992) revised the criterion slightly to $\\alpha_0 \\beta_0 \\la 0.15$, on the basis of an analysis of the flatness and stability of rotating discs. Tsuribe \\& Inutsuka (1999a), again using a semi-analytic approach, took account of the non-homologous nature of collapse, and obtained a somewhat different criterion, which can be approximated by $\\alpha_o \\la 0.55 - 0.65 \\beta_0$. This criterion was subsequently confirmed with simulations (Tsuribe \\& Inutsuka 1999b).  The standard model can be modified in several interesting ways.  (1) The standard model can be modified by adding an $m = 2$ azimuthal density perturbation, i.e. $\\rho({\\bf r}) = \\rho_0 \\left[ 1 + A {\\rm cos}(\\phi) \\right]$, where $\\phi$ is the azimuthal angle in spherical polar co-ordinates and $A$ is the fractional amplitude of the perturbation. This perturbation is a common basis for numerical explorations of collapse and fragmentation. The standard model was first simulated by Boss \\& Bodenheimer (1979), who invoked a perturbation with amplitude $A = 0.5$. They found that a cloud having $\\alpha_0 = 0.25$ and $\\beta_0 = 0.20$ collapsed and fragmented to form a binary system. This result was reproduced by Burkert \\& Bodenheimer (1993) using a more accurate code. Burkert \\& Bodenheimer (1993) also performed simulations with $\\alpha_0 = 0.26$ and $\\beta_0 = 0.16$ and with the amplitude of the perturbation set to $A = 0.5$ and $A = 0.1$. In the latter case they obtained not just a binary, but also a line of smaller fragments between the main binary components. However, Truelove et al. (1998) have demonstrated using an AMR code, that this is an artefact. If the collapse remains truly isothermal to high densities, and if the resolution is sufficiently high for the Jeans mass to be resolved at all times, the material between the binary components should not fragment. This has been confirmed by Boss et al. (2000), by Sigalotti \\& Klapp (2001), and by Kitsionas \\& Whitworth (2002), who followed the purely isothermal collapse to even higher density, using SPH with particle splitting.  (2) The standard model can be modified by changing the equation of state. Tohline (1981) and Miyama (1992) considered the collapse of clouds having an adiabatic equation of state, $P = K \\rho^\\gamma$, and concluded that the condition for fragmentation takes the form \\begin{eqnarray} \\label{crit_miyama} \\alpha_0 \\beta_0 ^{(4 - 3 \\gamma)} \\le f(\\gamma) \\,, \\end{eqnarray} where, for example, $f(5/3) \\simeq 0.064$. Alternatively, a barotropic equation of state can be adopted, in which the gas is isothermal at low densities (where star forming gas is expected to be thin to its own cooling radiation), and adiabatic at high densities (where the gas is expected to be optically thick to its cooling radiation). Collapse simulations using a barotropic equation of state of this form are presented by {Bonnell (1994)}, Bate \\& Burkert (1997), Boss et al. (2000), and Cha \\& Whitworth (2003). A barotropic equation of state of this type is also used in the simulations presented in this paper, and is discussed in Section \\ref{EOS}. It is clear that treatment of the equation of state has a profound influence on the outcome of collapse. For example, Bate \\& Burkert (1997) show that the collapse of an initially uniform-density cloud having $\\alpha_0 = 0.26$, $\\beta_0 = 0.16$, and an $m = 2,\\,A = 0.1$ azimuthal perturbation, {\\it does} produce a line of small fragments between the two main binary components, {\\it if} the barotropic equation of state described above is used (but not if the gas remains isothermal indefinitely). Perhaps the most telling result in this regard is reported by Boss et al. (2000) who simulated the collapse of a cloud with an $m = 2$, $A = 0.1$ perturbation, first using a barotropic equation of state, and then including radiation transport and an energy equation. The latter case produced a binary, whereas the former case did not, even though the variation of pressure with density was very similar in the two cases. The implication is that the subtle radiation-transport and thermal-inertia effects, which suddenly become important as isothermality gives way to adiabaticity, play a critical r\\^ole in determining the pattern of fragmentation.  (3) The standard model can be modified by imposing a density profile. Myhill \\& Kaula (1993), using a code which included radiation transport and an energy equation, showed that clouds with solid-body rotation, $\\alpha_0 = 0.16$, $\\beta_0 = 0.17$,  and an $m = 2$, $A = 0.1\\;{\\rm or}\\;0.5$ azimuthal perturbation do not fragment during the isothermal collapse phase, if the density profile is centrally peaked (specifically $\\rho \\propto r^{-n}$ with $n = 1$ or $n = 2$). However, Burkert, Bate \\& Bodenheimer (1997) repeated the simulation with $A = 0.1$ and $n = 1$, using a barotropic equation of state. They found that after the central primary protostar condenses out, a circumstellar disc forms around it, and this disc then fragments to produce companion protostars. A number of simulations have also been performed with Gaussian density profiles (e.g. Boss \\& Myhill 1995; Boss 1996; Burkert \\& Bodenheimer 1996; Boss et al. 2000), and also with exponential profiles (Boss 1993), but there does not appear to have been a systematic evaluation of the influence of these profiles on the outcome of collapse.  (4) The standard model can be modified by introducing differential rotation. Myhill \\& Kaula (1993), again using a code which included radiation transport and an energy equation, showed that clouds with $\\alpha_0 = 0.16$, $\\beta_0 = 0.17$, a centrally peaked density profile ($\\rho \\propto r^{-n}$ with $n = 1$ or $n = 2$), and an $m = 2$, $A = 0.1\\;{\\rm or}\\;0.5$ azimuthal perturbation, do fragment if they have sufficient differential rotation (in contrast with the result for solid-body rotation). The tendency for differential rotation to promote fragmentation has been confirmed by Boss \\& Myhill (1995) and Cha \\& Whitworth (2003).  (5) Finally, a number of authors have explored the effect of clouds having non-spherical initial shapes. Evidently the outcome here depends not only on $\\alpha_0$ (and $\\beta_0$, if there is rotation), but also on the aspect ratio of the initial configuration. Bastien (1993) simulated the collapse of isothermal, non-rotating cylindrical clouds and determined the mass per unit length required for fragmentation. This problem was revisited by Bonnell \\& Bastien (1991) and Bastien et al. (1991), and extended to polytropic cylinders by Arcoragi et al. (1991). Bonnell et al. (1991) explored the circumstances under which isothermal cylinders rotating about an axis perpendicular to their elongation fragment to produce binary systems, and Bonnell et al. (1992) considered isothermal cylinders rotating about an arbitrary axis. Bonnell \\& Bastien (1992) repeated this last study with cylinders having a density gradient along the symmetry axis, and showed that quite modest gradients were sufficient to produce binary systems with mass ratios in the range 0.1 to 1. Nelson \\& Papaloizou (1993) simulated the collapse of prolate spheroids, and showed that they fragment if the mass per unit length is sufficiently high (as for cylinders), but the binary components tend to be closer because there is less mass at the ends of a prolate spheroid. Boss (1993) simulated the collapse and fragmentation of a rotating, mildly prolate cloud with an exponential density profile, and showed that the conditions for fragmentation were rather more restrictive than those obtained by Miyama, Hayashi and Narita (1984). These results were extended to mildly prolate clouds having a Gaussian density profile and differential rotation by Boss \\& Myhill (1995); and to more elongated prolate clouds having a Gaussian density profile, but solid-body rotation, by Sigalotti \\& Klapp (1997). Boss (1996) simulated the collapse of rotating isothermal oblate spheroidal clouds and showed that fragmentation into multiple systems occurs provided $\\alpha_0 < 0.4$, almost independent of $\\beta_0$. Monaghan (1994) explored the effect of vorticity on the collapse and fragmentation of ellipsoidal clouds, using numerical simulations.      \\subsection{Outline of paper}  In this paper we pursue further our investigation of prestellar cores whose collapse is induced by a steady increase in the external pressure. We extend the investigation to the case of rotating cores, and focus on the formation of circumstellar discs around the central primary protostars (CPPs) and the subsequent fragmentation of these discs. Our main goal is to relate the properties of the star systems formed (multiplicities and mass ratios) to the dynamics of the collapse, and hence to the parameters $\\beta_0$ and $\\phi$, measuring -- respectively -- the initial rate of rotation and the rate of compression (see Eq.~(\\ref{phi})). Our model invokes initial conditions, and generates density and velocity profiles, very similar to those inferred from observations of real star forming cores. However, without performing a large (i.e. statistically robust) ensemble of simulations, we cannot know whether it also delivers multiple systems with statistics similar to those observed.  In Section \\ref{Constit} we describe the numerical method we use, the constitutive physics, and the initial and boundary conditions. In Section \\ref{grand_champs}, we present our results, with special attention to the large-scale velocity and density fields in the cores. Section \\ref{frag} discusses the evolution of the discs that form around the central primary protostars, with particular emphasis on the fragmentation process. Section \\ref{Summary} summarizes our main conclusions. In the Appendix, we develop a semi-analytical description of the key features discussed in Sections \\ref{grand_champs} and \\ref{frag}, and estimate the timescales influencing disc stability. We perceive this analysis as an integral part of our paper, and we have put it in an appendix purely so that those whose main interest is in the phenomenology of fragmentation can read about the results of our simulations without getting to grips with the more mathematical aspects, which give an approximate quantitative explanation for the dependence of disc stability on the initial core rotation and the rate of compression. ",
        "conclusions": "\\label{Summary}  We have investigated the effect of compression on the collapse and fragmentation of a rotating core.  Most of the conclusions of Paper I, concerning the large-scale dynamics of the collapsing core, are still valid. In particular, the increase in external pressure drives a compression wave into the core. The compression wave leaves in its wake a velocity field very similar to those recently inferred for prestellar cores, from observations of asymmetric molecular-line profiles. Tafalla et al. (1998) estimated inflow at $\\sim 0.10\\,{\\rm km}\\,{\\rm s}^{-1}$ in the outer layers of L1544, and Williams et al. (1999) found inflow at $\\sim 0.08\\,{\\rm km}\\,{\\rm s}^{-1}$ further in. Lee et al. (1999) detected inflow velocities ranging from 0.04 to $0.10\\,{\\rm km}\\,{\\rm s}^{-1}$ in several other prestellar cores. Our models generate comparable inflow velocities for $\\phi$ in the range 10 to 1, i.e slow to intermediate compression. Interestingly, in regions like Perseus and $\\rho$ Ophiucus, where star formation is triggered more violently, cores with significantly higher infall velocities have been observed. For example, Di Francesco et al. (2001) detected velocities of $\\sim 0.5\\;{\\rm to}\\;0.7\\,{\\rm km}\\,{\\rm s}^{-1}$ in the Class 0 protostars of NGC 1333 IRAS 4, suggesting that very fast compression may have occurred in this region.  The fundamental difference from Paper I is that because the cores we simulate here have rotation, the inflowing material does not all converge directly onto the central primary protostar (CPP). Instead, a large fraction of it first collects in an accretion disc around the CPP, and instabilities in the outer parts of this disc may then lead to fragmentation producing additional protostars. As noted by Larson (2002), the critical factors determining the stability of the outer disc are (i) its mass and density, and (ii) the speed with which it is assembled: {\\em a quickly assembled, massive, dense outer disc is unstable against fragmentation.}  We show numerically -- and, in the Appendix, semi-analytically -- that the density in the outer disc is larger for faster compression, and for larger $\\beta_0$. These effects may already have been observed. In cores belonging to the relatively quiescent star formation region Taurus, the inflow velocities are compatible with slow compression, and the density is close to the density of the singular isothermal sphere (SIS). Conversely, in cores belonging to more active star formations regions, the inflow velocities appear to be supersonic, and the densities are about one order of magnitude higher than the density of the SIS (Motte \\& Andr\\'e 2001; Andr\\'e et al. 2003). This accords with the predictions of the analysis in the appendix.  Except in the case of very rapid compression, the low angular momentum material arriving in the centre of the core goes directly into the CPP, and the high angular momentum material forms an accretion disc round the CPP. If the compression is more rapid, this has a number of effects which tend to render the disc more unstable.  First, the material accreting onto the outer parts of the disc arrives with higher density, and therefore the density in the outer disc -- following compression in the accretion shock at the edge of the disc -- is also higher. Second, material is delivered into the outer parts of the disc more rapidly, and therefore these outer parts become more massive. Third,  there is less time for the gravitational torques associated with symmetry-breaking instabilities in the disc (the same instabilities which lead to fragmentation) to redistribute angular momentum and thereby facilitate the continuing growth of the CPP. The combination of a massive, dense outer disc and a low-mass CPP makes the outer disc unstable against fragmentation, spawning secondary protostars with masses typically four or five times lower than the CPP.  For very rapid compression there is no CPP; all the material flows into the disc, and it is so concentrated towards the edge that it is more accurately described as a ring. The ring then fragments into two or three protostars of comparable mass.  For more rapid rotation ($\\beta_0 = 0.05$) the outer disc is even more massive in comparison to the CPP, even more extended, and therefore even more prone to fragment.  In the appendix we analyse the structure of the inflowing envelope, and its consequences for the stability of the central disc. This analysis explains why higher rotation (large $\\beta_0$) and more rapid compression (small $\\phi$) promote fragmentation and the formation of multiple protostars."
    },
    "0311/astro-ph0311169_arXiv.txt": {
        "abstract": "We report the discovery of the two-image gravitationally lensed quasar SDSS~J1335+0118. The object was selected as a lens candidate from the Sloan Digital Sky Survey. The imaging and spectroscopic follow-up observations confirm that the system exhibits two gravitationally lensed images of a quasar at $z=1.57$. The image separation is $1\\farcs56$.  We also detect an extended component between the two quasar images, likely the lensing galaxy. Preliminary mass modeling predicts the differential time delay $\\Delta t\\sim 30h^{-1}{\\rm day}$ assuming the redshift of the lens galaxy is $0.5$. ",
        "introduction": "Multiple images of a quasar produced by a foreground galaxy are useful for cosmological and astrophysical applications, including the measurement of the cosmological constant through the lensing rate \\citep{turner90,fukugita90,chae02}, the determination of global Hubble constant from time delays between images (Refsdal 1964, 1966), the study of the mass distributions of the lensing galaxies \\citep{kochanek91,metcalf01,chiba02,dalal02,schechter02,koopmans03,rusin03}, and the formation and evolution of galaxies \\citep{kochanek00,kochanek01,keeton01a,oguri02,ofek03}. A large homogeneous sample of lensed quasars with a well-understood selection function is essential in these statistical studies. In addition, lensed quasars can be a powerful tool to study quasar itself through microlensing, spectra, or variabilities of lensed quasars \\citep{chang84,grieger91,mediavilla98,mineshige99,yonehara99,wyithe00,yonehara01,lewis03}. Therefore, it is important to make a large sample of lensed quasars also to find more gravitational lens systems which are suitable for studying quasar.  The Sloan Digital Sky Survey (SDSS; \\cite{york00,stoughton02,abazajian03}) has the potential to provide such a lensed quasar sample. The SDSS is a survey to image $10^4{\\rm deg^2}$ of the sky as well as to obtain spectra of galaxies and quasars from the imaging data. The dedicated 2.5 meter telescope at Apache Point Observatory is equipped with a multi-CCD camera \\citep{gunn98} with five broad band centered at $3561$, $4676$, $6176$, $7494$, and $8873${\\,\\AA} \\citep{fukugita96b}. The imaging data are automatically reduced by the photometric pipeline \\citep{lupton01}. The astrometric positions are accurate to about $0\\farcs1$ for sources brighter than $r=20.5$ \\citep{pier03}. The photometric errors are typically less than 0.03 magnitude \\citep{hogg01,smith02}. The SDSS quasar selection algorithm is presented in \\citet{richards02}; the SDSS spectra cover $3800$--$9200${\\,\\AA} at a resolution of $1800$--$2100$. The final spectroscopic quasar sample is expected to comprise $10^5$ quasars, thus will contain $\\sim10^2$ lensed quasars given the typical lensing probability $0.1\\%$ \\citep{turner84}.  Indeed, several new gravitational lens systems have been discovered using SDSS data (\\authorcite{inada03a} 2003a,b,2004; \\cite{morgan03}; \\cite{johnston03}; \\cite{pindor04}; \\cite{oguri04}).  We report the discovery of the lensed quasar SDSS~J133534.79+011805.5 (SDSS~J1335+0118) in the SDSS. This quasar has already been identified in the Large Bright Quasar Survey (LBQS 1333+0133; \\cite{hewitt91}), but not as lensed. A search for lens candidates in the LBQS was carried out by \\citet{hewett98}, but those authors state that their search is not sensitive to lens systems with $\\leq 3''$ separations. The quasar was selected as a lens candidate in the course of an ongoing search for strongly lensed quasars in the SDSS. From the results of photometric and spectroscopic follow-up observations, we conclude that the quasar is lensed by an intervening galaxy. ",
        "conclusions": "}  We have reported the discovery of a new two-image gravitational lens SDSS~J1335+0118. The system was identified as a new lens candidate in the SDSS.  The image separation is $1\\farcs56\\pm0.002$ in the $i$-band image taken with Subaru Suprime-cam.  The spectroscopic observation with ESO NTT has confirmed both images have identical redshift $z=1.57$ ($z=1.57\\pm0.03$ and $z=1.57\\pm0.05$ for components A and B, respectively). We have also probably identified the lensing galaxy between two quasar components in the higher-resolution imaging data taken with Subaru Suprime-cam and Keck NIRC. The lens geometry is well reproduced with a simple mass model and reasonable model parameters. Assuming the redshift of the lens galaxy as $0.5$, we have predicted a differential time delay $\\Delta t\\sim 30h^{-1}{\\rm day}$.  \\bigskip We thank an anonymous referee for many useful comments. Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society. The SDSS Web site is http://www.sdss.org/.  The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions. The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, Princeton University, the United States Naval Observatory, and the University of Washington.  Part of the work reported here was done at the Institute of Geophysics and Planetary Physics, under the auspices of the U.S. Department of Energy by Lawrence Livermore National Laboratory under contract No.~W-7405-Eng-48.  This work is based in part on data collected at Subaru Telescope, which is operated by the National Astronomical Observatory of Japan, and partly based on observations collected at the European Southern Observatory, Chile under program 70.D-0469. 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. We thank the staffs of Subaru, Keck, and ESO NTT for their excellent assistance."
    },
    "0311/astro-ph0311396_arXiv.txt": {
        "abstract": "In the 1998-99 flight, BOOMERanG has produced maps of $\\sim 4 \\%$ of the sky at high Galactic latitudes, at frequencies of 90, 150, 240 and 410 GHz, with resolution $\\simgt 10'$. The faint structure of the Cosmic Microwave Background at horizon and sub-horizon scales is evident in these maps. These maps compare well to the maps recently obtained at lower frequencies by the WMAP experiment. Here we compare the amplitude and morphology of the structures observed in the two sets of maps. We also outline the polarization sensitive version of BOOMERanG, which was flown early this year to measure the linear polarization of the microwave sky at 150, 240 and 350 GHz. ",
        "introduction": "BOOMERanG is balloon-borne microwave telescope, sensitive at 90, 150, 240 and 410 GHz, with a resolution of $\\sim 10'$. The instrument was equipped with very sensitive spider-web bolometers (Mauskopf et al. 1997). The image of the sky is obtained by slowly scanning the full payload in azimuth ($\\pm 30^o$) at constant elevation. The scan center constantly tracks the azimuth of the lowest Galactic foreground region, situated in the southern hemisphere, in the Horologium constellation. Every day of the flight the instrument obtains a fully cross-linked map of about $45^o \\times 30^o$ of the sky. The instrument was flown in a long duration circum-Antarctic flight from Dec.28, 1998 to Jan.8, 1999, and has been described in Piacentini et al. (2002) and in Crill et al. (2003). The main target of the experiment was the detection of anisotropy in the Cosmic Microwave Background. The maps produced by the experiment have been published in de Bernardis et al. (2000), Masi et al. (2001), Netterfield et al. (2002), Ruhl et al. (2003). Due to the limited size and to 1/f noise, these maps do not contain information for angular scales $\\gg 5^o$ and have been filtered accordingly. The maps are calibrated to $\\sim 10 \\%$ in gain and to $\\sim 10\\%$ in beam FWHM. In the 150 GHz map, the signal is well above the noise, and maps taken at different scan speeds and in different locations are perfectly consistent, demonstrating the low level of systematic effects. The level of the noise is of the order of 50 $\\mu K$ per $7'$ pixel and of 20 $\\mu K$ per $28'$ pixel. At 90, 150 and 240 GHz the rms signal has the spectrum of CMB anisotropy, and does not fit any reasonable spectrum of foreground emission. The temperature fluctuations detected in the high latitude part of the 150 GHz map are remarkably gaussian (Polenta et al. 2002, De Troia et al. 2003). The angular power spectrum of the 150 GHz map in the multipole range $50 < \\ell < 1000$ has been computed in Netterfield et al. (2002) and Ruhl et al. (2003). Three peaks have been detected in the power spectrum, at multipoles $\\ell \\sim$ 210, 540, 845 (de Bernardis et al., 2002). Interstellar Dust contamination of the 150 GHz power spectrum has been shown to be less than 1\\% (Masi et al. 2001). These results fit the scenario of acoustic oscillations of the primeval plasma at horizon and subhorizon scales (Sunyaev and Zeldovich 1970, Peebles and Yu 1970). In the framework of adiabatic inflationary structure formation the cosmological parameters have been estimated from the measured power spectrum (Lange et al. 2001, Netterfield et al., 2002, de Bernardis et al. 2002, Ruhl et al. 2003). The three cosmological parameters best constrained by the BOOMERanG data are the curvature parameter $\\Omega = 1.03\\pm 0.05$  (the universe is nearly flat), the spectral index of the primordial density perturbations $n_s =1.02 \\pm 0.07$ (nearly scale invariant), the physical density of baryons $\\Omega_b h^2 = 0.023 \\pm 0.003$ (consistent with Big Bang Nucleosynthesis).  Here we compare the BOOMERanG maps to the maps recently obtained at similar frequency and resolution by the WMAP satellite (Bennett et al., 2003). Working from the advantage L2 point of the Sun-Earth system, in its first year of operation WMAP has produced full sky maps of the microwave sky at 22, 32, 41, 60, 94 GHz, with resolution of the order of $15'-30'$ and noise of $\\sim 35 \\mu K$ per pixel in 28' pixels. The maps are precisely calibrated (better than 1\\%). The power spectra obtained from these maps are fully consistent with the CMB power spectrum measured by BOOMERanG at 150 GHz, and with the adiabatic inflationary scenario: they allow the precise determination of most of the cosmological parameters; these estimates improve the precision of the BOOMERanG ones by a factor 2-3 for the parameters mentioned above; the full-sky coverage and polarization sensitivity allows the determination of previously poorly constrained parameters, like the reionization optical depth (Spergel et al. 2003).  \\noindent The comparison of the BOOMERanG and WMAP maps is carried out with three targets:  \\noindent $\\bullet$ Compare independent maps of the CMB with similar angular resolution, and confirm the detection of primordial structures in both experiments;  \\noindent $\\bullet$ Improve the BOOMERanG calibration using the precise calibration of WMAP, and demonstrate that 1\\% gain calibration is possible for the new, polarization-sensitive BOOMERanG-B2K survey;  \\noindent $\\bullet$ Infer the level and properties of foregorund contamination in forthcoming deep surveys of the CMB. ",
        "conclusions": ""
    },
    "0311/astro-ph0311113_arXiv.txt": {
        "abstract": "{We present a study of the star formation rate (SFR) for a sample of 16 distant galaxies detected by ISOCAM at 15$\\mu$m in the CFRS0300+00 and CFRS1400+52 fields. Their high quality and intermediate resolution VLT/FORS spectra have allowed a proper correction of the Balmer emission lines from the underlying absorption.  Extinction estimates using the H$\\beta$/H$\\gamma$ and the H$\\alpha$/H$\\beta$ Balmer decrement are in excellent agreement, providing a robust measurement of the instantaneous SFR based on the extinction-corrected $H\\alpha$ luminosity. Star formation has also been estimated exploiting the correlations between IR luminosity and those at MIR and radio wavelengths. Our study shows that the relationship between the two SFR estimates follow two distinct regimes: (1) for galaxies with SFR$_{IR}$ below $\\sim 100M_{\\odot}/yr$, the SFR deduced from H$\\alpha$ measurements is a good approximation of the global SFR and (2) for galaxies near of ULIRGs regime, corrected H$\\alpha$ SFR understimated the SFR by a factor of 1.5 to 2.  Our analyses suggest that heavily extincted regions completely hidden in optical bands (such as those found in Arp 220) contribute to less than 20\\% of the global budget of star formation history up to z=1. ",
        "introduction": "In the last few years, great steps toward a better understanding of the distant Universe have been done using a panchromatic approach. The mid IR windows opened by the camera ISOCAM (Cesarsky et al, 1996) have lead to the identification of the galaxies responsible for  the bulk of the infrared background, a population of dusty galaxies which evolves very rapidly in number (Elbaz et al., 1999, 2002). Luminous Infrared Galaxies (LIRGs, L$_{IR}\\ge 10^{11}L_{\\odot}$) detected by ISO are mainly dusty starbursts (Fadda et al., 2002) with large star formation ($> 50-100 M_{\\odot}$/yr) triggered by interactions, $\\sim$30-40\\% of detected galaxies show optical signs of interaction (Flores et al., 1999). Rigopoulou et al. (2000, hereafter R00) using the VLT studied the near IR spectra of a few LIRGs and claimed that the SFR measured from IR luminosity (hereafter $SFR_{IR}$) is at least 3 times higher than that based on $H\\alpha$ luminosity (hereafter SFR$_{H\\alpha}$). However, this claim is affected by large uncertainties on the SFR$_{H\\alpha}$, in absence of a proper extinction correction. Indeed extinction for such heavily dust-shrouded objects is very important, and only accurate estimates of the extinction can confirm or infirm the R00 claim.  Recently, Hammer et al. (2001) have analyzed intermediate resolution VLT/FORS spectra (R=3.5\\AA$\\;$ at rest) of 3 compact LIRGS at z$\\sim$0.6.  They derived accurate fluxes for Balmer emission lines (H$\\beta$, H$\\gamma$) after a proper removal of the underlying stellar absorption. Extinction estimates using observed H$\\beta$/H$\\gamma$ ratio were used to correct $H\\beta$ luminosities and derive SFR$_{H\\alpha}$ which are consistent with $SFR_{IR}$. These objects present strong metallic and Balmer absorption lines combined with intense emission lines, which indicate a particularly complex star formation history (Hammer et al, 2001). It is unclear yet if these properties are shared by the whole galaxy population at intermediate redshifts. Indeed, at z $\\ge$ 0.4, an important fraction of the galaxy population was identified to Balmer-strong galaxies with emission, in the field (Hammer et al, 1997) and in galaxy clusters (Poggianti et al, 1999).    \\\\ This letter presents a preliminary study of 16 LIRGs out of a sample of 90 observed at VLT and CFHT.  The data are described in Section 2. Section 3 summarizes the different techniques used for the data analysis.  The determination of the SFR is discussed in Sections 4 and 5. Throughout this paper we adopt H$_0$=50 km/s/Mpc and q$_0$=0.5 ($\\Omega_M=1,\\Omega_{Lambda}=0$).  \\begin{figure*}   % \\centering \\includegraphics[width=12cm]{fig1.eps} \\caption{ % Combined VLT FORS+ISAAC spectra, after masking sky emission lines, and smoothing the continuum of both galaxies (see Hammer et al. 2001;Gruel et al. in prep).  Zooms in the center shows details of our spectra (S/N$>$5 per resolution element in the continuum), including the underlying Balmer absorption.  Lower panel shows combined near IR spectra. } \\label{Fig1} \\end{figure*} ",
        "conclusions": "Extinction effects potentially affect a large cosmological sources and their consequences have been the subject of a long debate in observational cosmology. In this letter, we study a representative sample of dust enshrouded starbursts, which, at moderate redshifts, can be extensively studied in spectroscopy, at MIR and at radio wavelengths. We have compared the instantaneous SFR related to the optical nebular emission (SFR$_{H\\alpha}$) to the SFR estimated from the MIR emission measured by ISO , and, when available, from the radio emission These two methods give compatible results, and we denote the result by (SFR$_{IR}$).  H$\\alpha$ luminosities have been properly corrected for:(1) extinction corrections (using Balmer decrement ratio); (2) absorption line contamination (using moderate resolution spectroscopy) and (3) aperture correction (using photometric measurements). The latter correction could potentially lead to major errors. Indeed it assumes that the fraction of the emission line light which is sampled by the slit equals that of the continuum light derived from broad band photometry, which could be wrong.  However, our extinction estimate from the H$\\beta$/H$\\gamma$ ratio does not depend on such assumption and agrees well with our estimate from the H$\\alpha$/H$\\beta$ ratio. Moreover, for the 8 distant galaxies, the aperture correction is generally close to 1, and so we believe that uncertainties related to aperture correction are marginal. IR luminosities have been estimated by (1) using the correlation between MIR and far-IR luminosities and (2) using the correlation between radio and far-IR luminosities. A good agreement is found using the two methods.  Despite the large error bars on individual SFRs, we find a remarquable trend in the relationship between the two SFR estimates: (1) For starbursts with SFR$_{IR}$ below $90-130M_{\\odot}/yr$, SFR estimates based on H$\\alpha$ luminosities, properly corrected, agree well with SFR$_{IR}$, in agreement with analyses of local galaxies with low SFRs (Buat et al., 2002, Rosa-Gonzalez et al., 2002); (2) At larger IR luminosities, ${H\\alpha}$ emission lines underestimate the SFR by factors up to 2.5  at SFR$_{IR}$=$250M_{\\odot}/yr$. (3) few objects show properties strongly discrepant from the empirical relationship between SFR$_{H\\alpha}$ and SFR$_{IR}$ which could be due to their individual star formation history: for example, if the star formation was rapidly decreasing (increasing), one would expect SFR$_{IR}$ larger (smaller) than SFR$_{H\\alpha}$. Our result is in excellent agreement with the recent study of the IRAS sources detected by the SDSS (Hopkins et al, 2003, see their Fig 15) in the local universe. However it substantially differs from that of Franceschini et al (2003, hereafter F03), who have studied similar sources in the HDFS. Indeed in the F03's Fig. 8, IR and $H\\alpha$ estimates of the SFR are almost similar for LIRGs and ULIRGs, while for most of the less luminous IR galaxies, infrared provide larger values of the SFRs. Although the above results are somewhat in contradiction with our and Hopkins et al studies, one has to understand the major reasons of such a discrepancy. F03's Table 5 displays extinction corrected SFR(Halpha) values for 11 objects. For 7 of them, they have estimated dust extinction from the stellar continuum on the basis of the rest-frame V-K colors or on the comparison between H$\\alpha$ and rest-frame 2800AA fluxes; it is generally believed that the continuum can only provide gross estimates of the gas extinction, because it is strongly affected by the various and complex star formation histories.  For the 4 resting objects (HDFS25, 27, 53 and 55), the discrepancy might be related to the fact that F03 have not corrected H$\\beta$ fluxes for underlying absorption, leading to a systematic underestimate of H$\\beta$ emission flux. Moreover 2 of these objects (HDFS25, 55) present a low S/N for the H$\\alpha$ line (see R00 Fig. 1).  Using both SFR's, we have estimated an empirical polynomial relationship SFR$_{IR}$ = SFR$_{H\\alpha} \\times$ (A$\\times$ SFR $_{H\\alpha}$ + 1), where A=3.11$^{6.0}_{0.4}$ $10^{-3}$ (which fits at best the points in Figure 2b, weighted according to the error bars).  Applying our empirical relationship to the L$_{IR}$ estimated by Flores et al.  (1999) on field galaxies, one can derive in principle the H$\\alpha$ luminosity density corresponding to the IR luminosity density at z $\\le$ 1 (note that this result is insensitive to a change of cosmological constants). It results that the $H\\alpha$ star formation density corresponds to 83\\% to 89\\% of the IR star formation density, for the Flores et al. lower and upper limit, respectively.  This suggest that galaxies which actually dominates the cosmic star formation density at z$\\le$ 1 are not ULIRGs, but correspond to the much more abundant population of LIRGs and starbursts. For determining the SFR, MIR fluxes which can be measured rapidly by satellites (ISO, SIRTF) give, and at least for the intrinsically bright IR galaxies, more accurate estimates than cumbersome optical and NIR spectroscopy. However, the detailed spectra bring useful information on metal abundances and on the evolutionary history of the galaxies, as we will show in subsequent papers.  {\\it Acknowledgements:} We thank Rafael Guzm\\'an for enlightening discussions which greatly contributed to improve our paper. %"
    },
    "0311/astro-ph0311325_arXiv.txt": {
        "abstract": "We carry out a maximum-likelihood kinematic analysis of a sample of $\\nbhb$ blue horizontal branch (BHB) stars from the Sloan Digital Sky Survey presented in \\citet{halo1} (Paper~I).  Monte Carlo simulations and resampling show that the results are robust to distance and velocity errors at least as large as the estimated errors from Paper~I.  The best-fit velocities of the Sun (circular) and halo (rotational) are $\\vysolar \\pm \\vysolarerror \\kms$ and $\\vhaloval \\pm \\vhaloerror \\kms$ but are strongly covariant, so that $v_{\\odot} - \\vhalo = \\vysolarminushalo \\pm \\vysolarminushaloerr \\kms$. If one adopts standard values for the local standard of rest and solar motion, then the halo scarcely rotates.  The velocity ellipsoid inferred for our sample is much more isotropic [$(\\sigma_r,\\sigma_\\theta,\\sigma_\\phi) =(\\sigmar \\pm \\sigmarerror, \\sigmatheta \\pm \\sigmathetaerror, \\sigmaphi \\pm \\sigmaphierror) \\kms$] than that of halo stars in the solar neighborhood, in agreement with a recent study of the distant halo by \\citet{Sommer-Larsen_etal97}. The line-of-sight velocity distribution of the entire sample, corrected for the Sun's motion, is accurately gaussian with a dispersion of $\\sigmalosval \\pm \\sigmaloserror \\kms$. ",
        "introduction": "Among all visible components of the Galaxy, the stellar halo is the oldest and occupies the largest volume. It is therefore of interest both as a record of the early history of the Galaxy and as a tracer of the Galactic potential.  The kinematics of halo stars in the solar neighborhood are well established, as shown in Table~\\ref{t:localhalo}. For present purposes, the solar neighborhood extends as far as a few kpc from the Sun, but to a distance less than the Sun's own Galactocentric radius. Evidently, the systemic rotation $\\vhalo$ of the local halo is small or zero, and the velocity ellipsoid is distinctly elongated in the radial direction.  The more distant halo has been less well studied.  Proper motions are generally not measurable for distant stars, so that line-of-sight velocities must be analyzed under \\emph{a priori} assumptions about the symmetries of the velocity ellipsoid. The stars are also fainter, reducing the accuracy of the measured quantities.  Much work on the distant halo has been done using globular clusters, which have the advantage of being distinctive and bright, but have important drawbacks: there are only $\\sim 10^2$ of them, so that statistical uncertainties are large, and special conditions of formation or selective destruction by Galactic tides may have biased the properties of the cluster system compared to those of the field.  \\citet{Frenk_White80} analyzed the kinematics of a sample of 66 globular clusters under the assumption that the systemic rotation $\\vhalo$ is constant; they concluded that the residual velocities are consistent with being isotropic within broad uncertainties.  For a subsample of 44 metal-poor (``F'') clusters, they found $\\vhalo=42\\pm33\\kms$ and a root-mean-square (rms) residual line-of-sight velocity $\\sigmalos=124\\kms$.  In a similar analysis of clusters with $\\mbox{[Fe/H]}<-0.8$, \\citet{Armandroff89} found $\\vhalo=46\\pm29\\kms$ and $\\sigmalos =116\\pm10\\kms$.  The restriction to low metallicities follows the discovery by \\citet{Zinn85} that the more metal-rich clusters have a distribution and kinematics intermediate between those of the stellar halo and thin disk.  \\citet{Ratnatunga_Freeman89} sampled the halo field using relatively faint red giants at high Galactic latitude.  Their summary for the more metal-poor stars near the south Galactic pole concludes that $\\sigma_z$ is approximately constant at $\\sim 75\\kms$ out to $25\\kpc$, which would be consistent with the solar neighborhood (Table~\\ref{t:localhalo}) if the velocity ellipsoid is constant in cylindrical coordinates; however, for their most distant subsample near the SGP ($\\langle D\\rangle=17.1\\kpc$), they measured $\\sigmalos =42\\pm10$ (10 stars).  In a study similar to that reported here, \\citet{Sommer-Larsen_etal97} analyzed a sample of 679 blue horizontal branch stars drawn from multiple sources and spanning Galactocentric radii $\\sim 7-65\\kpc$.  With line-of-sight velocities and photometric distances, they found that the radial velocity dispersion (in spherical coordinates) falls rather sharply from its value in the solar neighborhood, $\\sigma_r=140\\pm10\\kms$, to an asymptotic value $89\\pm 19\\kms$ at $r\\gtrsim 20\\kpc$.  By requiring dynamical equilibrium in an assumed logarithmic Galactic potential, they concluded that the tangential velocity dispersion should rise from its local value to $\\sigma_\\theta\\approx\\sigma_\\phi\\approx 137\\kms$ at large $r$.  They did not claim, however, to have detected this rise directly in their data.  They further suggested that the shift from radial to tangential anisotropy with increasing Galactocentric radius might support theories in which the halo formed by heirarchical merging of smaller systems.  Because of the differences in kinematics cited above, it appears that local high-velocity, low-metallicity stars may not be entirely representative of the more distant halo. To study the more distant halo, we have isolated a sample of $\\nbhb$ blue horizontal branch (BHB) stars from the Sloan Digital Sky Survey \\citep[][hereafter Paper I]{halo1}.  In that paper, we estimated contamination from foreground stars and blue stragglers to be less than $10\\%$ for bright stars ($g<18$) and about $25\\%$ for faint stars ($g>18$).  An improved color-absolute magnitude relation was also derived, and these absolute magnitudes are used throughout this paper to derive distance moduli. This sample of BHB stars provides effective tracers of the Galactic halo, with no kinematic bias.  The sample may have selection bias, and is probably far from complete, but these biases do not affect the kinematic analyses presented in this paper other than degrading the statistics and increasing the errors. Using maximum-likelihood methods, we fit for the mean rotation and principle axes of the halo velocity ellipsoid on the assumption that these are constant in spherical or cylindrical coordinates.  We also fit for the orbital velocity of the Sun.  The present paper is organized as follows. We first discuss an overview of the kinematic analysis in \\S\\ref{s:ve_overview} and appendix~\\ref{s:maxlike}. The results for the kinematic parameters of the halo are presented in \\S\\ref{s:ve_results} and summarized in Table~\\ref{t:results}. Monte Carlo simulations which validate the analysis are presented in \\S\\ref{s:montecarlo}. \\S\\ref{s:discussion} summarizes our main conclusions and discusses the agenda for future work.  \\subsection{Terminology used in this paper}\\label{s:terminology} The following terminology and conventions are used throughout this work. A heliocentric right-handed Cartesian coordinate system $(x,y,z)$ is oriented so that $+x$ points towards the Galactic center (GC) and $+z$ points towards the North Galactic Pole, such that $+y$ is the direction of disk rotation.  Thus, $\\vsun = (v_{\\odot,x},v_{\\odot,y},v_{\\odot,z}) = \\vlsr + (U,V,W)$ where $(U,V,W)$ is the solar motion with respect to the local standard of rest (LSR). The velocity of the LSR is taken to be $\\vlsr = (0,220,0) \\kms$ \\citep{kerr_lyndenbell_1986}, and the canonical solar motion vector is $(U,V,W) = (10,5,7) \\kms$ \\citep[rounded to integer values]{dehnen_binney_1998}. Galactocentric spherical $(r,\\theta,\\phi)$ and cylindrical $(R,\\phi,z)$ coordinate systems will also be used. We make a rudimentary model of halo bulk rotation parametrized by a single velocity $\\vhalo$. A positive $\\vhalo$ indicates a flat halo rotation curve in the same sense as the disk rotation. Radial velocities to stars are measured as redshifts which are subsequently and immediately corrected for the Earth's motion around the Sun and the Earth's rotation about its own axis.  This heliocentric radial velocity $v_r$ contains a component due to the solar velocity $\\vsun$.\\footnote{Here the subscript ``$r$'' is used for components along the line of sight, but elsewhere it will be used for the radial component in spherical coordinates based on the Galactic Center.  The context should make our meaning clear.} The ``line-of-sight'' velocity $\\vlos$ will be defined as $v_r$ plus the projection of $\\vsun$ onto the line-of-sight to the star, minus the projection of its expected rotation component $-\\vhalo \\boldsymbol{\\hat{\\phi}}$ (if any).  That is, $\\vlos$ is the component of the purely random component of velocity in a model of the halo with a flat rotation curve, projected onto the line of sight.  We find $v_{\\rm los}$ conceptually useful, but one should not forget that it is contingent upon our assumptions concerning the rotation of the Sun and the halo, whereas $v_r$ is directly observed (modulo errors). Throughout, the distance from the Sun to the GC is assumed to be $8 \\kpc$ \\citep{reid_1993}. ",
        "conclusions": "\\label{s:discussion}  Applying maximum-likelihood methods to the line-of-sight velocities and photometric distances of BHB stars, we have modeled the kinematics of the relatively distant galactic halo.  (The median distance of the sample is about $25\\kpc$.) No significant anisotropy of the velocity ellipsoid is detected.  In a model aligned with spherical coordinates, for example, $\\boldsymbol{\\sigma} = (\\sigmar,\\sigmatheta,\\sigmaphi) \\kms$. However, the errors in $\\sigma_\\theta$ and $\\sigma_\\phi$ are much larger, $\\sim 15-20\\kms$, than that of $\\sigma_r$, because of the great distances of most stars: that is, the Galactocentric radius $\\scriptr$ of most stars is more than twice that of the Sun, so that the direction of the line of sight is not far from radial despite the broad angular range of our sample.  This result contrasts starkly and surprisingly with the radial anisotropy of halo velocities in the solar neighborhood ($\\boldsymbol{\\sigma} \\sim (150,100,100) \\kms$; see Table~\\ref{t:localhalo}). Our results are consistent, however, with the only other comparable study of distant halo field stars: \\cite{Sommer-Larsen_etal97}.  We are able to fit separately for the rotation speed of the halo ($\\vhalo$) and the orbital velocity of the Sun, although the two are strongly covariant.  The halo rotation is marginally consistent with zero at the one-sigma level, which is $\\sim 20\\kms$.  The difference $v_{\\odot,y}-\\vhalo=\\vysolarminushalo\\pm\\ \\vysolarminushaloerr \\kms$ is more tightly constrained than either quantity separately. This is consistent with a non-rotating halo if one accepts the IAU recommendation for the local standard of rest, $\\vlsr=220\\kms$ \\citep{kerr_lyndenbell_1986} and \\citet{dehnen_binney_1998}'s solar motion, $\\vsun-\\vlsr=(10,5,7)\\kms$, although some authors have advocated a substantially smaller \\citep{olling_merrifield98} or larger \\citep{uemura_etal00} value of the azimuthal component of $\\vlsr$.  (If the solar rotation \\emph{is} substantially less than $225\\kms$, then according to our results, the halo should counterrotate!) Our result is also statistically consistent with most of the studies cited in Table~\\ref{t:localhalo}, although the latter seem to prefer a slight positive halo rotation. Note that what those studies actually measure is the mean velocity of local halo stars relative to the Sun; the values for $v_{halo}$ in the fourth column of that table have been adjusted to $V_\\odot=225\\kms$ so that they can be compared directly to our result. Presumably, $\\vlsr$ will soon be determined with exquisite accuracy by a variant of the classical visual-binary method applied to stars orbiting Sgr~$\\rm A^*$, but accurate values are not in the public domain at the time of writing \\citep{ghez_etal03,salim_gould99}.  Of course, the measurement of $\\vhalo$ depends upon the assumption that the halo rotation curve is flat; line-of-sight velocities would be unaffected by adding a solid-body component to the angular velocity of the halo, although proper motions would be (see below).  These are the main results of our study.  It is likely that more science could be extracted from this sample, and completion of the SDSS will allow it to be enlarged.  For the present, we simply highlight certain interesting features of the data that have not yet been thoroughly explored.  The distribution of $\\vlos$ is plotted in figure~\\ref{f:vlos_dist} and beautifully follows a gaussian of width $\\sigmalos = \\sigmalosval \\pm \\sigmaloserror \\kms$. The uncertainty quoted here is one-sigma statistical, to which should be added a similar systematic uncertainty pending better quantification of our measurement errors.  An estimate of this systematic error is easily determined as $\\sigmalos - \\sqrt{\\sigmalos^2 - \\Delta v^2} \\approx 3.4 \\kms$ ($\\Delta v = 26 \\kms$ from Paper~I). The line-of-sight velocities (figure~\\ref{f:vlos_dist}) are about as gaussianly distributed as the number of stars would allow.  This is consistent with (although not required by) the isotropy of the velocity ellipsoid that we find, since a strongly anisotropic ellipsoid, even if gaussian along each individual axis, would not produce an exact gaussian in an average over many lines of sight.  Gaussianity could result from large observational errors, regardless of the intrinsic velocity distribution, if the errors were gaussian, but our errors seem not to be large enough to dominate the observed distribution (Paper~I).  By contrast, the distribution of transverse velocities of halo stars in the solar neighborhood is distinctly nongaussian, with a negative kurtosis \\citep{popowski_gould_1998}.  This is another indication that the local halo differs from the more distant halo probed by our BHB sample.  Another aspect of the data that could be thoroughly explored is the smoothness of the BHB phase-space distribution.  The notion of a stellar velocity ellipsoid implicitly presumes that phase space is smoothly populated.  If instead the distribution is very lumpy and the number of important lumps or streams is small, then statistical characterizations of \\emph{individual} halos, in particular that of the Galaxy, may not be meaningful.  In cosmology, however, statistical descriptions are unavoidable. Current theory holds that galactic halos have been built up by accretion of lesser systems \\citep{searle_zinn_1978,bullock_kravtsov_weinberg_2001}. Accreted gas joins galactic disks, and the details of its history are effaced by dissipation.  Stars that form in these systems before they are accreted (and presumably also their dark matter), however, behave collisionlessly and preserve a partial record of their origins.  The structure of the stellar and dark-matter halos should be related. \\citet{helmi_white_springel_03} conclude from a detailed analysis and scaling of cosmological simulations that the number of streams composing a halo similar to that of the Galaxy should be large, amounting to almost $10^4$ within $25\\kpc$, and that most streams joined the Galaxy early on, allowing more time for phase mixing.  The number of our BHB stars is much smaller than this, so that one might expect to find a smooth distribution. And yet we see structure in our sample that is consistent with the tidal stream of the Sagittarius Dwarf (Paper~I), as previously delineated in SDSS data \\citep{ivezic_etal_2000,yanny_etal_2000,ibata_etal_2001a, newberg_etal_2002} and in a survey of carbon stars \\citep{ibata_etal_2001b}.  Most of our results for the halo velocity ellipsoid are scarcely affected by excluding that part of the sample that lies close to the Sagittarius stream on the sky. In short, it is not yet clear whether the kinematics of the Galactic halo admit a statistical decription. Larger samples of stars will be necessary to clarify this.  In Paper~I, we used proper motions only as a check on the purity of our sample.  Calibration against the QSOs shows that the proper motions of individual BHB stars in our sample are smaller than measurement errors by factors $\\sim 3-5$, but statistical detection of the kinematics of the population as a whole should certainly be possible. Among other benefits, this would allow a direct constraint on the solid-body component of halo rotation and a further test of the isotropy of the velocity dispersions.  Systematic errors in the proper motion data present many complexities, however, so we defer their analysis to a later paper.  The quality of the USNO-B data should be appreciably better than that of the USNO-A, and one should probably use the former to determine the proper motions of SDSS objects.  Our BHB radial velocity data could be used to constrain the mass of the outer Galaxy.  We have not yet attempted a serious dynamical analysis, partly in order to do full justice to the kinematics, but also because mass measurements rely more heavily on aspects of the data that are less certain: namely, the photometric distances and the completeness as a function of magnitude. However, since no kinematic information was used in the selection of the sample, we may assume that the velocities fairly represent the whole population of BHB stars, whose spatial density is often described as a powerlaw in Galactocentric radius, $n\\propto \\scriptr^{-\\alpha}$, with exponent $\\alpha\\approx 3.5$ \\citep{kinman_suntzeff_kraft_1994}. In that case, to the extent that the dispersions are isotropic and constant with radius and rotation is negligible, and that the potential is spherical, Jeans' equations imply \\begin{eqnarray}\\label{e:Jeans} v_{\\rm circ} &\\approx& 187\\left(\\frac{\\alpha}{3.5}\\right)^{1/2}\\left(\\frac{\\sigma}{100\\kms}\\right),\\\\ M(\\scriptr) &\\approx& 2.0\\times 10^{11}\\left(\\frac{v_{\\rm circ}}{187\\kms} \\right)^{2}\\left(\\frac{\\scriptr}{25\\kpc}\\right)\\,M_\\odot. \\end{eqnarray} The value of $v_{\\rm circ}$ so obtained, with $\\sigma \\approx 100 \\kms$, is substantially lower than the circular velocity in the vicinity of the Sun.  Faced with similar evidence drawn from their own independent BHB sample, \\cite{Sommer-Larsen_etal97} concluded that the tangential velocity dispersion of the halo probably rises beyond $\\scriptr\\sim 20\\kpc$ so as to exceed the radial dispersion and be consistent with a constant circular velocity.  At such distances, however, the tangential dispersions are very poorly constrained by line-of-sight velocities without \\emph{a priori} assumptions about the potential (see Table~\\ref{t:results}).  Rather than make such assumptions, one may prefer to await further direct evidence."
    },
    "0311/astro-ph0311055_arXiv.txt": {
        "abstract": "We present a catalog of 49 spectroscopic strong gravitational lens candidates selected from a Sloan Digital Sky Survey sample of 50996 luminous red galaxies. Potentially lensed star-forming galaxies are detected through the presence of background oxygen and hydrogen nebular emission lines in the spectra of these massive foreground galaxies.  This multiline selection eliminates the ambiguity of single-line identification and provides a very promising sample of candidate galaxy-galaxy lens systems at low to intermediate redshift, with foreground redshifts ranging from 0.16 to 0.49 and background redshifts from 0.25 to 0.81. Any lenses confirmed within our sample would be important new probes of early-type galaxy mass distributions, providing complementary constraints to those obtained from currently known lensed high-redshift quasars. ",
        "introduction": "In its strongest form, gravitational lensing produces unmistakably distorted, amplified, and multiplied images of distant astronomical objects.  It is therefore not surprising that the majority of known galaxy-scale gravitational lens systems have been discovered through imaging observations. However, a small number of lenses have been discovered spectroscopically, with the spectrum of a targeted galaxy showing evidence of emission from a background source and follow-up imaging revealing lensing morphology.  The three most secure examples are the lensed quasars 2237+0305 \\citep{huch85} and SDSS J0903+5028 \\citep{john03} and the lensed Lyman-$\\alpha$ emitting galaxy 0047$-$2808 \\citep{w96, w98, w99}.  Several authors have made predictions for the frequency of lensed quasar discoveries in galaxy redshift surveys \\citep{koc92, mw00, mw01}. Others, inspired by the discovery of 0047$-$2808, have undertaken spectroscopic searches for lensed Lyman-$\\alpha$-bright galaxies \\citep{h00, wil00, hall00, behs}. The idea behind such searches is that a massive foreground galaxy should act as an effective gravitational lens of {\\em any} object positioned sufficiently far behind it and at small enough impact parameter, and any emission features from such lensed objects should be detectable in the spectra of the foreground galaxy.  Therefore a search for discrepant emission features in galaxy spectra can potentially lead to a sample of gravitational lens systems that would not be discovered in broadband imaging searches due to faintness of source relative to lens. This ``lenses-looking-for-sources'' approach is complementary to lens searches such as the recently completed Cosmic Lens All-Sky Survey \\citep[CLASS;][]{class1, class2} which proceed by targeting sources and looking for evidence of an intervening lens.  With its massive scale and quality of data, the spectroscopic component of the Sloan Digital Sky Survey (SDSS) provides an unprecedented opportunity for spectroscopic galaxy-galaxy gravitational lens discovery. This paper presents the first results of such a search within a sample of $\\sim$51,000 SDSS Luminous Red Galaxy (LRG) spectra \\citep[hereafter E01]{lrg}: a catalog of candidate lensed star-forming galaxies at intermediate redshift. The lens candidates presented in this paper were detected by the presence of not one but (at least) three emission lines in the LRG spectra identified as nebular emission from a single background redshift: \\oii\\ and two out of the three of \\hb, \\oiiia, and \\oiiib\\ (all emission-line wavelengths in this paper are vacuum values, in Angstroms).  This implies a maximum redshift of $z \\sim 0.8$ for any candidate lensed galaxies: at higher redshifts, [O~{\\sc iii}] emission moves off the red end of the SDSS spectrograph.  The lensing cross section of a particular foreground galaxy is lower for intermediate-redshift sources than high-redshift sources, and in this sense the lens search we describe here is at a disadvantage relative to searches for lensed high-redshift Lyman-$\\alpha$ emitters.  However, the identities of the emission lines in the sample we present here are absolutely unambiguous.  This cannot be said for spectroscopic Lyman-$\\alpha$ lens candidates, which are typically detected as single discrepant emission lines and are difficult to distinguish from lower-redshift emission or (in the case of a huge survey such as the SDSS) non-astrophysical spectral artifacts. (A forthcoming paper will present SDSS single-line [O~{\\sc ii}] and Lyman-$\\alpha$ lens candidates.) Our spectroscopic lens search based on multiple-line detection has further advantages. First, the source redshift of any lensed galaxies will be known from the outset.  Second, for a given limiting line flux, star-forming galaxies at intermediate redshift are more numerous on the sky than Lyman-$\\alpha$ emitters at high redshift \\citep[see][]{hip03, mai03}, and hence strong lensing events could be more frequent despite decreased lensing cross sections. The non-lensed source population should thus be more amenable to study and characterization \\citep[for example]{hogg98}, facilitating lens statistical analysis.  Finally, intermediate-redshift lensed galaxies would probe the mass distribution of the lens population in a systematically different manner than do high-redshift sources.  The search for new lens systems is motivated by the unique power of gravitational lensing to constrain the mass distribution of the lensing galaxy. \\citet{ml92} argued that strong galaxy-galaxy gravitational lenses in the optical band should be both ubiquitous and of great scientific interest, but their detection has proved difficult.  In addition to the optical Einstein ring 0047$-$2808 mentioned above, \\citet{h00} have published one more spectroscopic galaxy-lens candidate. {\\sl Hubble Space Telescope} ({\\sl HST}) imaging has also had some success in detecting lensed galaxies, as described by \\citet{cr02}, \\citet{rgo99}, and \\citet{fas03}, but spectroscopic confirmation remains a challenge in most cases. Lensed galaxies are especially useful for investigating the mass distribution of the lensing galaxy because unlike lensed quasars they are spatially resolved sources which provide more detailed constraints \\citep[for example]{kn92}. As with lensed quasars, the relative angular deflections between multiple images can constrain the total deflecting lens mass. But the relative distortions of multiple images can provide constraints as well, probing the slope of the radial mass profile and allowing more detailed connection with theories of galaxy formation and Cold Dark Matter (CDM). The promise of our candidate sample is enhanced by the fact that lens {\\em and} source redshifts for any confirmed lens systems will be known from the outset with great precision; this is generally not the case for lenses discovered in optical imaging or radio data.  The lack of sufficient lensing constraints on the radial slope of the galaxy mass profile in lensed quasar systems is the greatest source of uncertainty in the Hubble constant ($H_0$) values derived from the time delays between images \\citep{sch00}, which remain systematically low compared to the current best $H_0$ values from the distance ladder \\citep{fre01} and cosmic microwave background \\citep{spe03}; \\citet{koc02a, koc02b, koc03} highlights the importance of resolving this discrepancy. Any galaxy-galaxy lenses confirmed within our current list of candidates would help to solve this problem by expanding the sample for ensemble lensing studies of early-type galaxy mass profiles like that of \\citet{rus03} and allowing for the inclusion of distortion information as mentioned above. The generally lower lens and source redshifts of our candidate systems as compared to known lens systems are also of particular interest because the images of any strong lenses within our sample will tend to form at smaller {\\em physical} radii within the lens galaxy than in current lens systems \\citep[see][Table 1]{lensevol}, thus probing elliptical galaxy mass profiles in a systematically different manner\\footnote{The lowest-redshift lensed quasar currently known, RXJ 1131$-$1231, has $z = 0.658$ \\citep{slu03}.}. A large enough sample of strongly lensed star-forming galaxies could even be monitored directly for supernovae.  Multiple images of the same supernova would give time delays and $H_0$ measurements, as numerous authors have envisioned \\citep{ref64, pm00, holz01, goo02, ok03}. If we assume one supernova per galaxy per century as an order-of-magnitude estimate \\citep{db85}, a program to monitor $\\sim 50$ lensed galaxies for several years would have a good chance of success.  Numerous studies have addressed the inability of standard macroscopic galaxy mass models to predict the flux ratios observed between images of strongly lensed quasars. Some authors have suggested that this phenomenon may be due to the effects of substructure in the lensing galaxy on the scale of either CDM sub-halos \\citep{ms98, mm01, chi02, dk02} or stars \\citep{wms95, sw02}, while others have invoked macroscopically modified lens models as a possible explanation \\citep{ew03, mhb03}. A sizeable sample of strongly lensed galaxies would provide valuable insight into this problem: by virtue of the arcsecond-scale intrinsic angular size of the source, galaxy-galaxy lenses will be immune to the lensing effects of any substructure with lensing deflection angles on the scale of microarcseconds (stars) or milliarcseconds (sub-halos).  If the suppressed saddle point image fluxes seen in many quasar lenses are indeed due to macroscopic inadequacies in the standard lens models, similar effects should be seen in galaxy-galaxy lenses.  If on the other hand these effects are due to micro/millilensing by substructure, they should be absent from lensed galaxy systems. ",
        "conclusions": "We have presented a catalog of candidate strong galaxy-galaxy gravitational lens systems detected spectroscopically within the Sloan Digital Sky Survey.  These systems have known foreground {\\em and} background redshifts; only the detailed spatial alignment of foreground and background galaxies remains unknown. We plan to conduct integral-field spectrographic follow-up observations of these systems, which will allow us to resolve the spatial distribution of the background nebular line emission that we have detected. These observations should confirm a substantial number of lenses within our candidate sample, with many lensed galaxies at lower redshift than any other currently known lensed extra-galactic optical sources. Any lenses confirmed within our sample will be of considerable interest for the study of early-type galaxy mass distributions, and could have implications for lens-time-delay $H_0$ measurements and anomalous quasar-lens flux ratios. Our sample demonstrates the feasibility of the emission-line-based spectroscopic lens search technique within the SDSS and other redshift surveys, and we plan to extend the search to higher source redshifts in the near future. We have also developed and applied a method for abating the influence of night sky emission-line residuals in the 7000--9000-\\AA\\ range that allows us to detect [O~{\\sc iii}] and H$\\beta$ emission over the redshift range $z \\sim$ 0.4--0.8 without an excess of false-positive detections."
    },
    "0311/astro-ph0311263_arXiv.txt": {
        "abstract": "We present detections of the Sunyaev-Zeldovich Effect (SZE) at 150GHz and 275GHz for the X-ray luminous z=0.299 cluster 1E0657-67. These observations were obtained as part of an X-ray, weak lensing, and SZE survey of nearby X-ray bright clusters. The SZE maps were made with the ACBAR (150, 210, 275 GHz) bolometer array installed at the Viper telescope located at the South Pole. We also present preliminary results from a blind SZE cluster survey. ",
        "introduction": "We are undertaking a multiwavelength study of the physical properties of a sample of X-ray bright galaxy clusters. These clusters were selected from the REFLEX survey of southern X-ray clusters (~\\cite{2001A&A...369..826B}). They conform to the following criteria: L$_x > 4\\times 10^{44}$ erg/s [0.1-2.5 keV] and $\\delta < -44^{\\circ}$. In this study, we will combine new and archival X-ray data with weak lensing observations and new Sunyaev-Zeldovich Effect (SZE) measurements.  The rationale behind combining these multiwavelength observation is the fact that they probe the cluster components in different ways. For instance, X-ray and SZE data probe the gas density and temperature distributions, whereas weak lensing observations probe the total mass distribution. The combination of X-ray and SZE data provides three observables (X-ray surface brightness, SZE decrement, and projected temperature) which constrain the variation of density and temperature throughout the target cluster. Comparing weak-lensing mass distributions with X-ray and SZE gas distributions in the same clusters allows a detailed test of usual assumptions (hydrostatic equilibrium, spherical symmetry, and the lack of significant substructure) on which estimates for the amount and distribution of dark matter in clusters are based.  This study will attempt to:  1) characterize the density, temperature, and dynamical state of the gas in each cluster. 2) compare the mass estimates derived from the SZE and X-ray observations with the more direct weak lensing mass. This study is paramount to determine whether, and to what level, future cosmological parameter estimates from SZE and X-ray cluster surveys might be biased by non-thermal cluster physics and/or by non standard cluster geometries (e.g., \\cite{2001ApJ...560L.111H}).   \\subsection{Sunyaev-Zeldovich Effect Observations with ACBAR}  The SZE describes the inverse Compton scattering of Cosmic Microwave Background (CMB) radiation photons by hot electrons in the intra-cluster plasma.  There are two SZ effects: the kinetic and thermal effects.  The thermal effect is due to the random thermal motion of the cluster electrons and shifts lower energy CMB photons to higher energy while the kinetic effect is due to the bulk flow motion of the cluster. Since photon number is conserved in the process, the net effect of the thermal effect is to distort the CMB spectrum such that there is a null at about 215 GHz, an increment at shorter wavelengths, and a decrement at longer wavelengths. This spectral distortion allows multi-frequency instruments that span this null to separate the SZE from intrinsic CMB anisotropies. Except at frequencies close to the thermal null, the thermal effect dominates the kinetic effect by an order of magnitude.   We have performed cluster SZE observations with the Arcminute Cosmology Bolometer Array Receiver (ACBAR) mounted at the Viper telescope (~\\cite{2000ApJ...532L..83P}). This millimeter/submillimeter telescope consists of 4 mirrors arranged in an off-axis configuration. The radiation from an object collected by the 2.15m diameter primary and the secondary (which form an aplanatic Gregorian system) is reflected by a flat chopping tertiary mirror and focused by a hyperbolic condenser onto the focal plane of the telescope. Note that the tertiary mirror is located at the exit pupil of the system. Therefore, any motion of the chopping mirror is equivalent to a motion of the primary mirror. This allows us to perform fast raster scans of up to 3 $^{\\circ}$ in azimuth at a fixed elevation.  The ACBAR instrument (\\cite{run1}) consists of 16 micro-lithographed `spider-web' bolometers arranged in a $4\\times4$ array.  The detectors are spaced by $\\sim$ 16 arcmin and have beams that can be modeled as $\\sim$ 4.5 arcmin FWHM Gaussians (\\cite{kuo1}). A more extensive review of the instrument and its capabilities can be found in Runyan et al. (2003). For the 2001 season, observations were made at 150, 220, 275, and 345 GHz (band widths of 30, 30, 50, and 25 GHz respectively), with a column of four detectors per frequency.  These frequencies take full advantage of the atmospheric windows available at the Pole. They bracket the null in the thermal SZE spectrum and allow ACBAR to produce both SZE decrement and SZE increment images. For the 2002 observing season, ACBAR was modified so that each row performed observations at the same frequency. In addition, the number of 150GHz filters was doubled, by removing the 345 GHz channels.    During an observation, the chopping motion of the tertiary mirror sweeps each bolometer beam across the sky in the azimuth direction at an approximately constant elevation. The design of the telescope keeps the illumination pattern on the primary almost constant during the chopping; however, this is not the case over the secondary mirror. This introduces a systematic offset into the data because the emissivity over the secondary is not constant due to snow accumulation, ice, and temperature variations. To remove these offsets, we used a Lead-Main-Trail technique. The method is as follows: The main field sweeps (where the target is centered) are led and trailed by similar sets of raster scans (lead and trail scans) of the sky at the same elevation but spaced in azimuth by half a degree. The lead and trail observations are then subtracted from the main observation to correct for the offset. Note that the observing time in the lead and trailed stares is half the exposure time of the lead stare (typically, the Lead and Trail fields are observed for 30 sec while the Main field is observed for 60 seconds). This guarantees that the L-M-T observing sequence is performed fast enough so that we minimize the effects of any temporal variation in the offsets.  Note that for the cluster observations, we set the chopper throw to be $\\sim$ 1.5 degrees. With a sampling rate of $\\sim$ 300 Hz, this produced. $\\sim$ 200 samples per arcmin per stare. In order to make a 2d-map of a cluster, we take several stares (set of L-M-T field observations) spaced in elevation by 1 arcmin such that we cover the cluster in its entirety (e.g., for a $z=$0.05 cluster, we map a region that is 1 degree in elevation). In addition, we combine several (30 $\\rightarrow$ 50) of these raster scans in order to integrate long enough to build up signal to noise. ",
        "conclusions": ""
    },
    "0311/astro-ph0311277_arXiv.txt": {
        "abstract": "Multifrequency observations from the VLA, VLBA and OVRO Millimeter Array of a major radio outburst of \\cyg\\ in 2001 September are presented, measuring the evolution of the spectrum of the source over three decades in frequency, over a period of six days.  Following the peak of the flare, as the intensity declines the high-frequency spectrum at frequency $\\nu$ steepens from $\\nu^{-0.4}$ to $\\nu^{-0.6}$, after which the spectral index remains at this latter terminal value; a trend previously observed but hitherto not satisfactorily explained.  VLBA observations, for the first time, track over several days the expansion of a sequence of knots whose initial diameters are $\\sim 8$\\,milliarcseconds.  The light-crossing time within these plasmons is of the same order as the time-scale over which the spectrum is observed to evolve.  We contend that properly accounting for light-travel time effects in and between plasmons which are initially optically thick, but which after expansion become optically thin, explains the key features of the spectral evolution, for example the observed timescale.  Using the VLBA images, we have directly measured for the first time the proper motions of individual knots, analysis of which shows a two-sided jet whose axis is precessing.  The best-fit jet speed is $\\beta \\sim 0.63$ and the precession period is $\\sim 5$ days, significantly lower than fitted for a previous flare.  Extrapolation of the positions of the knots measured by the VLBA back to zero-separation shows this to occur approximately 2.5 days \\textit{after} the detection of the rise in flux density of \\cyg. ",
        "introduction": "\\label{sec:intro}  \\cyg\\ was first detected by \\citet{Gia67} using X-ray proportional counters during a 1966 rocket flight and rose to prominence with the detection of a giant radio flare in 1972 September \\citep{Gre72}.  It is a high-mass X-ray binary system located in the Galactic plane at a distance of $\\approx$10 kpc from the Earth \\citep{Pre00,Dic83}.  With Galactic co-ordinates $l=79.84$\\degree, $b=0.69$\\degree, it lies in or behind either the Perseus or the Outer Arm, depending on its exact distance from us.  The nature of the compact object is still uncertain.  The large interstellar extinction to the source precludes optical spectroscopy, making it difficult to obtain a reliable mass function and hindering identification of the companion star.  However, infrared spectra show strong, broad helium emission lines, and a lack of strong hydrogen emission.  Spectral analysis by \\citet{vanKer96} indicated that the companion is a Wolf-Rayet star of the WN7 subclass, although more recent observations by \\citet{Koc02} suggested that the subclass may in fact be WN8.  The orbital period, as inferred from X-ray and infrared flux modulations, is 4.8 hours \\citep[e.g.][]{Par72}.  This source occasionally undergoes huge radio outbursts where the flux density can increase up to levels of $\\sim20$\\,Jy.  Of order 50 outbursts with peak radio flux densities exceeding 1\\,Jy have been observed since the detection of the first such flare in 1972 September.  During one of these radio outbursts, the hard X-ray emission was found to correlate strongly with the radio emission, although the soft X-ray emission was anticorrelated with radio flux density \\citep{McC99}.  In quiescence, the soft-X-ray is found to be strongly correlated with the radio emission, and anticorrelated with the hard X-rays \\citep{Cho02}.  This soft-X-ray---radio correlation disappears immediately before and during periods of radio flaring \\citep{Wat94}.  During recent outbursts, jet-like structures have been observed at radio frequencies, although there has been some debate as to the morphology of the collimated emission.  Two-sided jets from a flare in 2000 September were detected on arcsecond scales in a north-south orientation with the Very Large Array (VLA) \\citep{Mar01}, with an inferred jet speed $v = 0.48c$. On milliarcsecond scales, a one-sided jet with the same orientation was imaged with the Very Long Baseline Array (VLBA) after a flare in 1997 February \\citep{Mio01}.  In this case, the derived jet speed was significantly higher ($v \\geq 0.81c$).  In this paper we analyse radio observations of the 2001 September outburst of Cygnus X-3 and describe some of the physical mechanisms responsible for the observed evolution of the source.  In \\S\\ref{sec:prev_work} we summarize the characteristics of previous radio flares of \\cyg.  In \\S\\ref{sec:VLA}, \\S\\ref{sec:ovro} and \\S\\ref{sec:vlba}, we present our observations from the VLA, Owens Valley Radio Observatory (OVRO) Millimeter Array, and the VLBA respectively, and in \\S\\ref{sec:vlba_analysis} we analyse the VLBA images and infer the speed, sidedness and precession parameters of the jet.  We go on to discuss the spectral evolution and the underlying physical mechanisms for the low-frequency turnover in \\S\\ref{sec:spectra}.  We then derive constraints on the magnetic field strength, electron number density and energy density in the source in \\S\\ref{sec:src_parms}, and describe the physics behind the high-frequency spectral evolution in \\S\\ref{sec:analysis}.  In \\S\\ref{sec:big_guys} the magnetic field strength and energy in the jet in this outburst of \\cyg\\ are compared to derived parameters of radio galaxies such as Cygnus\\,A and M\\,87. ",
        "conclusions": "We have presented time-sequenced milliarcsecond-resolution observations at three different frequencies of a major outburst of \\cyg, accompanied by simultaneous quasi-daily monitoring of the (unresolved) flux density of the source, with unprecedented spectral coverage of a large flare, ranging over three decades in frequency. The proper motions of individual plasmons have been directly tracked for the first time, and the expansion of those plasmons has been directly measured.  One main result is to securely determine the presence of a two-sided jet in \\cyg\\ on milliarcsecond scales following a major flare.  The southern jet appears to be approaching and the northern counterjet appears to be receding from us, being substantially de-boosted.  We derived a jet speed of $\\beta\\sim0.63$.  The jet appears to be precessing with a period of order 5.3\\,days, significantly shorter than that derived by \\citet{Mio01} for the 1997 flare.  It appears that the milliarcsecond-scale jet is not ejected immediately on the rise of the integrated flux density, but $\\sim 2.5$ days after this, a point which we explore in a forthcoming paper (Miller-Jones et al.\\ in prep.).  We believe that it is possible that the intrinsic characteristics of \\cyg\\ in the outbursts of, for example, 1997 and 2001, may have been different.  It is not clear that all outbursts of such sources follow the same patterns.  Certainly the flux densities and fading times are not the same for different events, as outlined above.  This variable, almost whimsical, nature of the source needs to be clarified.  Certain parameters would not be expected to change (e.g.\\ after accounting for precession, the inclination angle of the jet axis to the line of sight might be expected to be constant). Identical spectral, temporal, and \\textit{uv}-coverage of different events with similar radio lightcurves might help to determine whether the jet parameters are indeed constant for different flares.  Observations of the integrated flux density of the source showed an evolving spectrum similar to that seen in previous outbursts of \\cyg. The high-frequency spectrum at frequency $\\nu$ took the form $\\nu^{-\\alpha}$, where $\\alpha$ steepened from $\\sim 0.4$ to $\\sim 0.6$ and the spectral index then remained at this latter terminal value. We contend that the spectral evolution is the result of light travel-time effects in plasmons which were initially optically thick, but which became optically thin as they expanded.  Such a model can produce spectral evolution over timescales similar to those observed. Confirmation that opacity effects are important came from the integrated spectrum between 1.4\\,GHz and 74\\,MHz: this showed a turnover moving to lower frequencies with time, caused by the decreasing opacity (probably caused by either synchrotron self-absorption, free-free absorption or a low-energy cutoff in the electron spectrum).  The presence of an additional flat-spectrum core component which was responsible for the initial rise of the flare but which then faded with time was also postulated.  We derived an upper limit to the magnetic field in the plasmons of $0.15\\pm0.06$\\,G, consistent with previous estimates for similar systems.  We also placed lower limits on the minimum and maximum observed Lorentz factors of the emitting electrons, of $\\gamma=58$ and $\\gamma=321$.  Despite the unique nature of this system, \\cyg\\ does share many of the characteristics of other, better understood microquasars. Owing to its position within the Galaxy and the associated difficulties in imaging the system, \\cyg\\ still remains somewhat of an enigma.  Studying it can, however, enhance our understanding of the effects which we believe must play a part in the observed emission from microquasars.  Light travel-time effects are unavoidable in systems of this size evolving on similar timescales."
    },
    "0311/astro-ph0311041_arXiv.txt": {
        "abstract": "{We have carried out a radio-wavelength imaging survey of 45 bright-rimmed clouds (BRCs), using the Australia Telescope Compact Array to characterise  the physical properties in their ionised boundary layers. We detected radio emission from a total of 25 clouds and using a combination of Digitised Sky Survey and mid-infrared MSX 8$\\mu$m images classified the emission into that associated with the ionised cloud rims, that associated with embedded possible massive YSOs and that unlikely to be associated with the clouds at all. A total of 18 clouds display radio emission  clearly associated with the cloud rim and we determine the ionising photon flux illuminating these clouds and the electron density and pressure of their ionised boundary layers. Using a global estimate for the interior molecular pressure of these clouds we show that the majority are likely to be in pressure equilibrium and hence are currently being shocked by photoionisation-induced shocks. We identify those clouds where the predicted ionising photon flux is inconsistent with that derived from the observations and show that either the spectral types of the stars illuminating the BRCs are earlier than previously thought or that there must be additional ionising sources within the HII regions. Finally, we identify the radio sources embedded within the clouds with infrared stellar clusters and show that they contain late O and early B-type stars, demonstrating that a number of BRCs are intimately involved with high to intermediate-mass star formation. ",
        "introduction": "Bright rimmed clouds (BRCs) are found at the edges of HII regions and are the result of the interaction of OB stars with their environment. The surface layers of BRCs are photoionised and excited by nearby OB star(s), resulting in a bright optical rim located on the side of the cloud facing the OB star. The cloud material is often swept into a cometary morphology by the radiation pressure (e.g.~Bertoldi \\cite{bertoldi};  Lefloch \\& Lazareff \\cite{ll94}). Dense cores found at the head of the clouds may  shield the remaining cloud material from the UV radiation, resulting in a column or finger structure similar to that observed in the Eagle Nebula (Hester et al.~\\cite{hester96}; White et al.~\\cite{whiteeagle}). These phenomena lead to the other common names for bright-rimmed clouds such as elephant trunk nebulae, cometary globules or speck globules (Bertoldi \\cite{bertoldi}).  The photoionisation of the cloud surface has been linked to induced star formation within the clouds (e.g.~Elmegreen \\cite{e91}; Sugitani et al.~\\cite{sfo}; Sugitani et al.~\\cite{sfmo89}). In a process known as Radiative-driven Implosion (or RDI) shocks driven into the cloud by the photoionisation of its outer layers compress the molecular gas, forming dense cores at the head of the cloud and possibly triggering their collapse. The physical processes behind RDI are relatively well-understood and have been the subject of recent theoretical studies  (Bertoldi \\cite{bertoldi}; Bertoldi \\& McKee \\cite{bk90}; Lefloch \\& Lazareff \\cite{ll94}, \\cite{ll95}). The radiative-driven implosion of molecular clouds at the periphery of HII regions may amount to a cumulative total of some several hundred stars per HII region (Ogura et al.~\\cite{osp02}) and perhaps 15\\% or more of the low-to-intermediate mass stellar mass function (Sugitani et al.~\\cite{sfo}).  One of the key diagnostic indicators in the RDI models is the ratio of external to internal pressure. As the surface layers of the BRC become ionised a recombination layer known as the \\emph{ionised boundary layer} (or IBL) develops on the side of the cloud facing the ionising star. If the conditions within the IBL are such that the pressure within the ionised exterior gas is greater than or equal to the interior neutral cloud pressure, photoionisation-induced shocks and a D-critical ionisation front propagate into the cloud interior, compressing and heating the molecular gas (Bertoldi \\cite{bertoldi}; Lefloch \\& Lazareff \\cite{ll94}). On the other hand, if the BRC is overpressured with respect to the ionised boundary layer then the ionisation front stalls at the cloud surface until the increasing recombination within the IBL raises the ionised gas pressure to equilibrium with the interior cloud pressure (Lefloch \\& Lazareff \\cite{ll94}). As equilibrium is reached the evolution of the cloud follows the same path as the initially underpressured case.  The RDI models show that the evolution of a BRC critically depends upon the \\emph{duration} of its UV illumination.  The pressure balance between the interior and exterior of the BRC thus allows us to discriminate between clouds that are currently affected by photoionisation-induced shocks and those clouds that are as yet relatively unaffected. The physical conditions in the ionised boundary layer are crucial to understand the RDI process and any subsequent star formation, as it is the ionised boundary layer which traces the interaction between the UV radiation from the OB star(s) and the molecular gas of the bright-rimmed cloud. However, although at least 89 BRCs have been identified as potential star-forming regions via IRAS searches (Sugitani et al.~\\cite{sfo}; Sugitani \\& Ogura \\cite{so94}) only a handful of clouds have had the conditions in their ionised boundary layers  determined or their pressure balance investigated (Thompson et al.~\\cite{thompson03a}; White et al.~\\cite{whiteeagle}; Megeath \\& Wilson \\cite{mw97}; Lefloch et al.~\\cite{llc97}).  Radio continuum observations are a valuable probe of the conditions within the ionised boundary layer. From the radio free-free emission we can determine the ionised gas pressure, the ionising photon flux impinging upon the cloud and the electron density of the ionised boundary layer (Thompson et al.~\\cite{thompson03a}; Lefloch et al.~\\cite{llc97}). Data from the NRAO VLA Sky Survey (NVSS; Condon et al.~\\cite{nvss}) may be used for BRCs north of $\\delta = -40$\\degr\\ (e.g.~Thompson et al.~\\cite{thompson03a}) but no such resource\\footnote{The Sydney University Molonglo Sky Survey (SUMSS; Bock et al.~\\cite{bls99}) has comparable resolution and sensitivity to NVSS for $\\delta \\le -30$\\degr\\ but does not yet cover a substantial portion of the southern galactic plane.} currently exists for the BRCs in the southern hemisphere survey of Sugitani \\& Ogura (\\cite{so94}). As a first step toward establishing the physical conditions in the ionised boundary layers and pressure balances for a wide sample of BRCs we present an Australia Telescope Compact Array\\footnote{The Australia Telescope Compact Array is funded by the Commonwealth of Australia for operation as a National Facility  managed by CSIRO.} (ATCA) radio imaging survey of the BRCs from the Sugitani \\& Ogura (\\cite{so94}) catalogue.  The observations and data reduction procedure are described in Sect.~\\ref{sect:obsvns}. The data from the survey are presented and analysed in Sect.~\\ref{sect:analysis}, with particular reference to the likely nature of the detected emission (i.e.~does the emission emanate from the ionised boundary layer or from embedded sources?).  In Sect.~\\ref{sect:discuss} we discuss the implications of the survey and estimate the likely ratio between external and internal pressure of the BRCs  detected in the survey, using assumptions for the internal molecular gas pressure. Finally in Sect.~\\ref{sect:concln} we present a summary of our results and conclusions. ",
        "conclusions": "\\label{sect:concln}  We have carried out a radio-wavelength imaging survey of 45 bright-rimmed clouds (BRCs) with the Australia Telescope Compact Array, with the aim of measuring the physical conditions in their ionised boundary layers. We detected radio emission from a total of 25 clouds and using a combination of optical Digitised Sky Survey images and mid-infrared MSX 8$\\mu$m images classified the radio emission into 4 types. The four types are: Type 1 bright-rimmed clouds, Type 2 broken-rimmed clouds, Type 3 embedded radio sources and Type 4 unassociated radio sources.  Type 1 sources are those where the radio emission follows the bright optical rim of the cloud and correlates closely with the MSX 8$\\mu$m emission, which is a tracer of the photon-dominated surface layers of these clouds. Eighteen Type 1 clouds were detected in the survey and we evaluated the incident ionising photon flux, electron density and ionised gas pressure for their ionised boundary layers. With this sample we have more than doubled the existing number of BRCs with known ionised boundary layer conditions (e.g.~Thompson et al.~\\cite{thompson03a}; White et al.~\\cite{whiteeagle}; Megeath \\& Wilson \\cite{mw97}; Lefloch et al.~\\cite{llc97}). Type 2 sources are those clouds where the radio and 8$\\mu$m emission is positionally correlated, but the optical appearance of the cloud resembles the well known broken-rimmed globule CG 4 (Reipurth \\cite{reipurth83}; Gonz\\'alez-Alfonso et al.~\\cite{g-acr95}). Only one Type 2 source was identified (SFO 62).  Type 3 sources are strong radio emitters that are both positionally coincident with luminous mid-infrared sources and embedded within the bright-rimmed clouds rather than located at the edge. The high infrared and radio luminosities of these objects imply that they are embedded stellar clusters containing late O and early B-type stars. Four type 3 sources were detected in the survey. Type 4 sources are radio sources that are either not positionally coincident with the optical or mid-infrared emission tracing the clouds or whose morphology does not follow the optical or mid-infrared morphology of the clouds. It is possible that a number of Type 4 sources represent free-free emission associated with the clouds but observations of higher sensitivity and resolution than our survey are required to make a definitive statement.  We draw the following conclusions from our survey:  \\begin{enumerate}  \\item Comparing the pressure of the ionised gas in the ionised boundary layer to a global estimate of the  internal molecular pressure of  BRCs indicates that the majority of the clouds are  in pressure equilibrium. It is thus highly likely that photoionisation-induced shocks are currently propagating into these clouds, in line with the predictions of the radiative-driven implosion (RDI) models of Lefloch \\& Lazareff (\\cite{ll94}) and Bertoldi (\\cite{bertoldi}).  \\item The two clouds SFO 58b and SFO 61a have the lowest ionised boundary layer pressures and are thus the most probable clouds where the interior  molecular pressure could exceed the exterior ionised gas pressure. These clouds are presumably in the evolutionary phase prior to that of pressure balance between their interior and exterior (Lefloch \\& Lazareff \\cite{ll94}). In this phase the photoionisation-induced shocks are stalled at the ionised boundary of the cloud until the pressure in the ionised boundary layer rises and equilibrates with the interior molecular pressure of the cloud.  \\item In most cases the physical properties of the ionised boundary layers are consistent with their expected ionisation from the OB stars responsible for exciting the large HII regions at the edges of which the clouds are found. Where the predicted ionising flux illuminating the clouds is greater than that derived from the observations we have shown that differences of a factor of a few can be readily explained by projection effects. For a number of clouds the predictions and observations cannot be reconciled and we suggest that there are additional ionising sources present within the HII regions. A SIMBAD database search revealed several additional OB stars but their overall contribution to the ionising flux is still insufficient to match the observations. Either the spectral types of the known stars are earlier than estimated or there are additional ionising stars located within the HII regions.  \\item Using a multi-wavelength approach combining the radio, mid-infrared and optical data, we show that the ionising star of SFO 72 is not the O7.5 star HD 115455,  as originally thought, but is more likely to be the O9.5 star HD 115071, which lies some 21\\arcmin\\ to the west of SFO 72. The non-detection of this cloud in the radio suggests that it is located at least 93.5 pc in the foreground of HD 115455. An embedded cluster is currently breaking free from the molecular gas at the tip of SFO 72 and from our radio upper limit we constrained the earliest star within this cluster to be B1, consistent with the observations of Herbst (\\cite{herbst75}).  \\item The 4 embedded Type 3 radio sources  detected in the survey are shown to be embedded stellar clusters containing late O and early B-type stars. All of the embedded sources are indentified with infrared stellar clusters found in recent 2MASS surveys (Bica et al.~\\cite{bdb03}; Bica et al.~\\cite{bdsb03}; Dutra et al.~\\cite{dbsb03}) or in 2MASS images downloaded as part of this paper (see Fig.~\\ref{fig:sfo62_2mass}). Three more BRCs are also found to be associated with stellar clusters which have either disrupted their natal molecular cloud (SFO 62) or are on the verge of breaking free from the molecular gas (SFO 72 \\& SFO 82b). Each of these clusters is sufficiently luminous in the far-infrared or radio to contain at least early B-type stars and thus we have shown that bright-rimmed clouds are implicated in the formation of high to intermediate mass-stars. However, further observations and modelling are required to investigate whether this star formation may have been induced by the external action of the OB star or stars ionising the bright-rimmed clouds. \\end{enumerate}"
    },
    "0311/astro-ph0311088_arXiv.txt": {
        "abstract": "We find from high speed photometry of the Sloan Digital Survey cataclysmic variable candidate SDSS J161033.64--010223.3 that it has a double humped modulation, which is probably orbital, with a period of 80.52 minutes, and shows strong ZZ Cet type pulsations. The largest amplitude oscillations have periods near 607, 345, 304 and 221 s. The amplitudes of some of the oscillations are seen to vary on time scales of hours, indicating modes with unresolved multiplets in which the splittings are $\\sim$ 1 d$^{-1}$. This is only the second dwarf nova found to have non-radial pulsations of its white dwarf primary. ",
        "introduction": "The white dwarf accretors in cataclysmic variable (CV) stars have outer layers that differ from those of isolated white dwarfs. Not only is there an accumulation of hydrogen-rich material, leading ultimately to a nova explosion, the accretion process itself results in heating, mixing and compression which change the thermal structure and composition profile (Townley \\& Bildsten 2003). Observations that contribute towards probing these structures are not easy to interpret -- the chemical compositions of the outer layers of the primaries of CVs are complicated by the rapid settling rate of the heavier elements (Sion 1999), the chemical abundances in nova ejecta are the end product of a mixture of dredge-up from the core and nuclear reactions during the explosion, embedded in shock-driven hydrodynamic processes. It would be of greater value to probe the outer layers by more direct means, as is possible through the application of asteroseismology to single white dwarfs (e.g., Clemens 1993).  A start in this direction was made by the discovery of the first CV to show non-radial pulsations (i.e. a CV whose primary is a ZZ Cet star) -- namely the dwarf nova GW Lib (Warner \\& van Zyl 1998). GW Lib is a low mass transfer ($\\dot{M}$) system that has only had one observed outburst (in 1983) of very large amplitude. Its oscillation spectrum is complicated and has yet to be fully unravelled (van Zyl et al.~2000), and as such it remains to make a quantitative contribution to the determination of outer structure of an accreting white dwarf. Its predominant oscillation modes have periods that lie near 236, 376 and 650 s, with amplitudes in the range 5 -- 15 mmag, and evidence for other modes, including sum and difference frequencies, and some possible periods around 1000 s and 5000 s. As with large amplitude ZZ Cet stars, the oscillation spectra are not stationary, individual modes (which may in many cases be unresolved multiplets) varying greatly in amplitude on time scales of days, months and years.  Nevertheless, the discovery of further examples of this class of CV will surely be welcomed, and that is what we present here -- a second CV in which large amplitude non-radial oscillations of the primary can be studied by the brightness modulations that they cause. In Section 2 we describe what was previously known about the star, in Section 3 we present our photometric observations and their Fourier transforms and in Section 4 we add some general comments. ",
        "conclusions": "Clemens (1993) partitioned ZZ Cetis into a `hot' group and a `cool' group and shows that the distribution and range of periods is different in the two groups. A decade later, with far more ZZ Cet stars known, this division is less convincing; there are many stars that combine oscillation modes from both groups, as do GW Lib and SDSS1610 Clemens also pointed out that similarities in the ZZ Cetis, at least among the hotter ones, probably arise because all have a similar amount of hydrogen, $\\sim 10^{-4}$ M$_{\\odot}$ in their envelopes. The primary of a CV could have any amount of hydrogen, from zero to $\\sim 10^{-4}$ M$_{\\odot}$, depending on when it last had a nova eruption.  SDSS1610 probably has rare dwarf nova outbursts of amplitude similar to WZ Sge and would therefore reach $m_V \\sim 11$ at maximum. A search through archived sky patrol plates is therefore recommended, and a watch should be kept for future outbursts. The value of following the heating effects of an outburst on the eigenfrequencies of the primary has already been pointed out (Warner \\& van Zyl 1998)."
    },
    "0311/astro-ph0311331_arXiv.txt": {
        "abstract": "Energetic particles that undergo strong pitch-angle scattering and diffuse through a plasma containing strong compressible MHD turbulence undergo diffusion in momentum space with diffusion coefficient~$D_p$. If the rms turbulent velocity is of order the Alfv\\'en speed~$v_{\\rm A}$, the contribution to~$D_p$ from slow-mode eddies is~$\\simeq (2p^2 v_{\\rm A}/9 l)[\\ln (l v_{\\rm A}/D_\\parallel) + 2\\gamma - 3]$, where $l$ is the outer scale of the turbulence, $\\gamma\\simeq 0.577$ is Euler's constant, and~$D_\\parallel$ is the spatial diffusion coefficient of energetic particles, which is assumed to satisfy~$D_\\parallel \\ll l v_{\\rm A}$.  The energy spectrum of accelerated particles is derived for this value of~$D_p$, taking into account Coulomb losses and particle escape from the acceleration region with an energy-independent escape time.  Slow modes in the $D_\\parallel \\ll l v_{\\rm A}$-limit are an unlikely explanation for electron acceleration in solar flares to energies of 10-100~keV, because for solar-flare conditions the predicted acceleration times are too long and the predicted energy spectra are too hard.  The acceleration mechanism discussed in this paper could in principle explain the relatively hard spectra of gyrosynchrotron-emitting electrons in the~100-5000~keV range, but only if $D_\\parallel \\ll l v_{\\rm A}$ for such particles. ",
        "introduction": "\\label{sec:intro}  The first part of this paper treats stochastic particle acceleration by slow modes in compressible magnetohydrodynamic (MHD) turbulence assuming efficient pitch-angle scattering. The second part presents the energy spectra of accelerated particles taking into account Coulomb losses and assuming that particles escape the acceleration region on an energy-independent time scale. The third part of the paper shows that slow-modes in the efficient-pitch-angle-scattering limit are not a viable explanation for electron acceleration in solar flares in the 10-100~keV energy range, although in principle they could explain the electron spectra at 100-5000~keV inferred from microwave gyrosynchrotron radiation.  Throughout the paper it is assumed that energetic particles propagate through a magnetized plasma containing both ``large-scale'' turbulent compressible motions that vary over distances much larger than the energetic-particle gyroradius~$\\rho_g$ and ``small-scale'' waves that vary over distances~$\\sim \\rho_g$. The small-scale waves scatter particles, causing diffusion in physical space with mean free path~$\\lambda_{\\rm mfp}$ as well as diffusion in momentum space. The large-scale motions also cause spatial and momentum diffusion.  Only those large-scale plasma motions on scales $\\gg \\lambda_{\\rm mfp}$ are considered. Particle acceleration by velocities on scales $\\gg \\lambda_{\\rm mfp}$ has been considered previously by a number of authors [e.g., Kulsrud \\& Ferrari (1971) Bykov \\& Toptygin (1982, 1990, 1993), Ptuskin (1988), Dolginov \\& Silantev 1990, Katz \\& Stehlik (1991), Chandran \\& Maron (2003)]. In the present paper, the large-scale velocities are taken to be slow-mode eddies in strong anisotropic compressible MHD turbulence. Such eddies are elongated along the local magnetic field, and involve velocities directed primarily along the local magnetic field (Lithwick \\& Goldreich 2001, Cho \\& Lazarian 2002).  Part of the motivation for considering slow modes is that under certain conditions they can contribute significantly to momentum diffusion.  An additional motivation is for completeness; by comparison the effects of fast modes and Alfv\\'en modes are relatively well understood (e.g., Jokipii 1966, Kulsrud \\& Pearce 1969, Berezinskii et~al 1990, Miller et~al 1997, Schlickeiser \\& Miller 1998, Chandran 2000, Schlickeiser 2002, Yan \\& Lazarian 2002).  If the small-scale waves are Alfv\\'en waves, the transport equation describing the energetic-particle distribution function~$f$ is (Skilling 1975) \\begin{equation} \\frac{\\partial f}{\\partial t} + \\left[ \\frac{\\partial}{\\partial p^3} \\left(p^3 {\\bf w }\\right)\\right] \\cdot \\nabla f = \\nabla \\cdot \\left( D_\\parallel \\hat{\\bf b} \\hat{\\bf b} \\cdot \\nabla f\\right) + (\\nabla \\cdot {\\bf w}) \\frac{p}{3}\\, \\frac{\\partial f}{\\partial p} + \\frac{\\partial}{\\partial p^3}\\left[9 p^4 D_{\\rm pp}^{\\rm (s)} \\frac{\\partial f}{ \\partial p^3}\\right], \\label{eq:ski1} \\end{equation} where $p$ is momentum, $\\hat{\\bf b}$ is the magnetic-field unit vector, \\begin{equation} {\\bf w} = {\\bf u} + \\left\\langle \\frac{3}{2} (1-\\mu^2) \\frac{\\nu_+ - \\nu_-}{\\nu_+ + \\nu_-}\\right\\rangle v_{\\rm A} \\hat{\\bf b} \\label{eq:defw} \\end{equation} is an effective ``wave-frame'' velocity at which the cosmic rays are advected, ${\\bf u}$ is the large-scale plasma velocity, $v_{\\rm A}$ is the Alfv\\'en speed, $\\nu_+$ and $\\nu_-$ are, respectively, the pitch-angle scattering rates associated with waves traveling parallel to and anti-parallel to the magnetic field, $\\mu$ is the pitch-angle cosine, $\\langle \\dots \\rangle$ indicates an average over~$\\mu$, \\begin{equation} D_\\parallel = v^2 \\left\\langle\\frac{1-\\mu^2}{2(\\nu_+ + \\nu_-)}\\right \\rangle \\label{eq:defdp} \\end{equation} is the diffusion coefficient for particle motion along the magnetic field, \\begin{equation} D_{pp}^{\\rm (s)} = 4 \\gamma_L^2 m^2 v_{\\rm A}^2 \\left\\langle\\left( \\frac{1-\\mu^2}{2} \\right)\\frac{\\nu_+ \\nu_-}{\\nu_+ + \\nu_-}\\right\\rangle \\label{eq:defdpps} \\end{equation} is the momentum diffusion coefficient associated with the small-scale waves, and~$\\gamma_L$ is the relativistic Lorentz factor of an energetic particle.  Equations~(\\ref{eq:defw}) and (\\ref{eq:defdpps}) show that either~${\\bf w} \\neq {\\bf u}$ or $D_{ pp}^{\\rm s} \\neq 0$ (Schlickeiser 2002). Nevertheless, for simplicity the momentum diffusion associated with small-scale waves is ignored ($D_{pp}^{\\rm s} \\rightarrow 0$), and at the same time, inconsistently, it is assumed that ${\\bf w} = {\\bf u}$.  This approximation has been standard in studies of particle acceleration by velocities on scales~$\\gg \\lambda_{\\rm mfp}$ [Bykov \\& Toptygin (1982, 1990, 1993), Ptuskin (1988), Dolginov \\& Silantev 1990, Katz \\& Stehlik (1991)]; it reduces equation~(\\ref{eq:ski1})  to \\begin{equation} \\frac{\\partial f}{\\partial t} + {\\bf u } \\cdot \\nabla f = \\nabla \\cdot \\left( D_\\parallel \\hat{\\bf b} \\hat{\\bf b} \\cdot \\nabla f\\right) + (\\nabla \\cdot {\\bf u}) \\frac{p}{3}\\, \\frac{\\partial f}{\\partial p}. \\label{eq:ski2} \\end{equation}  Equation~(\\ref{eq:ski2}) is used to calculate the momentum diffusion coefficient arising from the large-scale turbulence.  It seems plausible that a nonzero $D_{pp}^{\\rm s}$ would approximately add to the momentum diffusion coefficient associated with the large-scale turbulence calculated from equation~(\\ref{eq:ski2}), but future work is needed to verify this. The spatial diffusion caused by the large-scale turbulence is neglected.  It is assumed that the rms large-scale turbulent velocity is comparable to~$v_{\\rm A}$ and that the ratio~$\\beta$ of thermal to magnetic pressure is~$\\lesssim 1$. The power spectrum of the slow modes is assumed to follow the theory of Lithwick \\& Goldreich (2001), which has received support from direct numerical simulations (Cho \\& Lazarian 2002).  It is also assumed that $v_{\\rm A} d_{\\rm min} \\ll D_\\parallel \\ll v_{\\rm A} l$, where $d_{\\rm min}$ is the length along the magnetic field of the smallest slow-mode eddies under consideration ($d_{\\rm min}$ is either the length of the eddies at the dissipation scale or several times~$\\lambda_{\\rm mfp}$, whichever is larger), and~$l$ is the length of the largest (outer-scale) eddies. Since the time for an eddy of length~$\\lambda_\\parallel$ (measured along the magnetic field) to be randomized is~$\\lambda_\\parallel/v_{\\rm A}$ (Lithwick \\& Goldreich 2001), the inequality $v_{\\rm A} d_{\\rm min} \\ll D_\\parallel \\ll v_{\\rm A} l$ implies that the largest eddies are randomized before a particle diffuses through them, and that a particle diffuses through eddies of length~$d_{\\rm min}$ before such eddies are randomized.  The assumption $D_\\parallel \\gg v_{\\rm A} d_{\\rm min}$ requires that the particle speed~$v$ satisfy $v\\gg v_{\\rm A}$ since $D_\\parallel \\sim v\\lambda_{\\rm mfp}$ and $d_{\\rm min} > \\lambda_{\\rm mfp}$.  The procedure used to calculate the momentum diffusion coefficient~$D_p$ arising from the large-scale velocities is to average over the slow-mode fluctuations, neglect spatial variations in the  averaged distribution function~$f_0$, and approximate the eddies on scales comparable to~$l$ as a uniform background field~${\\bf B} _0$. The eddies on scales~$\\ll l$ have velocities~$\\ll v_{\\rm A}$ and magnetic perturbations~$\\ll B_0$ and can be treated using quasilinear theory. It is found that \\begin{equation} \\frac{\\partial f_0}{\\partial t} = \\frac{1}{p^2} \\frac{\\partial }{\\partial p} \\left( p^2 D_p \\frac{\\partial f_0}{\\partial p}\\right), \\label{eq:fp0} \\end{equation} with \\begin{equation} D_p \\simeq \\frac{2 p^2 v_{\\rm A}}{9  l}\\; \\left[\\ln \\left(\\frac{l v_{\\rm A}}{D_\\parallel}\\right) + 2\\gamma - 3 \\right], \\mbox{ \\hspace{0.3cm} for $v_{\\rm A} d_{\\rm min} \\ll D_\\parallel \\ll v_{\\rm A} l$,} \\label{eq:dpintro} \\end{equation} where $\\gamma \\simeq 0.577$ is Euler's constant.  The corresponding acceleration time scale,~$p^2/D_p$, is of order the large-eddy turnover time $l/v_{\\rm A}$, and depends only weakly through the $\\ln(l v_{\\rm A}/D_\\parallel)$ term on~$p$ and particle species.  An approximate version of equation~(\\ref{eq:dpintro}) is obtained using phenomenological arguments in section~\\ref{sec:phen}.   Section~\\ref{sec:spectra} presents the time-dependent spectrum of accelerated particles assuming an ad-hoc model of particle escape from the acceleration region with an energy-independent escape time. Section~\\ref{sec:stst} gives the steady-state spectrum for electron acceleration taking into account Coulomb losses, again assuming an energy-independent escape time.  Section~\\ref{sec:sf} compares the predicted spectra and acceleration times to observations of electrons in solar flares.  The results of the paper are summarized in section~\\ref{sec:conc}.  Much of the notation used in the paper is defined in table~\\ref{tab:t1}.  \\begin{table*}[h] \\begin{tabular}{lc} \\hline \\hline & \\vspace{-0.25cm} \\\\ Notation & Meaning \\vspace{0.1cm} \\\\ \\hline \\vspace{-0.2cm} \\\\ $\\lambda_{\\rm mfp}$ & scattering mean free path of energetic particles \\\\ $\\lambda_{\\rm thermal}$ & Coulomb mean free path of thermal protons \\\\ $\\lambda_\\parallel$ & length of a turbulent eddy measured along the magnetic field \\\\ $\\lambda_\\perp$ &  width of a turbulent eddy measured across the magnetic field \\\\ $r_g$ & energetic-particle gyroradius \\\\ $f$ & energetic-particle distribution function \\\\ $N(E)$ & number of accelerated particles per unit energy \\\\ ${\\bf B} $ & magnetic field \\\\ $d_\\parallel$ & length along magnetic field of eddies at the dissipation scale \\\\ $d_{\\rm min}$ &  $d_\\parallel$ or several times~$\\lambda_{\\rm mfp}$, whichever is larger \\\\ $l$ & stirring scale (outer scale) of turbulence \\\\ $v_{\\rm A} $ & Alfv\\'en speed \\\\ $c_{\\rm s}$ & sound speed \\\\ $\\beta$ & ratio of thermal to magnetic pressure \\\\ ${\\bf u} $ & turbulent velocity \\\\ $v$ & energetic-particle velocity \\\\ $D_\\parallel$ & diffusion coefficient for motion along magnetic field \\\\ $D_p$ & momentum diffusion coefficient \\\\ $\\tau_{\\rm esc}$ & escape time scale \\\\ $\\tau_{\\rm acc}$ & acceleration time scale \\\\ $\\tau$ & $t/\\tau_{\\rm acc}$ \\\\ $\\chi$ & $\\tau_{\\rm acc}/\\tau_{\\rm esc}$ \\\\ $\\gamma$ & Euler's constant, $\\sim 0.577$ \\vspace{0.2cm} \\\\ \\hline \\hline \\end{tabular} \\caption{Definitions. \\label{tab:t1} } \\end{table*} ",
        "conclusions": "\\label{sec:conc}  The momentum diffusion coefficient~$D_p$ arising from slow modes in strong compressible anisotropic MHD turbulence is calculated for the case that the rms turbulent velocity is~$\\sim v_{\\rm A}$, pitch-angle scattering of energetic particles is very efficient ($d_{\\rm min} v_{\\rm A} \\ll D_\\parallel \\ll l v_{\\rm A}$), the energetic particle speed~$v$ is~$\\gg v_{\\rm A}$, and $\\beta \\lesssim 1$. It is found that \\begin{equation} D_p \\simeq \\frac{2 p^2 v_{\\rm A}}{9  l}\\; \\left[\\ln \\left(\\frac{l v_{\\rm A}}{D_\\parallel}\\right) + 2\\gamma - 3 \\right], \\mbox{ \\hspace{0.3cm} for $v_{\\rm A} d_{\\rm min} \\ll D_\\parallel \\ll v_{\\rm A} l$,} \\label{eq:dpsolve2} \\end{equation} where $\\gamma \\simeq 0.577$ is Euler's constant.  Aside from a possible weak momentum dependence associated with the $\\ln (l v_{\\rm A}/D_\\parallel)$ term, $D_p\\propto p^2$, implying that the acceleration time~$\\tau_{\\rm acc} = p^2/D_p$ is independent of~$p$. In addition, $\\tau_{\\rm acc}$ is of order~$l/v_{\\rm A}$ and approximately independent of particle species.  Slow modes in the $D_\\parallel \\ll l v_{\\rm A}$-limit are an unlikely explanation for electron acceleration in solar flares to energies of 10-100~keV, because for solar-flare conditions the predicted acceleration times are too long and the predicted energy spectra are too hard.  The acceleration mechanism discussed in this paper could in principle explain the relatively hard spectra of gyrosynchrotron-emitting electrons in the~100-5000~keV range, but only if $D_\\parallel \\ll l v_{\\rm A}$ for such particles."
    },
    "0311/astro-ph0311107_arXiv.txt": {
        "abstract": "We reanalyze the group of stars associated kinematically with Arcturus (Eggen 1971), and confirm that they constitute a peculiar grouping of metal-poor stars with similar apocentric radius, common angular momentum, and distinct metal abundance patterns. These properties are consistent with those expected for a group of stars originating from the debris of a disrupted satellite. The Arcturus group stands out clearly in the catalog of Beers et al. (2000) as well as in other compilations of metal-poor stars in the solar neighborhood.  Its angular momentum appears too low to arise from dynamical perturbations induced by the Galactic bar, and coincides with that of the kinematically-peculiar population of stars identified above and below the Galactic plane by Gilmore, Wyse \\& Norris (2002): the Arcturus group is in all likelihood the solar neighborhood extension of such population. Further analysis is needed to confirm our proposal that Arcturus, together with some of the brightest stars in the night sky, may have formed beyond the confines of the Galaxy. However, validating the extragalactic origin of the Arcturus group would lend support to the prevalence of accretion events envisioned in hierarchical models of galaxy formation, and would suggest that early mergers may have contributed a substantial fraction of the old and metal-deficient stars in the Galactic disk. ",
        "introduction": "\\label{sec:intro}  The discovery of substructure in the distribution of stars in the outer regions of galaxies like the Milky Way and M31 constitutes indisputable evidence that accretion events and mergers have played a role in building up their stellar halos (Ibata et al. 1994, Helmi et al. 1999, Majewski et al. 2003). Although the role of such events in the formation of the Galactic disk components is not yet fully understood, models of the thick disk component typically envision the tidal heating of an extant thin disk by a merging satellite (Quinn \\& Goodman 1986). Recent numerical simulations of disk galaxy formation also suggest that disrupted satellites might have contributed a significant fraction of the old stars in the disk of the Galaxy (Abadi et al. 2003a,b). As with the stellar halo, long-lived traces of such events may be preserved in the form of substructure in the orbital distribution of stars in the disk (Helmi et al 2003).  Identifying such debris ought to be easier in samples where contamination by the thin disk component is minimized, such as samples of metal poor/old stars, or else samples collected {\\it in situ} away from the disk or far from the Galactic center. The discovery of a ``ring'' of stars of almost certain extragalactic origin in the outer disk of the Milky Way by the SDSS (Newberg et al 2002, Yanny et al 2003), as well as recent report of a galaxy being cannibalized in the Milky Way disk (Martin et al 2003) provide examples that this process indeed seems to be at work in the Galaxy (Helmi et al 2003, Rocha-Pinto et al 2003).  Gilmore, Wyse \\& Norris (2002) have also recently argued for substructure in the Galactic thick disk. They postulate the existence of a distinct population of stars, with kinematics intermediate between the canonical thick disk and the stellar halo (mean rotational velocity of $\\sim 100 \\kms$) to explain the results of their spectroscopic survey of F--G stars located 0.5--5 kpc away from the Galactic plane.  Substructure in the disk has been noted since the early days of Galactic structure research (see, e.g., Eggen 1998 and references therein), and has usually been interpreted as signature of the gradual dissolution of loosely bound star clusters in the clumpy potential of the disk.  One example is the group of stars associated by Eggen (1971) with the bright star Arcturus, whose rotation speed lags the Local Standard of Rest (LSR) by roughly $\\sim 100 \\kms$. We discuss below that evidence for a population of stars with rotation speeds lagging the LSR by $\\sim 100 \\kms$ is also present in the metal-poor star catalog of Beers et al (2000; hereafter B00), and in the compilation of Gratton et al (2003, hereafter GCCLB).  In this {\\it Letter} we reassess the status of Eggen's Arcturus group by using modern distance, proper motions, radial velocities, and metallicity determinations available in the literature. We argue that the Arcturus group; the structure in the B00 and GCCLB compilations; as well as the kinematically-odd population of stars discovered by Gilmore, Wyse \\& Norris (2002), are all part of the same, dynamically coherent group of stars that might have originated in the disruption of a fairly massive satellite early during the Galaxy's assembly.   \\begin{figure*}[t] \\includegraphics[height=0.35\\textheight,clip]{f1a.eps} \\includegraphics[height=0.35\\textheight,clip]{f1b.eps} \\caption{Panel (a) shows the angular momentum distribution of stars in the B00 catalog (top dot-dashed histogram); those in the Arcturus group candidate list of Eggen (1996; 1998; shaded histogram at the bottom); and stars in the compilation of Gratton et al. (2003; solid histogram). The middle dot-dashed histogram correspond to all B00 stars on ``disk-like'' orbits, i.e., $|W|<50$ kms s$^{-1}$. There is evidence in all these samples (which favor metal-poor stars) for an ``excess'' of stars with angular momentum coincident with the Arcturus group at $J_z\\sim 950$ $\\kms$ kpc.  Panel (b) shows the rotational velocity distribution of all stars in the B00 catalog, decomposed into three canonical Galactic (Gaussian) components: halo, thick disk, and thin disk. The solid curve denotes the sum of all three components. Note the ``excess'' of stars lagging the LSR rotation by $\\sim 120 \\kms$, coincident with the Arcturus group. See text for full details. \\label{figs:jzhist}} \\end{figure*} ",
        "conclusions": "\\label{sec:disc}  The peculiar angular momentum of the Arcturus group, together with its moderate velocity dispersion and singular metal abundance, all point to a possible extragalactic origin for the group.  Age estimates are available for some stars in the group (Fuhrmann 2000, 2003), and suggest that it is perhaps as old as $10$-$12$ Gyr. This would imply that the group is the debris from an ancient accretion event.  This accretion event has probably contributed a significant number of metal-deficient stars to the solar neighborhood, as shown by the prominence of the group in samples that favor metal-deficient stars. The presence of this dynamically-coherent group may explain some of the variance in the results of dynamical studies of the thick disk. Indeed, some of these studies have concluded that the bulk of the thick disk lags the LSR by only $\\sim 30$--$50 \\kms$ (Beers \\& Sommer-Larsen 1995, Ojha et al 1996, Norris 1999), whereas others have favored a lag of $\\sim 100 \\kms$ instead (Wyse \\& Gilmore 1986, Chen et al 2000, Fuhrmann 2000, 2003), possibly reflecting the different contribution of stars in the group to each sample.  There are a number of inquiries that may further validate (or refute) our interpretation. For example, confirming the link between the Arcturus group and the debris identified by Gilmore, Wyse \\& Norris outside the plane would help to uphold the group's extragalactic origin. Most of the stars in the Arcturus group are located close to their orbital apocenter, where density enhancements are expected: we may have detected a ``shell'' similar to those observed around elliptical galaxies (Helmi et al 2003). Other distinct features of this shell interpretation would include a rapid decline in the importance of the group at increasingly large Galactocentric distance (shells are ``sharp'') and, of course, the presence of other shells at different Galactocentric radii.  It may also be possible to find other members of the progenitor of the Arcturus group by searching for stars on disk-like eccentric orbits with metal abundances consistent with the ``sequence'' outlined by the group in Figure~\\ref{figs:afeh}.  Finally, it is oddly gratifying to think of stars visible to the naked eye, such as Arcturus, as silent night-sky witnesses of the merging history of the Milky Way.   \\begin{figure}[t] \\begin{center} \\includegraphics[height=0.35\\textheight,clip]{f3.eps} \\end{center} \\caption{Abundance ratios for stars in the GCCLB sample. Open circles correspond to all stars in each sample, whereas filled circles indicate candidate stars of the Arcturus group. Note that these stars delineate a well define sequence in the abundance ratios that may be interpreted as resulting from a short episode of star formation where a system self-enriched to a metallicity of order one-third solar. The curve correspond to model 1 of Matteucci \\& Francois (1989). See text for further details. \\label{figs:afeh}} \\end{figure}"
    },
    "0311/astro-ph0311382_arXiv.txt": {
        "abstract": "We performed a systematic analysis of 54 Wide Field Camera (WFC) observations of ScoX-1 available in the BeppoSAX public archive. Observations span over the six years of BeppoSAX mission lifetime and include 690 hr of data. We searched for shifts and shape changes of the Z pattern in the color-color diagrams. We find that the Z pattern occupies most of the time the same locus in the color-color diagram. There are however a few exceptions, which are discussed in detail.  \\vspace{1pc} ",
        "introduction": "Sco X-1 is a low mass X-ray binary (LMXB), is the brightest extrasolar X-ray source in the sky and was discovered by R. Giacconi and collaborators in the seminal X-ray rocket experiment of June 1962. Neutron star LMXB are divided into two classes based upon the different morphology of their tracks in color-color diagrams (CD). The three branches of Z tracks are named  Horizontal (HB), Normal (NB) and Flaring (FB), while  atoll sources are characterized by an ``island'' and a ``banana'' state \\cite{HvdK89}. Track morphology is due to spectral variations on timescales of weeks, days or hours. In the case of ScoX-1, the complete Z track is generally traced out in a few hours to a day. Such variations are interpreted as an accretion sequence \\cite{H89}, where the mass accretion rate $\\dot{M}$ is assumed to change along the track (see however \\cite{bcbc03}). Secular shifts and shape changes of the Z track  were reported for several Z sources such as CygX-2 \\cite{Kuu96}, whose shifts were interpreted in terms of occultation of the emitting region by a precessing accretion disk,  or very recently for LMC X-2 \\cite{shk03}, in which secular shifts of 2.5-10\\% are seen. Conversely significant secular variations in the Z track of ScoX-1 were never observed \\cite{dvdk00}. However an inclination of 45 deg and a variable disk thickness could produce the same effect, besides reducing the HB extent, as discussed in \\cite{Foma01}. ",
        "conclusions": "Analysis of ScoX-1 CD, spanning over the six years of BeppoSAX lifetime, show that most of the time the track covers the same position. However there are few noticeable exceptions. The most evident is a shift of the vertex and FB positions towards soft color and hard color 6-13\\% lower than the more common position. This shifts are observed on timescales of years or even months. Moreover we find indications that shifts appear also on shorter timescale (days). Detailed analysis is in progress to confirm this latter point."
    },
    "0311/astro-ph0311457_arXiv.txt": {
        "abstract": "It is now established that several of the Local Group dwarf Spheroidal galaxies (dSphs) have large mass-to-light ratios. We consider the possibility that the dark matter in the haloes of dSphs is composed of massive black holes with masses in the range $10^{5}\\Msun$ to $10^7\\Msun$. We use direct $N$-body simulations to determine the long term evolution of a two-component dSph composed of a pressure-supported stellar population within a black hole dominated halo. The black holes are initially distributed according to a Navarro, Frenk \\& White profile. For black hole masses between $10^5\\Msun$ and $10^6\\Msun$, the dark matter halo evolves towards a shallower inner profile in less than a Hubble time. This suggests that black holes in this mass range might provide an explanation for the origin of the dark matter cores inferred from observations of Low Surface Brightness galaxy rotation curves. We compare the simulated evolution of the stellar population with observed data for the Draco dSph and find that dynamical heating generally leads to the rapid dispersal of the stellar population to large radii. The dependence of the heating rate on the black hole mass is determined, and an upper limit of $10^{5}\\Msun$ is placed on the individual black holes comprising the dark matter halo of Draco, if its present stellar distribution is representative of the initial one. We also present a simple scaling argument which demonstrates that the dynamical heating of an initially compact, though not self-gravitating, stellar distribution might produce a remnant distribution similar in extent to Draco after 10 Gyr, even for black hole masses somewhat in excess of $10^{5}\\Msun$. ",
        "introduction": "\\label{sec:intro}  The existence of dark matter in the Universe, which has an abundance six times that of baryonic matter, is inferred from sources as diverse as the disk rotation curves of spiral galaxies, gravitational lensing studies of galaxy clusters, the very high mass-to-light ratios of nearby dwarf galaxies and analysis of Cosmic Microwave Background fluctuations.  However, the actual composition of the dark matter in the Universe is as yet unknown ~\\citep{Jerry2003}. One method by which dark matter may be detected and its properties probed is by means of substructure lensing ~\\citep[e.g.][]{Metcalf2001,Li2002,Keeton2003}. Dark matter studies in the Local Group can also yield constraints on dark matter models via the mass density as inferred from kinematic measurements ~\\citep[e.g.][]{Mark2002}.  The standard cold dark matter (CDM) cosmology, though popular and well-tested on large scales, has some unresolved issues.  The two best known are the over-abundance of low-mass satellite haloes identified at the present epoch in cosmological simulations relative to the number of observed Local Group satellites ~\\citep[e.g.][]{Moore1999} and the much-debated inner profile of CDM haloes ~\\citep[e.g.][]{Blok2002,Swaters2003}.  The latter may be due to resolution issues in the simulations, where a range of inner density slopes are found by various authors~\\citep[see][ for a concise summary]{Suto}. Low mass haloes may also have intrinsically softer inner profiles than more massive systems ~\\citep{Massimo2003}. Observationally, there is considerable evidence that massive galaxies also do not have cusped haloes, but instead contain roughly uniform density cores~\\citep[e.g.][]{binev01,salucci01,bsd03}. The issue of the missing satellite haloes may be due to the fact that some of the satellite galaxies are very diffuse; a few bound stellar systems may therefore be undetected simply due to their low surface brightness ~\\citep{Mateo1998}.  It has also been suggested that the satellite crisis can be resolved if all the Local Group dwarf spheroidal galaxies (dSphs) lie in significantly more massive haloes than those implied by previous estimates, which were based only on their central velocity dispersions~\\cite[e.g.][]{Stoehr2002,Hayashi2003}. However, differences between the radial distribution of massive satellite haloes about their parent galaxies in simulations and the observed distribution of the Milky Way dSphs suggest that this cannot be the complete solution to this problem~\\citep[e.g.][]{DeLucia03}. Another possibility is that there may exist dark matter haloes which do not harbour a stellar population.  In order to address these issues, some alternatives to CDM have been proposed: these include warm dark matter and self-interacting dark matter ~\\citep[e.g.][]{SS2000,Gnedin_Ostriker}.  An interesting complete alternative to dark matter has also been suggested for the observed dwarf Spheroidals (dSphs) whereby special phase-space characteristics permit the long-lived remnant of a tidally disrupted satellite to assume the appearance of a dSph in projection ~\\citep{Klessen_Kroupa}; however, difficulties in reproducing the data on Draco have recently been noted in this model ~\\citep{Klessen2003}.   There has been considerable interest in cosmological CDM simulations in the past decade, with a variety of halo profile fits being presented by different authors.  Following a set of large cosmological simulations, ~\\cite{NFW96,NFW97} proposed a universal density profile to accommodate virialised CDM haloes of all masses, which has become commonly known as the NFW profile.  The density in the inner region is cusped as $r^{-1}$, whereas at larger radii it is proportional to $r^{-3}$, with the location of the transition being defined by a scale radius used to characterise the profile.  Several authors have subsequently confirmed the fit provided by the NFW profile, although others have also presented steeper inner slopes for the halo density ~\\citep[e.g.][]{Fukushige2001,Ghigna2000,Colin2003}.  A generalised density law has therefore been suggested, for which the original NFW profile is a special case, in order to account for such findings~\\citep[e.g.][]{Zhao1996}.  As long as the dark matter is cold and collisionless, the halo profile would not be expected to depend on the nature of the dark matter and, more specifically, it should be independent of the dark matter particle mass.  This remains true unless we consider particles whose masses are much greater than a solar mass, for which two-body dynamical effects would become relevant.  Such phenomena would first be observed in the lowest mass, lowest velocity dispersion systems such as the dSphs.  Much interest in recent years has been focused on the study of the Local Group satellites, not least because their relative proximity has led to the increased availability of high quality data ~\\citep{IH95,Mateo1998}, which place stricter constraints on the nature of the dark matter that appears to dominate their gravitating mass.  These systems have also received much attention on the modelling and simulation front ~\\citep[e.g.][]{Stoehr2002,Mark2002}. Amongst those with the highest mass-to-light ratios is the Draco dSph galaxy.  At a heliocentric distance of approximately $80\\kpc$, Draco has been studied extensively and a wealth of kinematic data exists; in particular, measurements of stellar velocities out to large radii are now available ~\\citep[]{Jan2002,Mark2004}.  Given that within the region probed by the stellar distribution it also appears to have suffered relatively little from the tidal effects of the Milky Way ~\\citep{Odenkirchen,Klessen2003,Mark2004}, unlike some of the other Local Group satellites such as the Sagittarius dwarf ~\\citep[e.g.][]{Helmi2001}, Draco is an ideal candidate with which to study the possible properties of dark matter.  Draco is also known to contain an old stellar population with an age of more than 10 Gyr; any model of this galaxy must therefore remain stable for at least this length of time in order to be compatible with this observation.  One conceivable candidate for the dark matter, either as the sole or major component, is an ensemble of massive black holes.  These could be primeval, or may have grown by accretion and mergers from a primeval population, or have origins in later non-linear structures. ~\\citeauthor{LO85} (1985: hereafter LO85) investigated the way in which a halo of massive black holes would affect the stellar dynamics in the Milky Way and obtained results that are generally consistent with observations.  Our present study investigates the dynamical effects of a dark matter halo composed purely of primordial massive black holes on the evolution of the stellar component of a dwarf galaxy.  It was shown in LO85 that for a dSph such as Draco, an upper limit for the individual black hole mass would be $2\\times10^6\\Msun$, based upon stability arguments for both the stellar population and the dark matter halo forming the galaxy.  More recently, \\cite{mao04} have suggested that massive black holes with masses in the range $10^5-10^6\\Msun$ may offer a viable explanation for the observed image flux ratios in strong gravitational lens systems. We use direct $N$-body simulations to trace the evolution of a model dSph whose initial structural parameters are chosen such that they are compatible with present-day observations of Draco.  The stellar population is treated as a cluster of tracer particles, whose evolution is dependent on the potential of the dark halo.  The distribution of the halo particles is chosen to follow an NFW profile, whilst the stellar population is given an initial distribution which is described by a Plummer law.  We assume the dark matter of the halo to consist entirely of massive black holes, formed prior to the formation of the stellar population. Their masses are chosen to cover a range whose maximum value has, for generality, been taken to be larger than the constraints imposed by LO85; in each simulation, all black holes have the same mass.  In Section~\\ref{sec:sim} we discuss the generation of our simulated models.  We present the results in Section~\\ref{sec:results} along with a discussion in Section~\\ref{sec:disc}.  In Section~\\ref{sec:init_cond} we discuss alternative initial conditions for the stellar distribution of dSphs and their implications for black hole haloes. Finally we present a summary and our main conclusions in Section~\\ref{sec:con}. ",
        "conclusions": "\\label{sec:con} The high mass-to-light ratios of certain Local Group dwarf galaxies suggest that they are among the most dark matter-dominated stellar systems known in the Universe. This makes them ideal laboratories in which to investigate the nature and properties of dark matter.  In addition, their relative proximity has allowed the accumulation of a wealth of high quality stellar kinematic data in recent years.  We present here the results of simulations where the modelled system is assumed to consist of a dwarf galaxy-sized stellar population, which initially resides in the centre of an extended, massive halo with an NFW profile.  The halo particles are assigned masses of between $10^{5}\\Msun$ and $10^7\\Msun$ and are modelled as massive black holes. The mass of the system is chosen to match that of the Draco dSph, a Galactic satellite for which evidence has recently been presented of the existence of a massive dark matter halo ~\\citep{Jan2001}.  The stellar distribution appears to have suffered only minor tidal distortion ~\\citep{Klessen2003,Mark2004}. This, together with the availability of stellar velocity measurements up to four Plummer radii \\citep{Jan2002,Mark2004}, makes Draco a useful test candidate for comparison with simulations.  The first question which our simulations address is whether the black hole population evolves too quickly to be consistent with the potential observed for Draco.  Perusal of the very short central relaxation times (Table~\\ref{tab:sim}) indicates that we might have expected the systems to undergo core collapse and re-expansion within a Hubble time, with subsequent dissolution in a Galactic tidal field, had the initial conditions been closer to a Plummer model ~\\citep{Lee1987}.  However, due to the inverted initial temperature profile of the NFW model, the evolution in our simulations is slower and initially in the opposite direction, towards a lower central density with a cored rather than a cusped density profile.  Only systems with black hole mass greater than $10^{6.5} \\Msun$ appear to evolve too rapidly towards core collapse.  For the lower end of the mass range investigated, evolution leads to a profile with a density profile that is more consistent with observations of dSphs ~\\citep{Jan2003,Magorrian2003} than the initial NFW model.  A more serious set of issues is raised by the evolution of the tracer populations.  We find that in all of our dSph-sized models, the central tracer density shows a very rapid decline in the first Gyr, corresponding to over-efficient heating by the black holes.  This effect can also be observed via the increase in the tracer velocity dispersion, which at 10 Gyr cannot be reconciled with the observed Draco dispersion. If the initial stellar distribution is chosen to match that of Draco, then only for black holes of mass below about $10^5\\Msun$ is the heating rate reduced to a level which could perhaps be compatible with the observations.  We conclude that the haloes of dSph galaxies cannot be composed of black holes more massive than $10^5\\Msun$ if their initial stellar distributions were similar to those observed at the present time.  We have also performed simulations in which the scale length was set to twice that used to model the halo of Draco.  The heating of the stellar population in these cases was somewhat reduced, and more than $10$ percent of the stellar population which initially occupied the central $500\\pc$ of the halo remained in this region after 10 Gyr for black hole masses of $10^6 \\,\\textrm{M}_{\\odot}$, as opposed to a loss of greater than 95 percent for the models with smaller $r_{\\rm s}$. However, the velocity dispersion of the evolved tracers within the same region is still uncomfortably large. Notwithstanding the question of how to suppress the heating of the stellar population, the evolution of the haloes of these models is potentially very interesting. As Fig.~\\ref{fig:vcirc} shows, the inner regions of the haloes evolve towards a more cored profile. Thus, these models may be able to explain the observations of LSB galaxies which appear to have large dark matter cores~\\citep{Blok2002}.  By constraining the amount of evolution allowed to take place for the stellar population within a given time scale, we are able to place a new upper limit on the mass of individual black holes which constitute the dark matter halo of Draco. If we require that the stellar count within $500\\pc$ of the halo origin decreases by no more than a factor of 2 within 5 Gyr, we obtain an upper mass constraint of approximately $10^{5}\\Msun$; the equivalent condition required over a time scale of 10 Gyr sets a limit of $10^{4.6}\\Msun$. These numbers call into question the hypothesis that massive black holes could be a significant component of the dark matter: black holes with masses below $10^5\\Msun$ are less interesting as dark matter candidates because they cannot explain either the heating of the Milky Way disk (LO85) or the origin of cores in dark haloes (Section~\\ref{subsec:draco}). Therefore, we have also investigated the possibility that a dSph whose initial stellar distribution was more compact than those of the present-day dSphs could evolve into an object resembling Draco as a result of heating. We have presented a simple scaling argument which implies that if the black holes have masses of $10^5\\Msun$ then an initial stellar distribution with scale radius $48\\pc$ would evolve into an object of similar extent to that of Draco in $10\\Gyr$; such a scenario may permit models with black hole masses somewhat larger than $10^5\\Msun$ and correspondingly more compact initial conditions. However, we note that this scenario may require very finely tuned initial conditions because $\\sqrt{N}$ noise in the black hole distribution may inflate the heating rate at early times.  In summary, we find that scenarios in which the dark matter is in the form of massive black holes, with masses between $10^5\\Msun$ and a few $\\times 10^5\\Msun$, may provide a viable channel for the production of cored, rather than cusped, haloes but only if the initial stellar distributions of the dSphs were considerably more compact than the currently observed ones. Further simulations will be required to determine if this picture is, in all details, consistent with the observed data on the Local Group dSphs.  \\subsection*{Acknowledgements} We thank Scott Tremaine for valuable comments on an earlier draft of this paper and Sverre Aarseth for many useful discussions and for the use of his NBODY codes. We thank Andy Gould, Bohdan Paczynski, Piotr Popowski, Paolo Salucci and F. Javier Sanchez for their critical comments on the feasibility of the models, and the anonymous referee for useful comments.  MIW acknowledges financial support from PPARC."
    },
    "0311/astro-ph0311511_arXiv.txt": {
        "abstract": "We present galaxy circular velocity functions (GCVFs) for 34 low redshift ($z \\la 0.15$) clusters identified in the Sloan Digital Sky Survey (SDSS), for fifteen clusters drawn from dark matter simulations of hierarchical structure growth in a $\\Lambda$CDM cosmology, and for $\\sim$29,000 SDSS field galaxies.  We find that the observed and simulated cluster GCVFs take the form of a power law.  The cumulative GCVFs of the simulated clusters are very similar across a wide range of cluster masses, provided individual subhalo circular velocities are scaled by the circular velocities of the parent cluster.  Once all sources of error are accounted for, the intrinsic scatter in the cumulative, scaled observed cluster GCVF is consistent with the simulations.  The slope of the observed cluster GCVF is $\\sim$-2.5, independent of cluster velocity dispersion.  The average slope of the simulated GCVFs is somewhat steeper, although formally consistent given the errors.  Using our highest resolution cluster, we find that the effects of baryons on galaxy rotation curves is to flatten the simulated cluster GCVF into better agreement with observations. Finally, we find that the field GCVF deviates significantly from a power law, being flatter than the cluster GCVF at circular velocities less than 200 $\\kms$, and steeper at circular velocities greater than 200 $\\kms$. ",
        "introduction": "\\label{sec:intro}  \\indent In  the widely  accepted  hierarchical clustering  model of  structure formation, adiabatic  fluctuations in  the initial density  field grow and  merge  into  increasingly   larger  dark  matter  halos.   Simple inflationary approaches model the initial power spectrum of fluctuations as Gaussian and  scale-invariant, while cosmological parameters and the nature of  the dark matter control its  subsequent evolution.  The resulting mass  function of dark  matter halos can be  predicted using {\\it  N}-body   simulations  \\citep{Sheth99,  Bode01,   Jenkins01, Reed03}  or analytic  methods  based  upon Press-Schecter  theory  \\citep{Press74, Bond91,  Bower91, Lacey93,  Sheth01}.  The  observed mass  function is therefore  an important  constraint on  fundamental properties  of the universe.  In the  following, we focus on the  mass function on galaxy scales.  Comparing predicted mass functions to observations requires mapping the observable properties of galaxies onto the masses of their dark matter halos.  In one technique, semi-analytic models (SAMs) are used to predict the luminosities of galaxies within halos of different masses, by parametrizing models of the complicated processes of gas cooling, star formation, stellar evolution, and feedback.  These techniques efficiently explore a wide variety of models.  However, their utility is limited by the extent to which the simple recipes capture the relevant physics.  An alternative to SAMs utilizes the empirical relationship between velocities and luminosities in both spiral and elliptical galaxies to transform the observed galaxy luminosity function into a galaxy circular velocity function (GCVF), which is closely related to the mass function \\citep{Cole89, Shimasaku93, Gonzalez00, Kochanek01}. While this approach avoids a number of assumptions required to assign luminosities to dark matter halos, baryons cannot be completely ignored.  The baryonic mass distribution affects a galaxy's circular velocity profile significantly, while also changing the equilibrium configuration of the host dark matter halo through adiabatic contraction \\citep{Blumenthal86}.  Baryons also determine the probability that a galaxy will be included in a given observational sample.  However, the circular velocity directly traces the total enclosed mass of a galaxy, which is dominated by dark matter at sufficiently large radii.  Therefore, models of the GCVF are less sensitive to inaccuracies in our understanding of star formation, stellar evolution, and feedback than those of the luminosity function. The observed GCVF therefore better constrains the mass function.  Recent comparisons between observed and theoretical field GCVFs indicate that low-mass halos with circular velocities $\\la$200$\\kms$ are over-predicted by simulations \\citep{Gonzalez00, Kochanek01}.  It is now possible to extend the analysis of the GCVF to galaxies within clusters.  The current generation of {\\it N}-body simulations \\citep{Ghigna98, Tormen98, Klypin99clus, Okamoto00, Springel01} has achieved the force and mass resolution necessary to capture the effects of tidal stripping, dynamical friction, and galaxy harassment, essential for predicting the GCVF in dense environments. \\citet{Moore99} used the Tully-Fisher relation to determine circular velocities for galaxies in the Virgo Cluster.  They compared the resulting distribution to that in a standard CDM simulation of a Virgo-like cluster, and found good agreement down to $\\sim$50 $\\kms$ without accounting for the effects of baryons on galaxy rotation curves.  Coupling dark matter merging trees with SAMs which include strong feedback, \\citet{Springel01} were able to reproduce a wide variety of cluster properties.  Although they did not directly compute the cluster GCVF, the agreement between a variety of observed and modeled cluster relations (morphology-density, luminosity function, Tully-Fisher, Faber Jackson) suggests that they would have found agreement for the cluster GCVF as well.  Although these works represent significant progress in testing simulations of cluster substructure against observations, larger samples of both simulations and observations are required to assess whether the results are typical, and over what cluster mass range.  Toward this goal, we determine the GCVFs for 34 nearby clusters from the Sloan Digital Sky Survey (SDSS).  To compare the observed clusters to theoretical predictions, we assemble a sample of 15 high resolution cluster simulations, from which we measure the theoretical cluster GCVF.  We also repeated the GCVF analysis for $\\sim$29,000 galaxies from the SDSS, to allow a comparison to the field.  Our method for determining the GCVFs for the SDSS sample is described in \\S{\\ref{sec:obs}}.  In \\S{\\ref{sec:sims}} we describe the simulations.  Our results are discussed in \\S{\\ref{sec:results}}. First we discuss the shape of the cluster GCVF as related to that of the cluster luminosity function (\\S{\\ref{sec:shape}}).  Second, we investigate the claim of \\citet{Moore99} that the velocity distribution of substructure on all mass scales is self-similar when scaled by the host halo velocity (\\S{\\ref{sec:selfsim}}).  Third, we present power-law fits to the cluster GCVF, quantify its dependence on cluster velocity dispersion, and make a comparison to dark matter cluster simulations (\\S{\\ref{sec:trends}}).  Fourth, we examine how the adiabatic contraction of dark matter halos due to baryons affects the comparison between observed and simulated GCVFs (\\S{\\ref{sec:hires}}).  Finally, we compare the observed cluster GCVF to that measured for the field (\\S{\\ref{sec:field}}).  A summary of our results can be found in \\S{\\ref{sec:summary}}.  Throughout we adopt $H_{{\\rm 0}} = 70\\kms {\\rm Mpc}^{-1}$, $\\Lambda = 0.7$, and $\\Omega = 0.3$. ",
        "conclusions": "\\label{sec:summary}  We constructed the cluster and field galaxy circular velocity functions using data from the Sloan Digital Sky Survey Early Data Release.  Circular velocities were determined using the inverse Tully-Fisher and Fundamental Plane relations on a galaxy-by-galaxy basis, rather than by performing a transformation of variables on a functional fit to the luminosity function.  For comparison to the observations, we also analyzed the GCVFs determined from 15 clusters simulated in the $\\Lambda$CDM concordance cosmology.  Our results can be summarized as follows:  1.  We found that, although the observed cluster luminosity function is a Schechter function,  the observed cluster GCVF can be described as a power law.  2.  We tested the self-similarity claimed by \\citet{Moore99} for our scaled simulated and observed cluster GCVFs.  The simulated GCVFs are indeed very similar to one another, and show only a small scatter.  The scatter in the observed cluster GCVFs is much larger, but can be explained by measurement error.  Thus, the intrinsic scatter in the observed cluster GCVF is similar to that predicted.  3.  The normalization of the observed cluster GCVF increases with cluster velocity dispersion, while the slope is constant.  The average slope of the observed cluster GCVF is somewhat flatter ($\\beta \\approx -2.5$) than that computed from $\\Lambda$CDM cluster simulations ($\\beta = -3.4$).  However, the errors on these average slopes are large, and overlap.  If the difference between the simulated and the observed clusters is real, it is most likely due to the fact that the simulations do not include the effects of baryons on galaxy rotation curves.  4.  We estimated the effects of baryons on the simulated GCVFs. Baryons can flatten the simulated GCVF enough to bring it into agreement with the GCVFs of observed clusters with similar velocity dispersions.  This agreement between the amount of cluster substructure predicted by N--body simulations and that observed in real clusters has been found previously by \\citet{Moore99} for an SCDM simulation of a Virgo-like cluster.  Our approach is different in that we treat early and late types distinctly, as well as account for the effects of baryons on galaxy rotation curves.  In addition, our simulated cluster has evolved in the concordance cosmology."
    },
    "0311/astro-ph0311401_arXiv.txt": {
        "abstract": "With a 1999 \\asca\\ observation, \\pgbal\\ became notable as the first Broad Absorption Line (BAL) quasar found to exhibit a typical radio-quiet quasar X-ray continuum underlying a large amount of intrinsic absorption.  We present a recent \\chandra\\ ACIS-S3 observation of \\pgbal\\ that demonstrates remarkable spectral and luminosity variability since that time. In addition to a decrease in the continuum normalization by a factor of $\\sim3.5$, the absorption column density has apparently increased substantially, and a strong feature in the \\feka\\ region has appeared. Concurrent \\hst\\ STIS data compared with archival \\hst\\ data from earlier epochs show evidence for variability of the continuum (up to a factor of $\\sim1.7$ in the ultraviolet), and in some absorption features of the \\CIV$\\lambda$1549 BAL since 1992; however, the \\OVI\\ BAL structure is consistent with a 1995 observation.  We also present evidence for Ly$\\beta$--\\OVI$\\lambda$1037.62 and Ly$\\alpha$--\\NV$\\lambda$1242.80 line-locked absorption systems, supporting the assumption that ultraviolet line pressure is driving the BAL outflow. Whereas ultraviolet BALs typically exhibit only modest equivalent-width variability over timescales of years, the dramatic X-ray variability of \\pgbal\\ suggests that X-ray spectral variability studies of \\balqs\\ have great potential for probing the physics of quasar winds. ",
        "introduction": "\\label{sec:intro}  Recent studies of the strong correlations between supermassive black hole (SMBH) masses and the properties of their hosts' galactic bulges \\citep[e.g.,][]{FerMer2000,GebhardtEtal2000} demonstrate clearly the intimate interrelation between the growth of SMBHs and their host galaxies.  This connection implies a physical mechanism for regulating the coeval development of these structures. During much of their accretion phases, SMBHs reveal themselves as luminous quasars \\citep[e.g.,][]{YuTre2002}, and thus quasar winds are likely to provide an important source of feedback during the coeval growth of SMBHs and their hosts' galactic bulges \\citep[e.g.,][]{Fabian1999}. These winds are directly observed in the population of Broad Absorption Line (BAL) quasars, which exhibits deep, broad absorption lines from high-ionization ultraviolet (UV) resonance transitions. BALs arise when our line of sight passes through thick material flowing outward from the nuclear region at speeds up to 0.03--0.3$c$.  Though \\balqs\\ typically compose $\\sim10\\%$ of the population in optically selected surveys, a recent analysis of the survey criteria for the Large Bright Quasar Survey, taking into account selection biases against the \\balq\\ population, concluded that the true fraction of \\balqs\\ in the radio-quiet quasar population is $22\\pm4\\%$ \\citep{HewFol2003}; \\citet{ReichardEtal2003} reached a similar conclusion based on the Sloan Digital Sky Survey \\balq\\ sample. Furthermore, spectropolarimetry, statistics from gravitational lensing, and theoretical arguments support this higher intrinsic \\balq\\ fraction \\citep[e.g.,][]{Goodrich1997,Chartas2000,KrVo1998}. Therefore, \\balqs\\ are not exotic objects but rather provide an opportunity to study powerful gas outflows that are likely present in typical quasars \\citep[e.g.,][]{WeMoFoHe1991}.  The presence of UV BALs is generally accompanied by extreme weakness in soft X-rays as a result of intrinsic absorption  with spectroscopically measured column densities typically in the range of (1--50)$\\times10^{22}$~\\cmsq\\ \\citep[e.g.,][]{GreenEtal2001,GaBrChGa2002}. Conversely, X-ray-weak Active Galactic Nuclei (AGNs) often show evidence for UV absorption, though not necessarily BALs \\citep*[e.g.,][]{BrLaWi2000}. While significant X-ray and UV absorption invariably occur in concert, the absorber properties do not always agree.  For example, the column density implied from the X-ray observation of the \\balq\\ PG~0946$+$301, \\nh\\,$\\gtrsim10^{24}$~\\cmsq\\ \\citep{MathurEtal2000}, is significantly larger than the best estimates from detailed UV absorption-line studies \\citep[\\nh$\\sim10^{20-22}$~\\cmsq;][]{AravEtal2001}. Furthermore, among \\balqs, there are no obvious trends relating the particulars of the UV absorption-line properties to the magnitude of the observed weakness in X-rays \\citep{GallagherEtal2003a}. Such a correspondence between X-ray and UV absorption properties might be expected based on theoretical quasar disk-wind models \\citep[e.g.,][]{MuChGrVo1995} and also on observations of Seyfert galaxies where high-resolution X-ray spectroscopy has revealed that the ionized absorbers seen in the X-ray regime often share kinematic properties with the UV-absorbing outflows \\citep[e.g.,][]{KaspiEtal2002,GabelEtal2003}. However, at present it is unclear if Seyfert ionized-absorber outflows are distinct from or belong to a continuum of winds that extend to the energetic BAL outflows in luminous quasars (though see Laor \\& Brandt 2002 for a discussion of this issue). \\nocite{LaoBra2002}  Though a high-resolution X-ray gratings observation of a \\balq\\ would offer a direct kinematic test of the connection between the velocity structure seen in X-ray and UV absorption, the X-ray faintness of this population makes such an observation expensive. Simultaneous X-ray/UV variability studies offer an alternative means of investigation. A clear demonstration of absorption variability in one wavelength regime without corresponding changes in the other would provide strong evidence that the two absorbers are distinct.  Conversely, corresponding changes in both X-ray and UV absorption would support an X-ray/UV absorber connection. Understanding the relationship between X-ray and UV absorption is essential for constraining important physical parameters, in particular the mass-outflow rate from the central engine.  The mass-outflow rate determines the kinetic influence of quasar activity on star formation in the host galaxy as well as the quasar contribution to the enrichment of the intergalactic medium with metals.  An excellent target for a variability study is the UV-bright, \\balq\\ \\pgbal\\ which has shown significant variability in both the UV continuum and X-ray in the past \\citep{GaEtal2001a}. This quasar is one of the most luminous low-redshift ($z=0.466$) Bright Quasar Survey objects with $M_{\\rm V}=-26.9$ (\\hnot$=70$\\hunits, $\\Omega_{\\rm M} = 0.3$, and $\\Omega_{\\Lambda} = 0.7$ are assumed throughout), and it has a broad-band optical polarization of $0.71\\%$ \\citep{BeScWeSt1990}. \\hst\\ spectra revealed broad, shallow \\CIV\\  absorption making this a \\balq\\ with a balnicity index of $\\approx 2980$\\kms\\ \\citep{JannuziEtal1998,BrLaWi2000}, well above the lower limit that defines the class \\citep{WeMoFoHe1991}. There are no spectra in the literature that cover the \\MgII\\ region, and so it is not known if \\pgbal\\ has broad \\MgII\\ absorption lines (see Boroson \\& Meyers 1992 for further discussion of \\MgII\\ \\balqs). \\nocite{BoMe1992} However, the lack of \\AlIII\\ absorption (which often accompanies \\MgII) suggests that \\pgbal\\ does not have low-ionization BALs. \\pgbal\\ was detected in soft X-rays with \\rosat\\ \\citep{BrLaWi2000} which is unusual for a \\balq\\ \\citep{KoTuEs1994,GrMa1996}. With the 1999 Oct 30 \\asca\\ data, it was the first \\balq\\ shown to have an X-ray continuum typical of radio-quiet quasars underlying intrinsic absorption.  \\pgbal\\ showed significant spectral and flux variability between the 1991 Nov 19 \\rosat\\ and the \\asca\\ observations \\citep{GaEtal2001a}.  With the high X-ray flux measured with \\asca\\ as well as its history of X-ray variability, \\pgbal\\ was a natural target for a coordinated X-ray/UV spectral study to investigate the relationship between absorption in the two regimes.  In this paper, we present the results from joint \\chandra/\\hst\\ observations of \\pgbal\\ taken on 2002 September 1.  For reference, the dates of the multiwavelength observations presented in this paper are presented in Figure~\\ref{fig:lc}. ",
        "conclusions": "Between the 1999 \\asca\\ spectra and the 2002 \\chandra\\ spectrum presented in this paper, the X-ray spectral shape and luminosity of \\pgbal\\ have varied dramatically. From the best-fitting partial-covering absorber models, the change in the continuum shape at low energies can be explained by a significant increase in the intrinsic column density (see Figure~\\ref{fig:contours}), while the underlying X-ray power-law photon index remained constant within the statistical errors.  This basic conclusion is not model-dependent; jointly fitting the \\asca\\ SIS and \\chandra\\ data with ionized-absorber models resulted in the same conclusion (see Table~{\\ref{tab2}}). Figure~\\ref{fig:4models} shows the ionized and partial-covering absorber models which best fit the data from both epochs. While the power-law photon indices are consistent between the two observations, the best-fitting 2002 continuum normalization is $\\sim30\\%$ of the 1999 level.  The absorption-corrected, rest-frame 2.0--10.0 keV luminosity, $L_{2.0-10.0}$, thus dropped by a factor of $\\sim3.7$ while the absorbing column density increased (see Table~{\\ref{tab1}). In addition, structure in the \\feka\\ region is now evident, consistent with a broad, neutral emission line with rest-frame EW=580--1570~eV.  A feature with this flux would not have been detectable in the 1999 \\asca\\ data given the high level of the continuum.  Comparison of the UV data from 2002 with the two previous epochs, 1992 and 1995, also reveals variability, but on much more modest scales than seen in the X-ray. The \\CIV\\ region observed in 1992 and 2002 shows evidence for small equivalent width decreases in individual absorption features, while the \\OVI\\ spectra taken in 1995 and 2002 do not show significant evidence for absorption feature changes. Over the course of a decade, the UV continuum luminosity varied as well; between 1992 and 2002 the UV continuum redward of \\Lya\\ became bluer and brighter.  Most of this increase in continuum luminosity could have occured between 1992 and 1995 as the 1995 continuum blueward of the \\Lya\\ emission line is within $\\sim25\\%$ of the 2002 level.  \\subsection{Implications of Variability for X-ray and UV Absorbers}  Given the present data quality, translating the significant empirical changes into physical understanding is challenging. Unfortunately, there are no UV observations simultaneous with the 1999 \\asca\\ observations, and so we cannot rule out significant changes in the UV absorption properties between 1999 and 2002. However, the concurrent lack of UV absorption-line or large-scale continuum variability from the 1995 GHRS to the 2002 STIS observations is worth noting.  The comparison of these data with the current STIS data is particularly relevant as 1995 is closer in time to the 1999 \\asca\\ observation and the continuum level in 1995 is more consistent with the 2002 level. As seen in Figure~\\ref{fig:uv}, there is no evidence for Ly edge absorption in the 2002 STIS G140L spectrum.  This puts an upper limit on the column density of neutral hydrogen obscuring the UV continuum of $\\nh\\lesssim 10^{17}$~\\cmsq\\ or perhaps a few times larger in the case of partial covering. However, the best-fit values for the X-ray absorbing column densities during both the \\asca\\ and \\chandra\\ epochs are $\\sim$5--6 orders of magnitude larger.  Either a highly ionized X-ray absorber or an X-ray absorber that does not obstruct the UV continuum could explain this discrepancy.  Though UV BAL studies are notoriously unreliable for constraining absorber column density if partial covering is not properly taken into account \\citep[e.g.,][]{Arav1997}, a column density increase by a factor of $\\sim7$ as indicated by our X-ray spectral fitting would likely have some effect on the UV absorption lines if the X-ray and UV absorption arise in the same gas.  The isolated, large-scale X-ray variability thus suggests the presence of an X-ray absorber distinct from the UV absorber in \\pgbal. However, without concurrent X-ray and UV data from more than one epoch, this claim remains speculative.  If the bulk of the X-ray absorption is distinct, it could arise in the ``shielding gas'' postulated in the disk-wind models of \\citet{MuChGrVo1995}.  In this picture, the broad emission lines and the BALs are produced in the same material at radii of a few light days with the shielding gas at smaller distances from the SMBH. This may be typical in \\balqs; the recent detection of relativistically blueshifted X-ray absorption lines in the \\balqs\\ \\apm\\ and \\pgone\\ (Chartas et al. 2002; Chartas, Brandt \\& Gallagher 2003), \\nocite{ChBrGaGa2002,ChBrGa2003} unmatched in velocity with UV absorption lines, also supports this picture.  The four UV absorption systems identified in \\S\\ref{sec:UV} demonstrate, if the line locking is real, that radiative forces play a role in the dynamics of the outflow.  In particular, Lyman series, \\OVI, and \\NV\\ line pressure was probably important for accelerating the gas, otherwise the systems would not have become locked when these wavelength coincidences occured. However, the Ly$\\alpha$--\\NV$\\lambda$1242.80 and Ly$\\beta$--\\OVI$\\lambda$1037.62 line pairs are only two line-locked pairs of many other possibilities among UV resonance lines, and so they may only contribute a small part of the total radiation force \\citep[e.g.,][]{KoVoMoWe1993}. If we speculate that the structure seen in BAL troughs is caused by systems locked into a complex network by many absorption-line pairs, then the UV BAL velocity structure might be expected to be extremely stable, even in the presence of a varying continuum.  Confirmation of a change in velocity of an X-ray BAL such as those detected in \\apm\\ or \\pgone\\ (if line locking plays such a significant role in UV BAL structure) would then suggest two intriguing possibilities: (1) the X-ray absorption is occuring in gas as it is being accelerated, before the system has become line locked, or (2) the X-ray absorbers are being driven by Compton pressure where line locking would not occur.  In the latter case, the X-ray absorbers would be distinct from the UV BAL gas.  \\subsection{Iron Emission}  The combination of the apparent width ($\\sigma=0.46^{+0.73}_{-0.32}$~keV) and equivalent width (EW=$1050^{+520}_{-471}$~eV) derived from modeling the Fe emission line in \\pgbal\\ is notable.  Intrinsically broad lines are typically associated with emission from the inner accretion disk; however, the EWs of such features are expected to be $\\sim150$~eV \\citep[e.g.,][]{MaPePiSt1992}.  Empirically, the mean for such lines has been found to be higher, $\\sim220$~eV \\citep[e.g.,][]{GeorgeEtal2000}, though still significantly less than we observe for \\pgbal. Therefore, if the line in \\pgbal\\ is intrinsically broad, this is hard to reconcile with its large EW.  Furthermore, a time-delay scenario in which a previously strong continuum has illuminated material whose emission is now more apparent as the continuum level has decreased is unlikely with the small scales ($\\sim$ light hours) of the inner disk. The Fe emission from that region would react on short time scales to a diminution of the direct continuum, and so a broad, high-EW line from this region would be unlikely.  In contrast, large EW Fe lines are common features of Seyfert~2 X-ray spectra with Compton-thick obscuration to the direct continuum \\citep[e.g.,][]{GeoFab1991,MatBraFab1996,TuGeNaMu1997}. In this case, reflection of direct X-rays off optically thick material usually associated with the torus results in a Compton-reflection hump and narrow, fluorescent \\feka\\ emission with large EW. Therefore, an alternative picture to an intrinsically broad line is that the apparent breadth of the line arises from unresolved, narrow \\feka\\ lines from many ionization states.  Given the limited statistics and CCD resolution, this scenario is plausible from the positive residuals above rest-frame 6.4~keV in Figure~\\ref{fig:chandra}, though the apparent lack of a strong reflection continuum (see \\S\\ref{sec:xanal}) is problematic unless the optically thick reflector is highly ionized.  Given the larger spatial scales on which the narrow emission lines are believed to arise, a time-delay scenario makes more sense.  If this is the case, then the current \\feka\\ emission is reflecting the previously strong level of the continuum as was seen in the Seyfert galaxies NGC~2992 and NGC 4051 during low continuum states \\citep[e.g.,][]{WeaverEtal1996,GuainazziEtal1998}. In this situation, the \\feka\\ line strength would be expected to diminish over time; therefore X-ray monitoring of this source has the potential to constrain the spatial scale of any reflecting medium in this luminous quasar.  \\subsection{\\boldmath\\aox} \\label{sec:aox}  X-ray spectroscopy of \\balqs\\ has indicated that they have typical radio-quiet quasar X-ray continua and spectral energy distributions in which soft X-rays are depressed due to intrinsic absorption. The power in the X-ray regime relative to the UV continuum is parameterized using the quantity \\aox=0.384~$\\log(f_{\\nu,\\rm2keV}/f_{\\nu,\\rm2500\\AA})$, where the flux densities, $f_{\\nu}$, are measured at the rest-frame frequencies as labelled. To calculate $f_{\\nu,\\rm2500\\AA}$, the flux density at rest-frame 2200~\\AA\\ (observed-frame 3225~\\AA) was estimated by extrapolating a short distance ($\\sim55~\\AA$) along the continuum from the long-wavelength end of the \\hst\\ STIS spectrum.  From that point, the flux density at rest-frame 2500~\\AA\\ was calculated using the Sloan Digital Sky Survey composite quasar spectral slope, $\\alpha_{\\nu}=-0.44$ \\citep{VandenBerk2001}. We chose this extrapolation to avoid potentially significant contamination from a complex of blended \\FeII\\ emission lines around 2500~\\AA, and this value for the spectral index is consistent with the observed UV continuum. The rest-frame 2~keV (observed-frame 1.36~keV) flux density is determined from a local continuum fit to the \\chandra\\ data. Based on these simultaneous X-ray and UV data of \\pgbal, the observed value of \\aox=$-2.07\\pm0.05$ is within the normal range for \\balqs\\ \\citep[$-$2.6 to $-$1.9; e.g.,][]{GallagherEtal2003a,GreenEtal2001}. However, correcting the X-ray spectrum for the intrinsic absorption from the best-fitting partial-covering absorber gives a ``corrected'' \\aoxc=$-1.79$. In contrast, the value calculated from the absorption-corrected 1999 \\asca\\ spectra and the 1992 \\hst\\ FOS data, \\aoxc$_{1999}=-1.55$, suggested much stronger (a factor of $\\sim 4.2$) X-ray emission relative to the UV continuum, though the 1992 UV data were not simultaneous with the \\asca\\ observation.  As can be seen in Figure~\\ref{fig:uv}, the continuum flux densities at rest-frame 2200~\\AA\\ (observed-frame 3225~\\AA) in both epochs are very similar, and so the decrease in \\aox\\ is dominated by the decrease in X-ray flux. For comparison, normal radio-quiet quasars of this luminosity typically have \\aox$\\sim-1.62\\pm0.1$ \\citep[see  eq. 4 of][]{VignaliEtal2003b}.  The change in \\aoxc\\ of $-0.24$ over the $\\sim3$-yr time span is thus still within the spread for normal radio-quiet quasars, though \\aoxc$_{2002}$ is at the weak end of the distribution. At this point, it is unclear whether the 1999 \\asca\\ spectra or the 2002 \\chandra\\ data are more representative of the ``normal'' state of the quasar continuum in \\pgbal.  If the 1999 \\asca\\ data are more representative, then \\pgbal\\ might return to this emission level at any time.  At least one \\balq, PG~1254$+$047, was argued to be intrinsically X-ray-weak in addition to suffering from heavy X-ray absorption \\citep{SaHa2001}.  Though the X-ray data quality in that particular case was poor (only $\\sim44$ photons), it is quite possible that the strength of the X-ray continuum of that source varies as well. Probing the timescale and interplay of X-ray continuum and absorption variability requires additional monitoring observations of sources with demonstrated variability such as \\pgbal.  \\subsection{Comparison of the Bolometric and Eddington Luminosities}  In contrast to the conclusion drawn by \\citet{WeMoFoHe1991} that all quasars are likely to have BAL regions, others have suggested that the BAL phenomenon may instead be a manifestation of an evolutionary phase, perhaps of enhanced accretion rate \\citep[e.g.,][]{HaMoTeMc1984,BeckerEtal2000}.  To investigate this issue for \\pgbal, we estimated the bolometric luminosity as a fraction of the Eddington luminosity, $L_{\\rm bol}/L_{\\rm Edd}$, to determine if \\pgbal\\ is an outlier in the $L_{\\rm bol}/L_{\\rm Edd}$ distribution of unabsorbed radio-quiet quasars. From $I$-band imaging, \\citet{FyMoTh2001} measure $I=18.8$ for the apparent magnitude of the luminous elliptical host galaxy which corresponds to an absolute magnitude of $M_I=-23.2$  (in the observed-frame Cousins $I$ band).  After converting to the rest-frame $V$-band luminosity \\citep[including the appropriate $K$ and color-corrections for elliptical galaxies from Table 3 of][]{FukShiIch1995}, we obtain an estimate of the black hole mass of $M_{\\rm BH}$=8.8$\\times10^{8}\\msun$ \\citep[using eq. 7 from][]{Laor2001}. To estimate the black hole mass from the optical spectroscopic properties of the quasar, we utilize $L_{\\rm 3000 \\AA}$ and H$\\beta$ FWHM and eq. 3 in \\citet{Laor1998} to obtain $M_{\\rm BH}$=6.8$\\times10^{8}\\msun$.  Both of these mass estimates have significant systematic uncertainties, perhaps as large as a factor of $\\sim3$ \\citep[e.g.,][]{Laor1998,Vestergaard2002}. However, it is notable that these two completely independent methods give comparable estimates of the black hole mass.  Using the bolometric correction from the 2500~\\AA\\ monochromatic luminosity for radio-quiet quasars of 6.3$\\pm3.1$ \\citep[as calculated empirically by][]{ElvisEtal1994}, we estimate a bolometric luminosity for \\pgbal\\ of $L_{\\rm bol}=(7.7\\pm3.8)\\times10^{46}$~\\lumin\\ based on $L_{\\rm2500\\AA}$ measured from the 2002 STIS spectrum as described in $\\S$\\ref{sec:aox}.  Combining these values with the log average of the black hole mass estimates, $M_{\\rm BH}$=7.7$\\times10^{8}\\msun$, gives an Eddington ratio of $L_{\\rm bol}/L_{\\rm Edd}\\sim0.8$. Comparison of this rough value with those of similarly luminous quasars in the sample accumulated by \\citet[][see their Figure~8]{WooUrr2002} indicates that \\pgbal\\ does not have an anomalously high Eddington ratio at present.  Similarly, \\citet{YuaWil2003} found that the Eddington ratios of a sample of $z\\sim2$ BAL and non-\\balqs\\ were consistent.   \\subsection{Future Work}  The lack of variability in the velocity structure of UV BALs on timescales of years is one of the most serious challenges to the current disk-wind models \\citep[e.g.,][]{deKool1997}.  While changes in the depths of individual absorption-line features have been seen, this EW variability is typically ascribed to changing ionization structure within the absorbing gas \\citep[e.g.,][]{BarlowEtal1992}. Such a change in ionization is consistent with the decrease in the depth of \\CIV\\ absorption features we detected during an epoch when the UV continuum increased in normalization as well as became bluer. In contrast, changes in velocity of individual absorption components are almost unknown. However, the most comprehensive absorption-line variability investigation to date, the 3-yr study by \\citet{Barlow1993}, was not long enough for a conclusive test; the hydrodynamic wind models of \\citet{PrStKa2000} generated instabilities on rest-frame timescales of $\\sim3$~yr. With longer time baselines, velocity increases have been claimed for at least one \\balq, Q1303$+$308 \\citep{VilIrw2001}, and an extended study is warranted to investigate this issue further. With high-quality UV coverage now extending from 1992 to 2002, almost seven years in the quasar rest-frame, \\pgbal\\ is an excellent candidate for continued monitoring of its BAL structure.  The large gaps in coverage of \\pgbal\\ over the last decade, illustrated in Figure~\\ref{fig:lc}, have limited our ability to interpret its multiwavelength variability. Additional epochs of simultaneous X-ray/UV monitoring will contribute to understanding the connection between X-ray and UV absorption in \\balqs. The detection of four narrow, probably line-locked absorption systems indicates that high resolution UV spectroscopy covering these features would enable a more detailed investigation into the partial covering properties of these UV absorption systems, for comparison with the partial covering of the X-ray gas. Furthermore, if the bulk of the X-ray absorbing gas lies at smaller radii than the UV BAL gas, then the X-ray absorbers might be expected to vary, in ionization state, column density, and velocity, on more rapid timescales than the UV BALs. There are now examples of \\balqs\\ with evidence for X-ray absorption variability: \\pgbal, \\apm, and perhaps \\pgone\\ \\citep{Chartas2000,ChBrGa2003}. These three objects compose a significant fraction of the handful of \\balqs\\ that have X-ray data adequate for variability studies; this demonstrates that the X-ray regime may be a fertile area for exploring details of quasar winds inaccessible in the UV.  Furthermore, to interpret X-ray variability accurately, the hydrodynamic disk-wind models \\citep[e.g.,][]{PrStKa2000} may need to be extended to smaller radii ($<10^{15}$~cm) with higher resolution, and they must also more explicitly include the effects of Compton and X-ray edge and line pressure.  For the highly ionized gas expected at small radii,  these will dominate the radiation force over UV line pressure \\citep{CheNet2003a,CheNet2003b}. More appropriate models for low-energy absorbers which include the potentially substantial opacity from velocity-broadened X-ray BALs are also required.  This line opacity could be comparable to the edge opacity in certain cases. Additional items for an X-ray observer's wish list from theoretical modeling are specific predictions for the strength, profiles, and variability properties of iron emission and absorption lines."
    },
    "0311/astro-ph0311292_arXiv.txt": {
        "abstract": "I review recent progress in understanding the formation of the first stars and quasars. The initial conditions for their emergence are given by the now firmly established model of cosmological structure formation. Numerical simulations of the collapse and fragmentation of primordial gas indicate that the first stars formed at redshifts $z\\simeq 20 - 30$, and that they were predominantly very massive, with $M_{\\ast}\\ga 100 M_{\\odot}$. Important uncertainties, however, remain. Paramount among them is the accretion process, which builds up the final stellar mass by incorporating part of the diffuse, dust-free envelope into the central protostellar core. The first quasars, on the other hand, are predicted to have formed later on, at $z\\sim 10$, in more massive dark matter halos, with total masses, $\\sim 10^8 M_{\\odot}$, characteristic of dwarf galaxies. ",
        "introduction": "How and when did the first sources of light form in the universe? Within the framework of modern cosmology, we have learned that the first stars and quasars formed at the end of the so-called ``dark ages'', a few $10^{8}$~yr after the big bang (e.g., Barkana \\& Loeb 2001; Miralda-Escud\\'{e} 2003). The cosmic dark ages began $\\sim$~500,000~yr after the big bang when the photons of the cosmic microwave background (CMB) were emitted. At this time, the CMB photons redshifted into the infrared wavelength band so that they were no longer energetic enough to ionize hydrogen atoms, and could henceforth propagate freely. From the viewpoint of a human observer, the universe descended into a state of complete darkness, that lasted until the cosmic Renaissance of first light occurred much later on.   At these early times, the universe did not yet contain any gravitationally bound structures, and its state was simple enough to allow a description in terms of precise, linear physics (e.g., Hu \\& Dodelson 2002). We know from observations of the highest redshift quasars and galaxies that the dark ages must have ended prior to $z\\sim 7$ (e.g., Hu et al. 2002; Fan et al. 2003). Whereas we can directly probe the state of the universe at $z\\sim 1000$, by measuring CMB anisotropies, and at $z\\la 6$, by observing quasar absorption lines and Lyman-$\\alpha$ emission from galaxies, we know little about the formative epoch in between. Thus, this era is the final frontier of observational cosmology, and the crucial question is (e.g., Barkana \\& Loeb 2001; Loeb \\& Barkana 2001; Miralda-Escud\\'{e} 2003): {\\it How and when did the dark ages end?}  Within the now firmly established model of cosmological structure formation (e.g., Spergel et al. 2003), the first stars are predicted to have formed at redshifts $z\\simeq 20 - 30$ in dark matter (DM) halos of total mass $\\sim 10^{6}M_{\\odot}$ (e.g., Couchman \\& Rees 1986; Ostriker \\& Gnedin 1996; Tegmark et al. 1997; Abel et al. 1998; Yoshida et al. 2003a). Numerical simulations of the collapse and fragmentation of metal-free gas, assumed to contain no dynamically significant magnetic fields yet, have indicated that the first stars, the so-called Population~III (e.g., Bond 1981), were predominantly very massive, with $M_{\\ast} \\ga 100 M_{\\odot}$ (e.g., Bromm, Coppi, \\& Larson 1999, 2002; Abel, Bryan, \\& Norman 2000, 2002; Nakamura \\& Umemura 2001). Many uncertainties, however, remain. Two key problems are: (i) The physics of accretion from a dust-free envelope, which determines the final stellar mass, is not yet well understood (e.g., Omukai \\& Palla 2001, 2003; Tan \\& McKee 2003). (ii) A complete theory of primordial star formation would allow us to predict the exact functional form of the Population~III initial mass function (IMF). Again, a reliable determination is still eluding us (see, e.g., Nakamura \\& Umemura 2001, 2002; Omukai \\& Yoshii 2003 for first attempts). I will address both issues in \\S 3.  The first quasars, on the other hand, powered by accretion onto supermassive black holes (SMBHs), are expected to have formed later on, at $z\\sim 10$, in more massive, dwarf-sized systems (e.g., Haiman \\& Loeb 2001; Bromm \\& Loeb 2003a). Numerical simulations of high-redshift SMBH formation indicate that these could have formed only {\\it after} a significant episode of previous star formation (e.g., Bromm \\& Loeb 2003a; Di Matteo et al. 2003). The long-standing question (e.g., Silk \\& Rees 1998) of what came first, stars or quasars, would then be answered in favor of the former.  The formation of the first luminous sources had important implications for the evolution of the intergalactic medium (IGM). Massive Population~III stars were efficient producers of ionizing photons (e.g., Tumlinson \\& Shull 2000; Bromm, Kudritzki, \\& Loeb 2001b; Schaerer 2002, 2003; Tumlinson, Shull, \\& Venkatesan 2003). The contribution from these stars, thus, could have been important in reionizing the IGM early on, at $z\\ga 15$, as may be required by the recent {\\it Wilkinson Microwave Anisotropy Probe (WMAP)} data (Kogut et al. 2003; Spergel et al. 2003) on the optical depth to Thomson scattering (e.g., Cen 2003a,b; Ciardi, Ferrara, \\& White 2003; Haiman \\& Holder 2003; Holder et al. 2003; Kaplinghat et al. 2003; Sokasian et al. 2003a,b; Wyithe \\& Loeb 2003a,b). The second key feedback effect concerns the production of the first heavy elements, and their subsequent dispersal by supernova (SN) explosions into the pristine IGM (e.g., Gnedin \\& Ostriker 1997; Madau, Ferrara, \\& Rees 2001; Mori, Ferrara, \\& Madau 2002; Bromm, Yoshida, \\& Hernquist 2003; Wada \\& Venkatesan 2003). Once the IGM has been metal-enriched above a minimum level, star formation is predicted to shift from a predominantly high-mass (Population~III) to a lower-mass (Population~II) mode (e.g., Omukai 2000; Bromm et al. 2001a; Bromm \\& Loeb 2003b; Schneider et al. 2003a). The nature of this fundamental transition, whether gradual or sudden, will depend on the detailed IGM enrichment history (e.g., Schneider et al. 2002; Mackey, Bromm, \\& Hernquist 2003; Scannapieco, Schneider, \\& Ferrara 2003; Ricotti \\& Ostriker 2004; Yoshida, Bromm, \\& Hernquist 2004).  Very recently, novel empirical probes have begun to test our theoretical ideas on star and galaxy formation in the early universe. Next to the {\\it WMAP} data on the ionization state of the IGM, these probes include determinations of the chemical abundance pattern in extremely metal-poor Galactic halo stars (e.g., Christlieb et al. 2002), measurements of the near IR cosmic background radiation (e.g., Santos, Bromm, \\& Kamionkowski 2002; Salvaterra \\& Ferrara 2003), and the prospect of detecting high-redshift gamma-ray bursts (GRBs) with the {\\it Swift} satellite to be launched in 2004 (e.g., Lamb \\& Reichart 2000; Bromm \\& Loeb 2002).  In this review, I will briefly summarize the current state of key aspects in this rapidly evolving field (see Bromm \\& Larson 2004 for a more detailed discussion). ",
        "conclusions": "The first stars and quasars, their formation and their cosmological implications, are a fascinating subject at the frontier of modern cosmology. Up to now, most of our efforts have been theoretical, but we are approaching the point where observations can test our ideas. This is very significant, as we are bound to learn important lessons about the physical state of the universe at the end of its dark age. These empirical probes are provided by both high-redshift observations, and complementary to this, by the `near-field cosmology' (Freeman \\& Bland-Hawthorn 2002) of scrutinizing the chemical abundance patterns in extremely metal-poor stars in our local cosmic neighborhood. One can only wonder what exciting discoveries are awaiting us, and it is a great privilege to be a part of this grand endeavor."
    },
    "0311/astro-ph0311547_arXiv.txt": {
        "abstract": "We investigate the nature of the diffuse Galactic emission in the Wilkinson Microwave Anisotropy Probe (WMAP) temperature anisotropy data.  Substantial dust-correlated emission is observed at all WMAP frequencies, far exceeding the expected thermal dust emission in the lowest frequency channels (23, 33, 41 GHz).  The WMAP team (Bennett \\etal) interpret this emission as dust-correlated synchrotron radiation, attributing the correlation to the natural association of relativistic electrons produced by SNae with massive star formation in dusty clouds, and deriving an upper limit of 5\\% on the contribution of Draine \\& Lazarian spinning dust at K-band (23 GHz). We pursue an alternative interpretation that much, perhaps most, of the dust-correlated emission at these frequencies is indeed spinning dust, and explore the spectral dependence on environment by considering a few specific objects as well as the full sky average. Models similar to Draine \\& Lazarian spinning dust provide a good fit to the full-sky data. The full-sky fit also requires a significant component with free-free spectrum uncorrelated with \\Halpha, possibly hot ($\\sim 10^6$K) gas within $30\\degree$ of the Galactic center.   \\emph{Subject headings: } cosmic microwave background --- diffuse radiation --- dust, extinction --- ISM: clouds --- radiation mechanisms: thermal --- radio continuum: ISM ",
        "introduction": "\\label{sec_intro}  The first year of data from the Wilkinson Microwave Anisotropy Probe (WMAP) has sparked a revolution in the study of cosmology and the early universe (Bennett \\etal\\ 2003, Hinshaw \\etal\\ 2003, Spergel \\etal\\ 2003). In addition, it has provided full sky maps of the interstellar medium at $23-94\\GHz$ at a sensitivity of $\\sim 200\\microK$ per $7'$ pixel at resolutions much higher than previous full-sky observations. The quality of the data has rekindled the debate about the origin of ``anomalous'' dust-correlated microwave emission (Kogut \\etal\\ 1996, de Oliveira-Costa 1997,1998,1999,2002,2003; Finkbeiner \\etal\\ 2002, 2003; Banday \\etal\\ 2003; Lagache 2003) dubbed ``Foreground X'' by de Oliveira-Costa \\etal\\ (2002).  The WMAP team relied on a Maximum Entropy Method (MEM) analysis to determine that spinning dust\\footnote{The idea of grain rotational emission was first discussed by Erickson (1957). More recently, after the discovery of the population of ultra small grains, Ferrara \\& Dettmar (1994) noted that the rotational emission from such grains may be observable, but their treatment incorrectly assumed Brownian thermal rotation of grains.} with the Draine \\& Lazarian (1998) cold neutral medium (CNM) spectrum accounts for less than 5\\% of the emission in any WMAP waveband.  We show that for different models of spinning dust, these constraints do not apply.  Indeed, in most of the Galaxy, spinning dust emission (or some other mechanism such as magnetic dipole emission, Draine \\& Lazarian 1999) explain the data just as well.  We emphasize that the WMAP team's primary objective was to prevent Galactic foreground emission from contaminating their cosmological results based on the CMB anisotropy.  They have done a superb job of modeling the foregrounds, regardless of whether the physical interpretation of the subtracted temperature anisotropy signal is correct. The WMAP polarization data also provide important cosmological constraints as well as additional information about foregrounds. For example, if the Bennett \\etal\\ hard synchrotron component can be shown to be $\\sim30$\\% polarized, then spinning dust would become an unlikely explanation (see Lazarian \\& Draine 2000).  On the other hand, if it is unpolarized, the synchrotron would be ruled out. Something between the two extremes is likely, and this will help us determine the relative contributions of synchrotron and spinning dust.  As the polarization data are not yet public, discussion of polarization will be deferred to a later paper. ",
        "conclusions": "We have performed an analysis of the WMAP temperature anisotropy data, using free-free, soft synchrotron, FDS thermal dust, and spinning dust as templates.  We find:  1.  The FDS thermal (non-magnetic) dust prediction is good on average over the high-latitude sky at 94 GHz, and even in dark cloud LDN1622 and \\HII\\ region \\LPH.  2.  Free-free is fit by $T = 5500$K ionized gas (rather cool, but see e.g. Heiles \\etal\\ 2000), plus a haze toward to the Galactic center with free-free spectrum.   This is the first detection of emission from this component.  Further study is needed to test whether this component is indeed free-free or perhaps a very hard synchrotron component.  3.  The residual dust-correlated emission is well fit by $T^3_{dust} \\times$ dust column density over most of the sky.  Such a model, along with free-free, soft synchrotron, and FDS thermal dust, accounts for $\\sim98\\%$ of the variance in the WMAP data.  4.  There is no conflict between the general idea of spinning dust emission and the WMAP data.  In spite of the Bennett \\etal\\ (2003) statement that the Draine \\& Lazarian CNM model accounts for less than 5\\% of the ISM emission (which is true), spinning dust remains a reasonable hypothesis both morphologically and spectrally.  5.  The Green Bank Galactic Plane Survey GPA data at 8 and 14 GHz (Langston \\etal\\ 2000) rule out hard synchrotron as a majority component in the diffuse ISM in the WMAP bands (see companion paper, Finkbeiner, Langston, \\& Minter 2003).  This leaves spinning dust as a viable alternative.  The WMAP data alone cannot establish the relative levels of spinning dust, magnetic dust, and hard synchrotron emission, since any of these theories can be adjusted to explain the data. However, objects such as LDN1622 (Figure 3) are clearly not dominated by synchrotron in the WMAP bands, and are well fit by a superposition of spinning dust models. Therefore, the dominance of spinning dust at 23 GHz is a good working hypothesis.  Further data (including surveys at lower frequencies, better WMAP signal to noise, and WMAP polarization data coming soon) will help lift the remaining degeneracies and determine the identity of the dust-correlated microwave emission from``Foreground(s) X.''    I am indebted to Bruce Draine, Ed Jenkins, and David Schlegel for encouragement and helpful advice. Roger Hildebrand called the WMAP spectrum of LDN1622 to my attention, and provided other insights. Lyman Page, Chuck Bennett, and Gary Hinshaw provided helpful comments on the proper use of the WMAP data in this work.  Comments from an anonymous referee substantially improved the text. Some of the results in this paper were derived using HEALPix (G\\'orski, Hivon, and Wandelt 1999). This research made use of the NASA Astrophysics Data System (ADS) and the IDL Astronomy User's Library at Goddard\\footnote{http://idlastro.gsfc.nasa.gov/}.  DPF is a Hubble Fellow supported by HST-HF-00129.01-A. \\newpage"
    },
    "0311/astro-ph0311221_arXiv.txt": {
        "abstract": "{ We propose an extension of  the semi-analytical solutions derived by Lin et al. (1965) describing the two-dimensional homologous collapse of a self-gravitating rotating cloud having uniform density and spheroidal shape, which includes magnetic field (with important restrictions) and thermal pressure. The evolution of the cloud is  reduced to three  time dependent ordinary equations  allowing to conduct a quick and preliminary investigation of the cloud dynamics during the precollapse phase, for a wide range of parameters. We apply our model to the collapse of a  rotating and magnetized oblate and prolate isothermal core. Hydrodynamical numerical simulations are performed and comparison with the semi-analytical solutions is discussed. Under the assumption that all cores are similar, an apparent cloud axis ratio distribution is calculated from the sequence of successive evolutionary states for each of a large set of initial conditions. The   comparison with the observational distribution of the starless dense cores belonging to the catalog of Jijina et al. (1999) shows a  good agreement for the rotating and initially prolate cores (aspect ratio $\\simeq 0.5$) permeated by an helical magnetic field ($\\simeq 17-20 \\mu$G for a density of $\\simeq 10^4$ cm$^{-3}$). ",
        "introduction": "The stars form through the gravitational collapse of dense cores located in the molecular clouds. These condensations are complex and involve gravity, hydrodynamics, magnetic  and thermal processes. %  The core dynamics has been  investigated numerically by several authors. Although the numerical approach undoubtedly allows a deep understanding of the cloud dynamics, it does not replace the analytical studies, first because analytical work gives more explicit descriptions  and second because the complexity of numerical methods keeps increasing, making  a full exploration of the space parameters  and their use by non numerical experts difficult.  Analytically, most of the efforts have been dedicated to the description of  the isothermal gravitational collapse (Larson 1969, Penston 1969a, 1969b, Hunter 1977, Shu 1977, Whitworth \\& Summers 1985, Bouquet et al. 1985). These solutions which assume self-similarity are very useful to get a physical hint on how the collapse proceeds and to give strong benchmarks to the numerical studies. However in spite of its importance, very few analytical tools are available to describe the dynamics of the precollapsing or starless phase. The present paper is an attempt to fill this gap.   In this paper, we first present  new solutions of the gravo-magnetic condensation. They are an extension of  solutions obtained by Lin et al. (1965) for the gravitational part and simultaneously by Aburihan et al. (2001) and by the author (Hennebelle 2001) for the MHD part. We then  use these solutions to study the phenomenology of the collapse  of oblate and prolate, isothermal, magnetized and non magnetized clouds and compare the result with the evolution obtained with numerical simulations. The solutions allow to calculate the cloud shape evolution. With the assumption that all cores are similar, distributions of apparent axis ratio are predicted. We then compare these distributions with the apparent axis ratio distribution of the starless dense cores belonging to the catalog of Jijina et al. (1999).   Indeed  the geometry of the dense cores is an important issue since it depends on the physical processes that drive their evolution. Previous investigations by Myers et al. (1991) and Ryden  (1996) conclude that the observational data are more consistent with axisymmetric prolate cores than with axisymmetric oblate cores. The same conclusion is reached by the recent study of Curry (2002) with a different sample of cores and by Hartmann (2002) who shows that the major axis of the cores in the Taurus cloud is preferentially aligned with the filament in which they are embedded. Two recent studies, (Jones et al. 2001 and Goodwin et al. 2002) analyse the observational apparent axis ratio distribution derived from the catalog of Jijina et al. (1999) and conclude that the core apparent axis ratio distribution is compatible with  triaxal cores but being more nearly oblate than prolate. From a theoretical point of view, the equilibrium solutions of a rotating and magnetized cloud (Mouschovias 1976, Tomisaka et al. 1988) present an oblate shape except if there is a substantial toroidal magnetic component   (Tomisaka 1991) or if the structure of the poloidal magnetic field is such that it compresses the cloud (Stahler \\& Curry 2001). Curry (2000)  also shows that two-dimensional solutions of the isothermal Lane-Emden equation present a prolate configurations. On the other hand, recent studies show (Nakamura et al. 1993, Fiege \\& Pudritz 2000a, 2000b) that prolate cores can be formed from a magnetized  filament (with poloidal and toroidal field) through a combination of the gravitational and the magnetic sausage instabilities. Fiege \\& Pudritz  (2000c) also demonstrate that the axis ratio of the prolate magnetized cores is compatible with the observational apparent axis ratio distribution. \\\\    In the second Section of the paper, we derive the solutions and reduce the self-gravitating MHD equations to three time dependent ordinary equations. We discuss their mathematical and physical properties, their potential and intrinsic weaknesses. In the third Section, we solve  these  ordinary equations numerically and describe the phenomenology of the collapse of  an isothermal dense core. We also compare the semi-analytical solutions with the results of the hydrodynamical numerical simulations of an isothermal dense core. The fourth Section is devoted to the comparison  between the theoretical and observational apparent axis ratio distributions. A discussion and summary are given in the fifth part. ",
        "conclusions": "New solutions of the gravo-magnetic contraction have been presented. They describe the evolution of a uniform density cloud with spheroidal shape and can be applied in various situations. These solutions allow to explore a wide set of parameters and constitute a useful complement to numerical simulations. They can also be used as benchmarks in problems involving gravitational collapse.  We have explored the phenomenology of the gravitational contraction of isothermal dense cores described by these semi-analytical solutions with special emphasis on the axial and  radial  rebounds. Numerical simulations have been performed to compare these solutions with the evolution of an isothermal initially uniform cloud and it has been found that they  can reproduce  some aspects of cloud evolution  during the precollapse phase. The most important disagreement is the cloud outer part which is significantly influenced by the rarefaction wave not taken into account in the semi-analytical solutions.  For each of a large set of initial values, we have calculated the apparent axis ratio distribution resulting from the evolution of the core and tentatively, compared it to the observational one derived from the catalog of Jijina et al. (1999).  We find that the axis ratio distribution of a collapsing rotating, prolate (initial aspect ratio $\\simeq 0.5$) magnetized ($\\simeq 17-20 \\mu$G  for a density of $\\simeq 10^4$ cm$^{-3}$)  filament is in  good agreement with the observational distribution (see Fig.~\\ref{stat_bo}). Definite conclusion cannot be drawn at this stage however  since our model is very simplified and (mhd) numerical studies are required to confirm the results. Also   the observational core sample is not homogeneous and the statistical significance  remains poor. Finally   alternative distributions of evolving or static cores can  fit the observational data as well. Nevertheless, we believe that prolate cores permeated by an helical magnetic field deserve  further investigations."
    },
    "0311/astro-ph0311367_arXiv.txt": {
        "abstract": "{  We report a detailed analysis of the coronal abundance of 4 stars with varying levels of activity and with accurately known photospheric abundances. The coronal abundance is determined here using a line flux analysis and a full determination of the differential emission measure. Photospheric abundance values are taken from literature works. Previous coronal abundance determinations have generally been compared to solar photospheric abundances; from this a number of general properties have been inferred, such as the presence of a coronal metal depletion with an inverse First Ionization Potential correlated with activity level. Here we show that, when coronal abundances are compared with real photospheric values for the individual stars, the resulting pattern can be very different. Also, we present evidence that, in some cases, the coronal metal abundance may not be uniform in the corona; in particular it can vary with the temperature of the emitting plasma. ",
        "introduction": "Following the initial results from the ASCA-SIS instrument \\citep[e.g.][]{whi94,dra96}, which showed that the low-resolution X-ray spectra of active coronal sources where often best fit with a low metal abundance (typically about 0.3 the solar photospheric abundance) the issue of the metal abundance of the coronal plasma has become one of the main areas of interest for X-ray spectroscopy of active stars. Due to the limited spectral resolution of the ASCA spectra, the determination of abundance for elements other than Fe was often not accurate \\citep{fav03}. The availability of the XMM-\\emph{Newton} and \\emph{Chandra} high-resolution X-ray spectra of coronal sources has allowed to observe individual, resolved spectral lines of a number of elements in the corona of active stars, and thus to address the issue of the abundance of a number of elements. Unfortunately up to now the results have not been as conclusive as initially expected. One reason is that often the same data analyzed with different approaches result in significantly different value for the elements abundances \\citep[see][for a number of examples]{fav03}.  One of the issues most often addressed by coronal abundance studies is whether a bias with the First Ionization Potential (FIP) of the element is present. Such ``FIP effect'' is found in the solar corona, where the abundance of elements with low FIP ($< 10$~eV; e.g. Mg, Si, Fe) is observed to be enhanced, in the corona, with respect to their photospheric values, while elements with high FIP ($\\ga 10$~eV; e.g.  O, Ne, Ar) remain at abundances similar to the photospheric values. Such behavior is observed in the upper solar atmosphere at temperatures higher than $\\sim 1$~MK, both in the non-active regions and in full-disk spectra \\citep{lam95,feld00}.  EUVE observations showed the same FIP bias to be present in stars with low activity levels, such as $\\alpha$~Cen~AB \\citep[G2V+K1V,][]{drake97}, and possibly (only limited knowledge of their photospheric abundance was available at that time) $\\xi$~Boo~A \\citep[G8V,][]{lam99} and $\\epsilon$~Eri \\citep[K2V,][]{lam96}.  EUVE spectra of more active stars did not show any FIP effect, with however a rather low coronal Fe abundance with respect to the solar photosphere, which suggested some metal depletion in the corona of the active stars. \\citet{feld00} review some of the results on solar and stellar coronal abundances previous to \\emph{Chandra} and XMM-\\emph{Newton}.  \\begin{table*} \\caption{Observation log}\\label{obslog} \\begin{center} \\begin{footnotesize} \\begin{tabular}{lccccrcrrrr} \\hline \\hline {Object} & {HD \\#} & {SpT} & {Instrument} & {\\gl \\gl (\\AA)} & {$t_{\\rm exp}$ (ks)} & {Date} & $d$ (pc) & $L_{\\rm X}$ & $L_{\\rm X}$/$L_{\\rm bol}$ & $L'_{\\rm X}$ \\\\ \\hline Procyon & 61421 & F5IV & \\emph{Chandra}/LETG & 3--175 & 69.6 & 6 Nov 1999 & 3.50 & 8.89E27 & 3.39E-7 & 9.41E26\\\\ Procyon & 61421 & F5IV & \\emph{Chandra}/LETG & 3--175 & 69.7 & 7 Nov 1999 & `` & `` & `` & ``\\\\ $\\epsilon$~Eri & 22049 & K2V & \\emph{Chandra}/LETG & 3--175 & 105.3 & 21 Mar 2001 & 3.22 & 1.79E28 & 1.42E-5 & 9.65E27\\\\ $\\lambda$~And & 222107 & G8III/? & \\emph{Chandra}/HETG & 1.7--27 & 81.9 & 17 Dec 1999 & 25.8 & 2.74E30 & 3.25E-5 & 1.73E30\\\\ V851~Cen & 119285 & K2IV-III/? & XMM/RGS & 4--38 & 29.2 & 7 Mar 2003 & 76.2 & 6.46E30 & 3.01E-4 & 4.47E30\\\\ \\hline \\end{tabular} \\end{footnotesize} \\end{center} Note: Distances ($d$) from \\citet{hippa}. $L_{\\rm X}$ (erg s$^{-1}$)  in the range 5--100~\\AA\\ (0.12--2.4 keV), the range covered by the ROSAT/PSPC. $L'_{\\rm X}$ (erg s$^{-1}$) in the range 6--20~\\AA\\ (0.62--2.1 keV), common to RGS, HETGS and LETGS. \\end{table*}  The new generation of X-ray satellites XMM-\\emph{Newton} and \\emph{Chandra} offers access to high resolution spectra in a wavelength region where numerous spectral lines from different elements are formed at coronal temperatures. This permits to measure the fluxes of individual lines, in principle allowing a much better determination of the coronal thermal structure and abundances. Initial results on coronal abundances determined from data from the two observatories have been reviewed by \\citet{dra02} and \\citet{fav03}. These new results seem to point towards a scenario with stars with low activity levels showing, when their coronal and photospheric abundances are compared, a solar-like FIP effect \\citep[case of $\\alpha$~Cen,][]{raa03} or no FIP bias \\citep[for Procyon,][]{raa02}. Active stars (like HR~1099, II~Peg, AR~Lac or AB~Dor), with a quite different coronal thermal structure, dominated by material emitting at temperatures of $\\sim 10$~MK, show a pattern characterized by a general depletion of elements with low FIP (when they are compared with solar photospheric values) while the elements with high FIP remain at solar (or sometimes higher) photospheric values \\citep{dra01,bri01,huen01,huen03,sanz03} although photospheric abundances of other elements than Fe are not known for HR~1099. \\citet{aud03} analyzed XMM-\\emph{Newton} observations of 5 active binary systems, concluding that low-FIP elements are subject to larger depletion in more active stars, while high-FIP elements remain at similar levels (although photospheric abundances were available only for one of the objects in their sample, namely $\\lambda$~And).  In general, the comparison of stellar coronal abundances with solar photospheric values is of little use, given the wide range of abundances found among late type stars \\citep[e.g.][]{cay01}.  Most of the stars for which coronal abundances have been determined have no published photospheric abundances, so that a true test of e.g. FIP-related effects cannot be done. This is specially true for active stars, for which the large rotational velocity complicates the measurements of photospheric abundances. In most cases the available photospheric abundances are limited to Fe \\citep[e.g.][]{ran93}, and even these in many cases are likely to be affected by systematic errors.  Moreover, different methods to determine the distribution of material with temperature in the corona (the so called ``emission measure distribution'', EMD) sometimes produce contradictory results, depending on the analysis technique employed \\citep{fav03,sanz03}.  In present work we determine the coronal thermal structure and abundances of four stars with different activity levels (in terms of $L_{\\rm X}/L_{\\rm bol}$) observed with relatively high statistics with \\emph{Chandra} and XMM-\\emph{Newton} (see Table~\\ref{obslog}), and for which there are accurate photospheric abundances measurements available, with the aim to explore the relation between activity levels and coronal relative abundance patterns. The chosen stars are Procyon (one of the stars with lowest $L_{\\rm X}/L_{\\rm bol}$ values), the intermediate activity star $\\epsilon$~Eri, the long-period active binary $\\lambda$~And and the active binary V851~Cen. These four objects are among the few stars for which both high-resolution coronal spectra with good statistics, and measurements of photospheric abundances for several elements, are available.  The observations are described in Sect.~2, while a brief description of the analysis techniques is presented in Sect.~\\ref{sec:analysis}. The scientific results for each of the star are discussed in Sect.~\\ref{sec:results}, followed by the discussion and conclusions in Sect.~\\ref{sec:conclusions}.  \\begin{figure*} \\begin{center} \\includegraphics[angle=90,width=0.4\\textwidth]{proc_lc.ps} \\includegraphics[angle=90,width=0.4\\textwidth]{eeri_lc.ps} \\includegraphics[angle=90,width=0.4\\textwidth]{land_lc.ps} \\includegraphics[angle=90,width=0.4\\textwidth]{v851cen_lc.ps} \\end{center} \\caption{Light curves of Procyon (LETGS),  $\\epsilon$~Eri (LETGS), $\\lambda$~And (HETGS) and V851~Cen (XMM/MOS 2) with 1-$\\sigma$ error bars. The upper axis reports the photometric phase ($\\epsilon$~Eri) and orbital phase ($\\lambda$~And and V851~Cen, with $T_0$ corresponding to secondary star located behind the primary star). \\label{lcurves}} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusions}  The results found in the present sample show that the issue of coronal abundances in stars other than the Sun is complex and still poorly understood. Solar coronal abundance is by itself a complex issue, with the actual values depending on which part of the corona is observed (with coronal holes, active regions and flares showing quite different abundance patterns). Results obtained from full disk spectra (thus comparable to stellar spectra) show enhancements by at least a factor of 3 in elements with ${\\rm FIP}\\la 10$~eV \\citep{lam95}. Young solar active regions show coronal abundance compatible with the photospheric values, but they soon develop a FIP bias that reaches average coronal values in two to three days \\citep{she96,wid01}. The FIP effect that is observed in the solar corona seems to disappear with increasing activity level of the star, as it has been shown in cases like II~Peg \\citep{huen01}, AR~Lac \\citep{huen03}, AB~Dor \\citep{sanz03}, $\\lambda$~And and V851~Cen.  A solar-like FIP effect appears to be present in $\\alpha$~Cen~A \\citep{raa03}, a star very similar to the Sun in many physical parameters, as well as in other solar analogs with low levels of activity, such as $\\pi^1$~UMa and $\\chi^1$~Ori \\citep{gud02}, although only the Fe photospheric abundance is known for these two objects (so that abundance patterns assume a solar-like mix). A mid-activity level star, $\\epsilon$~Eri, could be showing the progressive disappearance of the FIP effect, with elements like Mg and Fe already showing photospheric abundances in the corona, and leaving a marginal coronal abundance enhancement only for Ca and Ni.  Finally, Procyon seems to be quite different from the Sun, with no FIP-related trends evident, although there could be an increment of Fe abundance with the coronal temperature relative to other elements with higher FIP (see Sect.~\\ref{sec:procyon}). The FIP effect detected in the solar corona above quiet regions grows with temperature, in a range similar to that covered here for Procyon \\citep{feld00}.  The interpretation of the metal depletion in the corona of active stars that has been given in the past from initial results is somewhat in contradiction with the present results.  Many of the stars for which a ``metal abundance deficiency syndrome'' (MAD) has been claimed do not have measurements of the photospheric abundance available (or they are limited to Fe abundance determinations), and published work is limited to the comparison of stellar coronal patterns with solar photospheric abundances (a rather arbitrary practice). The case of $\\lambda$~And is quite remarkable, with compatible coronal and photospheric abundances for all elements tested, and a definitely non-solar photospheric abundance pattern. V851~Cen, with similar coronal thermal structure (Fig.~\\ref{emds}) shows a marginal coronal abundance depletion. Other cases like II Peg \\citep{huen01} or AR Lac \\citep{huen03} show also moderate depletion when they are compared with the respective photospheric values, but strong metal depletion seems to be clear only for AB~Dor \\citep[][and references therein]{sanz03}.  However AB~Dor is a fast rotator, and the only available determination of photospheric abundances, carried out by \\citet{vil87}, needs further confirmation given the difficulty for the measurements in a spectrum so much affected by rotational broadening.  While the absence of FIP effect in active stars, at least at the solar level, is well supported at this point, the presence of a general ``MAD effect'' or an ``inverse FIP effect'' increasing with activity is not yet clearly established, given the discrepancies found between stars with quite similar levels of activity and thermal coronal structure like AB~Dor, $\\lambda$~And and V851 Cen.  Moreover, the abundance pattern could also be changing in different parts of the corona, depending on the coronal temperature, as present in Procyon. A larger sample of stars with both photospheric and coronal abundances, determined in a homogeneous way, would be necessary in order to establish a general trend between coronal abundances and activity levels."
    },
    "0311/astro-ph0311151_arXiv.txt": {
        "abstract": "We report on new simultaneous phase resolved spectroscopic and polarimetric observations of the polar (AM Herculis star) V834~Cen during a high state of accretion. Strong emission lines and high levels of variable circular and linear polarization are observed over the orbital period.  The polarization data is modelled using the Stokes imaging technique of Potter et al. The spectroscopic emission lines are investigated using the Doppler tomography technique of Marsh \\& Horne and the Roche tomography technique of Dhillon and Watson.  Up to now all three techniques have been used separately to investigate the geometry and accretion dynamics in Cataclysmic Variables. For the first time, we apply all three techniques to simultaneous data for a single system. This allows us to compare and test each of the techniques against each other and hence derive a better understanding of the geometry, dynamics and system parameters of V834 Cen.  All three techniques are consistent with an interpretation in which a ballistic stream extends to a minimum of $\\sim$ 40 degrees in azimuth around the white dwarf before becoming threaded by the magnetic field lines. Interestingly, the observed ballistic Doppler velocites do not show a reduced $v_{y}$ component, as found in Doppler imaging of other AM Her systems. Furthermore, the secondary star in V834 Cen shows more He II (4686 \\AA \\ ) emission on its leading inner face, as opposed to the trailing face like in other AM Her systems. We propose that the accretion shock preferentially illuminates the leading face of the secondary star. In addition, the ballistic stream does not obscure the leading face of the secondary from the accretion shock, and, in fact, our Doppler maps show that the ballistic stream is a strong He II (4686 \\AA \\ ) emission source in itself and thus adds to the illumination of the leading face of the secondary. ",
        "introduction": "The standard picture of a cataclysmic variable (CV) is a binary system in which there is a Roche-lobe filling red dwarf (the secondary star) and an accreting white dwarf (the primary), see \\eg Warner 1995 for a review of Cataclysmic Variables. In AM Her systems, also known as polars, the white dwarf primary has a sufficiently strong magnetic field to lock the system into synchronous rotation and to prevent an accretion disc from forming.  Instead, the material overflowing the Roche lobe of the secondary initially continues on a ballistic trajectory until, at some distance from the white dwarf, the magnetic pressure overwhelms the ram pressure of the ballistic stream. From this point onwards the accretion flow is confined to follow the magnetic field lines of the white dwarf, see \\eg Cropper 1990 or Warner 1995 for a full review of magnetic CVs.  The now supersonic accreting material suddenly becomes subsonic at a shock region, which forms at some height above the white dwarf surface. The shock-heated material reaches temperatures of $\\sim$ 10-50 keV and is ionised. The hot plasma cools by two different mechanisms as it settles onto the white dwarf, resulting in a density and temperature stratification in the post-shock region. The two cooling mechanisms are X-ray cooling, in the form of bremsstrahlung radiation and, with sufficient magnetic fields, cyclotron cooling in the form of optical/IR cyclotron radiation.  V834~Cen was discovered as an X--ray source (E1405-451 = H1405-45) with the HEAO 1 and Einstein satellites by Jensen et al. (1982), and Mason et al. (1993) identified it with a stellar object. It was later given the classification of an AM Her system when Tapia (1982) and Bailey et al. (1983) reported strong linear and circular phase dependent polarization. Extensive polarimetric and spectroscopic observations by Cropper et al. (1986) and Rosen et al. (1987) followed.  In 1987, Wickramasinghe, Tuohy \\& Visvanathan detected the presence of strongly phase-dependent absorption features in the spectrum of V834 Cen. They interpreted the features as Zeeman components of H$\\alpha$, originating from a cloud of cool gas surrounding the accretion shock. Puchnarewicz et al. (1990) and Schwope \\& Beuermann (1990) also reported the presence of pure photospheric Zeeman absorption lines during a low state. Ferrario et al. (1992) also reported observations of V834~Cen when it was in a low state. They detected Zeeman absorption lines from the photosphere of the white dwarf and cyclotron emission features from the accretion shock, indicating that there was still some accretion from the companion star. These authors were therefore able to unambiguously determine the magnetic field strength and structure of V834~Cen, which, in turn, makes it an ideal object for Stokes imaging (Potter, Hakala \\& Cropper 1998).  Mauche (2002) reported on EUVE (Extreme Ultraviolet Explorer) observations of V834~Cen during two epochs of different accretion states. He gave a qualitative explanation for the EUV light curves by invoking a simple model of accretion from a ballistic stream along field lines of a tilted magnetic field centered on the white dwarf. During the higher accretion state, accretion would occur over a broader range of azimuths than during the lower accretion state.  In this paper, a new version of Stokes imaging, which was described in Potter, Hakala \\& Cropper (1998) in its original form, is applied to polarimetric observations of V834~Cen. Here we also describe our new, more robust, methodology for Stokes imaging, which allows objective mapping of the cyclotron emission regions in magnetic cataclysmic variables in terms of their location, shape and size.  In addition, we apply Doppler mapping, of Marsh \\& Horne (1988), to spectroscopic data taken simultaneously and four months later than the polarimetry. The Doppler solutions are then modelled using a single particle trajectory code and compared to the results derived from Stokes imaging. Finally we apply Roche tomography, of Dhillon \\& Watson (2001) to the same spectroscopic data in order to gain knowledge of the irradiation pattern of the inner irradiated surface of the companion star and compare the results with those obtained from the other two techiques. ",
        "conclusions": "We have presented new photopolarimetric and spectroscopic observation of the AM Herculis binary V834~Cen during a high state of accretion. The observed high levels of polarization over the whole orbit are modelled using Stokes imaging. This technique predicts a single emission region, visible for the whole orbit, to be responsible for the observed polarized variations. The location of the region on the white dwarf is consistent with recent work on EUV observations by Mauche (2002). Both analyses place the main emission region approximately more than 40 degrees and up to 70 degrees in azimuth from the line of centers of the two stars.  Our Doppler tomographic analysis shows several components associated with the secondary star, the ballistic stream and parts of the magnetically confined stream. Interestingly, the ballistic stream (at least the earlier parts) do not show a reduced Doppler $v_{y}$ velocity as is commonly seen in other polars. We therefore reason that the earlier part of the ballistic stream is relatively free from magnetic influences of the white dwarf. Magnetic Doppler trajectories are also calculated and shown to be generally coincident with the magnetic signatures found in the Doppler maps. We also calculated the location of the footprints of the magnetic field lines on the surface of the white dwarf and showed that these are coincident with the prediction for the location of the cyclotron emission region from Stokes imaging.  Roche tomography of the narrow component seen in the HeII 4686 \\AA \\ emission line reveals two asymmetries in the emission distribution on the Roche surface of the secondary. Namely, there is generally more emission on the inner surface and also on the leading face of the secondary. We argue that this is still consistent with irradiation of the secondary, as the lack of a magnetic curtain at the beginning of the ballistic stream reduces its obscuring ability. This is not only supported by the Doppler tomogram, but also by Stokes imaging results, which show that most of the accreting material comes along field lines that intersect the later parts of the ballistic stream (see Fig.~\\ref{alltogethernow})."
    },
    "0311/astro-ph0311198_arXiv.txt": {
        "abstract": "{The force-free coronal loop model by \\cite{Tit:Dem-99} is found to be unstable with respect to the ideal kink mode, which suggests this instability as a mechanism for the initiation of flares. The long-wavelength ($m\\!=\\!1$) mode grows for average twists $\\Phi\\ga3.5\\pi$ (at a loop aspect ratio of $\\approx$\\,5). The threshold of instability increases with increasing major loop radius, primarily because the aspect ratio then also increases. Numerically obtained equilibria at subcritical twist are very close to the approximate analytical equilibrium; they do not show indications of sigmoidal shape. The growth of kink perturbations is eventually slowed down by the surrounding potential field, which varies only slowly with radius in the model. With this field a global eruption is not obtained in the ideal MHD limit. Kink perturbations with a rising loop apex lead to the formation of a vertical current sheet below the apex, which does not occur in the cylindrical approximation. ",
        "introduction": "\\label{intro} Magnetic loops are the elementary building block of the solar corona and of low-plasma-beta environments of astrophysical objects in general. Understanding their instabilities is one of the fundamental problems in corona physics. Due to the small plasma beta ($\\beta=P_\\mathrm{kin}/P_\\mathrm{mag}\\sim10^{-3}\\dots10^{-2}$ in the inner solar corona), the magnetic configuration of stable loops must be nearly force free. Any instability must be caused by currents flowing mainly along the loops. A sigmoidal (S- or inverse-S) shape of the projection of loops onto the solar surface is regarded as a signature of such currents, and indeed, a strong sigmoidal shape of soft X-ray loops correlates with eruptive activity \\citep{Canf:al-99}.  The stability of current-carrying force-free (or nearly force-free) fields was extensively studied for toroidal geometry in fusion research \\cite[see, e.g.,][]{Biskamp-93} and for cylindrical geometry with fixed ends in the astrophysical context. The latter is considered as an approximation to coronal loops with large aspect ratio which includes the effect of photospheric line tying \\cite[e.g.,][]{Hoo:Pri-81, Vell:al-90,Miki:al-90,Einaudi-90,Hood-92,Bat:Hey-96,Baty-01,Gerr:al-02}. It was found that the stability is mainly controlled by the total twist, \\begin{equation} \\Phi = \\frac{lB_\\phi(r)}{rB_z(r)}\\,, \\label{eq_Phi} \\end{equation} where $l$ is the length of the current-carrying flux system, $r$ is the (minor) radius, and $B_z$ and $B_\\phi$ are the axial and azimuthal field components, respectively. In addition, the radial profile of the twist, the effect of line tying, and the ratio of radius and twist scale length ($\\lambda=2{\\pi}l/\\Phi$) are important. For example, the uniformly twisted force-free toroidal or periodic cylindrical configuration is kink-unstable for $\\Phi>\\Phi_\\mathrm{c}=2\\pi$, but line tying raises the instability threshold to $\\approx2.5\\pi$ \\citep{Hoo:Pri-81,Ein:VHov-83}. Localizing the current within a certain radius, $a$, and embedding it in a potential field also raises the threshold (as any substantial radial variation of the twist). Configurations of this type were found to be stable for $\\Phi_\\mathrm{max}\\la5\\pi$, where $\\Phi_\\mathrm{max}$ is the peak value of the twist in the configuration \\citep{Miki:al-90,Bat:Hey-96}. The instability threshold of embedded loops rises strongly if their radius becomes very small, i.e. if $\\delta=2{\\pi}a/\\lambda\\la1$ \\citep{Baty-01}. For cylindrical line-tied force-free configurations it was also shown that the longest-wavelength kink mode (azimuthal wave number $m\\!=\\!1$) becomes unstable before any other mode in the ideal MHD limit \\citep{VdLin:Hoo-99}.  The stability of arched magnetic loops has not yet been investigated. However, numerical studies of the injection of twist by slow vortex motions at the footpoints of an initially current-free loop-shaped flux bundle indicate that raising the twist of a loop beyond a critical value leads to destabilization \\citep{Amar:al-96,Toe:Kli-03}. The critical twist lies in the range $2.5\\pi<\\Phi_\\mathrm{c}<2.75\\pi$ for the specific initial configurations studied. The driving by photospheric vortices generally influences a large surrounding volume in addition to the twisted loop. It also causes an increase of the loop length and width with increasing twist. Clearly, a stability consideration of loops free from such additional influences is needed.  In this paper we study the stability of the loop model by \\cite{Tit:Dem-99}, cited as T\\&D in the following (see their Fig.\\,2 for a schematic). This approximate, cylindrically symmetric, force-free equilibrium consists of a toroidal ring current of major radius $R$ and minor radius $a$, whose outward-directed Lorentz self-force is balanced with the help of a field by two fictitious magnetic charges of opposite sign which are placed at the symmetry axis of the torus at distances ${\\pm}L$ to the torus plane. That axis lies below the photospheric plane $\\{z\\!=\\!0\\}$ at a depth $d$. The resulting field outside the torus is current-free and contains a concentric magnetic X line between the torus and its centre. A toroidal field component created by a fictitious line current running along the symmetry axis is added. Its strength controls the twist of the field in the torus, and it turns the X line into a hyperbolic flux tube (HFT; for a definition, see \\citeauthor{Tito:al-02}, \\citeyear{Tito:al-02}). The existence of the HFT is generic to such force-free configurations with a net current. The accuracy of the obtained equilibrium improves with decreasing parameters $a/R$ and $a/L$. Its section in the volume $\\{z\\!>\\!0\\}$ is a model of a coronal magnetic loop.  T\\&D have investigated the stability of the torus with respect to global expansion (growing perturbations ${\\delta}R>0$) and found instability for sufficiently large radii, $R\\ga\\sqrt{2}\\,L$. \\cite{Rous:al-03} confirmed this recently by numerical simulations, which also suggested that the instability threshold lies near $5L$ and that the slow decrease of the surrounding toroidal field with distance from the generating line current leads to saturation of the instability and prevents the loop from erupting. In the following, we investigate the stability of the configuration with respect to ideal kink modes in dependence of the twist in its coronal part. ",
        "conclusions": "\\label{concl} The magnetic loop equilibrium by \\cite{Tit:Dem-99} is kink-unstable for twists $\\Phi>\\Phi_\\mathrm{c}$, with $\\Phi_\\mathrm{c}\\approx3.5\\pi$ at a loop aspect ratio $R/a\\approx5$. The instability threshold rises with rising aspect ratio. No indication of sigmoidal shape at slightly subcritical twists was found. The unstable kink mode with upward displacement of the loop apex forms a vertical current sheet at the HFT underneath the loop, which has no counterpart in kink-unstable cylindrical loop models but corresponds to the central element of the ``standard flare model'' \\cite[e.g.,][]{Shibata-99}. The surrounding potential field eventually terminates the exponential growth of kink modes in the considered configuration, apparently inhibiting a global eruption within the framework of ideal MHD. However, the configuration may nevertheless suffer a global eruption if (a) a surrounding field which decreases more rapidly with height is employed, (b) magnetic reconnection is permitted to occur in the formed current sheets, or (c) magnetic reconnection with neighbouring loops is triggered in a multiple-loop configuration."
    },
    "0311/astro-ph0311145_arXiv.txt": {
        "abstract": "We provide a ``Baedecker'' or travel guide to the directions on the sky where the dark matter annihilation signal may be expected. We calculate the flux of high energy $\\gamma$-rays from annihilation of neutralino dark matter in the centre of the Milky Way and the three nearest dwarf spheroidals (Sagittarius, Draco and Canis Major), using realistic models of the dark matter distribution. Other investigators have used cusped dark halo profiles (such as the Navarro-Frenk-White) to claim a significant signal. This ignores the substantial astrophysical evidence that the Milky Way is not dark-matter dominated in the inner regions. We show that the annihilation signal from the Galactic Centre falls by two orders of magnitude on substituting a cored dark matter density profile for a cusped one. The present and future generation of high energy $\\gamma$-ray detectors, whether atmospheric Cerenkov telescopes or space missions like GLAST, lack the sensitivity to detect any of the monochromatic $\\gamma$-ray annihilation lines. The continuum $\\gamma$-ray signal above 1 GeV and above 50 GeV may however be detectable either from the dwarf spheroidals or from the Milky Way itself. If the density profiles of the dwarf spheroidals are cusped, then the best prospects are for detecting Sagittarius and Canis Major. However, if the dwarf spheroidals have milder, cored profiles, then the annihilation signal is not detectable. For GLAST, an attractive strategy is to exploit the wide field of view and observe the Milky Way at medium latitudes, as suggested by Stoehr {\\em et al.} This is reasonably robust against changes in the density profile. ",
        "introduction": "The foremost candidate for the cold dark matter (CDM) composing galactic haloes is the lightest neutral supersymmetric particle, namely the neutralino \\cite{Jungman}. If so, then neutralino pair annihilation may lead to observable consequences, in particular the emission of high energy $\\gamma$-radiation \\cite{Bergstrom98}. The possibility that such $\\gamma$-rays may be identified by forthcoming atmospheric Cerenkov telescopes (ACT) such as VERITAS \\cite{veritas} or by satellite-borne detectors like GLAST \\cite{glast} has excited considerable recent interest \\cite{baltz,tyler,ts,stoehr}.  It is clearly of importance to identify the best places to search for such an annihilation signal. Inspired by the highly cusped models based on numerical simulations of dark halo formation \\cite{navarro,moore}, a number of investigators have suggested that the centre of the Milky Way may be the optimum target. For example, Bergstrom {\\it et al.} \\cite{Bergstrom98} have shown that if the dark matter density is cusped as $1/r$ at small radii, then the $\\gamma$-ray flux would be detectable for typical neutralino properties in the minimal supersymmetric extension of the Standard Model. Inspired also by the persistence of substructure in numerical simulations, a number of authors \\cite{Bergstrom99,ts,crm} have argued that a substantial enhancement in the $\\gamma$-ray signal can be expected from such `clumps'. In these calculations, the inner regions of the substructure are also usually assumed to be cusped. However, even within the framework of the cusped models favored by cosmological simulations, these conclusions have been contested as being overly optimistic \\cite{stoehr}.  More awkwardly, there is a substantial body of astrophysical evidence that the halo of the Milky Way is not cusped at all \\cite{evansidm}. First, the microlensing optical depth towards the Galactic Center is very high. Particle dark matter does not cause microlensing, whereas faint stars and brown dwarfs do. The total amount of all matter within the Solar circle is constrained by the rotation curve, so this tells us that lines of sight towards the Galactic Center are not dominated by particle dark matter. More specifically, haloes as strongly cusped as $1/r$, normalised to the local dark matter density as inferred from the stellar kinematics in the solar neighbourhood, are ruled out by the high microlensing optical depth \\cite{be}. Second, the pattern speed of the Galactic bar is known to be fast from hydrodynamical modelling of the motions of neutral and ionised gas. If dark matter dominates the central regions of the Milky Way, then dynamical friction will strongly couple the dark matter to the Galactic bar and cause it to decelerate on a few bar rotation timescales \\cite{debs}. It is now largely accepted by astronomers that bright galaxies like the Milky Way do not have cusped dark haloes today, with some investigators suggesting that feedback from star formation may provide a resolution with cold dark matter theories \\cite{bgs}.  In fact, there is {\\em no} observational evidence whatsoever that any nearby galaxy has a cusped dark halo profile. The rotation curves of low surface brightness and dwarf spiral galaxies have been the subject of a long controversy \\cite{mcgaugh,bosch}. The effects of beam smearing mean that the H~I rotation curves of many dwarf spirals are broadly compatible with both cores and cusps. However, the H~II rotation curves for at least some dwarf spirals are not compatible with cusps \\cite{blais}. Most dwarf spheroidals (dSphs) do not contain gas and so the structure of the dark halos must be inferred from stellar motions. Very recently, the survival of kinematically cold substructure in the Ursa Minor dSph has been used to argue against a cusped halo \\cite{kwge}. Hence, even at the least massive and most dark matter dominated end of the galaxy mass spectrum, the predictions of cold dark matter theories concerning halo structure seem to disagree with the observations. Nonetheless, for the three nearest dSphs -- Draco, Sagittarius and Canis Major -- there is no direct evidence either for or against central cusps in their dark matter distribution.  Given this weight of evidence, it seems very prudent to use both cored and cusped halo models to estimate the range of the expected $\\gamma$-ray annihilation signal. In this paper, we examine four possible locations -- the Galaxy Center and the centers of the three nearest dark-matter dominated dwarf spheroidals (Draco, Sagittarius and Canis Major). In Section~II we use the most recent data on the velocity dispersion of the dSphs to constrain a variety of dark halo models, and in Section~III evaluate the $\\gamma$-ray flux from neutralino self-annihilations. Sections~IV and V summarise the expected contribution from the background and the criterion for detection respectively. The results for second generation ACTs and for GLAST are given in Section~VI. ",
        "conclusions": "If the dark matter present in the Universe is composed at least in part by the lightest supersymmetric particle, then this could manifest itself via $\\gamma$-ray emission from pair annihilations.  It is clearly important to estimate the likely magnitude of the neutralino annihilation signal. It is also important to identify the likely locations and spectral r\\'egimes in which the signal should be sought. This paper has provided new estimates of the signal towards the Galactic Center and the nearby dwarf spheroidals using a variety of models.  There have been a number of recent calculations predicting that the neutralino annihilation flux from the inner Galaxy will be detectable with forthcoming satellites like GLAST and with second generation atmospheric Cerenkov telescopes (ACTs) \\cite{Bergstrom98,stoehr}. These calculations assume that the cusped Navarro-Frenk-White (NFW) models for the Milky Way halo hold good. This assumption is in contradiction with a substantial body of astrophysical evidence about the inner Galaxy \\cite{evansidm,be,debs}. In any case, even if the Milky Way halo was originally of NFW form, the formation of the disk and bulge will have reprocessed the primordial dark matter distribution \\cite{bgs}. In contradiction with earlier results, we do not find the prospects of detecting the annihilation flux from the Galactic Center to be particularly promising. In particular, the $\\gamma$-ray line coming from the the $\\gamma\\gamma$ and $Z\\gamma$ final states is not detectable either with second generation ACTs or with the GLAST satellite. {\\em We caution that many of the recent estimates of high flux are sensitively dependent on the assumptions made regarding the innermost structure of the dark halo. Even the best numerical simulations have difficulty in resolving structures on scales less than 1 kpc, and so the inner profile is always found by extrapolation.}  The high mass-to-light ratios of the Local Group dwarf spheroidals (dSphs) makes them attractive targets.  Cusped profiles like NFW are not presently ruled out for dSphs like Sagittarius or Draco. It may be that the visible dwarf galaxy lies entirely within the central parts of a cusped dark matter halo. If so, then the optimum targets are the Sagittarius and Canis Major dSphs. The detection of monochromatic lines is still extremely difficult, but the GLAST satellite may detect excess continuum $\\gamma$ ray emission. This is of course a less distinctive feature than a sharp line. In particular, if the Sagittarius or Canis Major dSphs have a strongly cusped dark halo profile ($\\rho \\sim r^{-1.5}$ or $\\rho \\sim r^{-1}$), then some regions of supersymmetric parameter space can be ruled out. Again, however, this conclusion only holds good if the dark halo profile is cusped. Using a cored isothermal-like model for the dark halo, even the Sagittarius and Canis Major dSphs may be invisible to GLAST and second generation ACTs.  Unlike Bergstrom {\\it et al.} \\cite{Bergstrom98}, we do not find the Galactic Center to be a promising location. Partly, this is because we believe that the Milky Way does {\\em not} have a strongly cusped profile based on the available astrophysical evidence \\cite{evansidm,be,debs}. Partly, this is because Bergstrom et al. chose a generous overall normalisation anyhow --- they used the NFW model corresponding to the maximum flux which satisfies two weak constraints on the mass and the rotation curve. Accordingly, the local dark matter density is as high as $\\sim 0.6$ GeV cm$^{-3}$ in their model. When the circular velocity curve of such a halo is combined with that for the disk and bar, then it necessarily violates the constraint on the Galactic rotation curve in the inner parts. One important caveat of our results --- however --- is that the possible effects of a central black hole are not included in our calculations. Here, we merely note that the observability of any expected signal depends on the manner in which the black hole grows \\cite{gs,merritt}.  Stoehr {\\it et al.} \\cite{stoehr} have also recently emphasised that the $\\gamma$-ray emission from the Galactic Center may have been overestimated by the use of too strongly cusped profiles. They suggest that the galaxy at moderate latitudes ($|b| > 10^\\circ)$ may also be a good target for detecting the continuum emission (they do not study the line emission). This is not really an option for ACTs with their small field of view. However, {\\it it is an attractive possibility for GLAST, as the continuum emission is detectable irrespective of uncertainties in halo structure.} For ACTs, the best targets remain the Sagittarius and Canis Major dSphs.  Very recently, the Large Magellanic Cloud (LMC) has been suggested as another likely target \\cite{olinto}. Judging from \\cite{vdm}, the average mass to light ratio of the LMC within 8.9 kpc is only $\\sim 3$ (as opposed to $\\sim 100$ for the compact dSphs). This is an upper limit to the central mass to light ratio. In other words, much as in the Milky Way, dark matter dominates the outer parts of the LMC and is responsible for the asymptotic flatness of the ratio curve.  However, the central parts of the LMC are dominated by the luminous bar and disk. The assumption that the dark halo dominates the gravitational potential everywhere is therefore not valid. Hence, the procedure used in \\cite{olinto} of fitting the rotation curve to a NFW dark halo is flawed. The gravitational potential of the gas and stellar disk and bar simply cannot be ignored in the central regions."
    },
    "0311/astro-ph0311373_arXiv.txt": {
        "abstract": "Under the assumption that the concordance $\\Lambda$ cold dark matter (CDM) model is the correct model, we test the cosmic microwave background (CMB) anisotropy data for systematic effects by examining the band pass temperature residuals with respect to this model. Residuals are analysed as a function of angular scale $\\ell$, galactic latitude, frequency, calibration source, instrument type and several other variables that may be associated with potential systematic effects. Our main result is that we find no significant systematic errors associated with these variables. However, we do find marginal evidence for a trend associated with galactic latitude indicative of galactic contamination. ",
        "introduction": "The cosmic microwave background (CMB) power spectrum is a particularly potent probe of cosmology.  As long as the systematic errors associated with these observations are small, the detected signal has direct cosmological importance. The ever-tightening network of constraints from CMB and non-CMB observations favours a concordant $\\Lambda$ cold dark matter (CDM) model that is commonly accepted as the standard cosmological model (Table~\\ref{concord}).  Since the anisotropy power spectrum is playing an increasingly large role in establishing and refining this model, it is crucial to check the CMB data for possible systematic errors in as many ways as possible.  Systematic errors and selection effects are notoriously difficult to identify and quantify.  Individual experimental groups have developed various ways to check their CMB observations for systematic effects \\citep[e.g.][]{kogut96,miller02}, including the use of multiple calibration sources, multiple frequency channels and extensive beam calibrating observations. Internal consistency is the primary concern of these checks.  Testing for consistency with other CMB observations is another important way to identify possible systematic errors.  When the areas of the sky observed overlap, this can be done by comparing CMB temperature maps \\citep[e.g.][]{ganga94a,lineweaveretal95,xu01}. When similar angular scales are being observed one can compare power spectra \\citep[e.g.][Figure 11]{sievers02}.  A prerequisite for the extraction of useful estimates for cosmological parameters from the combined CMB data set is the mutual consistency of the observational data points \\citep{wangetal02a}; the best-fit must also be a good fit. \\citet{wangetal02a} and \\citet{sievers02} have recently explored the consistency of various CMB observations with respect to power spectrum models and concluded that the CMB fluctuation data is consistent with several minor exceptions.  Although individual observational groups vigorously test their data sets for systematic errors, the entire CMB observational data set has not yet been collectively tested.  Here we check for consistency of the concordance model (Table~\\ref{concord}) with respect to possible sources of systematic error.  Under the assumption that the concordance model is the correct model (i.e. more correct than the best-fit to the CMB data alone), we explore residuals of the observational data with respect to this model to see if any patterns or outliers emerge.  We attempt to identify systematic errors in the data that may have been ignored or only partially corrected for.  With only a few independent band power measurements the usefulness of such a strategy is compromised by low number statistics. However, we now have hundreds of band power measurements on scales of $2<\\ell<2000$ from over two dozen autonomous and semi-autonomous groups. There are enough CMB fluctuation detections from independent observations that subtle systematic effects could appear above the noise in regression plots of the data residuals.  This is particularly the case when one has a better idea of the underlying model than provided by the CMB data alone.  The history of the estimates of the position of the CMB dipole illustrates the idea.  Once a relatively precise direction of the dipole was established, the positional scatter elongated in the direction of the galactic centre could be distinguished unambiguously from statistical scatter and more reliable corrections for galactic contamination could be made \\citep[Figure 2]{lineweaver97}.  We aim to ascertain whether the use of the concordance model as a prior can help to separate statistical and systematic errors in the CMB anisotropy data.  In \\S2 we discuss constraints on cosmological parameters, the current concordance model and how simultaneously analysing combinations of independent observational data sets can tighten cosmological constraints. Our analytical methodology is detailed in \\S3.  In \\S4 possible sources of systematic uncertainty are discussed. In \\S5 and \\S6 our results are discussed and summarised. ",
        "conclusions": "Over the past 10 years, successive independent and semi-independent data sets have extended the angular scale, calibration precision and freedom from galactic contamination of the CMB power spectrum. Each CMB measurement contains useful cosmological information and no data set is immune to contamination. It is therefore important to compare data sets and check for systematics. We have collectively tested the full CMB data set for inconsistencies with the concordance model and our results indicate that the model is consistent with the data although a need to slightly dampen power in the model at high $\\ell$ is indicated (Figure~\\ref{lresid}).  We have explored residuals of the observational data with respect to the concordance model to see if any patterns emerge that may indicate a source of systematic error. We have found little significant evidence for inter-experiment inconsistencies other than a trend associated with galactic latitude that may be an indication of low-level galactic contamination of CMB observations made closer to the galactic plane (Figure~\\ref{gallatresid}). A more detailed comparison of CMB fields of observations with galactic dust and synchrotron maps will be necessary to clarify the source of this trend."
    },
    "0311/astro-ph0311003_arXiv.txt": {
        "abstract": "In this paper we investigate the possibility that there is a first phase of partial ionisation due to X-rays produced by black hole accretion in small-mass galaxies at redshifts $7 < z < 20$. This is followed by complete reionisation by stellar sources at $z \\simeq 7$. This scenario is motivated by the large optical depth to Thompson scattering, \\taue$\\simeq 0.17 \\pm 0.04$, recently measured by WMAP. But it is also consistent with the observed Gunn-Peterson trough in the spectra of quasars at $z \\sim 5-6$. We use a semianalytic code to explore models with different black hole accretion histories and cosmological parameters.  We find that ``preionisation'' by X-rays can increase the intergalactic medium (IGM) optical depth from \\taue$\\approx 0.06$ given by stellar sources only, to $0.1 \\simlt \\tau_{\\rm e} \\simlt 0.2$, if a fraction of baryons $10^{-5} \\simlt \\omega_{\\rm ac} \\simlt 10^{-4}$ is accreted onto seed black holes produced in the collapse of low metallicity, high mass stars before $z \\simeq 15$.  To be effective, preionisation requires a non-negligible star formation in the first small-mass galaxies in which seed black holes are formed.  By $z \\sim 20-25$ the IGM is re-heated to $10,000$ K and the ionisation fraction is about 20\\%.  The increase of the IGM Jeans mass is effective in reducing star formation in the smaller-mass haloes. Large values of \\taue are obtained in models with top-heavy stellar initial mass function only if pair-instability supernovae are not important. Seed black holes are assumed to accrete at near the Eddington limit with a duty cycle that decreases slowly with increasing time.  Alternatively, a moderate fraction of the black holes must be ejected from the host galaxy or exist without merging into the supermassive black holes in galactic centres.  The model predicts that dwarf spheroidal galaxies, if they are preserved fossils of the first galaxies, may host a mass in black holes that is 5-40\\% of their stellar mass. The redshifted X-ray background produced by this early epoch of black hole accretion constitutes about $5-10$\\% of the X-ray background in the 2-50 keV bands and roughtly half of the currently estimated black hole mass density was formed at early times.  Moreover, in most models, the photons from the redshifted background are sufficient to fully reionise \\GII at redshift $z \\sim 3$ without any additional contribution from quasars at lower redshifts and the temperature of the mean density intergalactic medium remains close to $10^4$ K down to redshift $z \\sim 1$. ",
        "introduction": "The mean transmitted flux along the line of sight of the furthest quasars found by the Sloan Digital Sky Survey \\citep[\\eg,][]{Becker:01} shows that the ionisation fraction of the intergalactic medium (IGM) is rapidly decreasing approaching $z \\sim 6.3$. Assuming, at that redshift, an approximately constant stellar emissivity of ionising radiation, the IGM would have been reionised at \\zrei$\\sim 7$ by stellar sources \\citep[][ and others]{Gnedin:00,Djorgovski:01,Songaila:02,Fan:03}.  But a much earlier ionisation epoch is implied by the WMAP observations \\citep{Bennet:03, Kogut:03}. If the IGM had a sudden transition from neutral to completely ionised, producing the optical depth to Thompson scattering \\taue$=0.17 \\pm 0.04$ measured by WMAP, \\zrei$ \\simeq 17$ for the best fit $\\Lambda$CDM cosmological model \\citep{Spergel:03}.  In a companion paper \\citep[paper~I,][]{RicottiOI:03}, we have tried to explain these apparently contradictory results by considering the contribution to reionisation due to zero-metallicity (\\pop3) stars. Since such stars are thought to be massive ($M_* > 5-10$ M$_\\odot$) \\citep[\\eg,][]{OmukaiN:98, Abel:00, Bromm:02, Abel:02, NakamuraU:02} and hot \\citep{Tumlinson:00}, their efficiency for UV emission is larger than for \\popII and \\popI stars that have a Salpeter initial mass function (IMF).  A time dependent efficiency of UV emission, $\\epsilon_{\\rm UV}$, produced by the transition from a top-heavy IMF at high redshift to a Salpeter IMF at low redshift, could possibly explain the large \\taue measured by WMAP and would be consistent with the Sloan data at $z \\simlt 6.3$ \\citep{ChiuFO:03}.  An upper limit of the UV emissivity, $\\epsilon_{\\rm UV}^{\\rm max} \\simlt 2-3\\times 10^{-3}$ can be derived by assuming the maximum efficiency of energy production by thermonuclear reactions in massive stars. For a Salpeter IMF, $\\epsilon_{\\rm UV} \\sim 2-4 \\times 10^{-4}$, depending on the stellar metallicity \\citep{Tumlinson:00}.  Therefore, assuming a constant (maximal) escape fraction of ionising photons from the galaxies, \\fesc$\\simeq 1$, a redshift dependent IMF would increase the UV emission efficiency from low to high redshift by about 10-20 times.  Several authors have tried to reproduce the high optical depth measured by WMAP with the increased ionising photon emission by zero-metallicity stars. \\cite{Cen:03b} has used a semianalytic calculation to indicate that the IGM may have experienced two distinct epochs of reionisation. \\cite{HaimanH:03} have investigated different reionisation histories showing that optical depths consistent with WMAP can be achieved assuming very massive \\pop3 stars. \\cite{WyitheL:03} show that, if feedback regulates star formation in early low-mass galaxies as observed in present-day dwarfs, \\pop3 stars forming with a heavy IMF are required to match WMAP data. \\cite{CiardiFW:03}, using numerical simulation, show that zero metallicity stars with a moderately heavy IMF (characteristic stellar mass of 5 M$_\\odot$) and \\fesc$=20$\\% are able to generate an optical depth of 0.17, consistent with WMAP data. \\cite{Sokasian:03}, also using numerical simulations, showed that zero metallicity stars with a top-heavy IMF and \\fesc$\\sim 100$\\% is needed to produce an optical depth consistent with WMAP data.  In paper~I we have found (consistent with the work of the other investigators mentioned above) that, assuming the most extreme properties for zero metallicity stars (top heavy IMF and \\fesc$\\sim 100$\\%), the maximum optical depth to Thomson scattering produced by \\pop3 stars is \\taue$ \\sim 0.13$, marginally consistent with WMAP data. In order to produce this large value of \\taue, massive zero-metallicity (\\pop3) stars must be the dominant population until redshift $z \\simeq 10$. But we have noted that, if the ratio of metal atoms to ionising photons produced by \\pop3 stars is normal, the metal enrichment of the ISM prevents metal-poor stars from being produced for long enough to ionise more than a small fraction of the IGM, \\ie that the \\pop3 phase is very rapidly self-limiting. In addition, if pair-instability SNe are important, the mechanical energy input by SN explosions produces strong outflows in galaxies with masses $M_{\\rm dm} \\simlt 10^{9}$ M$_\\odot$, reducing their star formation and delaying reionisation. These two effects will lower the redshift of reionisation and \\taue to values inconsistent with WMAP but still consistent with the Sloan quasar data. Thus we have found it essentially impossible to reproduce the WMAP and Sloan results using stellar sources of UV ionisation.  In summary, the reason for the partial disagreement with the conclusions of previous works on the ability of \\pop3 stars to produce the large optical depth observed by WMAP, is motivated by our findings concerning the self-termination of \\pop3 stars by the enhanced metal pollution and the negative feedback on star formation from the enhanced energy input by SN explosions, that inevitably follow from the assumption of a top-heavy IMF. In previous works these concerns have not been addressed.  However in paper~I we also note that, if most \\pop3 stars collapse into black holes (BHs), the metal pollution problem is alleviated and the mechanical feedback from SN/hypernova explosions will be reduced as well.  An interesting consequence of this scenario is a copious production of BHs from the first stars \\citep{Schneider:02}.  For this reason it is possible that the secondary radiation from accretion onto seed BHs might be a more important source of ionising radiation than the primary radiation from the same stars. In particular because, as we will show, accretion on seed BHs is not sensitive to the early termination of the \\pop3 epoch by metal pollution.  Motivated by these results, in this paper we study the partial ionisation of the IGM by an early X-ray background produced by accreting BHs, followed by more complete stellar reionisation by \\popII stars at \\zrei$ \\approx 7$. An hybrid model where UV from \\pop3 stars and X-ray emission from accreting BHs are both important in producing a large optical depth is also viable if self-termination of \\pop3 stars happens late (\\eg, $z \\sim 10$) and if we assume extreme parameters for \\fesc and for the global \\pop3 stars UV emissivity. But, given the difficulties in modelling properly the transition from \\pop3 to \\popII stars, the predicted importance of UV from \\pop3 stars is very uncertain (almost a free parameter at this point). As an example, in the present paper we show a model where the transition from \\pop3 to \\popII stars is at redshift $z \\sim 15$. In this case we show that UV radiation from \\pop3 stars has negligible effect on the Thompson optical depth.  For a fixed emissivity of ionising radiation, X-rays are less efficient in reionising the IGM, because most of their energy goes into heat instead of ionisation. The inclusion of secondary electrons can boost the ionisation efficiency by about a factor of ten if the electron fraction of the IGM is \\xe$<10$ \\%, but roughly $1/3$ of the primary electron energy is always converted into heating of the IGM \\citep{ShullVan:85}. It will be shown that the partial ionisation of the IGM must begin at redshifts $z \\simgt 15$ in order to have an important effect on \\taue.  This can only happen if the X-ray sources (\\ie, accreting BHs) form in small-mass haloes with masses $M_{\\rm dm} \\simlt 10^8$ M$_\\odot$.  Previous work on the formation of the first galaxies \\citep{RicottiGSa:02, RicottiGSb:02} has used cosmological simulations with radiative transfer to show that star formation in the first small-mass galaxies is reduced by feedback effects but is not suppressed. It is therefore plausible that a substantial production of seed BHs from the first stars took place in the first small-mass galaxies.  Published studies have investigated the effects of X-rays on early galaxy formation using a semianalytic approach \\citep{Oh:00, Venkatesan:01} and cosmological simulations \\citep{Machacek:03}. But since those works were completed before WMAP results, they did not focus on scenarios that could produce \\taue$\\simeq 0.17$.  The study presented in this paper (paper~IIa) is carried out using a semianalytic code based on principles similar to those adopted by \\cite{Chiu:00}. In \\cite*{RicottiOG:03} (paper~IIb) we present the results of cosmological hydrodynamic simulations that include radiative transfer. We discuss the observational signatures of X-ray preionisation compared to stellar reionisation, including calculations of the expected amplitude of secondary anisotropies of the CMB and the redshifted 21cm signal in absorption and emission against the CMB. The semianalytic models presented in this paper allow us to explore a larger parameter space and help us to interpret the numerical result of the full cosmological simulations. This new code, presented in Appendix~\\ref{ap:A}, calculates the evolution of the filling factor of \\HII regions and the ionisation and thermal history of the IGM outside the \\HII regions, produced by an X-ray background.  The semianalytic code has been tested and calibrated using the results of the cosmological simulations (\\cf, paper~I). Note that some cosmological simulations presented here include the effects of SN explosions, using the recipe discussed in \\cite{Gnedin:98a}, in contrast to earlier papers in this series \\citep{RicottiGSa:02, RicottiGSb:02}.  This paper is organised as follows. In \\S~\\ref{sec:rat} we present qualitative arguments showing why and for which models X-ray preionisation can be more effective in increasing \\taue than reionisation by stellar sources. In \\S~\\ref{sec:xsour} we estimate the accretion rate onto seed BHs and we compute the integrated X-ray energy required to explain the large \\taue measured by WMAP. In \\S~\\ref{sec:res_sa} we summarise the characteristics of the semianalytic code for reionisation and we show its results. In \\S~\\ref{sec:conc} we conclude this work with a summary and a discussion of some of the observable effects of early black hole X-ray heating. A full description of the semianalytic code is given in Appendix~\\ref{ap:A}. ",
        "conclusions": "\\label{sec:conc}  In a previous paper (paper~I), we showed that even with a top heavy mass function \\pop3 stars could not reionise the universe to a significant degree if they end their lives as pair instability supernovae. The reason is that such events contaminate the high density regions to such an extent that further formation of extremely low metallicity stars becomes impossible before significant ionisation has occurred.  But, if these same stars implode to massive black holes, then accretion onto them at near Eddington rates will produce a pervasive X-ray background which can quite efficiently partially preionise the bulk of the cosmic baryons before $z=15$ and give an optical depth to Thompson scattering (\\cf, \\tab~\\ref{tab:2}) consistent with WMAP results.  The primary reason for the increased efficiency is the long mean free path of X-rays which both increases the fraction of photons that escape the parent halo and ionises preferentially in the low density regions where recombination rates are low. In addition X-rays only partially ionise the IGM to 10-20\\%, therefore the recombination rate is reduced with respect to the reionisation scenario produced by UV sources. Thus we find that the average number of times a baryon is photoionised in this picture is about unity for \\taue$\\simeq 0.15$ while it is about ten in the case when stellar UV is used (paper~I). Moreover, since the X-rays preionise the low density regions of the IGM at early times, the required number of ionising photons from \\popII stars to have reionisation at redshift $z \\simeq 6$ is much reduced.  But we might ask, could the X-rays arise as easily from X-ray binaries or could SN remnants \\citep[\\cf,][]{Oh:00}?  These other sources are certainly helpful but they will generate somewhat less energy in X-rays than accretion can do. The ratio of X-ray energy to massive star formation expressed as an efficiency is $\\epsilon_{\\rm SN}=E_{51} 10^{51} {\\rm erg}/143 M_\\odot c^2 \\sim 4 \\times 10^{-6}$ for these sources (assuming \\popII with Salpeter IMF and neglecting Compton cooling of the X-ray emitting hot gas) but we see from $\\epsilon_{\\rm ac}=0.15 \\beta (M_{\\rm SMBH}/M_{\\rm bulge})$ that the comparable number from accretion is $2.25 \\times 10^{-5}$.  What are the observable consequences of the scenario we are proposing (in addition to generating a large \\taue)? \\begin{enumerate} \\item The power spectrum of the cross correlation of temperature and polarisation of the CMB is sensitive to the visibility function $g(z)=e^{-\\tau_{\\rm e}}d \\tau_{\\rm e}/d \\eta$ (\\cf, \\fig~\\ref{fig:mod3}-\\ref{fig:mod4}). The Planck satellite should be able to distinguish between the visibility function produced by an early X-ray partial ionisation (peaking at large z) or the one expected for reionisation by stellar sources (peaking at lower z). \\item Since the assumed spectrum (\\fig~\\ref{fig:sed}) has a positive slope $(d\\ln L_\\nu /d\\ln \\nu) >3$ for much of the relevant range of frequencies there are ever fresh ionising photons available and the probability per unit time that a neutral hydrogen atom is ionised does not decrease with time. This produce softer photons that, even if the ionisation efficiency of X-rays decreases when the ionisation fraction exceeds 10-20\\%, enhance the ionisation fraction in the voids to values larger than 20 \\% because of the contribution from the redshifted photons of the X-ray background.  In most models, the photons from the redshifted X-ray background can fully reionise \\GII at redshift $z \\sim 3$ without any additional contribution from quasars at lower redshifts. Due to the heating rate by \\GII ionisation the temperature of the mean density intergalactic medium remains close to $10^4$ K down to redshift $z \\sim 1$ and the IGM equation of state has index $(\\gamma-1) \\sim 0.3$. \\item The maximum accretion rate onto seed black holes as a function of redshift is limited by the observed magnitude of the $\\gamma$ ray background at $z=0$. This limits the ability of late preionisation models to produce large values of \\taue and contribute substantially to the mass growth via accretion of seed black holes. In models that produce optical depth \\taue$\\sim 0.17$ consistent with WMAP, the redshifted X-ray background produced by the early epoch of black hole accretion constitutes about $5-10$\\% of the hard X-ray background in the 2-50 keV bands. \\item The model predicts that dwarf spheroidal galaxies, if they are preserved fossils of the first galaxies, would host a mass in black holes that is 5-40\\% of their stellar mass (\\cf, \\tab~\\ref{tab:2}). The larger values are produced in the early preionisation scenarios and it might be possible to rule them out on the base of observations from dynamical considerations. \\item The IGM is reheated to $T =10^4$ K before the ionisation fraction exceeds 10\\% producing typical values of the Compton distortion parameter $y \\sim 4 \\times 10^{-7}$.  The observational upper limit on this parameter is $y < 1.5 \\times 10^{-5}$, too large to constrain our models. Star formation in the smaller mass haloes is reduced by the increase in the IGM Jeans mass following this early reheating. \\item The redshifted 21cm signal in emission and absorption against the CMB can be used to discriminate between preionisation scenarios and reionisation from UV sources. This will be discussed in paper~IIb, where we will also show that preionisation by X-rays produces CMB secondary anisotropies on small angular scales that are easy to recognise when compared to models of inhomogeneous early reionisation by stellar sources with the same \\taue. \\end{enumerate}  What are the requirements of the X-ray preionisation models in terms of seed black hole production from the first stars and accretion?  Since the time scale for accretion, $t_{\\rm Ed}=10^8$ yr, is comparable to the Hubble time at $z=30$ the maximum X-ray emissivity at high-redshift is proportional to the total mass of the seed BHs. For this reason, in order to get a larger \\taue in early preionisation scenarios, a top-heavy IMF is favoured over a Salpeter IMF. But this is true only if super-massive stars mostly end their lives collapsing directly into BHs without exploding as SNe or as pair-instability SNe.  We have shown models where the accretion onto seed black holes is near the Eddington rate at early redshift. This might perhaps be difficult to achieve.  But it is worth noticing that it is not strictly necessary that the accretion is at the Eddington rate initially. Although, if the accretion rate is less efficient than that, a larger stellar mass fraction producing seed black holes, $f_{\\rm BH}$, is needed to accommodate for this (see \\eq~(\\ref{eq:ombh})).  In this paper we have explored a range of models with different accretion histories onto seed black holes.  But we note that the models are not all equally plausible scenarios.  The very early preionisation scenario (M5) is not the most favourable for several reasons: (i) as noted by \\cite{MadauR:03}, accretion in small-mass haloes can be difficult because of the substantial gas photoevaporation in these galaxies (ii) if seed BHs do not accrete at near the Eddington rate the mass fraction, $f_{\\rm BH}$, of seed BHs is extreme also assuming top-heavy \\pop3 stars (iii) the number of recombinations per ionising photon in this model is larger than in the intermediate preionisation model and therefore \\taue is reduced.  The late preionisation model (M3) is also not favourable because it produces a low \\taue if we constrain the maximum black hole accretion rate using \\eq~(\\ref{eq:maxac}) to be consistent with the observed $\\gamma$-ray background at $z=0$.  We therefore conclude that preionisation starting at $z \\sim 15-20$ (\\eg, model M4) is the most plausible scenario to reproduce \\taue measured by WMAP. In this case $\\sim 10$\\% of the X-ray background in the 2-50 keV band is produced by the early generation of mini-quasars and roughtly half of the currently estimated black hole mass density was formed at early times (even if only a fraction of them may have merged into the SMBHs in the galactic centres).  In paper~IIb we study this scenario using hydrodynamic cosmological simulations with radiative transfer and we discuss observations that could probe this model.  We also note in conclusion that this picture while consistent with the central values of WMAP results could not easily produce \\taue at the upper end of the presently allowed values. If \\taue$ \\simgt 0.2$, X-ray preionisation is not likely the dominant responsible mechanism and other scenarios need to be considered (\\eg, decaying particles, non Gaussian perturbations, spectral index $n_{\\rm s}>1$, primordial BHs, unconventional recombination).  \\subsection*{ACKNOWLEDGEMENTS}  MR is supported by a PPARC theory grant. Research conducted in cooperation with Silicon Graphics/Cray Research utilising the Origin 3800 supercomputer (COSMOS) at DAMTP, Cambridge.  COSMOS is a UK-CCC facility which is supported by HEFCE and PPARC. MR thanks Martin Haehnelt and the European Community Research and Training Network ``The Physics of the Intergalactic Medium'' for support.  The authors would like to thank Andrea Ferrara, Nick Gnedin, Martin Haehnelt, Piero Madau and Martin Rees for stimulating discussions and Mark Dijkstra for noticing some incorrect numbers in the first draft of this paper. MR thanks Erika Yoshino support.  \\appendix"
    },
    "0311/astro-ph0311529_arXiv.txt": {
        "abstract": "{High resolution spectra of 34 halo population dwarf and subgiant stars have been obtained with VLT/UVES and used to derive sulphur abundances from the $\\lambda \\lambda 8694.0, 8694.6$ and $\\lambda \\lambda 9212.9, 9237.5$ \\SI\\ lines. In addition, iron abundances have been determined from 19 \\FeII\\ lines and zinc abundances from the $\\lambda \\lambda 4722.2, 4810.5$ lines. The abundances are based on a classical 1D, LTE model atmosphere analysis, but effects of 3D hydrodynamical modelling on the [S/Fe], [Zn/Fe] and [S/Zn] ratios are shown to be small.  We find that most halo stars with metallicities in the range $-3.2 < \\feh < -0.8$ have a near-constant $\\sfe \\simeq +0.3$; a least square fit to \\sfe\\ vs. \\feh\\ shows a slope of only $-0.04 \\pm 0.01$. Among halo stars with $-1.2 < \\feh < -0.8$ the majority have $\\sfe \\simeq +0.3$, but two stars (previously shown to have low $\\alpha$/Fe ratios) have $\\sfe \\simeq 0.0$. For disk stars with $\\feh > -1$, \\sfe\\ decreases with increasing \\feh . Hence, sulphur behaves like other typical $\\alpha$-capture elements, Mg, Si and Ca. Zinc, on the other hand, traces iron over three orders of magnitude in \\feh, although there is some evidence for a small systematic Zn overabundance ($\\znfe \\simeq +0.1$) among metal-poor disk stars and for halo stars with $\\feh < -2.0$. Recent measurements of S and Zn in ten damped Ly$\\alpha$ systems (DLAs) with redshifts between 1.9 and 3.4 and zinc abundances in the range $-2.1 < \\znh < -0.15$ show an offset relative to the \\szn\\ -- \\znh\\ relation in Galactic stars. Possible reasons for this offset are discussed, including low and intermittent star formation rates in DLAs. ",
        "introduction": "\\label{introduction}  Sulphur is generally regarded as an $\\alpha$-capture element. The work on Galactic stars by Fran\\c{c}ois (\\cite{francois87}, \\cite{francois88}) supported this view by showing that [S/Fe] increases from zero at solar metallicities to a plateau level of about +0.5 dex in the metallicity range $-1.8 < \\feh < -1.0$. This is an analogous behaviour to those of other $\\alpha$-elements, Mg, Si, and Ca - see Norris et al. (\\cite{norris01}) and Carretta et al. (\\cite{carretta02}), who both find \\alphafe\\ to be nearly constant at a level of +0.4\\,dex in the metallicity range $-4 < \\feh < -1$ and then to decrease towards the solar \\alphafe\\ ratio for $\\feh > -1$. The standard interpretation is that this trend arises from the time delay in the production of about two thirds of the iron by supernovae (SNe) of Type Ia relative to the near-instantaneous release of the $\\alpha$-elements by Type II SNe. In this connection it should be noted that the metallicity at which Type Ia SNe start to contribute (the ``knee'' of the [$\\alpha$/Fe] - [Fe/H] trend) seems to depend on the orbital properties of the stars. Halo stars with $\\feh \\sim -1$ belonging to the outer halo tend to have lower \\alphafe\\ than stars in the inner halo (Nissen \\& Schuster \\cite{nissen97}; Stephens \\& Boesgaard \\cite{stephens02}; Gratton et al. \\cite{gratton03}). Stephens \\& Boesgaard find a slope of $-0.15$ for the mean \\alphafe\\ vs. \\feh\\ for 53 stars in the metallicity range $-3.8 < \\feh < -0.6$, but this is for a sample with special kinematics, i.e. stars with large maximum distances from the Galactic center and/or the Galactic plane or stars having extreme retrograde motion.  Recent observations of the $\\lambda \\lambda 8694.0, 8694.6$\\,\\AA\\ \\SI\\ lines in spectra of  metal-poor stars by Israelian \\& Rebolo (\\cite{israelian01}) have, however, challenged the view that sulphur is an $\\alpha$-element. Their data, obtained with the 4-m William Herschel Telescope on La Palma, suggest that \\sfe\\ increases linearly with decreasing \\feh\\ to a level as high as $\\sfe \\sim +0.7$ at $\\feh = -2.0$. The study of Takada-Hidai et al. (\\cite{takada02}), based on Keck High Resolution Echelle Spectrograph (HIRES) observations, supports a quasi linear dependence of \\sfe\\ on \\feh\\ although in their case \\sfe\\ reaches only $+0.5$ dex at $\\feh = -2.0$.  As a possible explanation of the high value of \\sfe\\ in metal-poor stars, Israelian \\& Rebolo (\\cite{israelian01}) proposed that very massive supernovae with exploding He-cores and a high explosion energy make a significant contribution to the early chemical evolution of galaxies. According to Nakamura et al. (\\cite{nakamura01}) these hypernovae overproduce S with respect to O, Mg and Fe. With this intriguing possibility in mind a more thorough investigation of sulphur abundances in halo stars seems worthwhile.  A clarification of the trend of S abundances is also much needed in deciphering the chemical enrichment of damped Ly$\\alpha$ systems (DLAs), widely regarded as the progenitors of present-day galaxies at high redshift. The importance of sulphur stems from the fact that, unlike most other heavy elements, S is not depleted onto dust. Consequently, observations of the relatively weak \\SII\\ triplet resonance lines at $\\lambda\\lambda 1250, 1253, 1259$\\,\\AA\\ yield a direct measurement of the abundance of S in DLAs. Another element for which this is the case is Zn, and indeed most of our current knowledge of the chemical evolution of the universe at high redshift is based on surveys of [Zn/H] in DLAs (e.g. Pettini et al. \\cite{pettini99}; Prochaska \\& Wolfe \\cite{prochaska02}). {\\it If} S is an $\\alpha$-capture element, then its abundance relative to Zn (assumed to be an iron-peak element) could be used as `a chemical clock' to date the star-formation process at high $z$. Specifically, if a major star formation episode in a DLA occurred within $\\approx 0.5$\\,Gyr prior to the time when we observe the galaxy, we would expect to measure an enhanced [S/Zn] ratio, and {\\it vice versa}. Measurements of [S/Zn] in DLAs have been relatively scarce until recently, but are now becoming available in increasing numbers thanks the Ultraviolet and Visual Echelle Spectrograph (UVES) on the ESO Very Large Telescope (VLT). The VLT/UVES combination affords high spectral resolution and high efficiency over most of the optical spectrum, from blue and near-UV wavelengths to the far red. However, without a secure knowledge of the behaviour of S in metal-poor Galactic stars we clearly stand little chance of interpreting the situation at high $z$.  In the present paper we report on a survey of sulphur abundances in 34 metal-poor dwarf stars based on high resolution observations of the $\\lambda \\lambda 8694.0, 8694.6$ and $\\lambda \\lambda 9212.9, 9237.5$\\,\\AA\\ \\SI\\ lines obtained with UVES. The $\\lambda \\lambda 4722.2, 4810.5$\\,\\AA\\ \\ZnI\\ lines are also included in our spectra, allowing us to study the Galactic evolution of both sulphur and zinc.  \\begin{figure} \\vspace{-3cm} \\hspace{-0.5cm} \\resizebox{10.0cm}{!}{\\includegraphics{n0063.f1.eps}} \\vspace{-2.5cm} \\caption{Portion of the VLT/UVES spectrum of \\object{HD\\,110621} ($V = 9.9$, $\\feh = -1.66$) in the far red. Top panel: the spectrum before removal of the numerous telluric \\water\\ lines which occur in this region. Bottom panel: the spectrum after division by that of the B-type star \\object{HR\\,5488} using the IRAF task {\\tt telluric}. The \\SI\\ triplet, a \\MgII\\ line and the Paschen $\\zeta$ \\HI\\ line of \\object{HD\\,110621} can all be clearly seen. Note that the spectrum shown has not been corrected for the radial velocity of the star (which introduces a shift of $+6.4$\\,\\AA ), and that the wings of the \\HI\\ line have been removed by the continuum fitting. Excessive noise in the stellar spectrum is seen where strong \\water\\ lines have been removed.} \\label{fig:telluric} \\end{figure} ",
        "conclusions": "\\label{summary} One of the motivations for carrying out the present survey of S and Zn abundances in Galactic stars is the important role played by these two elements in deciphering the chemical enrichment and star formation histories of damped Ly$\\alpha$ systems. Up to now measurements of the sulphur abundance in Galactic stars have been sparse and doubts have been raised as to whether S is a typical $\\alpha$-capture element formed in Type II supernovae (Israelian \\& Rebolo \\cite{israelian01};  Takada-Hidai et al. \\cite{takada02}). The lack of data for sulphur, coupled with insufficient knowledge of the nucleosynthetic channels for the production of Zn, has led some (e.g. Prochaska et al. \\cite{prochaska00}) to question the validity of using these elements, and in particular their ratio, in the interpretation of the abundance patterns seen in DLAs.  The new data presented here have clarified the situation as regards Milky Way stars. By targeting the \\SI\\ $\\lambda\\lambda 9212.9, 9228.1, 9237.5$ triplet we have overcome the limitations of most earlier studies and probed the abundance of sulphur with higher precision and to lower metallicities than had been possible previously. We find that the trend of [S/Fe] as a function of \\feh\\ is very similar to those of other typical $\\alpha$-capture elements, Mg, Si and Ca. [S/Fe] is nearly constant at a level of $\\sfe \\simeq +0.3$\\,dex in the metallicity range $-3.2 < \\feh < -1.2$, starts to decrease at $\\feh \\simeq -0.7$, and reaches a solar ratio at $\\feh \\sim 0$. In the halo-disk transition region, at metallicities $-1.2 < \\feh < -0.7$, there is significant scatter in the values of \\sfe. More stars in this interval should be studied to look for possible correlations between stellar kinematics and \\sfe . Precise abundance analyses by Fuhrmann (\\cite{fuhrmann98}); Gratton et al. (\\cite{gratton00}); and Feltzing et al. (\\cite{feltzing03}) have shown a clear separation in \\mgfe\\ between thin and thick disk stars and it will be important to establish if this is also the case for \\sfe.  As far as zinc is concerned, we confirm the results of several earlier surveys which have shown that, to a first approximation, Zn tracks Fe over three orders of magnitude in \\feh. There is some evidence for a small overabundance of Zn ($\\znfe \\simeq +0.10$) for metal-poor disk stars and halo stars with $\\feh < -2$. However, there may be systematic errors in \\znfe\\ at a level of $\\pm 0.1$\\,dex due to the difficulty in analyzing the $\\lambda \\lambda 4722.2, 4810.5$ \\ZnI\\ lines in the solar spectrum. Non-LTE effects on the derived zinc abundances are also a potential problem that should be addressed. New studies of the nucleosynthesis of zinc in supernovae and by the $s$-process in AGB stars would help in understanding the reasons why Zn behaves like an Fe-peak element and the origin of the 0.1\\,dex offset at low metallicities.  When we turn to damped Ly$\\alpha$ systems, however, the Galactic pattern of S and Zn abundances does not seem to apply. The sample is still small (accurate measurements of [S/Zn] are available for only ten DLAs) and there is scatter in the data but, taken at face value, there is no obvious $\\alpha$-element enhancement in DLAs. We can not identify any systematic effect in the data nor in their analysis which would mask such an enhancement, if it were there, and conclude that its absence is probably real. Presumably, it is indicative of a bursting history of star formation in the galaxies giving rise to damped systems. However, it is important to realise that a variety of recent observations suggest that the $\\alpha$-enhancement exhibited by metal-poor stars of the Milky Way may in fact be the exception rather than the rule---it is generally not seen in dwarf spheroidal and dwarf irregular galaxies (Venn et al. \\cite{venn03}; Shetrone et al. \\cite{shetrone03}; Tolstoy et al. \\cite{tolstoy03}; Aloisi et al. \\cite{aloisi03}), in old stars of the Large Magellanic Cloud (Hill et al. \\cite{hill00}), in DLAs (this paper), in some Galactic halo stars with large orbits (Nissen \\& Schuster \\cite{nissen97}), and in the globular cluster \\object{Pal\\,12} which may originally have been part of the \\object{Sgr\\,dSph} galaxy (Cohen \\cite{cohen03}). The challenge now is to incorporate this rapidly growing set of abundance measurements into a comprehensive picture of the chemical evolution of galaxies."
    },
    "0311/astro-ph0311235_arXiv.txt": {
        "abstract": "We outline the theory and practice of measuring the four Stokes parameters of spectral lines in emission/absorption observations. We apply these concepts to our Arecibo HI absorption line data and present the results. We include a detailed discussion of instrumental effects arising from polarized beam structure and its interaction with the spatially extended emission line structure. At Arecibo, linear polarization (Stokes $(Q,U)$ has much larger instrumental effects than circular (Stokes $V$). We show how to reduce the instrumental contributions to $V$ and to evaluate upper limits to its remaining instrumental errors by using the $(Q,U)$ results. These efforts work well for opacity spectra but not for emission spectra. Arecibo's large central blockage exacerbates these effects, particularly for emission profiles, and other telescopes with weaker sidelobes are not as susceptible. We present graphical results for 41 sources; we analyze these absorption spectra in terms of Gaussian components, which number 136, and present physical parameters including magnetic field for each. ",
        "introduction": "In February 1999 we used the Arecibo\\footnote{The Arecibo Observatory is part of the National Astronomy and Ionosphere Center, which is operated by Cornell University under a cooperative agreement with the National Science Foundation.} telescope to begin a series of Zeeman-splitting measurements of the 21-cm line in absorption against continuum radio sources. Heiles \\& Troland (2001a, 2001b; papers I and II) reported on a by-product of this survey, namely the Stokes $I$ data from which spin temperatures and other information were derived. The present paper focuses on the technical aspects of processing the polarized spectral data and evaluating the instrumental errors. We also present the derived magnetic fields for 41 sources, which have 136 Gaussian components.  There has been much discussion of HI Zeeman splitting measurements because the polarized sidelobes, interacting with the angular structure of the HI emission, can produce instrumental effects. Even in the emission-absorption measurements presented in this paper, these effects can in principle be serious. Therefore, we discuss these effects in considerable detail. The bottom line is that instrumental effects are evaluated for each source independently and, generally, are negligible for the opacity spectra. However, for the expected emission spectra they are not negligible.  \\S \\ref{stokestheory} outlines the basic theoretical concepts involving Stokes parameters and emission/absorption lines. \\S \\ref{stokespractice} discusses the least squares fitting process required to extract the expected emission profile and the absorption spectrum by combining calibrated on- and off-source data.  \\S \\ref{generaldiscussion} discusses the physical reasons for and contributors to instrumental effects that arise from polarized beam structure.  The next three sections deal with instrumental contributions to the {\\it opacity} spectra. We begin by treating Stokes $V$ in  \\S \\ref{problemevaluation}, discussing the empirical least-squares evaluation and elimination of the two most basic instrumental effects, namely the trigonometric dependences on parallactic angle $PA$ and $2PA$ (squint-like and squash-like dependences, respectively).  We use the terms ``squint-like'' and ``squash-like'' for these empirically-determined dependencies because they include contributions from far-out sidelobes; in contrast, we use the terms ``squint'' and ``squash'' (sometimes preceded by the clarifier ``true'') for  the contributions from only the primary beam and first sidelobe\\footnote{The terms squint and squash normally refer to only to the main beam, not including the first sidelobe}. We then do the same in \\S \\ref{empiricalqu} for the linearly polarized Stokes parameters $(Q,U)$. We reach the important conclusion that instrumental effects in linear polarization are about ten times larger than in circular. This allows us to make an independent estimate of instrumental effects for $V$ using $(Q,U)$, as we discuss explicitly in \\S \\ref{usingqu}.  Our final discussion of instrumental effects (\\S \\ref{emissionone}) addresses the reliability of Zeeman splitting results for {\\it emission} profiles. Here we independently evaluate ``true'' squint and squashand also the squint-like and squash-like contributions. The difference between these is the contributions from the far-out sidelobes. The far-out sidelobe contribution is large for the Arecibo telescope. This makes the instrumental effects quite serious for emission profiles. Accordingly, we do not discuss the Zeeman splitting results for the emission profiles.  Finally, \\S \\ref{results} presents the profiles for all sources and a tabular list of the results for the Gaussian components. We select a good sample for statistical analyses, successfully compare with previous literature, and point out yet another source of uncertainty in the derived magnetic fields. ",
        "conclusions": "\\label{vproblems}  \\label{generaldiscussion}  This section discusses the most important instrumental contributions to instrumental polarization. Polarized opacities are small, so we must consider systematic effects at very low levels.  The most important contributors include the following: \\begin{enumerate}  \\item There is instrumental coupling between the polarized Stokes parameters and Stokes $I$, which creates replicas of the Stokes $I$ line in the polarized Stokes parameter spectra. These couplings are described by the Mueller matrix (Heiles et al 2001b) and are corrected for, but the corrections are imperfect.  Experience teaches us that this is usually the dominant instrumental contribution to the measured Stokes $V$ spectra, typically amounting to a few tenths of a percent. Fortunately, in deriving Zeeman splitting its presence is unimportant: observers normally call this a ``gain error'' and use a least squares technique to remove it.  \\item The Mueller matrix is derived from observations of small-diameter continuum sources, so it applies to beam center. However, it changes within the telescope beam. For Stokes $V$ the primary effect is  ``beam squint'', for which the $V$ beam has positive and negative lobes on opposite side of beam center. This arises from the two polarizations pointing in slightly different directions. Beam squint interacts with the first spatial derivative of the Stokes $I$ profile to produce a false contribution. We assume that the beam squint is fixed with respect to the feed, so its false contribution varies periodically with the parallactic angle $PA$.  Beam squint is theoretically predicted to occur for Stokes $V$ but not for $Q$ and $U$. However, Arecibo has significant beam squint in $Q$ and $U$ as measured by the response to the first spatial derivative in Stokes $I$; it is probably produced by the significant aperture blockage. Figures 12, 13, and 14 of Heiles et al (2001a) show that the effective beam squint for both $Q$ and $U$ is about ten times that for $V$.  Beam squint correlates with the angular structure in the {\\it emission} line to produce unreal features in the {\\it polarized} Stokes emission and absorption spectra. Experience shows that for Stokes $V$, beam squint interacting with the first derivative of HI intensity is the most serious instrumental problem (Heiles 1996). For conventional telescopes this contribution can be measured and removed rather accurately. For the Arecibo telescope, however, the polarized beam pattern changes partly systematically, partly erratically with hour angle, making the correction impractical.  We can estimate the magnitude of this contribution from the measured angular derivatives of the Stokes $I$ expected profile and the approximately-known magnitude of the beam squint (\\S \\ref{twopointfive}). Also, we can determine the approximate contribution by fitting for the change in the apparent polarized opacity spectrum as the polarized beam rotates with respect to the sky with hour angle (\\S\\ref{problemevaluation}). However, because Arecibo's beam characteristics, including squint, change while tracking, the squint contribution varies not only with $PA$ but also has a semi-random component.  We least-square fit equation  \\ref{trigeq} for the systematic squint component, as discussed below, but the semi-random component remains.  \\item For Stokes $Q$ and $U$, the primary effect for the polarized beam is ``beam squash'', for which the polarized beams have a four-lobed cloverleaf pattern, with two positive lobes on opposite side of beam center and two negative ones rotated $90^\\circ$. This arises from the two polarizations having slightly different beamwidths. Beam squash interacts with the second spatial derivative of the Stokes $I$ emission profile to produce a false contribution that varies periodically with twice the parallactic angle $2PA$.  Beam squash is theoretically predicted to occur for Stokes $Q$ and $U$, but not for $V$. However Arecibo, with its significantly blocked aperture, violates this rule. As with beam squint, Figures 12, 13, and 14 of Heiles et al (2001a) show that Arecibo's effective beam squash for both $Q$ and $U$ is about ten times that for $V$.  For  Stokes $V$, most telescopes have negligible beam squash, so our experience with them does not apply to Arecibo. We can estimate the magnitude of the squash contribution from the measured angular derivatives of the Stokes $I$ expected profile and the approximately-known magnitude of the beam squash (\\S \\ref{twopointfive}). Also, we can approximately determine the contribution by fitting for the change in the apparent polarized spectrum as the polarized beam rotates with respect to the sky with hour angle (\\S\\ref{problemevaluation}); this works well for the opacity spectrum. However, it does not work well for the expected emission spectrum because of the incomplete $PA$ coverage (\\S \\ref{emissionone}).  For Stokes $Q$ and $U$ we can estimate the squash contribution from the measured spatial derivatives and the known beam squash. However, we cannot determine it by fitting for the change with $2PA$ because real linear polarization also varies with $2PA$.  \\item All four Stokes parameters are affected by ``far-out sidelobes'', i.e.\\ sidelobes outside the first sidelobe. These arise primarily from ordinary diffraction and, secondarily, from surface inaccuracies. Far out sidelobes for Arecibo are particularly strong because of the severe aperture blockage. And they contribute disproportionately to the squint and squash response. They are likely to be polarized comparably to the first sidelobe, and the HI spatial structure within them changes more because the angular differences are larger; this magnifies their importance. The first sidelobe itself serves as an illustration: for a uniform extended source in Stokes $I$, it contributed 0.34 as much as the primary beam in 2000\\footnote{Arecibo's surface was readjusted and the focus point changed in 2002. The current fraction is much lower, $\\lesssim 0.10$.}); yet in Stokes $V$ it contributed almost {\\it twice} as much to the observed squint response as does the primary beam! (see Table 1 and Figure 14 of Heiles et al 2001a).  Far-out sidelobes can, like squint and squash, be expressed as a Fourier series in $PA$. The first term is squint-like and the second is squash-like. Below, we use the terms {\\it squint-like} and {\\it squash-like} to denote the contribution from the telescope, i.e.\\ from the primary beam and all sidelobes, and unless the context dictates otherwise we use the terms {\\it squint} and {\\it squash} to denote the contributions from the primary beam and first sidelobe\\footnote{To be precise, true squint and squash refer only to the main beam without the first sidelobe. But Heiles et al (2001a) measured the influence of both---for Arecibo, the first sidelobe is large because of the large aperture blockage. For this reason, in this paper we stretch the definition of ``true''.}.  \\end{enumerate}  Contribution (1) produces only a replica of the Stokes $I$ opacity profile; in Zeeman splitting this is routinely removed by least-squares fitting and has no damaging effect. Accordingly, we do not consider it further.  The  contributions (2), (3), and (4) exhibit sinusoidal dependences on either $PA$ (squint-like) or $2PA$ (squash-like). There are two independent ways to evaluate these contributions, each with its own problem. The {\\it empirical} way (\\S \\ref{problemevaluation}) is observing a range of PA and performing a least squares fit on the results to directly evaluate (and remove) the sinusoidal dependences; the problem, especially for the emission profile, is the incomplete $PA$ coverage. The other way (\\S \\ref{twopointfive}) {\\it predicts} the instrumental effects using the angular derivatives of the HI structure in the sky, together with the already known polarization structures of the Arecibo beam; the problem is the neglect of the far-out sidelobe contribution. We turn to a detailed discussion of these issues in the next few sections.    We discuss the measurement of polarized Stokes parameter profiles $(Q,U,V)$ from emission/absorption line observations  toward background continuum sources. For each of these Stokes  parameters, we derive the {\\it opacity profile} and the {\\it emission  profile} expected if the continuum source were absent.  We give  special emphasis to the evaluation of instrumental effects.  A  principal motivation for our study is detection of the Zeeman  effect in Stokes $V$ profiles. Arecibo suffers with respect to  other telescopes because its large central blockage produces  large sidelobes. The sidelobes are polarized, and their  interaction with spatially extended 21-cm emission produces most  of the instrumental contributions to the polarized profiles. These contributions are particularly serious for emission  profiles. But Arecibo's large collecting area compensates for  this problem when measuring opacity profiles. As a result, our  Stokes $V$ opacity profiles are generally reliable so we can make  21-cm Zeeman effect measurements in the Galactic CNM.  In \\S \\ref{stokestheory} we outline the basic theoretical concepts involving Stokes parameters in emission/absorption line observations. Here we derive the fundamental equation \\ref{eqns0}. This equation relates the $observed$ Stokes parameter profiles (on and off-source) to the physically significant polarized opacity  profiles and polarized emission profiles. In \\S \\ref{stokestheory} we also explain the contribution of the Zeeman effect to the Stokes $V$ profiles.  This contribution is always proportional to the  derivative of the Stokes $I$ profile (emission or absorption),  regardless of the line opacity.  Therefore, we can always derive  line-of-sight magnetic fields by fitting the Stokes $V$ profile  to the derivative of the Stokes $I$ profile.  In \\S \\ref{stokespractice} we present equation \\ref{eqnten}. This  is the fundamental equation that we least squares fit (independently for each spectral channel) to derive the Stokes  $V$ opacity profile and the Stokes $V$ emission profile for each  source. This equation includes an instrumental error term which  is extensively discussed in later sections. In \\S  \\ref{stokespractice} we also present the related discussion for Stokes $Q,U$ profiles, including equation \\ref{skymatrix6}. This equation, the analog of equation \\ref{eqnten} for Stokes $V$, includes a rotation matrix to account for the parallactic angles of the observations.  Much of the paper is concerned with instrumental effects arising from polarized beam structure.  This structure interacts  with spatially extended line emission.  As a result, instrumental  effects appear in the polarized emission profiles and, to a much smaller extent, in the polarized opacity profiles. In \\S  \\ref{generaldiscussion} we describe the nature and the effects of  various types of polarized beam structure.  One type is ``beam  squint''. Beam squint interacts with the first spatial derivative  of the Stokes $I$ profile to make instrumental contributions to  the Stokes $Q,U,V$ profiles. Unfortunately, the beam squint  contribution to the Stokes $V$ profiles can mimic the Zeeman effect. If beam squint remains fixed relative to the telescope  feed system, then its contributions to the polarized Stokes  profiles vary with the parallactic angle $PA$.  Another type of  polarized beam structure is``beam squash''. Beam squash interacts  with the $second$ derivatives of the Stokes $I$ profile to make  instrumental contributions to Stokes $Q,U,V$. If beam squash  remains fixed relative to the telescope feed system, then its  contributions vary as $2PA$. For prime focus telescopes, beam squint is theoretically  expected only for Stokes $V$, and beam squash is expected only  for Stokes $Q,U$. However, Arecibo breaks these rules, having  beam squint and squash in all polarized Stokes parameters. The  final type of polarized beam structure arises outside the primary  beam and first sidelobe, in the ``far-out sidelobes''.  Far out  sidelobes at Arecibo are particularly strong because of the large  aperture blockage.  We define all contributions to the polarized  Stokes profiles as ``squint-like'' and ``squash-like'' if they are  functions of $PA$ and $2PA$, respectively. In practice, squint-like contributions may arise from true beam squint (in the  primary beam and nearest sidelobe) and, also, from the far-out sidelobes.  The same is true for beam squash.  In \\S \\ref{problemevaluation} we treat instrumental  squint-like and squash-like contributions to Stokes $V$ opacity  profiles.  These profiles are of particular interest for the  Zeeman effect.  Equation \\ref{trigeq} expresses these instrumental  contributions to the fundamental fitting equation \\ref{eqnten} as sinusoidal functions of $PA$ (squint-like) and $2PA$ (squash-like). We fit equation \\ref{eqnten} in three ways: (a) with neither  squint-like nor squash-like contributions included, (b) with  squint-like contributions only and (c) with both squint and  squash-like contributions. Results of these three types of fits  are Stokes $V$ opacity profiles and emission profiles with (a) no  instrumental effects removed, (b) squint-like effects removed and  (c) both squint- and squash-like effects removed, respectively.  Fits of type (c) are usually best for the Stokes $V$ opacity  profiles, allowing us to remove squint and squash-like contributions accurately. In \\S \\ref{3C138examplev}, we illustrate these concepts for 3C138 which has a clearly-detected  Zeeman effect. In this section, we also develop a matrix  representation of the coupling coefficients between the fitted  Stokes $V$ opacity profile for a given source and the squint and  squash-like instrumental contributions.  From this matrix, we  derive a profile of $\\delta \\tau'_{V,2}(\\nu)$ for each source.   This profile represents the maximum possible instrumental contributions to the Stokes $V$ opacity profile after squint and squash-like contributions have been removed by the fitting  process.  An example for 3C138 is shown in the bottom panel of  Figure 9.  Here the profile of $\\delta \\tau'_{V, 2}(\\nu)$ is insignificant compared to the Zeeman effect in the Stokes $V$ opacity profile.  For most sources, the profile $\\delta \\tau'_{V,2}(\\nu)$ is very small compared to that of the Stokes $V$ opacity  profile, another reason to expect that the latter are reliably  determined.  \\S \\ref{empiricalqu} treats squint-like and squash-like contributions to the linear polarization Stokes $Q,U$ opacity  profiles. Least squares fits for squash-like contributions are  not possible: squash-like effects have the same $2PA$ dependence  as true linear polarization, so the two cannot be distinguished.  However, true linear polarization of the HI opacity profile  should be very small, particularly for small sources like 3C454.3  (angular size 14 milliarcsec). Therefore, the apparent linear polarization we measure in the 3C454.3 opacity profile must be  instrumental.  We judge that the same conclusion holds for all  sources, leading to a second method of estimating instrumental  effects in Stokes $V$ opacity profiles.  The Arecibo beam squint  and squash are known to be about 10 times greater in Stokes $Q,U$  than in Stokes $V$. Therefore, we can estimate instrumental  polarization in Stokes $V$ opacity profiles by dividing the  apparent linear polarization in the Stokes $Q,U$ profiles by 10. Since squint and squash-like effects have already been removed  from Stokes $V$ opacity profiles, this estimate applies to  instrumental effects (including those from the far-out sidelobes)  that are not a function of $PA$ or $2PA$.  This technique is  described in  \\S \\ref{usingqu}, and an illustrative profile for  3C138 is shown in the bottom panel of Figure 9.  Our final discussion of instrumental effects (\\S \\ref{emissionone}) addresses the reliability of the Stokes $V$  emission profiles.  For these profiles, fits of type (c) usually  suffer from large covariance between the $PA$ and $2PA$-dependent  fit parameters. As a result, the fit parameters are poorly  determined, and the type (c) fits yield very noisy Stokes $V$  emission profiles.  An example for 3C138 is shown in the bottom  panel of Figure 5. In effect, we are unable to remove reliably  the squint- and squash-like contributions from the Stokes $V$  emission profiles so we cannot derive the Zeeman effect from  them. Covariance between the fit parameters would be much less if  we had a larger range of $PA$ included in the data sets for our sources.  In \\S \\ref{emissionone} we also independently evaluate  the contributions of $true$ beam squint and squash (from the main  beam and the first sidelobe) to the Stokes $V$ emission profiles.   We compare these contributions to the empirically-fitted squint  and squash-like contributions (from the main beam and all  sidelobes). The difference between these two is the contribution  from the far-out sidelobes alone. We find that the far-out  sidelobes contribute most of the instrumental effects at Arecibo.   We expect these contributions to be mainly squint or squash-like, but there might also other contributions that do not depend on $PA$ or $2PA$; such contributions contribute additional uncertainties. {\\bf NOTE THIS CHANGE FROM ORIGINAL!}  For all these reasons, Zeeman effect results derived from Arecibo Stokes $V$ emission profiles are unreliable.  \\S \\ref{results} presents the Stokes $I$ and $V$ opacity profiles for all sources, along with profiles of possible  instrumental effects in the latter.  Profiles for 3C138 are shown  in Figure 9, profiles for other sources others are provided  electronically. This section also presents a tabular list of  parameters, including line-of-sight magnetic field strengths, for  the CNM Gaussian components in the Stokes $I$ opacity profiles.  We select a sample 69 Gaussian components for which $\\Delta B_{||}$ (uncertainties in the derived $B_{||}$) should be Gaussian distributed and, additionally, for which $\\Delta B_{||} < 10$ $\\mu$G. {\\bf NOTE THIS CHANGE FROM ORIGINAL!}  We will subject this sample to a future statistical analysis of magnetic field strengths in the CNM.  In \\S  \\ref{results}, we also discuss another source of uncertainty in  derived magnetic field strengths, the unknown sequential order along the  line-of-sight of the various CNM velocity components.  However,  this source of error has no significant effect upon the  qualitative or statistical properties of the magnetic field  measurements."
    },
    "0311/astro-ph0311079_arXiv.txt": {
        "abstract": "The Liverpool Telescope, situated at Roque de los Muchachos Observatory, La Palma, Canaries, is the first 2-m, fully instrumented robotic telescope. It recently began observations. Among Liverpool Telescope's primary scientific goals is to monitor variable objects on all timescales from seconds to years. An additional benefit of its robotic operation is rapid reaction to unpredictable phenomena and their systematic follow up, simultaneous or coordinated with other facilities. The Target of Opportunity Programme of the Liverpool Telescope includes the prompt search for and observation of GRB and XRF counterparts. A special over-ride mode implemented for GRB/XRF follow-up enables observations commencing less than a minute after the alert, including optical and near infrared imaging and spectroscopy. In particular, the moderate aperture and rapid automated response make the Liverpool Telescope excellently suited to help solving the mystery of optically dark GRBs and for the investigation of currently unstudied short bursts and XRFs. \\vspace{1pc} ",
        "introduction": "Robotic telescopes are, due to their efficiency and flexibility, opening up numerous possibilities in many important branches of astrophysics \\cite{Bode}. Distinct advantages of robotic telescopes are: \\begin{itemize} \\item{ the rapid response to Targets of Opportunity (ToO) alerts,} \\item{the efficiency in the systematic monitoring of variable objects on all timescales from seconds to years,} \\item{ simultaneous and coordinated observations with other facilities, either ground or space based,} \\item{conditions (seeing, photometricity, etc.) dependent observations, and} \\item{effective small scale surveys and routine tasks.} \\end{itemize} They are invaluable in the study of variable astronomical objects of various types, including short and unpredictable phenomena, such as GRBs, XRFs, supernovae, dwarf novae and gravitational lensing; all types of binary systems and variable stars; and reaching from distant sources such as active galactic nuclei to the study of comets and near-Earth asteroids in our Solar system.  With first light achieved at the end of July 2003, the Liverpool Telescope became the world's largest operational robotic telescope and joins the international efforts in the study of exciting new fields of research in time domain astrophysics. \\begin{figure} \\centering \\includegraphics[height=5.2cm]{g_fig1.eps} \\caption{The Liverpool Telescope at Roque de los Muchachos, La Palma, Canaries is a 2-m fully robotic altitude-azimuth design telescope with fully openable enclosure, which minimises the slewing time to the position of a ToO and enable observations starting less than a minute after the alert.} \\label{figLT} \\end{figure} ",
        "conclusions": "The unique r\\^ole that the robotic telescopes can play in the study of variable and transient sources has long been recognized. With the start of observations, the LT has great potential for contributing to many interesting areas of time variable astrophysics. The short response time, moderate aperture, excellent site and LT's instrumentation including optical camera, infrared imager and spectrograph, promise an interesting and valuable scientific harvest from the ToO follow-up observations, especially afterglows of GRBs and XRFs. In collaboration with satelites (Swift, HETE-2, INTEGRAL) and other ground-based facilities, including the Faulkes Telescopes (clones of the LT), we expect to be able to follow-up around 1 in 6 GRB/XRF immediately after the alert and we plan to monitor scientifically interesting afterglows also into their later stages.  \\vspace{0.5cm} {\\it Acknowledgments} \\\\ The Liverpool Telescope is funded via EU, PPARC, JMU grants and the benefaction of Mr. A. E. Robarts. A.G. acknowledges the receipt of the Marie Curie Fellowship from the EU."
    },
    "0311/astro-ph0311553_arXiv.txt": {
        "abstract": "We investigate chemical enrichment in Damped Lyman $\\alpha$ (DLA) systems in the hierarchical structure formation scenario using a semi-analytic model of galaxy formation.  The model developed by Nagashima, Totani, Gouda \\& Yoshii takes into account various selection effects on high-redshift galaxies and can show fundamental observational properties of galaxies, such as luminosity functions and number-magnitude/redshift relations. DLA systems offer the possibilities of measuring metal abundance more accurately than faint galaxies. For example, recent measurements of zinc abundance can provide good evidence for understanding the processes of metal pollution and star formation in DLA systems because zinc is virtually unaffected by dust depletion.  Here we focus on this advantage for observation in order to explore the metallicity evolution in DLA systems at high redshifts. We can consistently show the metallicity evolution for reasonable models which also reproduce fundamental properties of local galaxy population. This result suggests that the chemical evolution of DLA systems can be consistently reconciled with the observational features of typical galaxies.  We also investigate other properties of DLA systems (column density distribution and mass density of cold gas), and find that star formation in massive galaxies should be more active than that in low-mass ones. This is consistent with the results by Nagashima et al. and Cole et al. in which the star formation timescale is set by reproducing cold gas mass fraction in local spiral galaxies. Finally we discuss host galaxies associated with DLA systems. We conclude that they primarily consist of sub-$L^{*}$ and/or dwarf galaxies from the observations. ",
        "introduction": "The absorption-line systems observed in the spectra of background quasars have been proven to be a strong probe into the physical conditions of the Universe at high redshifts: abundances of neutral gas and metals, kinematic properties, etc. In particular, Damped Lyman $\\alpha$ Absorption (DLA) systems, which have been associated with cold gas in protogalactic disks \\citep{WTSC}, provide some advantages in the investigation of various characteristics of primordial galaxies similar to faint ones at high redshifts. For example, we can obtain abundant data in the early universe from quasar spectra which are less affected by observational limits caused by faintness of absorbers and chemical abundance, etc. Therefore, they have been studied extensively at redshifts $0\\la z \\la 4.5$ to explore the physical processes in galaxies: star formation rates, abundances of metals and dust \\citep{FP93, L96, PSKH, PESB, P00, PW99, PW00}, column density distribution \\citep[e.g.][]{SW00, RT00} and kinematics \\citep{PW97, PW98}. These observational features would provide important constraints on clarifying the evolutionary link between DLA systems and typical galaxies.  In particular, in DLA systems, we can measure abundances of various elements (Fe, Si, Ni, Mn, Cr, Zn, etc.) and thereby find good clues to the metallicity evolution of galaxies in the early universe.  For example, relative elemental abundances are one of the most important complements to this information because abundance ratios of elements produced by supernovae in different proportions can provide independent insights into the chemical evolution. However, it has been difficult to establish precisely some abundance patterns (e.g. [$\\alpha$/Fe]) in DLA systems because the level of depletion of refractory elements onto dust grains may be uncertain. For this reason, the zinc abundance has recently been recognized as a good tracer of the metallicity evolution instead of \\ion{Fe}{2} because zinc is one of the undepleted elements.  Recent observations have generally reported that the absolute Zn-abundance is $\\sim 10 \\%$ of the solar value with no clear evolution over the redshift \\citep{PSKH, PESB, PW99, VBCM}.  The low abundance and the apparent lack of redshift evolution have been considered interesting problems for DLA systems as proto-galaxies, in particular, at low redshifts where abundances much closer to the solar ones are expected.  This advantage of Zn-measurement confirms that DLA systems so far remain one of the best probes at high redshifts into star formation history.   Recent theoretical studies focus on the nature of DLA systems and physical relations with galaxies observed at high and low redshifts. \\citet{Ka96} studied the formation of DLA and Lyman-limit systems in a hierarchical clustering scenario using a hydrodynamical simulations of a CDM universe($\\Omega=1$). They concluded that the observed \\ion{H}{1} column  density distributions can be reproduced within a factor of 2 if DLA systems are relatively massive and dense protogalaxies. They also showed that their model could reproduce the other properties of DLA observations. However, these results remain somewhat inconclusive. The calculations were restricted to a single CDM cosmology ($\\Omega=1$ CDM). In addition, they could not include galactic disks embedded in halos with circular velocities $V_{\\rm c} \\la 100$ km s$^{-1}$ caused by limited resolutions ($M \\sim 10^{11}~h^{-1}~M_{\\odot}$). Gardner et al.(1997a,b, 2001) resolved these problems by using a relation between absorption cross-section $\\sigma$ and halo circular velocity $V_{\\rm c}$ obtained from numerical simulations. Assuming that the $\\sigma - V_{\\rm c}$ relation can be applied for previously unresolved halos, they presented numerical results for various cosmological models. As a result, the observed abundance of DLA systems can be reproduced if their extrapolation procedure applied for less-massive halos ($V_{c} \\ga 50-80$ km s$^{-1}$). However, the validity of introducing their correction method remains uncertain. Therefore, their calculations could not still include the contribution of DLA systems below limited resolutions, especially at redshifts $0 \\le z \\le 3$. Haehnelt, Steinmetz \\& Rauch (1998, 2000) focussed on the kinematics to explore the nature of DLA systems in SPH simulations with high resolutions ($\\sim$ a few kpc). They concluded that the high abundance of protogalactic clumps can reproduce the observed velocity width distribution and asymmetries of absorption lines found by Prochaska and Wolfe (1997,1998). Their conclusion suggested that the {\\it majority} of DLA systems are not large, rapidly rotating disks with $V_{\\rm c} \\ga 100$ km s$^{-1}$ but protogalactic clumps with the typical circular velocity of DLA halos $V_{\\rm c} \\sim 100$ km s$^{-1}$. Their numerical simulations presented that their models can reproduce observational properties of DLA systems and provided interesting insights to explore the nature of DLA systems.   Recently a different approach, semi-analytic modelling, has been applied with a view to deciphering the clues to the formation process of galaxies in the hierarchical clustering scenario (e.g., Kauffmann, White, \\& Guiderdoni 1993; Cole et al. 1994, 2000; Somerville \\& Primack 1999; Nagashima, Gouda \\& Sugiura 1999; Nagashima et al. 2001, hereafter NTGY; Nagashima et al. 2002). These approaches have some advantages. For example, semi-analytic modelling can study the effect of star formation or supernovae feedback on galaxy evolution even under simple recipes. It does not also suffer from resolution limitations. This approach can also take into account merging the histories of dark halos based on the power spectrum of the initial density fluctuation, and have successfully provided galaxy formation models for reproducing observational properties of galaxies such as luminosity functions, the Tully-Fisher relation, the relation between \\ion{H}{1} gas mass fraction and luminosities, and so forth. To understand the galaxy formation process, it is also important to focus on observational features of high redshift galaxies.  For example, the faint galaxy number counts are good evidence for constraining the key process of galaxy formation in addition to the local feature.  NTGY confirms, using a similar semi-analytic model, that the fundamental properties of local galaxies can be reproduced by their semi-analytic galaxy formation model in a cosmological constant-dominated flat universe ($\\Lambda$CDM model). Their model shows good agreement with the observational features containing those of galaxies at high redshifts such as galaxy counts, taking into account high-redshift selection effects caused by cosmological dimming of surface brightness, absorption by intergalactic \\ion{H}{1} gas, and internal dust absorption.  Therefore, it is valuable to examine how improvements in the above model affect our predictions for the metallicity evolution of cold gas in high-redshift galaxies. We investigate here the star formation history of DLA systems in our semi-analytic model, and whether we can consistently reproduce the metallicity evolution for the models which best match the observational constraints on local galaxy population. We also address what kinds of physical process, such as star formation timescale, play key roles in the evolution of DLA systems within the framework of the semi-analytic approach investigated here. Furthermore, we study other properties of DLA systems: \\ion{H}{1} column density distribution, mass density of cold gas and so on. Finally we discuss what kinds of galaxies are associated with DLA systems.  In \\S 2, we briefly describe the semi-analytic model used here.  In \\S 3, we show the evolution in chemical enrichment of DLA systems. In \\S 4, we present other properties of DLA systems in our calculation. Host galaxies of DLA systems are also discussed in this section. Finally we summarize our conclusions in \\S 5. ",
        "conclusions": "We investigated the metallicity evolution of DLA systems in a hierarchical galaxy formation scenario using a semi-analytic galaxy formation model given by NTGY, which has been found to show good agreement with many properties of galaxies.  In the previous theoretical work, it has been claimed that the metallicity of DLA systems is too high to reproduce the observed ones \\citep[e.g.][]{spf01}.  We find that, in contrast to previous work, DLA systems in our model have low metallicity, $\\sim 1/10 Z_{\\odot}$, consistent with observational data when we use the same model as that reproducing many observations of local and high redshift galaxies given by NTGY. This conclusion is given by setting the star formation timescale as follows: (1) nearly constant against redshift, and (2) longer timescale in lower circular velocity systems.  Our results suggest $\\tau_{*} \\propto V_{\\rm circ}^{\\alpha_{*}}$ and $\\alpha_{*}=-2$, which is entirely consistent with NTGY. The first assumption (1) leads that there is a large amount of cold gas at high redshift. The second (2) indicates that a stronger $\\alpha_{*}$ dependence causes low star formation rate in low mass galaxies that contribute more to the total number of DLA system. These points mainly led to the result that DLA systems have low metallicity ($Z \\sim 1/10 Z_{\\odot}$) consistent with observational data.  We can also reproduce column density distributions {\\it even} under the assumption that the radial distribution of \\ion{H}{1} column density follows the exponential profile. In addition with our results for the average circular velocities and absolute luminosities, our calculation indicates that DLA systems are primarily composed of less massive systems rather than present-day $L^{*}$ galaxies. In a subsequent paper, we intend to investigate in detail host galaxies of DLA systems in our model, including comparison with observational properties of dwarf galaxies and/or the other types of galaxies.  In further analyses, the following effects might be considered:  First, in our model, the metal enrichment of cold gas is tightly coupled with star formation because we assume here that the metal ejected by stars is completely mixed with cold gas and finally becomes distributed throughout the disk.  However, it is preferable to consider that galactic disks generally have metallicity gradients: outer regions may be metal-poor contrary to the metal-rich inner regions \\citep{P97, T98}. Second, the metal distribution might not be consistently diffuse.  \\citet{CK00} pointed out that typical DLA systems have a multi-phase medium composed of cold and warm \\ion{H}{1} gas induced by the observational fact that the spin temperatures of \\ion{H}{1} 21cm transition, which correspond to the kinetic temperature, appear to be higher ($T \\ga 1000$K) than those typical ($T \\sim 100-200$K) of the Milky Way or nearby spiral galaxies.  Moreover, \\citet{KGC01} and \\citet{LBS00} found that the large amount of neutral gas is in the warm phase in two DLA systems at low redshifts from the spectra well fitted by multi-components with various spin temperatures.  The inhomogeneous gas may arise from the wide range of metallicity in present DLA systems.  Moreover, local observations show that dusty systems generally have larger abundance. In practice, dusty DLA systems systematically fail to be observed because dust in DLA systems dims background quasars. These problems are clearly issues that deserve further investigation.  In the future, the combination of accurate dust-correction using various elements and detections of more DLA systems especially at $0 \\la z \\la 1$ are expected to reveal the origin of DLA systems and the star formation history at high redshifts."
    },
    "0311/astro-ph0311286_arXiv.txt": {
        "abstract": "We present the results of a numerical simulation in which star formation proceeds from an initially unbound molecular cloud core. The turbulent motions, which dominate the dynamics, dissipate in shocks leaving a quiescent region which becomes gravitationally bound and collapses to form a small multiple system. Meanwhile, the bulk of the cloud escapes due to its initial supersonic velocities. In this simulation, the process naturally results in a star formation efficiency of $\\sim 50$\\%. The mass involved in star formation depends on the gas fraction that dissipates sufficient kinetic energy in shocks. Thus, clouds with larger turbulent motions will result in lower star formation efficiencies. This implies that globally unbound, and therefore transient giant molecular clouds (GMCs), can account for the low efficiency of star formation observed in our Galaxy without recourse to magnetic fields or feedback processes. Observations of the dynamic stability in molecular regions suggest that GMCs may not be self-gravitating, supporting the ideas presented in this letter. ",
        "introduction": "The formation of stars in giant molecular clouds appears to be a very inefficient process. Global estimates of the star formation efficiency, based on the mass in molecular clouds and the Galactic star formation rate(Scalo 1986; Evans 1999), are of order a few percent while small-scale estimates for stellar clusters approach 50\\% (Lada 1992; Lada \\& Lada 2003).  The dispersion in estimates can, to some extent, be decreased under the assumption that molecular clouds are long-lived entities that exist for many dynamical times, although this necessitates a supporting mechanism to balance gravity (McKee \\etal 1993).  As star formation from molecular gas involves the gravitational contraction through many orders in magnitude in size, it is commonly assumed that the largest scale objects associated with star formation, molecular clouds, are themselves bound. In this case, the problem is how does a large bound region ($10^{4} - 10^{6}$ \\solmas ) only permit a small fraction of its mass to undergo gravitational collapse. Various mechanisms, such as magnetic fields and feedback from young stars, have been evoked over the years in order to support molecular clouds and explain the low efficiency of star formation (Shu, Adams \\& Lizano 1987; Franco, Shore \\& Tenorio-Tagle 1994). In this letter, we explore an alternative scenario whereby molecular clouds are not bound on the large scale, circumventing the need for any additional supporting mechanisms.  In a recent paper, Elmegreen (2000) has collected observational evidence implying that GMCs are short-lived objects, with dispersal times comparable to their crossing times ($t_{dyn} \\sim 3 \\times 10^{7}$years). The formation of stars would then have to occur in a short burst ($\\sim t_{dyn}$), as has been suggested from recent observations (Hartmann, Ballesteros-Paredes, \\& Bergin 2001; Hartmann 2002), rather than an ongoing process. This implies that molecular cloud formation is dynamical, and self-gravity need not play a dominant role. Such a scenario has recently been advanced by Pringle, Allen \\& Lubow (2001) where the passage of a spiral arm triggers the agglomeration of unbound molecular gas and the ensuing shock dissipates sufficient kinetic energy to allow self-gravity to (locally) form stars. The observed molecular clouds could then be transient structures with lifetimes comparable to their dynamical times (Elmegreen 2000; Pringle et al 2001; Larson 2003). In this picture it would seem likely that these regions are globally unbound and highly dynamic, with only a small fraction achieving gravitational instability before the cloud dissipates. This is possible if the internal bulk motions of the cloud are stronger than the self-gravity of the region. The supersonic nature of the motions produces shocks which dissipate kinetic energy, eventually leading to regions becoming gravitationally unstable. The efficiency of the star formation would then be governed by the strength of the internal kinetic energy at the point of the cloud's formation, along with how much support can be removed in shocks. This would naturally lead to different clouds having different star formation efficiency and potentially explain the variance in the observed estimates.  Observational estimates of the mass, and therefore the energies, of the molecular gas is inherently difficult and uncertain. Heyer, Carpenter \\& Snell (2001) have investigated the dynamic stability of molecular regions in the outer Galaxy. Results suggest that the molecular material is unbound until the mass of the region approaches $10^{5}$\\solmas. It is however noted in the paper that obtaining a value for the hydrogen mass via measurements of CO flux involves the assumption that the gas is virialised. If the region in question is then unbound, the subsequently derived value for the mass of molecular hydrogen is an over-estimate of the region's true mass. As a result, regions will appear more bound than they actually are. This places serious doubt on the generally accepted idea of globally bound GMCs.  The ability of turbulence to character the interstellar medium (ISM) along with its role in star formation in GMCs has been extensively studied by a number of authors. For a complete review we point the reader to the article by Mac Low \\& Klessen (2003). It is now well established that large scale turbulent motions can produce the filamentary structure that is commonly observed in star forming regions (V{\\'a}zquez-Semadeni, Ballesteros-Paredes \\& Rodriguez 1997; Ostriker, Gammie \\& Stone 1998; Ballesteros-Paredes, Hartmann \\& V{\\'a}zquez-Semadeni 1999; Padoan \\& Nordlund 1999) and that even MHD turbulence decays on a dynamical timescale (Mac Low et al 1998; Stone, Ostriker \\& Gammie 1998) such that a driving mechanism is required if the clouds are to exist for many dynamical times. These studies highlight that transient objects can form from large scale flows in the ISM, supporting the work of Elmegreen (2000) that star forming regions are short lived in nature. Recent work on subregions within GMCs has also shown that driven turbulence, and more importantly the scale of the driving, can effect the local efficiency of the star formation (Klessen, Heitsch \\& Mac Low 2000; Klessen 2001). The periodic boundary conditions in this last study make it difficult to find estimates of the global star formation efficiency of transient molecular regions, since the mean mass density cannot decrease. Any newly formed stars can thus continue to accrete mass indefinitely. Global calculations have tended to study clouds that are only marginally unbound and hence do not provide information on transient GMCs (Bate, Bonnell \\& Bromm 2003; Bonnell, Bate \\& Vine 2003)  It is the picture of GMC formation presented by Pringle et al (2001) that we use as the motivation for this letter. They suggest that molecular clouds are formed via the accumulation of small parcles of material. In this manner, we interpret the internal motions of the cloud as a result of the original random motions of the constituent parcles of gas. The internal motions therefore need not be turbulent in the strict definition and so we use the notion of ``turbulence'' loosely in this paper. The formation mechanism for molecular regions supported here is distinctly different from that proposed by Ballesteros-Paredes et al (1999). Since we are concerned with a transient object and one which is not part of some large scale flow, we do not require boundary conditions on our model. Futhermore, since the object in study is transient, we do not include any feedback effects into the internal motions as is consistent with the ideas presented by Elmegreen (2000).  We present here the first numerical simulation in which the turbulent gas is both initially unbound and free to expand, showing that this does result in star formation but with a relatively low efficiency. In section 2 we describe the details of the simulation technique with a discussion of the evolution and efficiency of the star formation given in section 3. We summarise and discuss the implications of this result in section 4.  \\begin{figure} \\centerline{\\psfig{figure=Lagcurves.ps,width=2.9truein,height=2.9truein}} \\caption{\\label{lagradii} The evolution of Lagrangian radii enclosing a fixed fraction of the total mass. The solid lines are from the original simulation and the dashed lines are from the period where the smallest binaries are merged. The vertical dotted line denotes the time in the simulation at which the first protostar is formed.} \\end{figure}  \\begin{figure*} \\centerline{\\psfig{figure=picture1.ps,width=1.8truein,height=1.8truein} \\psfig{figure=picture2.ps,width=1.8truein,height=1.8truein} \\psfig{figure=picture3.ps,width=1.8truein,height=1.8truein}} \\centerline{\\psfig{figure=picture4.ps,width=1.8truein,height=1.8truein} \\psfig{figure=picture5.ps,width=1.8truein,height=1.8truein} \\psfig{figure=picture6.ps,width=1.8truein,height=1.8truein}} \\caption{\\label{evolpics} The panels show logarithmic column density snapshots from the simulation, following the evolution of the molecular cloud. The maximum densities in the first four panels are 0.21, 0.17, 0.13 \\& 3.8 ${\\rm{g\\;cm^{-2}}}$ respectively and the last two panels have a maximum density 100${\\rm{g\\;cm^{-2}}}$, taken to arbitrarily represent the protostars. Minimum column densities are 0.001${\\rm{g\\;cm^{-2}}}$ for the first four panels and 0.01${\\rm{g\\;cm^{-2}}}$ for the last two. The turbulent velocities injected into the cloud at the beginning of the simulation quickly erases the uniform density and imposes a characterising filamentary structure. The kinetic energy is larger than the potential causing the cloud to expand, with the outer layers escaping the self-gravity of the system. Left behind by the escaping mass is a small dense region. The gas at the centre of the dense region, it is able to rid itself of kinetic energy via shocks with the shrouding material. Eventually gravity manages to dominate and produce a dense core $(t \\sim 4t_{ff})$ which fragments to form a 3 body hierarchical system with one protostar being ejected (visible at the top right of the last panel in the figure).} \\end{figure*} ",
        "conclusions": ""
    },
    "0311/astro-ph0311415_arXiv.txt": {
        "abstract": "We report the discovery of \\psr, a 2.95~millisecond pulsar in a nearly circular 1.53~day orbit.  Its narrow pulse width of 43\\,$\\mu$s allows pulse arrival times to be determined with great accuracy.  We have spectroscopically identified the companion as a moderately hot ($T \\approx 8500$\\,K) white dwarf with strong absorption lines.  Radial velocity measurements of the companion will yield the mass ratio of the system.  Our timing data suggest the presence of Shapiro delay; we expect that further timing observations, combined with the mass ratio, will allow the first accurate determination of a millisecond pulsar mass.  We have measured the timing parallax and proper motion for this pulsar which indicate a transverse velocity of $140^{+80}_{-40}$\\,km\\,s$^{-1}$.  This pulsar's stunningly narrow pulse profile makes it an excellent candidate for precision timing experiments that attempt to detect low frequency gravitational waves from coalescing supermassive black hole binaries. ",
        "introduction": "Binary radio pulsars provide a rich laboratory for a wide range of physical inquiry, including neutron star masses and the evolution of binary systems.  The pulsars with the best-determined masses all have eccentric relativistic orbits, spin periods of tens of milliseconds, and masses clustered very tightly around 1.35\\,M$_{\\odot}$ \\citep{tc99, bok+03}.  It is generally thought that millisecond pulsars ($P \\le 10$\\,ms) should have greater mass than these longer-period pulsars due to a larger accreted mass.  There is indeed a statistical suggestion that this is the case, but to date no precise millisecond pulsar mass measurements exist.  Currently, the best-determined millisecond pulsar masses are $m_{\\rm p} = 1.57^{+ 0.12}_{- 0.11}\\,{\\rm M}_\\odot$ for PSR~B1855+09 \\citep{nss03} and $m_{\\rm p} = 1.58 \\pm 0.18\\,{\\rm M}_\\odot$ for PSR~J0437$-$4715 \\citep{vbb+01}.  \\psr\\ was discovered during the Swinburne High Latitude Pulsar Survey, a recently-completed pulsar search using the 13-beam multibeam receiver on the 64-m Parkes radio telescope \\nocite{jac03} (Jacoby 2003; Jacoby et al. in preparation).  This survey is similar to, and an extension of, the highly-successful Swinburne Intermediate Latitude Pulsar Survey \\citep{ebvb01, eb01, eb01b}.  The pulsar was discovered in three adjacent survey beams observed on 2001 January 25 and confirmed on 2001 May 26.  Due to the pulsar's strong broad-band scintillation, the initial position inferred from the discovery data was in error by nearly a full beamwidth; it was not until we obtained a phase-connected timing solution that we realized the correct position.  Among the properties that can make a pulsar useful for precision timing experiments are a bright, narrow pulse profile, extremely stable rotation, proximity to the earth, and a companion star which can be detected and studied.  \\psr\\ appears to possess all of these desirable traits.  Nearly all of the pulsar's flux comes in a single narrow, sharp peak with full-width half-maximum (FWHM) of only 43\\,$\\mu$s, less than 1.5\\% of the 2.95\\,ms pulse period (Fig.~\\ref{fig:1909_profile}).  The brightness of the pulsar varies by at least a factor of 30 in the 21-cm band. The mean flux density of our observations is 3\\,mJy, but this is biased by the fact that we tend to observe longer when the pulsar is bright. Extremely high precision timing observations are possible during episodes of favorable scintillation. ",
        "conclusions": "\\label{sec:conclusions}  Pulse timing precision is often limited by systematic errors that distort the shape of pulsar profiles due to the interstellar medium and imperfect polarimetric and instrumental response.  Intrinsically narrow pulses are less susceptible to these errors, as are pulsars with small dispersion measures.  \\psr\\ is likely to become a cornerstone of future timing array experiments to detect low-frequency gravitational waves. It is noteworthy that \\psr\\ is far from other high-precision pulsars on the sky and therefore probes a unique line of sight.  In the shorter run, further timing of the pulsar, especially with CPSR2, should yield an accurate inclination and a reasonably precise companion mass.  Furthermore, in a modest spectroscopic observing campaign on an 8-m class telescope, the radial velocity orbit of the companion can be measured precisely. These combined results should allow a very reliable pulsar mass measurement, and advance our understanding of binary pulsar evolution and recycling."
    },
    "0311/astro-ph0311309_arXiv.txt": {
        "abstract": "{ We analyze H$\\alpha$ observations of the recurrent nova  \\object{T CrB} obtained during the last decade. For the first time the H$\\alpha$ emission profile is analyzed after subtraction of the red giant contribution. Based on our new radial velocity measurements of the H$\\alpha$ emission line we estimate the component masses of \\object{T~CrB}. It is found that the hot component is most likely a massive white dwarf. We estimate the inclination and the component masses to be $i\\simeq67^\\circ$, $M_{\\rm WD}\\simeq1.37\\,\\pm0.13M_\\odot$ and $M_{\\rm RG}\\simeq1.12\\,\\pm0.23M_\\odot$, respectively. The radial velocity of the central dip in the H$\\alpha$ profile changes nearly in phase with that of the red giant's absorption lines. This suggests that the dip is most likely produced by absorption in the giant's wind.  Our observations cover an interval when the H$\\alpha$ and the $U$-band flux vary by a factor of $\\sim6$, while the variability in $B$ and $V$ is much smaller. Based on our observations, and archival ultraviolet and optical data we show that the optical, ultraviolet and H$\\alpha$ fluxes strongly correlate. We argue that the presence of an accretion disc can account for most of the observed properties of \\object{T~CrB}. ",
        "introduction": "\\object{T Coronae Borealis} is a well known recurrent nova with an orbital period of 227\\fd57 (Kenyon \\& Garcia \\cite{kg}; Fekel et al. \\cite{fek}). The system contains  a M4.5\\,III giant (Murset \\& Schmid \\cite{mur}), which fills its Roche lobe and transfers mass toward a hot component (HC). The optical spectrum of \\object{T~CrB} is dominated by the red giant (RG) with superimposed Balmer emission lines, which are thought to originate in an accretion disc (AD) around the HC (Selvelli et al. 1992; hereafter \\cite{sel}). The high excitation \\ion{He}{ii}\\,$\\lambda$4686 emission line is sometimes also visible (Iijima \\cite{iij}). The H$\\alpha$ emission line is well visible all the time and shows variability with the orbital period as well as large variations on a time scale of years (Anupama \\& Prabhu \\cite{anu}; Zamanov \\& Marti \\cite{zm}). The photometric variability in the long wavelength range is dominated by the ellipsoidal variability of the RG (Shahbaz et al. \\cite{sha}; Belczynski \\& Mikolajewska \\cite{bm}; Yudin \\& Munari \\cite{yum}). At short wavelengths \\object{T~CrB} shows strong flickering activity, with the same parameters as those of normal cataclysmic variables, in spite of the large difference in the size of the AD (Zamanov \\& Bruch \\cite{zam}). Apart from the recurrent nova outbursts, \\object{T~CrB} shows other periods of enhanced luminosity. The amplitude of these variations increases toward the shorter wavelengths and reaches $\\sim2$ mag in the $U$-band (Zamanov \\& Zamanova \\cite{zz}; Hric et al. \\cite{hric}). The equivalent width of the line H$\\alpha$ (EW(H$\\alpha$)) also greatly increases  and reaches $\\sim30-35$ \\AA\\  (Anupama \\& Prabhu \\cite{anu}; Zamanov \\& Marti \\cite{zm}). This activity is not well understood.   \\begin{figure*}[t] \\includegraphics*[width=17cm]{ms2956f1.eps} \\caption[]{H$\\alpha$ emission line profiles before (left panels) and after (right panels) the subtraction of the giant spectrum. The first number in each of the right panels shows the date of observations in YYMMDD format, and the second the orbital phase. Note that these phases were calculated with the orbital period of of Fekel et al. (\\cite{fek}), but the zero phase refers to the inferior conjunction of the RG.} \\label{spec} \\end{figure*}   For more than 30 years the HC in \\object{T~CrB} has been considered as a main sequence star (Kraft \\cite{kraft}; Kenyon \\& Garcia \\cite{kg}). Investigations over the last 10 years, however, suggest that the accreting object must be a white dwarf (WD). First, \\cite{sel}  analyzed long-term ultraviolet (UV) spectral observations of \\object{T~CrB} and concluded that the HC is a WD. Belczynski \\& Mikolajewska (\\cite{bm}) analyzed the ellipsoidal variability of the RG and suggested the presence of a WD. In a study of recurrent novae with giant secondaries, Anupama \\& Mikolajewska (\\cite{am}) also found that the HCs in these systems are WDs embedded in the dense wind of the RG. Finally, Hachisu \\& Kato (\\cite{kato}) successfully modeled recurrent novae outbursts as thermonuclear runaway on a massive WD. However, no system parameters determination based on the radial velocities (RVs) of the two components has been published since Kraft (\\cite{kraft}). Recently, Hric et al. (\\cite{hric}) re-analyzed the RVs of Kraft (\\cite{kraft}) and concluded that the hot component in \\object{T~CrB} is a WD. However their analysis is still based on Kraft's seven measurements of the HC RV.  In this paper we present an analysis of spectral H$\\alpha$ and photoelectric $UBV$ observations of \\object{T~CrB} obtained during  the last decade. This is the first investigation of this object where the RG contribution is subtracted and the \"cleaned\" H$\\alpha$ profile is measured. Based on our H$\\alpha$ radial velocity measurements we obtain a new orbital solution. We also discuss the long-term variability of \\object{T~CrB}. ",
        "conclusions": "The radial velocities of the H$\\alpha$ emission line of the recurrent nova \\object{T~CrB} were precisely measured after subtraction of the red giant spectrum, thus obtaining the unbiased H$\\alpha$ profile. Based on our 38 new radial velocity  measurements we obtain a new solution for the system parameters. The solution indicates that at a system inclination $i>65^\\circ$ the mass of the hot component is below the Chandrasekhar limit and thus the hot component in \\object{T~CrB} is most likely a massive white dwarf. The most likely values of the system inclination and the component masses are $i\\simeq67^\\circ$,  $M_{\\rm WD}\\simeq1.37\\,\\pm0.13M_\\odot$ and $M_{\\rm RG}\\simeq1.12\\,\\pm0.23M_\\odot$.  The radial velocity of the central dip of the H$\\alpha$ emission line profile changes almost in phase with that of the red giant. This suggests that the central dip originates from the giant's wind.  During our observations \\object{T~CrB} showed a period of activity when the system brightness increased by $\\sim2.0$, 1.0 and 0.3 mag in $U$, $B$ and $V$ bands, respectively. The H$\\alpha$ equivalent width also increased by a factor of $\\sim6$ reaching $\\sim30$\\,\\AA\\ and the line had a complex multi-component structure. The phasing of the modulation of H$\\alpha$ emission line radial velocity shows that the emitting region is gravitationally connected to the hot component of the system and most probably related to the outer part of an accretion disc. We argue that a number of other characteristics of  \\object{T~CrB} are also in support of this supposition. The long-term variability is then assumed to be caused by variations of the mass accretion rate, possibly due to changes of the mass transfer rate from the companion or disc instability. This is also supported by our success in modeling the UV-optical-IR SED of \\object{T CrB} with a three-component model: an optically thick accretion disc, a nebular continuum and a red giant. However, to construct a consistent picture and to understand the underlying physical processes in \\object{T~CrB} better, simultaneous UV and optical spectral observations during future active phases are needed. Space observatories offering simultaneous optical/UV spectroscopic capabilities like Kronos (Peterson et al. \\cite{pet}) and WSO/UV (Wamsteker \\& Ponz \\cite{wam}) would be a good solution."
    },
    "0311/astro-ph0311623_arXiv.txt": {
        "abstract": "In the current paradigm of star formation magnetic fields play a very central role. Indeed, they probably help or even channel the initial gravitational collapse of the parent molecular cloud.  But their most spectacular effect is certainly the production of ubiquitous, powerful, bipolar self-collimated jets.  Also, in the last decade, evidence has come that in such systems a magnetosphere links the protostar to its surrounding accretion disc.  I will first briefly review the magnetohydrodynamic models of jets from young stars.  Then, I will discuss some constraints on the magnetospheric interaction since accretion must proceed despite strong stellar magnetic fields.  Finally, I will introduce the 3Reconnection X-wind\" scenario that leads to episodic jet formation with efficient removal of protostellar angular momentum. ",
        "introduction": "Collimated ejection of matter is widely observed in several astrophysical objects: inside our own galaxy from all young forming stars (Hartigan et al. \\cite{hart95}) and some X-ray binaries (Mirabel and Rodriguez \\cite{mir99}), but also from the core of active galaxies (Cao and Jiang \\cite{cj99}, Jones et al. \\cite{jones00}).  All these objects share the following properties: jets are almost cylindrical in shape; the presence of jets is correlated with an underlying accretion disc surrounding the central mass; the total jet power is a sizeable fraction of the accretion power.  Most of observed images show jets that are extremely well collimated {\\em already  close to the source}, with an opening angle of only some degrees.  On the other hand, the derived physical conditions show that jets are highly supersonic.  Indeed, emission lines require a temperature of order $10^4$ K, hence a sound speed $c_s \\sim 10$ km/s while the typical jet velocity is $v_j\\sim 300$ km/s.  The opening angle $\\theta$ of a ballistic hydrodynamic flow being simply $\\tan \\theta = c_s/v_j$, this provides $\\theta\\sim 5^o$ for YSOs, nicely compatible with observations.  Thus, observed jets could well be ballistic.  Note that this qualitative argument is not changed by recent observations of rotation in jets, since the rotation velocity is of the same order than the sound speed (Testi et al. \\cite{testi02}).  Therefore, the fundamental question is {\\em how does a physical system produce an unidirectional supersonic flow~?} This implies that confinement must be intimately related to the acceleration process. To date, the only physical process proved to be capable of accelerating plasma along with a self-confinement relies on the action of a large scale magnetic field (carried along by the jet).  Such a magnetic field is assumed to arise from either advection of interstellar (fossil) magnetic field or from local dynamo (or both).  There are three different situations potentially capable, at all evolutionary stages\\footnote{The evolutionary scenario of a young forming star goes from Classes 0, I (embedded source) to Class II (optically visible Classical T~Tauri star) and Class III (Weak line T~Tauri star: no signs of accretion/ejection).  The correlation between accretion and ejection remains valid from Class 0 to Class II phases (Cabrit et al. \\cite{cabrit92}, Hartigan et al. \\cite{hart95}, Bontemps et al. \\cite{bontemps96}).}, of driving magnetized jets from young forming stars:\\\\ $\\bullet$ {\\bf The protostar alone}: stellar winds extract both energy and mass from the protostar itself (Sauty et al. \\cite{sauty02}). \\\\ $\\bullet$ {\\bf The accretion disc alone}: ``disc winds'' are produced from a large radial extension in the disc and are fed with both mass and energy provided by accretion (Casse and Ferreira \\cite{cf2000a}).\\\\ $\\bullet$ {\\bf The interaction zone between the disc and the protostar}: ``X-winds'' are disc-winds produced in a tiny region around the magnetopause between the disc and the protostar (Shu et al. \\cite{shu94}, Lovelace et al. \\cite{love99}).  Pure stellar wind models are less favoured because observed jets carry far too much momentum.  In order to produce an observed jet, a protostar should be either much more luminous or rotating faster than observations show (DeCampli \\cite{dC81}, Konigl \\cite{konigl86}).  Unless otherwise proven, this leaves us with either disc-winds or X-winds.  Note that what distinguishes these last two scenarii is mainly belief: belief that no significant large scale magnetic field could be found in the innermost disc region (up to a few au).  It is however possible to observationaly discriminate between them (Shang et al. \\cite{shang02}, Garcia et al. \\cite{garcia01b}).  Nevertheless, the disc-wind scenario is usually favoured because it offers an \"universal\" paradigm able to explain jets from several astrophysical objects without relying on the central object. ",
        "conclusions": ""
    },
    "0311/astro-ph0311352_arXiv.txt": {
        "abstract": "We present two new non-parametric methods for quantifying galaxy morphology: the relative distribution of the galaxy pixel flux values (the Gini coefficient or $G$) and the second-order moment of the brightest 20\\% of the galaxy's flux ($M_{20}$). We test the robustness of $G$ and $M_{20}$ to decreasing signal-to-noise and spatial resolution, and find that both measures are reliable to within 10\\% for images with average signal-to-noise per pixel greater than 2 and resolutions better than 1000 pc and 500 pc, respectively. We have measured $G$ and $M_{20}$, as well as concentration ($C$), asymmetry ($A$), and clumpiness ($S$) in the rest-frame near-ultraviolet/optical wavelengths for 148 bright local ``normal\" Hubble type galaxies (E-Sd) galaxies, 22 dwarf irregulars, and 73 $0.05 < z < 0.25$ ultra-luminous infrared galaxies (ULIRGs). We find that most local galaxies follow a tight sequence in $G-M_{20}-C$, where early-types have high $G$ and $C$ and low $M_{20}$ and late-type spirals have lower $G$ and $C$ and higher $M_{20}$. The majority of ULIRGs lie above the normal galaxy $G-M_{20}$ sequence, due to their high $G$ and $M_{20}$ values. Their high Gini coefficients arise from very bright nuclei, while the high second-order moments are produced by multiple nuclei and bright tidal tails. All of these features are signatures of recent and on-going mergers and interactions. We also find that in combination with $A$ and $S$, $G$ is more effective than $C$ at distinguishing ULIRGs from the ``normal\" Hubble-types. Finally, we measure the morphologies of 49 $1.7 < z < 3.8$ galaxies from HST NICMOS observations of the Hubble Deep Field North.  We find that many of the $z \\sim 2$ galaxies possess $G$ and $A$ higher than expected from degraded images of local elliptical and spiral galaxies, and have morphologies more like low-redshift ULIRGs. ",
        "introduction": "The evolution of the physical structure of galaxies is one of the keys to understanding  how matter in the universe assembled into the structures we see today.  The most accessible tracer of a galaxy's physical structure is its morphology, i.e. the organization of its light (stars and dust), as projected into our line of sight and observed at a particular wavelength. As we examine more distant galaxies, we find that galaxy morphologies become increasingly chaotic. The disk and spheroidal structures abundant in the local universe disappear at early times in the universe (e.g. Abraham et al. 1996, Abraham \\& van den Bergh 2002). The emergence of the local Hubble sequence of spiral and elliptical galaxies at late times is one of the predictions of the hierarchical picture of galaxy assembly.  While the first morphological studies sought to describe the variety of galaxy shapes and forms, the goal of present-day morphological studies is to tie the spatial distribution of stars to the formation history of the galaxy. A major obstacle to this goal has been the difficulty in quantifying morphology with a few simple, reliable measurements. One tack is to describe a galaxy parametrically, by modeling the distribution of light as projected into the plane of the sky with a prescribed analytic function. For example, bulge-to-disk light ratios may be computed by fitting the galaxy with a two-component profile, where the fluxes, sizes, concentrations, and orientations of the bulge and disk components are free parameters (Peng et al. 2002, Simard et al. 2002). This B/D ratio correlates with qualitative Hubble type classifications, although with significant scatter.  Unfortunately, there is often a fair amount of degeneracy in the best-fitting models and B/D ratios, and structures such as compact nuclei, bars, and spiral arms introduce additional difficulty in fitting the bulge and disk components (e.g. Balcells et al 2003). A related approach is to fit a single Sersic profile to the entire galaxy (Blanton et al. 2003a). Profiles with high Sersic indices are interpreted as bulge-dominated systems, while low Sersic indices indicate disk-dominated systems.  However, not all bulges have high Sersic index values - some are exponential in nature (Carollo 1999), so not all objects with bulges will produce intermediate or high Sersic indices. Both the one-component and multiple-component fitting methods assume that the galaxy is well described by a smooth, symmetric profile - an assumption that breaks down for irregular, tidally disturbed, and merging galaxies.  Non-parametric measures of galaxy morphology do not assume a particular analytic function for the galaxy's light distribution, and therefore may be applied to irregulars as well as standard Hubble type galaxies.  Abraham et al. (1994, 1996) introduced the concentration index $C$ (which roughly correlates with a galaxy's B/D ratio) and Schade et al. (1995) put forward rotational asymmetry $A$ as a way to automatically distinguish early Hubble types (E/S0/Sa) from later Hubble types (Sb/Sc) and classify irregular and merging galaxies. Subsequent authors modified the original definitions to make $C$ and $A$ more robust to surface-brightness selection and centering errors (Wu 1999, Bershady et al. 2000, Conselice et al. 2000).   The third quantity in the ``$CAS$'' morphological classification system is a measure of a galaxy's residual clumpiness $S$, which is correlated with a galaxy's color and star-formation rate (Isserstedt \\& Schindler 1986; Takamiya 1999; Conselice 2003). Other more computer-intensive approaches to galaxy classification such as artificial neural networks and shapelet decomposition have also been applied to local and distant galaxies.  Artificial neural networks are trained by an astronomer on a set of galaxies of known morphological type and use a combination of size, surface-brightness, concentration, and color to classify galaxy types (Odewahn et al. 1996, Naim et al. 1997).   ``Shaplets''  deconstruct each galaxy's image into a series of Hermite polynomials (Refregier 2003,  Kelly \\& McKay 2004). The eigen-shapes produced by shapelet decomposition are often difficult to interpret by themselves, and the additional step of principle component analysis is performed to classify galaxies.  While $CAS$ is perhaps the most straightforward of the non-parametric methods, it is not without its weaknesses.  Because concentration is measured within several circular apertures about a pre-defined center, it implicitly assumes circular symmetry, making it a poor descriptor for irregular galaxies.  Asymmetry is more sensitive to merger signatures than concentration, but not all merger remnant candidates are highly asymmetric, and not all asymmetric galaxies  are mergers (e.g. dusty edge-on spirals).  Finally, the clumpiness determination requires one to define a galaxy smoothing length, which must be chosen carefully to avoid systematic effects dependent on a galaxy image's point spread function (PSF), pixel scale, distance, and angular size.  Also, the bulges of highly concentrated galaxies give strong residuals which are not due to star-forming regions and must be masked out when computing $S$.  In this paper, we examine two new non-parametric ways of quantifying galaxy morphology which circumvent some of the problems with the ``$CAS$'' system.  We use the Gini coefficient, a statistic used in economics to describe the distribution of wealth within a society. It was first adapted for galaxy morphology classification by Abraham et al. (2003) to quantify the relative distribution of flux within the pixels associated with a galaxy.  It is correlated with concentration, but does not assume that the brightest pixels are in the geometric center of the galaxy image. We also define a new indicator, $M_{20}$, which describes the second-order moment of the brightest 20\\% of the galaxy.  While similar to the concentration index, $M_{20}$ is more sensitive to merger signatures like multiple nuclei and does not impose circular symmetry.  In \\S 2, we modify Abraham's definition of the Gini coefficient in order to make it applicable to distant galaxies and we define $M_{20}$. In \\S 3, we test the robustness of these statistics to decreasing $S/N$ and resolution, and find that at average $S/N$ per galaxy pixel $> 2$ and spatial resolutions $<$ 500 pc, they are reliable to within 10\\%.  We also compare the robustness of $G$ and $M_{20}$ to $CAS$. In \\S 4, we compare the ability of $G$ and $M_{20}$ to classify local Hubble-type and merging galaxies to the $CAS$ system. Finally in \\S 5, we examine the near-ultraviolet/optical morphologies of 49 $ 1.7 < z < 3.8$ Lyman break galaxies and attempt to classify these LBGs as ellipticals, spirals, or merger candidates. ",
        "conclusions": "We have re-defined the Gini coefficient in Equation 6 as a statistic for measuring the distribution of flux values within a galaxy's image, and introduced $M_{20}$ (Equation 8), the second-order moment of the brightest 20\\% of the galaxy's flux.  These two indices are complementary, non-parametric morphology measures.  We have tested robustness of $G$ and $M_{20}$ to decreasing $S/N$ and resolution and found them to change by less than 10\\% for average $S/N$ per pixel $\\geq$ 2 and resolutions better than 500 pc. At worse resolutions, $C$, $A$, and $M_{20}$ have systematic biases which are a function of Hubble type, while $S$ becomes unreliable. $G$, on the other hand, appears to be remarkably stable at low resolutions and therefore is a powerful tool for classifying the morphologies of high-redshift galaxies.  We have measured $C$, $A$, $S$, $G$, and $M_{20}$ from the near-UV/optical images of 170 local E-S0-Sa-Sb-Sc-Sd-dI galaxies, 73 $z\\sim 0.2$ ULIRGs, and 49 $1.7 < z < 3.8$ Lyman break galaxies.  We find that:  1)Normal Hubble-type galaxies follow a tight $G-M_{20}-C$ sequence.  Early-type and bulge-dominated systems have high Gini coefficients and concentrations and low second-order moments as a result of their bright and compact bulges. Shallower surface brightness profiles, spiral arms, and off-center star clusters give late-type disks lower Gini coefficients and concentrations and higher second-order moments.  2) In combination with $A$ and $S$, $G$ is more effective than $C$ at distinguishing ULIRGs from normal Hubble types. We also find that most ULIRGs lie above the $G-M_{20}$ sequence and can be identified by by their higher $G$ and $M_{20}$ values. The high Gini coefficients arise from very bright compact nuclei, while multiple nuclei and bright tidal features produce large second-order moments.  3) ULIRGs with double and multiple nuclei have a statistically different distribution in morphology space than single nuclei ULIRGs.  ULIRGs with double/multiple nuclei typically have higher second-order moments and asymmetries and slightly lower concentrations than single nuclei ULIRGs.  Singly-nucleated ULIRGs are more likely to possess low asymmetries and low second-order moments, and often have higher concentrations and Gini coefficients than ULIRGs in well-separated galaxy pairs.  4) Many of HDFN galaxies at $z \\sim 2$ have higher rest-frame $B$-band Gini coefficients and asymmetries than expected for local elliptical and spiral galaxies degraded to the same resolution. Instead, these objects are most similar in morphology to local ULIRGs.  Our revised Gini coefficient has proven itself to be a highly robust and unbiased non-parametric morphological indicator for $z > 2$ galaxies observed at HST NICMOS resolution, and therefore has opened a window into the morphologies and assembly of the earliest galaxies. At lower redshifts, and in combination with $M_{20}$, $A$, and $S$, the Gini coefficient allows us to more precisely classify galaxy morphologies and identify merger candidates. In our next paper, we analyze a suite of hydrodynamical galaxy merger simulations to predict the evolution of merging galaxies in G-M-C-A-S morphology space. These simulations will explore a range of merger mass ratios, orbital parameters, and star-formation feedback efficiencies, and will trace the spatial distribution of dark matter, gas, and old and new stars as a function of time (Cox et al. 2004).  We would like to thank T.J. Cox and P. Jonsson for their valuable input and careful reading of this manuscript.  We also gratefully acknowledge C. Conselice for his comments and access to his morphology analysis code, M. Dickinson for use of the NICMOS Hubble Deep Field North observations, and the anonymous referee for their useful suggestions. Support for J.L. was provided by NASA through grant number 9515  from the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA contract NAS 5-26555. J.P. acknowledges support from NSF through grant AST-0205944 and NASA through NAG5-12326. P.M. acknowledges support by NASA through grants NAG5-11513 and GO-09425.19A from the Space Telescope Science Institute.  This work is based on observations made with the NASA/ESA Hubble Space Telescope, obtained from the data archive at the Space Telescope Science Institute, and observations from the Sloan Digital Sky Survey. Funding for the Sloan Digital Sky Survey (SDSS) has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society. The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions. The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "0311/astro-ph0311487_arXiv.txt": {
        "abstract": "We present a cosmologically-motivated model in which the seeds of supermassive black holes form out  of the  lowest angular  momentum  gas in  proto-galaxies at  high redshift.   We  show that  under  reasonable  assumptions, this  leads naturally  to a  correlation between  black hole  masses  and spheroid properties,   as  observed  today.    We  assume   that  the   gas  in early-forming, rare-peak haloes has a distribution of specific angular momentum similar to  that derived for the dark  matter in cosmological N-body simulations.   This distribution has a  significant low angular momentum   tail,  which implies   that  every proto-galaxy should contain  gas that ends up in  a high-density disc. In haloes  more massive  than a critical  threshold of $\\sim  7 \\times 10^7 \\Msun$  at $z \\sim  15$, the discs are  gravitationally unstable, and  experience  an  efficient  Lin-Pringle viscosity  that  transfers angular momentum outward and allows  mass inflow.  We assume that this process continues until the first massive stars disrupt the disc.  The seed black holes created in  this manner have a characteristic mass of $\\sim 10^5  \\Msun$, roughly independent of the  redshift of formation. This serves as a lower bound for black-hole masses at galactic centers today.  The  comoving density of mass  in black hole  seeds grows with time, tracking the continuous  production of critical mass haloes, and saturates  when cosmic  reionization acts  to prevent  gas  cooling in these low-mass  systems.  By  $z \\sim 15$,  the comoving  mass density becomes comparable  to that inferred from observations,  with room for appropriate additional luminous growth during a later quasar accretion phase.  The  hierarchical merger process  after $z \\sim  15$ naturally leads  to a  linear correlation  between black-hole  mass  and stellar spheroid  mass, with  negligible black  hole masses  in disc-dominated galaxies.  The formation  of massive seeds at high  redshifts, and the relatively important role  of mergers in the buildup  of today's black holes, are key elements in the proposed scenario. ",
        "introduction": "\\label{sec:intro}  Detailed studies  of gas  and stellar kinematics  near the  centers of present-day galaxies  have revealed  that almost all  galaxy spheroids (ellipticals and  bulges of  spirals) host super-massive  black holes, with masses $\\Mbh  = (10^6 - 10^{9} )\\Msun$, or  $\\sim 10^{-4}$ of the total  stellar mass  of  their parent  galaxies (e.g.   \\citealt{KG01, FM00,M98}).   Although most of  these black  holes are  not associated with quasar activity today,  presumably they represent the now-dormant counterparts  to the quasar-powering  engines known  to exist  at high redshift \\citep{FETAL01}.   Indeed, under reasonable  assumptions, the luminosity function  of Active Galactic Nuclei (AGN)  can be explained by modeling mass  accretion onto black holes that  have the same range of masses as observed in  local spheroids (see e.g. \\citealt{SW03} and references  therein).   The  standard   expectation  is  that  once  a primordial  galaxy is  populated with  a ``seed''  black hole  at some early time, the black hole can grow via the accretion of available gas and, in  the process,  give off light  in proportion to  its Eddington luminosity.  The seed population is an  important ingredient in any model that aims to explain  AGN or  local black  holes, but the  origin and  nature of these objects remain  unknown.  In this paper, we  provide a model for producing rather  massive ($\\sim 10^{4-5} \\Msun$) seed  black holes at early times from high density, low angular momentum gas in rare halos. We  argue that  massive, early  forming  seeds help  explain the  high redshift AGN population,  and, unlike lower mass seeds,  may give rise to the observed correlations  between host galaxy properties and black hole masses without the need for fine-tuned feedback.  Consider,  for  example  the  analysis of  \\citet{BETAL03},  who  have estimated a mass of $M_{\\rm bh}  = (2-6) \\times 10^9\\Msun$ for a black hole powering a $z=6.4$ quasar  found in the Sloan Digital Sky Survey. For the \\LCDM cosmology adopted below, the age of the universe at this redshift  is  $840$  Myr.   Even   the  remnant  of  a  very  massive, early-forming star of  $\\sim 100 \\Msun$ would have  to accrete mass at the Eddington limit  for $17$ e-folding times in order  to grow to the observed  mass.  For  a standard,  thin-disk accretion  model  with an efficiency $\\sim 0.1$, the e-folding  Salpeter time is $\\sim 45$ Myrs, so   an   intermediate-mass  seed   would   require   $770$  Myrs   of accretion\\footnote{ This  timescale is much larger  than the estimated lifetimes of quasars at lower redshift (see \\citealt{M03} for a recent review).}   in order  to reach  to observed  mass, corresponding  to a formation  redshift  for the  seed  of $z_f  \\gsim  40$.   This is  an extremely early redshift in  the context of \\LCDM structure formation, since  the minimum  halo mass  for ${\\rm  H}_2$ cooling  at  this time represents  a  $\\ga  5  \\sigma$  fluctuation.   While  this  does  not completely rule out  the possibility that the seed  was small, it does require an  extremely rare event.  A  $10^5 \\Msun$ seed,  on the other hand, would require $\\sim 10$ e-folding times of accretion, and thus a more  recent formation  redshift  $z_f \\gsim  11$.   Of course  these numbers are sensitive to certain assumptions, and may be influenced by allowing  low efficiency  accretion  or super-Eddington  luminosities. Nevertheless, it  is generally easier to understand  the population of $z \\gsim 6$  quasars if massive seeds are present  (see also Haiman \\& Loeb 2001 for a related discussion).  Perhaps  the  most  intriguing  clues  to the  nature  of  black  hole formation and  growth are the  strong correlations between  black hole masses and global  properties of their host spheroids  --- scales that differ by  many orders of magnitude.   The black hole  mass is tightly correlated with  the central velocity dispersion of  the spheroid, and shows a significant correlation  with the spheroid luminosity or mass. The theoretical  challenge is to connect  the galaxy-formation process originating  on scales  of megaparsecs  with the  black-hole formation processes on  sub-parsec scales.  In  addition, the epoch  of spheroid formation is  likely much  later than the  time when the  first quasar black holes begin  to shine, so an additional  challenge is to connect these  two  components in  time,  in a  way  that  maintains, or  even creates, the observed local correlations.  Several  scenarios have been  presented for  the formation  of massive black hole seeds.   The general frameworks range from  those that rely on the coalescence of massive stellar remnants that sink to the galaxy centers \\citep{QS90, HB95, L95, EETAL01} to one that models black hole growth   via  ballistic   particle  capture   during   bulge  collapse \\citep{AGR01, AGR03}.   Other models are based  on the hydro-dynamical evolution of  super-massive objects formed directly  out of primordial gas \\citep{HR93, LR94, EL95a,  HNR98,G01, BL03}.  Our model is related most closely to this third type of scenario.  The key obstacle  in any model of black hole  formation is clearly the centrifugal  barrier.   In  the  standard  cosmological  framework  of hierarchical   structure  formation,   proto-haloes   acquire  angular momentum by tidal torques  exerted by the background fluctuation field \\citep{H49,  P69, BE87,  PDH02a,  PDH02b}.  Haloes  obtain a  specific angular momentum that  is on average $10-100$ times  smaller than that required for centrifugal support at  the virial halo size (with a spin parameter  of $\\lambda \\sim  0.01-0.1$, see  below).  This  is roughly independent of mass  and time.  Thus, even after  cooling, most of the gas can contract  only to radii 10-100 times  smaller than the initial halo radius before it is halted in circular orbits.  The implied scale (at  $z=0$) is  $\\sim$kpc, many  orders of  magnitude larger  than the corresponding Schwarzschild  radius, $\\sim 10^{-5}  (M/10^9 \\Msun)$pc. The situation may get even  worse at earlier times because the typical collapsing mass $M_\\star(a)$ is smaller  in proportion to a high power of the expansion factor $a  \\equiv (1+z)^{-1}$.  This implies that the ratio    of   the   centrifugally    supported   size    ($\\propto   a M_\\star(a)^{1/3}$) to  the Schwarzschild radius  ($\\propto M_\\star(a)$ if  the black  hole mass  is  a constant  fraction of  the halo  mass) increases dramatically at early times,  unless the sites of black hole formation are  rare haloes much more massive  than $M_\\star(a)$.  This line of  argument indicates that  the seed black holes  originate from material at  the very lowest end of  the angular-momentum distribution within rare,  massive haloes at  high redshift.  This  material should then perhaps lose all its  angular momentum via an effective mechanism of  angular momentum  transfer,  which  would be  enhanced  in such  a high-density environment.  \\citet{EL95a} considered the  sites of black hole formation  to be the very  high density discs  which form  in those  rare haloes  that have extremely  low  total  angular  momentum.   Assuming  a  standard  CDM scenario, they used  an analytic model to predict  the distribution of halo spins and estimated the  abundance of haloes with low enough spin for efficient  viscosity \\citep{EL95b}. They argued that  there may be enough low-spin haloes to explain the black hole abundance.  They then adopted the  common view that  viscosity in a  differentially rotating gaseous disc is the mechanism that transfers angular momentum outward, thus  allowing  gas   to  contract  into  a  black   hole.   In  their calculations, they considered hydrodynamic (or perhaps hydro-magnetic) turbulence  to be the  viscous mechanism  responsible for  the angular momentum  loss,   and  appealed  to   the  $\\alpha$-disc  prescription \\citep{SS73}.  Our approach  is related, but  qualitatively different.  We  propose a scenario where  seed black holes  form from material with  low angular momentum in  all haloes  that are massive  enough to host  an unstable self-gravitating disc.  The required high-density discs can only occur in  haloes that collapse  at high  redshift, $z  \\ga 10-20$,  when the angular momentum  distribution of the (unprocessed) gas  should be the most similar  to that of the  dark matter, and  before reionization or some other feedback mechanism \\citep[e.g.,][]{DS86, DW03} prevents the growth  of  cold  discs.   The  crucial  point  is  that  every  halo, regardless  of its  total  angular  momentum, is  expected  to have  a distribution of specific angular momentum including a tail of low-spin material.  The haloes we consider are  massive enough so that the baryons trapped in their potential wells can cool and condense.  In the absence of any other effect that  may alter the angular momentum  content of baryons, the  lowest  spin   baryonic  material  in  each  system   ends  in  a high-density central  region forming  a proto-galactic disc. As an illustrative fiducial model, we assume that the angular-momentum distribution in  each halo  matches the ``universal''  distribution of \\citet{B01} (B01  hereafter). If  the angular  momentum is  approximately preserved  during baryonic infall into  the disk, the  B01 angular momentum  distribution implies that the inner disk becomes  self gravitating with the surface density inversely  proportional to  the radius.   We show  below that  if this behavior extends below the inner $\\sim  2\\% $ of the disc radius, then enough  mass  may acquire  a  high  enough  density for  viscosity  to efficiently redistribute the  angular momentum of the disc. For the final contraction of  disc material into a black  hole we evaluate the angular-momentum loss by the effective kinematic viscosity that arises from  gravitational  instabilities  \\citep{LP87}.  This  mechanism  is somewhat more efficient than  that from a typical $\\alpha$-disc model. This viscosity-driven  mass inflow cannot  continue indefinitely; once the  disc  significantly fragments,  or  once  it  is heated  or  even disrupted by feedback, the inflow stops.  The semi-quantitative scenario presented here can be regarder as an illustration of a general model based on the assumption that proto-galactic gaseous discs develop high density inner regions dominated by low angular momentum material. Our current analysis is basically a feasibility study of such a model, and an attempt to recover its robust predictions and to constrain the model parameters by theoretical considerations as well as by the observations.  The outline of  this paper is as follows.   In \\S~\\ref{sec:profile} we discuss the  angular momentum  distribution in proto-galaxies  and the implied   surface   density   profile   for  the   cold   discs.    In \\S~\\ref{sec:viscosity} we investigate  the effect of kinetic viscosity on the evolution  of the disc.  In \\S~\\ref{sec:mass}  we calculate the resulting   seed   black   hole    mass   and   mass   function.    In \\S~\\ref{sec:density} we  calculate the redshift evolution  of the seed black  hole  mass  density.   In \\S~\\ref{sec:correlation}  we  make  a preliminary attempt  to address the correlation between  the masses of black  holes  and galactic  stellar  spheroids,  to  be pursued  in  a subsequent paper.   In \\S~\\ref{sec:conc} we  summarize our conclusions and  discuss our results.   When necessary,  we assume  a $\\Lambda$CDM cosmology, with a  Hubble constant $h=0.7$ ($H_0 =  100 h$ km s$^{-1}$ Mpc$^{-1}$),   universal  mass   density  $\\Omega_{\\rm   M}   =  0.3$, cosmological  constant $\\Omega_{\\Lambda} =  0.7$, and  rms fluctuation amplitude $\\sigma_8 = 0.9$. ",
        "conclusions": "\\label{sec:conc}  We presented a physical model for the production of massive seed black holes  at  high  redshifts,  with  a characteristic  mass  $\\sim  10^5 \\Msun$. We argued that this model provides a useful starting point for explaining the  observed relation between the masses  of today's black holes and  the spheroidal components  of galaxies.  Massive  seeds may also  help  explain the  existence  of  the  supermassive black  holes associated with luminous AGN at $z \\gsim 6$.   Our model  for seed formation  is based on  the general idea  that the mass   in  every  early-collapsing   halo  is   expected  to   have  a {\\textit{distribution}}  of specific angular  momentum, with  at least some  fraction of  its  gas  having a  much  smaller specific  angular momentum than  is characterized  by the global  spin parameter  of the halo.   After the low-spin  gas cools,  it can  fall into  the central region and form  a dense disc, where certain  viscous processes become effective and may potentially lead to black hole formation.  In our exploration of this  idea, we assumed that the specific angular momentum of the gas is similar to that measured for the dark-matter in N-body simulations.   In this case,  if angular momentum  is conserved during the  gas infall, the central gas  discs become self-gravitating and Toomre  unstable in  all the haloes  above a critical  mass.  This critical mass  has a median value  of $\\sim 7 \\times  10^7\\Msun$ at $z \\sim 15$,  and it is  expected to be  scattered log-normally with  a ln standard deviation  of $\\sim 0.9$.   The sites of seed  production are thus much more massive than  the typical collapsing mass at that time, and are therefore associated  with high-sigma peaks in the fluctuation distribution.   Gravitational  instabilities  in  the  disc  drive  an effective viscosity that redistributes angular momentum.  As a result, material from the low angular  momentum tail of the distribution loses its  angular momentum  on a  viscous timescale  and contribute  to the formation of a  black hole.  We associate this  viscous timescale with the  timescale  for disruption  and/or  heating  of  the disc  due  to substantial star  formation.  We  expect this timescale  to be  on the order of $1-30$ Myr, but its exact value is not well determined.  We found  that the  typical seed black  holes, which form  in critical mass haloes,  have a characteristic  mass of $\\sim 10^5  \\Msun$, quite insensitive to the epoch of formation.  This predicts a lower bound to possible masses of  massive black holes, which can  be confronted with observations today.  This scenario should operate continuously as additional haloes grow to above the critical mass and  create new potential sites for black-hole formation.   The  process  is  expected  to stop,  however,  when  the Universe becomes sufficiently reionized that ${\\rm H}_2$ molecules are destroyed and  cannot provide efficient cooling.   We therefore expect seed black holes to form before $\\zre \\sim 13-21$.  We showed that the comoving density  of mass in seed black holes grows gradually  with time, in  association with  the birthrate  of critical mass  haloes.  This density  becomes comparable  to the  observed mass density of supermassive black holes  by $z \\sim 15$, near the redshift of reionization.  Depending on the adopted viscous timescale, there is room for an order of magnitude  growth in black hole masses via quasar powering accretion.  We find that  a viscous timescale of $t\\sim 5-10$ Myr provides quite reasonable results.  We explored  the development of a correlation  between black-hole mass and spheroid  mass in today's  galaxies by tracking our  early forming massive  seed  population  through  a  halo merger  tree.   Under  the assumption that spheroids  form via major mergers, we  showed that the merger process leads naturally to a linear relation between black-hole mass and spheroid mass.  This result is mainly due to the formation of {\\textit{massive}}  seeds  only  in  rare, high-peak  haloes  at  high redshifts,  and  is not  too  sensitive to  the  details  of our  seed formation model.  Starting with $\\sim 1-100\\msun$ seeds would not lead to  a  linear  correlation  via  mergers  alone;  additional  feedback associated with AGN accretion is  needed.  The seed black holes in our model have masses that are  proportional to the virial masses of their host haloes  at formation.  We demonstrated that  this linear relation is preserved  if black  holes merge in  concert with  the hierarchical growth of haloes and spheroids.  One  potentially interesting ramification  of our  scenario is  that a population  of massive  black holes  at high  redshift might  serve to power  small  quasars,  fueled  perhaps  by  low-angular-momentum  gas liberated by  merger events that occur  after the seeds  are in place. \\citet{MAD03} discussed how a population of ``mini-quasars'' like this might be  responsible for  reionizing the universe  at $\\zre  \\sim 15$. However,   in  their   model,  the   mini-quasar   power-sources  were intermediate-mass $\\sim 100 \\Msun$ black holes. Under similar assumptions,  the massive seed black holes  of our model would produce  even more ionizing  photons per site, and  perhaps lead more naturally to an early  reionization.  In our current scenario the reionization epoch was an input parameter detached from the black-hole formation  process, but, in  principle, the  accretion onto  the seeds themselves could  be responsible for the  reionization that eventually inhibits further black-hole formation.  Our quantitative  predictions should not be  interpreted too literally. For   example,   our  fiducial   model   assumes   that  the   initial angular-momentum distribution of the gas mirrors that of a typical CDM halo, and  that it  is preserved during  disk formation.  There  is no evidence in local  galaxies for discs with inner  densities as high as invoked by our  model for proto-galaxies.  This might  be explained by processes  that  have altered  the  gas angular-momentum  distribution preferentially  at  late  epochs,   such  as  feedback  effects  (e.g. \\citealt{MD02}), but one should clearly assign a substantial degree of uncertainty  to  our  model  assumption.   In order  to  address  this uncertainty, at least qualitatively, we explored the effect of varying the initial angular  momentum distribution.  We found that  if the low angular  momentum tail  is as  minor as  in an  exponential disc  of a constant-density core  ($\\xi =2$), the resulting seed  black holes are too small to be of relevance.  An intermediate case, with $\\xi = 1.3$, can produce sensible black-hole seeds, comparable to our fiducial case ($\\xi  =  1$), provided  that  the  viscosity  timescale is  slightly longer, $20 \\le t \\le 30$  Myr.  For $\\xi=1.3$, the critical halo mass becomes $\\sim 10^8 \\Msun$ at $z=15$,  and the inner $\\sim 10\\%$ of the disc  mass manages  to  lose  its angular  momentum  by viscosity  and produce a  black hole of  mass $\\sim 10^6  \\Msun$.  This seed  mass is larger than in the fiducial case, but the seeds in this case are rarer events (see figure \\ref{fig:fig4}).   Therefore, in order to recover a similar total mass density in  black holes, the viscosity must operate for a  longer duration.  The  actual low angular-momentum tail  of the gas  distribution   remains  an  unknown,  which   will  hopefully  be constrained  by  future,  more realistic  hydrodynamical  simulations, properly including feedback and other important physical effects.  In conclusion, while the model presented here is clearly idealized, it does  provide a useful  example for  the kind  of mechanism  needed to connect  cosmological  structure   formation  on  galactic  scales  to black-hole formation on much smaller scales.  Observations suggest the presence of  massive black holes  at high redshifts and  therefore the need for  the formation of massive  seed black holes  at still earlier times.  According to our  scenario, the observed correlations of black hole  and galaxy properties  are directly  linked to  the hierarchical nature of structure formation itself,  and owe their existence to rare systems forming in association with the first light."
    },
    "0311/astro-ph0311164_arXiv.txt": {
        "abstract": "{ Numerical hydrodynamical modelling of supernova-driven shell formation is performed with a purpose to reproduce a giant HI ring (diameter 1.7~kpc) in the dwarf irregular galaxy Holmberg~I (Ho~I). We find that the contrast in HI surface density between the central HI depression and the ring is sensitive to the shape of the gravitational potential. This circumstance can be used to constrain the total mass (including the dark matter halo) of nearly face-on dwarf irregulars. We consider two models of Ho~I, which differ by an assumed mass of the dark matter halo $M_{\\rm h}$. The contrast in HI surface density between the central HI depression and the ring, as well as the lack of gas expansion in the central hole, are better reproduced by the model with a massive halo of $M_{\\rm h}=6.0\\times 10^9~M_{\\odot}$ than by that with a small halo of $M_{\\rm h}=4.0\\times 10^8~M_{\\odot}$, implying that Ho~I is halo-dominated. Assuming the halo mass of $6.0\\times 10^9~M_{\\odot}$, we determine the mechanical energy required to form the observed ring equal to $(3.0 \\pm 0.5) \\times 10^{53}$~ergs, equivalent $300\\pm 50$ Type~II supernovae. The inclination of Ho~I is constrained to $15^\\circ-20^\\circ$ by comparing the modelled HI spectrum and channel maps with those observed. ",
        "introduction": "The interstellar medium of gas-rich dwarf irregular galaxies (dIrr) is usually dominated by features described as shells, holes, or rings. The number of such structures varies from a few (IC~10, Wilcots \\& Miller \\cite{Wilcots}) to several tens (Holmberg II, Puche et al. \\cite{Puche}) and their sizes are within several hundred parsecs. However, there are a few dIrr's the HI morphology of which is totally dominated by a single ring structure of the size comparable to their optical extent (Ott \\cite{Ott2}). Among them is Holmberg~I, a member of the M81 group of galaxies (distance $\\sim$~3.6~Mpc). Ott et al. (\\cite{Ott}) have found that most of its HI content (75\\%) is localized within a giant ring of 1.7~kpc diameter. They suggested that strong stellar winds and supernova (SN) explosions might be responsible for the peculiar HI morphology in Ho~I.  In the present paper we test this hypothesis via detailed numerical hydrodynamical modelling. Two model galaxies are considered: a fast rotating galaxy with a massive halo and a slowly rotating one with a small halo. We attempt to determine the amount of SNe required to reproduce a ring-like morphology of the distribution of the atomic hydrogen in Ho~I. We show that modelling of the HI spectrum, HI channel maps, and the radial HI profile may constrain the inclination angle and consequently the dark matter content of Ho~I.  In Sect.~\\ref{origin} the origin of a ring-like HI morphology in Ho~I is discussed. In Sect.~\\ref{model} the numerical hydrodynamical model for simulating the supernova-driven shell dynamics is formulated. The results of simulations are presented in Sect.~\\ref{results} and Sect.~\\ref{discus}. Our conclusions are summarized in Sect.~\\ref{concl}. ",
        "conclusions": "\\label{discus} Ho~I has a ring-like HI distribution, which is moderately non-axisymmetric with respect to its morphological center (see Fig.~\\ref{Fig1}). The ellipticity of the ring makes it rather difficult to draw any quantitative comparison with the observations on the basis of two-dimensional axisymmetric simulations. However, the basic qualitative features of the gas distribution and velocity structure in Ho~I, such as the size of the ring, the contrast in the gas surface density (hereafter, $\\Sigma_{\\rm g}$) between the central depression and the ring, and the lack of gas expansion in the central hole can be modelled in the two-dimensional case.  The lack of gas expansion in the central depression  reported by Ott et al. (\\cite{Ott}) argues that the shell has suffered a blowout. Otherwise, one would expect the shell to be still in the Sedov expansion phase with some piled-up gas in the z-direction. However, we do not see any HI along the line of sight to the central hole down to a column density of $\\sim 6.0 \\times 10^{19}$~cm$^{-2}$. The measured contrast in $\\Sigma_{\\rm g}$ between the central depression and the ring (a factor of ~ 15) is also indicative of a blowout, since the `limb brightening' effect may account for a factor of 2 difference in $\\Sigma_{\\rm g}$ between the projected rim of a three-dimensional shell and its center. Hence, the occurrence of a blowout (implied by the measured contrast in $\\Sigma_{\\rm g}$ between the central depression and the ring) at the time when the shell has expanded out to its present size of 1.7~kpc is the main criterion for the selection between the two models.  As discussed in Sect.~\\ref{model1} and \\ref{model2}, the blowout scenario of Ho~I is better reproduced in model~1 at $t=55 \\pm 5$~Myr. At a first glance, the radial gas distribution in model~2 at 73 Myr seems by eye to better fit the overall observational data. However, it is not only the measured size of the ring (1.7 kpc), but also the measured contrast in $\\Sigma_{\\rm g}$ between the central depression and the ring (a factor of ~ 15) that are not reproduced for this case. We would like to note here that the estimated decrease in HI column density near the central depression caused by the external UV background is approximately $2\\%$ at $t>45$~Myr in model~2. Hence, the external UV radiation field cannot reconcile the measured contrast of ~ 15. Although the overall radial gas distribution in model 2 at 73 Myr fits better the measured profile, it fails to reproduce both the size of the ring and the measured contrast in the gas surface density between the central depression and the ring.  On the other hand, our best model is not optimized to fit the observations. At the time when the ring has attained its present size of ~1.7 kpc, the maximum gas surface density exceeds that of Ho~I by approximately $5~M_{\\odot}$~pc$^{-2}$ (see Fig.\\ref{Fig8}, the solid line). At the same time, the width of the modelled ring is smaller than that of the observed ring. This is most probably due to the two-dimensional axisymmetric nature of our simulations attempting to reproduce a moderately non-axisymmetric structure, or azimuthal averaging of the inclined ring with vertical walls. Three-dimensional non-axisymmetric modelling of Ho~I is required for a fine tuning of the model radial gas distribution.  There are a few assumptions inherent to the model that need further justification: \\\\ {\\it The assumed cooling function}.  The cooling function below $10^{4}$~K depends on the degree of ionization of the gas. We made a few test runs with the cooling function of Spaans \\& Norman (\\cite{Spaans}) truncated at the lower temperatures, i.e. the cooling was set to zero at $T<10^4$~K, and found that the influence of the cooling at $T<10^4$~K on our results was minimal, which is due to the fact that thermal energy of the shell is mostly lost at higher temperatures as mentioned earlier.\\\\ {\\it Injection of energy by supernova explosions}. In the present simulations we use a constant wind approximation (Mac Low \\& Ferrara \\cite{MF}). However, our simulations have shown that the injection of thermal energy due to discrete SN explosions produce similar results. Indeed, the frequency of SN explosions in an instantaneously born stellar cluster remains nearly constant, if a Salpeter initial mass function and an upper stellar mass of $100~M_{\\odot}$ are assumed. This finding was also confirmed in numerical simulations of Mac Low \\& McCray (\\cite{MM}). \\\\ {\\it Star formation history}. In the present simulations we have assumed a single starburst localized in the central region of Ho I that provides the energy for 30~Myr. Constant star formation in Ho~I can be ruled out because of the low total $H\\alpha$ luminosity of $4.3 \\times 10^{38}$ ergs s$^{-1}$ (Miller \\& Hodge \\cite{Miller}), which implies a SFR of only $0.004~M_{\\odot}$~yr$^{-1}$. Indeed, assuming the constant SFR of $0.004~M_{\\odot}~$yr$^{-1}$ and an age of Ho~I of 10~Gyr, one obtains an estimate on the total stellar mass of $4 \\times 10^7 M_{\\odot}$, which is 2.5 times lower than that actually measured for Ho~I (Ott et al. \\cite{Ott}). Variable star formation is plausible. However, in dIrr's variations have a burst-like nature (see e.g. Searle et al. \\cite{Searle}). Our model corresponds to an explosive energy input following evolution of a single OB association with a total mass of $\\sim 3 \\times 10^4~M_{\\odot}$. A comparison of the photometry presented in Ott et al. (\\cite{Ott}) with the STARBURST99 synthesis models (Leitherer et al. \\cite{Star}) shows that the B-band magnitude as well as the U$-$B color of the optical emission within the giant HI hole is in agreement with a stellar cluster 55~Myr of age (predicted by model~1) and a mass of $\\sim 2\\times 10^{5} M_{\\odot}$ (Salpeter IMF, Z=0.004 metallicity, upper and lower mass cutoff: $100~M_{\\odot}$ and $1~M_{\\odot}$, respectively). Such a stellar population can easily provide the mechanical energy input to Ho~I needed to satisfy model~1. The discrepancy of the observed stellar mass with our simulations (about 1 order of magnitude) is most likely due to confusion with older stellar populations.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{H4449F11.ps}} \\caption{Azimuthally averaged radial gas distribution obtained for a time-varying mechanical luminosity with {\\bf a)} model~1 and {\\bf b)} model~2 at three different phases of the shell expansion. The {\\it measured} radial gas distribution of Ho~I is plotted with the filled squares. } \\label{Fig11} \\end{figure}  Nevertheless, we examine if variable star formation can affect our conclusions. We assume that the mechanical luminosity produced by SN explosions is time-dependent, for instance with a Gaussian shape. Its maximum is centered at 15~Myr after the beginning of the energy input phase, with a FWHM of 8.5~Myr. This time-dependent mechanical luminosity is supposed to mimic the energy input from successively born stellar clusters within a single stellar association. To be consistent with the simulations in Sect.~\\ref{results}, we set the energy input phase to last for 30~Myr and the total energy input equivalent to 300 and 150 SNe for model~1 and model~2, respectively. Figure~\\ref{Fig11} shows the azimuthally averaged gas distributions obtained at the same phases of the shell expansion as in Figs.~\\ref{Fig8} and \\ref{Fig10}. As is obviously seen, model~1 with variable star formation reproduces the observed distribution and the contrast in $\\Sigma_{\\rm g}$ between the central depression and the ring (the solid line, Fig.~\\ref{Fig11}a) even better than the same model with a single starburst approximation (the solid line, Fig.~\\ref{Fig8}). Model~2 clearly fails to reproduce the measured contrast in $\\Sigma_{\\rm g}$ for both star formation frameworks. Our simulations also show that multiple, successively born smaller clusters are less effective in creating a bubble structure than a single starburst with the same total ejected energy. \\\\ {\\it Numerical resolution}. We made a few test runs with a higher resolution of 5 pc. Although the higher resolution simulations provide more details on the gas flow dynamics and the development of Rayleigh-Taylor and Kelvin-Helmholtz instabilities (Figs.\\ref{Fig4} and \\ref{Fig9}), the integrated images such as the HI spectrum, the channel maps, and the radial gas distributions (Figs.~\\ref{Fig8}, \\ref{Fig10}, and \\ref{Fig11}) remain very similar to those derived with a lower resolution of 10 pc."
    },
    "0311/astro-ph0311508_arXiv.txt": {
        "abstract": "This note addresses possible applications of the Tikhonov regularization to image reconstruction of gravitational lens systems. Several modifications of the regularization algorithm are discussed. Our illustrative example is the close quadruple gravitational lens QSO 2237+0305 (Einstein Cross). The restored image of the lens is decomposed into two parts -- the quasar components and the background galaxy. ",
        "introduction": "Over the last decade gravitational lensing has become a powerful tool to probe the Universe. The number of lenses discovered so far is just below one hundred and is growing quickly. Gravitational lensing comes in a variety of shapes and sizes. One of the most interesting examples of this phenomenon is systems of close multiple quasar images. For a typical lens galaxy the image separation is of order of an arc second which happens to be the resolution of some of the best ground-based telscopes. Hence, what the observer will see is the source copies blurred and overlapped and possibly superimposed on the lens galaxy. To use the \"nature telescope\" one needs to restore the light distribution preferably splitting it into parts belonging to the \"background\" and to the quasar images. The problem addressed in this note is the problem of the correct reconstruction of complex gravitational lens images. ",
        "conclusions": "The application of the Tikhonov regularization to the image reconstruction of gravitational lens systems seems to be an effective way of tackling some of the common problems shared by conventional deconvolution algorithms. We have considered several modifications of the regularization method. The illustrative example of the Einstein Cross lens system was used to compare different versions of the algorithm. Although the results look quite encouraging the further work needs to be done to select the appropriate stabilizer function and to produce the pipe-line version of the software."
    },
    "0311/astro-ph0311022_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1} In the very earliest stage of star formation, the young stellar object is surrounded by a collapsing envelope containing mainly dust and a large variety of molecules. Recently modeling of the dust continuum emission from the envelopes around high-mass (van der Tak et al. $1999$, $2000$) and low-mass protostars (J{\\o}rgensen et al. $2002$) has been used to establish their physical parameters and as input for radiative transfer modeling in order to establish their chemical structure. The chemistry in these envelopes is complex and has several chemical characteristics with different diagnostics (van Dishoeck $2002$). Although molecular clouds exposed to external UV and X-rays have been studied for decades, very little is known about the influence of central UV radiation and X-rays on the chemical evolution of protostellar envelopes.  In recent years it has become clear that YSOs are strong X-ray emitters. The observed X-ray luminosities range from approximately $10^{28.5}$ to $10^{33}$\\,erg\\,s$^{-1}$ (Feigelson \\& Montmerle $1999$, Preibisch $1998$). The emission may originate from accretion, from magnetic activity or, in massive stars, from stellar wind instabilities. Since the emission starts in the deeply embedded phase, its onset and initial intensity cannot be observed directly due to very high extinction. It is therefore important to understand the influence of high-energy radiation on the chemistry in the envelopes around young stars in order to use molecules as tracers of X-rays and UV radiation. For this purpose we use detailed thermal balance and gas-phase chemical models to study the chemical evolution. In addition, we have started a JCMT program to search for chemical tracers such as CN, CO$^{+}$ and NO in the envelopes of both high- and low-mass stars, as well as intermediate mass YSOs. ",
        "conclusions": ""
    },
    "0311/gr-qc0311065_arXiv.txt": {
        "abstract": "It is well known that radiative corrections evaluated in nontrivial backgrounds lead to effective dispersion relations which are not Lorentz invariant. Since gravitational interactions increase with energy, gravity-induced radiative corrections could be relevant for the trans-Planckian problem. As a first step to explore this possibility, we compute the one-loop radiative corrections to the self-energy of a scalar particle propagating in a thermal bath of gravitons in Minkowski spacetime. We obtain terms which originate from the thermal bath and which indeed break the Lorentz invariance that possessed the propagator in the vacuum. Rather unexpectedly, however, the terms which break Lorentz invariance vanish in the high three-momentum limit. We also found that the imaginary part, which gives the rate of approach to thermal equilibrium, vanishes at one loop. ",
        "introduction": "In recent years modified dispersion relations which break Lorentz invariance have appeared in different contexts of gravitational physics.  On the one hand, they appeared in several works which address the so-called \\emph{trans-Planckian problem} \\cite{Jacobson91,Jacobson93,Jacobson99}. In black hole physics the modes responsible for Hawking radiation reach arbitrarily high energies near the black hole horizon when measured by a free-falling observer. This observation has lead some authors to study the robustness of Hawking radiation when modifying (from the outset) the dispersion relation beyond the Planck scale \\cite{Unruh95,BroutEtAl95,CorleyJacobson96,Helfer03}. A similar problem appears in inflationary cosmology: the modes responsible for the large scale structure had length scales much smaller than the Planck length in the early stages of inflation. Therefore it is of interest to determine to what extent the properties of the fluctuation spectrum are sensitive to modifications of the dispersion relation at the Planck scale \\cite{MartinBrandenberger00,Niemeyer00,NiemeyerParentani01,MartinBrandenberger03}. In these works, the introduction of a nontrivial dispersion relation has been originally suggested by formal analogies with condensed matter \\cite{Unruh81,Unruh95,Visser98,GarayEtAl99,VisserEtAl01}. However, the main physical motivation is that unknown effects of quantum gravity (such as radiative corrections) might introduce nontrivial dispersion relations in these contexts. Following 't Hooft's \\cite{tHooft96} observation that strong gravitational interactions in the near horizon region might alter the semiclassical description of black hole evaporation, a first dynamical realisation of this line of sought has been pursued in Refs.~\\cite{Parentani01,Parentani02}.  On the other hand, there have been hints from high energy astroparticle physics that we could already see the effects of a violation of the Lorentz symmetry. In particular modified dispersion relations have been used in order to provide an explanation for observations of ultrahigh energetic cosmic rays (with energies higher than $10^{20}\\,\\mathrm{eV}$) beyond the Greisen-Zatsepin-Kuzmin (GZK) cutoff \\cite{TakedaEtAl98,Kifune99,Amelino-CameliaPiran01a,Amelino-CameliaPiran01b}. As in the trans-Planckian problem, several dispersion relations have been exploited, inspired by various approaches to quantum gravity \\cite{AlfaroEtAl99,AlfaroEtAl02}, string theory \\cite{KosteleckySamuel89}, non-conmutative field theory \\cite{CarrollEtAl01}, variation of couplings \\cite{KosteleckyEtAl02} and multiverses \\cite{Bjorken03}. Modified dispersion relations often break the energy degeneracy for a given three-momentum: the particle energy can become helicity dependent \\cite{AlfaroEtAl99,AlfaroEtAl02} and light can become birefringent \\cite{CarrollEtAl90,KosteleckyMatthew01}. A popular form for these modified dispersion relations below the Planck scale is, in the cosmological rest frame, \\begin{equation} \\label{Eq1} E^2 = m^2 + |\\vect p|^2 + \\sum_{n \\ge 3} \\eta_n \\frac{|\\vect p|^n}{\\Mp^{n-2}}, \\end{equation} where $\\Mp$ is Planck's mass and $\\eta_n$ are coefficients of order 1 if the Planck mass is the relevant scale. Jacobson \\emph{et al.}\\ have studied the constraints on the possible values of the parameters $\\eta_n$ based on current astrophysical observational data \\cite{JacobsonEtAl02,JacobsonEtAl03a,JacobsonEtAl03b,JacobsonEtAl03c}. It is important to notice that both the real and imaginary parts $\\eta_3$ are already observationally constrained to be much smaller than 1. This shows that the ``natural'' assumption that the $\\eta_n$ should be of order 1 is perhaps too naive. The possibility of quadratic modifications of the dispersion relation is also strongly constrained \\cite{CarrollEtAl90,KosteleckyMatthew01}.  As we learned from particle physics, Lorentz invariance is a key element in renormalized quantum field theory \\cite{Hatfield,WeinbergQFT}. Nevertheless, it might be broken by quantum gravity at a fundamental level, or, alternatively, it may be broken in an effective way in nontrivial backgrounds only. Indeed in backgrounds which possess a preferred reference frame, radiative corrections to self-energies might contain terms which effectively break the (local) Lorentz invariance ---\\ie, the Lorentz invariance in the tangent plane. Let us emphasize that this second possibility does not imply any kind of fundamental breaking of the Lorentz symmetry or any new physics. Rather by ``effective breaking of the Lorentz invariance'' we mean that radiative corrections to the dispersion relation may contain terms which depend on vector or tensor fields characterizing the background.  In this paper we explore a model in the framework of this second alternative. We consider the propagation of a scalar particle immersed in a thermal bath of gravitons. Our aim is to determine how the inertia of the thermal bath affects the propagation of the scalar particle through radiative corrections. To this end, we shall compute the thermal corrections to its self-energy. The real part of the self-energy gives the thermal mass shift and might introduce as well a nontrivial dispersion relation \\cite{DonoghueEtAl85}. The imaginary part instead gives the rate of approach to thermal equilibrium \\cite{Weldon83}.  In fact it is well known that quantum effects evaluated in nontrivial backgrounds induce non Lorentz invariant terms in the dispersion relation. For instance, the self-energy of a charged fermion immersed in a thermal bath of photons has been computed in Refs.~\\cite{Weldon82,DonoghueHolstein83,DonoghueEtAl85,Weldon99}, and the effects of the thermal bath on the speed of light have been studied in Refs.~\\cite{Tarrach83,Barton90,LatorreEtAl95}. Let us mention also that electromagnetic interaction in curved spacetimes also leads to modifications of the dispersion relations. The effect of QED vacuum polarization on the speed of light in nontrivial spacetimes has been studied in Refs.~\\cite{DrummondHathrell80,Shore02,Shore02b,Shore03,Shore03b}. However, most calculations done so far do not concern gravitational interactions. Gravity-induced corrections should be dominant at energies approaching the Planck scale, hence relevant for the situations concerned with the trans-Planckian problem. Additionally notice that gravitational interactions are universal and not limited to charged particles.  Two important points concerning our approach should be mentioned. On the one hand, since gravitational interactions are non-renormalizable, the system of the scalar field coupled to gravity should be conceived as an effective field theory \\cite{Weinberg79,WeinbergQFT}. Although the full theory is non-renormalizable, low-energy predictions which do not depend on the Planck scale behavior of gravity can be extracted \\cite{Donoghue94a,Donoghue94b,Donoghue95}. On the other hand, since thermal corrections do not affect the ultraviolet properties of the theory, thermal field theory fits well with the effective theory approach. At a technical level, we shall incorporate the thermal effects through the real-time description of thermal field theory \\cite{LandsmanWeert87,DasThermal,LeBellac}, which can be seen as a particular application of the closed-time path (CTP) method in field theory \\cite{Schwinger61,Keldysh65,ChouEtAl85}.    We should also emphasize that we shall compute only the lowest order corrections. Hence, we deal with free gravitons; \\ie, we neglect the backreaction of the scalar field on the metric perturbations. The modifications to the metric propagator can be computed both in the framework of stochastic gravity \\cite{CalzettaHu94,HuMatacz95,HuSinha95,CamposVerdaguer96, CalzettaEtAl97,MartinVerdaguer99a,MartinVerdaguer99c,MartinVerdaguer00,HuVerdaguer03, BarrabesFrolovParentani00} or in the large-$N$ limit of quantum gravity \\cite{Tomboulis77} (both approaches can be seen to be equivalent \\cite{RouraThesis,RouraVerdaguerInPrep}). A computation along these lines (but in the semiclassical approximation) has been carried out by Borgman and Ford \\cite{BorgmanFord03}. They study the fluctuations in the focusing of a bundle of geodesics, propagating in a spacetime with metric perturbations induced by a thermal scalar field. To our knowledge, the higher order corrections to the self-energy have not yet been computed in the quantum field theoretical framework.  Gravity-mediated modifications of the dispersion relation have been also considered by Burgess \\emph{et al.}~\\cite{BurgessEtAl03}. Working in the context of brane-world scenarios, these authors compute the change to the dispersion relation of photons and fermions generated by interaction with four-dimensional effective gravitons. In that case the primary source of violation of the Lorentz symmetry is the modification of the dispersion relation of effective four-dimensional gravitons induced by certain extra-dimensional configurations. Similarly to our case, this modification is communicated to brane-bound particles through gravitational interactions.  The plan  of the paper is the following. In Sec.~II we introduce the action for the system considered ---namely, a scalar field coupled to gravity. In Sec.~III we compute perturbatively the leading correction to the self-energy of the scalar field at zero temperature. In Sec.~IV we assume that the scalar field and the gravitons are in thermal equilibrium at a given temperature, and compute the corrections to the retarded self-energy. Finally in Sec.~V we analyze the results and discuss the physical significance of the terms which break Lorenz invariance. The Appendixes contain reference material as well as some technical details of the calculation. In Appendix \\ref{app:FeyRules} we present the Feynman rules for the scalar field coupled to linearized gravity. In Appendix \\ref{app:Int} we give the results of the calculation of some integrals in dimensional regularization. In Appendix \\ref{app:TFT} we give a brief account of the real-time approach to thermal field theory. Appendix \\ref{app:ABCD} is mostly technical and is devoted to the calculation of certain integrals that appear in the calculation of the thermal contribution to the self-energy.  Throughout this paper we shall use a system of units with $\\hbar=c=k_\\mathrm B=1$. The signature of the metric will be $(-,+,+,+)$. ",
        "conclusions": "At zero temperature, the position of the pole of the propagator gives the energy of the state and hence defines the dispersion relation, according to the K\\\"allen-Lehmann specral representation \\cite{WeinbergQFT,Hatfield,Peskin}. The position of the pole is found to be \\begin{equation} \\begin{split} (p^0)^2  &= m^2 +|\\vect p|^2 + \\Sigma\\TO(-m^2) = m^2 + |\\vect p|^2, \\end{split} \\end{equation} where the second equality is a consequence of the renormalization condition \\eqref{RenorCond}. The dispersion relation is clearly Lorentz invariant, as expected.   Similarly, at finite temperature it can be shown \\cite{DonoghueEtAl85} that the ``effective'' dispersion relation of the particle, given by the location of the poles of the retarded propagator, determines the inertial properties of the particle. The location of the poles is given by \\begin{equation} \\label{DispRel} (p^0)^2 - |\\vect p|^2 = m^2 + \\Re \\SigmaR(p^0,\\vect p). \\end{equation} The thermal mass is obtained by setting $\\vect p=\\vect 0$: \\begin{equation} \\mth^2 = m^2 + \\Re \\SigmaR(\\mth,\\vect0). \\end{equation} In a Lorentz invariant situation one would simply have $(p^0)^2 = \\mth^2+ |\\vect p|^2$, but in general there can be additional dependence on the three-momentum $\\vect{p}$ on the right-hand side, \\begin{equation} \\label{RelDispTh} (p^0)^2 = \\mth^2+|\\vect p|^2 + \\mathcal F(\\kappa,T,\\mth,\\vect p). \\end{equation} The Lorentz-breaking additional term in the dispersion relation leads to modifications of the group velocity of the particles $\\vect v = \\mathrm dp^0/\\mathrm d\\vect p$: \\begin{equation}\\label{GroupVelocity} \\vect v = \\frac{\\vect p}{p^0} + \\frac{1}{2p^0}\\derp{\\mathcal F}{\\vect p}. \\end{equation}     Let us find the explicit form of the thermal mass and Lorentz-breaking terms both in the low and high temperature regimes. Equation \\eqref{DispRel} can be solved perturbatively: \\begin{equation} \\label{DispRelGen} (p^0)^2 = m^2 + |\\vect p|^2 +  \\Re \\Sigma_\\mathrm R(E_\\vect p,\\vect p) + O(\\kappa^4), \\end{equation} where we recall that $E_\\vect p = \\sqrt{ m^2 + |\\vect p|^2}$. At low temperatures the modified dispersion relation, according to Eqs.~\\eqref{ReSigmaT} and \\eqref{DispRelGen}, is approximately given by \\begin{equation} \\begin{split} \\label{RelDispT} (p^0)^2 &\\approx \\mth^2 +|\\vect p|^2 -   \\kappa^2 \\sqrt{\\dfrac{ m^5 T^3 }{{2\\pi^3}}} \\expp{-m/T} \\left( \\dfrac{|\\vect p|^2}{3m^2+4|\\vect p|^2} \\right), \\end{split} \\end{equation} where the leading contribution to the thermal mass is \\begin{equation} \\label{PhysMassT} \\mth^2 \\approx m^2 - \\frac{1}{6}  \\kappa^2 m^2 T^2. \\end{equation} In the high three-momentum limit the term which modifies the dispersion relation becomes a constant and can be reabsorbed in the thermal mass. The group velocity is modified according to \\Eqref{GroupVelocity}: \\begin{equation} \\vect v \\approx \\frac{\\vect p}{p^0} \\left( 1 - \\frac{  {\\kappa }^2 m^{9/2} {T  }^{3/2} \\expp{-m/T} } {{\\sqrt{2 \\pi^3}}\\,{\\left( 3m^2 + 4|\\vect p|^2  \\right)}^2} \\right). \\end{equation} Notice that at low temperatures the Lorentz breaking term carries a Boltzmann factor $\\expp{-m/T}$. This is due to the mentioned fact that the nontrivial momentum dependence comes from the thermal scalar particles whose abundance is exponentially suppressed at low temperatures. Analogously to what happens in electrodynamics \\cite{DonoghueEtAl85}, the effect of the graviton bath only shows up in the thermal mass.   At high temperature, $T\\gg m$, the modified dispersion relation is found to be \\begin{equation} \\label{RelDispHighT} \\begin{split} (p^0)^2 &\\approx \\mth^2 + |\\vect p|^2  \\\\ &\\quad+  \\frac{\\kappa^2 m^2 T^2}{48} \\left[ \\frac{ E_\\vect p}{|\\vect p|} \\ln \\left(\\frac{ 2E_\\vect p+ |\\vect p|}{ 2E_\\vect p- |\\vect p|}\\right) - 1 \\right], \\end{split} \\end{equation} and the thermal mass is \\begin{equation} \\mth^2 \\approx m^2 - \\frac{11}{48} \\kappa^2 m^2 T^2. \\end{equation} At high three-momentum the modification of the dispersion relation can be also reabsorbed in the thermal mass. The group velocity is given by \\begin{equation} \\label{SpeedHighT} \\begin{split} \\vect v \\approx \\frac{\\vect p}{p^0} \\bigg[ 1 &+ \\frac{1}{96} \\frac{ \\kappa^2 T^2 m^4}{(4m^2+3|\\vect p|^2)|\\vect p|^2} \\\\ &- \\frac{1}{96} \\frac{\\kappa^2 T^2  m^4}{|\\vect p|^3 E_\\vect p} \\ln \\left( \\frac{ 2 E_\\vect p + |\\vect p|}{2 E_\\vect p - |\\vect p|} \\right) \\bigg]. \\end{split} \\end{equation} At high temperature the terms which break the Lorentz symmetry are no longer exponentially suppressed. Notice also that Eq. (\\ref{RelDispHighT}) shows that there is no modification to the dispersion relation for massless scalars.  The Lorentz breaking term is negative in the low-temperature case, which implies that the speed of propagation is lowered with respect to the standard relativistic case. In contrast, the Lorentz breaking corrections have a positive sign at high temperatures, so that the speed of propagation is increased. However, it is always lower than the speed of light, as can be seen by expanding \\Eqref{SpeedHighT} in the ultrarelativistic limit: \\begin{equation} \\vect v \\approx \\frac{\\vect p}{|\\vect p|} \\left(1 - \\frac{m^2\\left[ 48 + {\\kappa }^2 T^2 \\left( \\ln 3 -1 \\right) \\right] }{96 {|\\vect p|}^2} + \\cdots \\right)  . \\end{equation}   We should emphasize that the above breaking of the Lorentz symmetry is only effective: if one applies a Lorentz transformation simultaneously to the particle and the bath, all results are Lorentz invariant. This can be seen can by introducing the unit vector $l^\\mu$ which gives the four-velocity of the thermal bath. Then the energy and the three-momentum of the particle with respect to the bath are \\begin{equation} p^0 = -l_\\mu p^\\mu, \\qquad |\\vect p| = \\sqrt{(l_\\mu p^\\mu)^2+p_\\mu p^\\mu}\\, . \\end{equation} In this article for obvious reasons of simplicity we worked in the bath rest frame where $l^\\mu=(1,0,0,0)$.  For low values of the momentum the effective dispersion relation can be expanded in powers of $\\vect p$. For instance in the low-temperature case, for momenta satisfying $|\\vect p|<m$, the effective dispersion relation \\eqref{RelDispT} can be expanded as \\begin{equation} \\begin{split} \\label{RelDispTExp} (p^0)^2 \\approx \\mth^2 +|\\vect p|^2 &-   \\dfrac{\\kappa^2 m^{5/2} T^{3/2}\\expp{-m/T} }{\\sqrt{2\\pi^3}}  \\frac{|{\\vect p}|^2}{3m^2} \\\\ &+ \\dfrac{\\kappa^2 m^{5/2} T^{3/2}\\expp{-m/T} }{\\sqrt{2\\pi^3}} \\frac{2|{\\vect p}|^4}{27m^4} + \\cdots . \\end{split} \\end{equation} It is worth mentioning the difference between this expansion and that of Eq.~\\eqref{Eq1}. In a fundamental approach to the Lorentz symmetry breaking one expects each additional power of momentum to be suppressed by increasing powers of the Planck mass, whereas in the effective breaking approach we have pursued there are several energy scales. Hence there is more freedom in the possible values of the suppression factor. For instance, in the low-temperature regime ---\\ie, in \\Eqref{RelDispTExp}--- the $|\\vect p|^4$ term is suppressed by $(T/m)^{3/2}e^{-m/T} \\ll 1$ in Planck units. The fact that in the latter approach there are more energy scales also explains why modifications to the $|\\vect p|^2$ term are present in Eq.~\\eqref{RelDispTExp} while they are not included in Eq.~\\eqref{Eq1}. Notice also that the corrections that we found only contain even powers of the momentum, and thus are in agreement with the results of Ref.~\\cite{Lehnert03}.  In the present universe the effects we discussed are completely negligible when applied to electrons or protons, both because they are proportional to the Planck length square and because they are exponentially suppressed. To obtain relevant effects, one should consider Planckian temperatures. However, in this case perturbation theory around a flat-spacetime approximation probably fails. Therefore the results obtained here should be considered an indication that effective violations of Lorentz invariance indeed occur, and that gravitational interactions cannot be neglected when exploring highly energetic regions of dispersion relations.   Whereas the real part of the self-energy gives the change in energy of the particle and hence the dispersion relation, the imaginary part accounts for the dissipative effects and gives the thermalization rate. As we have already commented, at the order $\\kappa^2$ the imaginary part of the self-energy is zero. We do not see the thermalization effect since all processes which could contribute to it, such as the one shown in Fig.~\\ref{fig:cut} (on the left), are kinematically forbidden. In order to account for thermalization effects we should compute the imaginary part of the self-energy to order $\\kappa^4$. At this order, the particle can exchange momentum with the thermal bath through processes such as the one to the right of Fig.~\\ref{fig:cut}, which would correspond to the Compton scattering in electrodynamics. Notice here a significant difference with respect to the vacuum case: while at zero temperature the self-energy is purely real to all orders of perturbation theory (because the particle is stable), at finite temperature it will acquire an imaginary contribution to order $\\kappa^4$.     \\begin{figure} \\hfill \\includegraphics{cut1.eps} \\hfill \\includegraphics{cut2.eps} \\hfill\\mbox{} \\caption{Diagrams such as the one depicted on the left would contribute to order $\\kappa^2$ to the thermalization rate at finite temperature, and hence to the imaginary part of the self-energy. However all those one-vertex diagrams are identically zero because of kinematical reasons. The first non-vanishing contribution would be due to Compton-like diagrams such as the one depicted on the right, and hence it would be a $\\kappa^4$ contribution.} \\label{fig:cut}  \\end{figure}  Let us end with a short summary of the main points of the paper. We have illustrated with a particular example how local Lorentz symmetry is effectively violated because of the interactions with a nontrivial ensemble of metric fluctuations, even if Lorentz symmetry holds at a fundamental level. We have also shown how this effective violation can be addressed at low energies within standard physics. The main quantitative result of the paper are Eqs.~\\eqref{RelDispT} and \\eqref{RelDispHighT}, which explicitly show the modifications of the dispersion relation. As in the electromagnetic case, this effect is exponentially suppressed at low temperatures. Moreover, the modifications of the dispersion relation are suppressed when the three-momentum of the massive particle (defined in the heat bath rest frame) is much larger than the temperature and the mass. This last result is somewhat unexpected since the gravitational coupling grows with the energy, unlike the electromagnetic case. Therefore, in spite of the derivative coupling, no violation of Lorenz invariance is found in the high-momentum limit, at least at one loop. Finally we have also shown that gravitational interactions do not generate a four-derivative term in the scalar field action."
    },
    "0311/astro-ph0311572_arXiv.txt": {
        "abstract": "We use the Karachentseva (1973) ``Catalogue of Very Isolated Galaxies'' to investigate a candidate list of $>\\,$100 very isolated early-type galaxies. Broad-band imaging and low resolution spectroscopy are available for a large fraction of these candidates and result in a sample of 102 very isolated early-type galaxies, including 65 ellipticals and 37 S0 galaxies. Many of these systems are quite luminous and the resulting optical luminosity functions of the Es and early-types (E+S0s) show no statistical differences when compared to luminosity functions dominated by group and cluster galaxies. However, whereas S0s outnumber Es 4:1 in the CfA survey, isolated Es outnumber S0s by nearly 2:1. We conclude that very isolated elliptical galaxies show no evidence for a different formation and/or evolution process compared to Es formed in groups or clusters, but that most S0s are formed by a mechanism (e.g., gas stripping) that occurs only in groups and rich clusters. Our luminosity function results for ellipticals are consistent with very isolated ellipticals being formed by merger events, in which no companions remain.  CHANDRA observations were proposed to test specifically the merger hypothesis for isolated ellipticals. However, this program has resulted in the observation of only one isolated early-type galaxy, the S0 KIG 284, which was not detected at a limit well below that expected for a remnant group of galaxies. Therefore, the hypothesis remains untested that very isolated elliptical galaxies are the remains of a compact group of galaxies which completely merged. ",
        "introduction": "Because the two-point correlation function of galaxies is very steep \\citep[e.g.,][]{peebles93}, the best place to find a galaxy is next to another galaxy. Another way of saying this is that truly isolated galaxies are exceptionally rare in the Universe. For example, \\citet{tully87} finds no completely isolated galaxies at all within the local supercluster; all are members of small or large bound groups or loose associations. And yet the mythical ``field population'' of galaxies continues to be referenced in the literature as a comparison to various cluster, group, compact group and interacting galaxy studies \\citep[e.g.,][]{zabludoff96,koopmann98,toledo99,christlein03}. Still, it should be possible to locate some galaxies which are very isolated relative to the much larger number which are members of clusters, small bound groups or loose and still unbound associations. If the case can be made that a potentially isolated galaxy has not experienced a merger or interaction with another galaxy for a time much longer than the timescale for the physical process under study (e.g., 10$^8$ yrs for a starburst; 10$^7$ yrs for an AGN active phase; few $\\times 10^7$ yrs for spiral galaxy density wave generation), then these very isolated galaxies provide an excellent baseline comparison sample for that property. For example, \\citet{haynes80,haynes84a} and \\citet{haynes84c} used an isolated galaxy sample as a baseline for the \\ion{H}{1} properties of galaxies; \\citet{adams80} have shown that very isolated galaxies are deficient in radio continuum emission compared to other galaxies; and \\citet{koopmann98} have shown that spirals in the Virgo cluster have different structural properties than more isolated spirals. In all of these cases, the use of an isolated galaxy comparison sample allows us to infer something about the effects that an external environment can have on the internal characteristics of a galaxy. For this reason, a large sample of very isolated galaxies is extremely valuable, providing an important comparison sample, which facilitates studies of environmental effects on galaxies.  Based upon a visual inspection of all $\\sim 30,000$ bright (m$_B \\leq 15.7$) galaxies catalogued by \\citet{zwicky57}, \\citet{karachentseva73} listed over 1000 Zwicky galaxies that are very isolated. Because few redshifts of these galaxies were available at the time of Karachentseva's work, she based her isolation criterion on the observed angular sizes and distances between galaxies. Thus, this sample is representative, not complete, because other isolated galaxies would not be included if they were projected onto foreground or background galaxies of comparable angular size (see below). Surprisingly, Karachentseva's (1973) list contains over 100 galaxies which she classified morphologically as early-types (Es and S0s).  The possible existence of very isolated early-type galaxies is unexpected based upon the typical environment for such systems \\citep[e.g.,][on the morphology/density relationship]{dressler84,oemler92}. And yet recent theoretical \\citep[e.g.,][]{barnes92,barnes85,athanassoula97} and observational work \\citep[e.g.,][]{mulchaey98,mulchaey99,zabludoff03} suggest that very isolated elliptical galaxies could be the final outcome in the evolution of dense groups of galaxies that have completely merged, leaving only a single, large elliptical galaxy behind. The detection of the isolated elliptical NGC 1132 as a diffuse X-ray source by \\citet{mulchaey99} at a luminosity (L$_x \\approx 5 \\times 10^{42}$h$^{-2}_{70}$ ergs s$^{-1}$) comparable to the L$_x$ of elliptical-dominated groups of galaxies is significant new evidence in favor of the merger hypothesis for forming ellipticals in general. Thus, the number, detailed structure, and X-ray properties of very isolated elliptical galaxies provide new tests of the merger hypothesis for elliptical galaxies. Are these true elliptical galaxies or do they possess systematic differences from the ellipticals found in rich groups and clusters? Similar questions could be asked about very isolated S0 galaxies, whose detailed histories are even more poorly understood. Proposed scenarios include ram-pressure stripped spirals \\citep{quilis00}, mergers of large galaxies with small companions \\citep{mihos95b,bekki98}, and early starbursts, followed by passive evolution \\citep{welch03}. Only two of these possible histories can account for isolated S0 galaxies.  In this Paper, we begin the study of very isolated early-type galaxies by scrutinizing a list of candidate very isolated early-type galaxies drawn from \\citet{karachentseva73}, presenting a final list of very isolated E \\& S0 galaxies and computing the optical luminosity function (LF) for them. In addition, we present a single CHANDRA ACIS imaging analysis of one of the most luminous galaxies in this sample, a first attempt at testing the merger hypothesis for the formation of very isolated ellipticals by detecting diffuse X-ray emission.  In the following section we describe the \\citet{karachentseva73} sample and present the sample of isolated early-type galaxies identified from it through further scrutiny. In Section 3, we present the optical luminosity functions of the ellipticals and the S0s in the sample. Section 4 presents the CHANDRA and optical imaging and spectroscopy results for the one isolated early type galaxy observed by CHANDRA: KIG 284. Section 5 includes a brief discussion and conclusions drawn from this work. Throughout this work we adopt the NASA Extragalactic Database (NED) abbreviation for the Karanchentseva (1973) galaxies (KIG) as well as H$_0 = 70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "In this paper we have presented a sample of bright (m$_{Zw} \\leq 15.7$), very isolated, early-type galaxies that can be used to investigate how early-type galaxies are formed. Starting with the KIG Catalogue of $>1000$ very isolated galaxies selected by \\citet{karachentseva73} from the CGCG, we used the POSS and deeper images to scrutinize more than 200 KIG galaxies as potential Es and S0s. Table \\ref{tbl:objlist} presents the basic data, including luminosities and recession velocities, for the 65 ellipticals and 37 S0 galaxies in the very isolated early-type galaxy sample that resulted from this detailed examination process. We emphasize that this sample is representative and not complete because it does not include very isolated galaxies which, by chance, are projected close to foreground or background galaxies that happen to have a similar angular size. We have used this sample to construct luminosity functions (LFs) for these galaxies and have compared them to the LFs of \\citet{marzke94b} for ellipticals and S0s found throughout the CfA redshift survey region (and therefore, biased heavily towards ellipticals and S0s in rich galaxy regions). After appropriate scaling, we find that the LFs for ellipticals and S0s from \\citet{marzke94b} are excellent matches to the very isolated E and S0 LFs. However, the relative numbers of Es and S0s in the \\citet{marzke94b} and the KIG samples are quite different. In the CfA survey S0s outnumber Es $\\sim$ 4:1 but in the KIG Catalogue (see Table \\ref{tbl:objlist}), ellipticals outnumber S0s nearly 2:1.  We interpret these LF results to mean that very isolated environments are just as likely to form very luminous ellipticals as are dense groups and rich clusters. Since it is now thought that many or perhaps all elliptical galaxies formed by mergers of disk galaxies \\citep{barnes92,mihos95a}, our LF results for KIG galaxies are not inconsistent with the hypothesis that very isolated Es formed by mergers as well. This is a particularly interesting suggestion for the most luminous ellipticals, whose merger histories in clusters and dense groups could be quite different from the mergers that occasionally occur in sparser environments. If the very luminous Es in the KIG sample formed through mergers, then they should be the best cases in the KIG sample for being the merger remnants of entire dense groups of galaxies.  We intended to further test the merger hypothesis for forming very isolated ellipticals specifically by searching for extended X-ray emission around a few of the most luminous KIG Es. Luminous diffuse X-ray emission would be expected if these ellipticals were the final stage in the evolution of poor, compact groups of galaxies \\citep{mulchaey98}. \\citet{mulchaey99} have reported the detection of one isolated elliptical, NGC 1132, at L$_x = 5 \\times 10^{42}$h$^{-2}_{70}$ ergs s$^{-1}$ and \\citet{mulchaey98} have detected a few elliptical-dominated poor groups of galaxies at comparable L$_x$. Thus, we had hoped to observe several luminous KIG ellipticals at or below this sensitivity limit to characterize the X-ray properties of isolated Es. However, only one of our four proposed targets was observed with CHANDRA/ACIS-S and we have now shown this galaxy, KIG 284, to be a luminous S0 and not an elliptical. Therefore, our single non-detection at L$_x < 10^{41}$h$^{-2}_{70}$ ergs s$^{-1}$, a full factor of ten below the NGC 1132 detection, does not constrain significantly the origins of the KIG isolated ellipticals or S0s. And since our upper limit on KIG 284 is comparable to the L$_x$ observed for some S0 galaxies \\citep[e.g., NGC 1553;][]{blanton01}, we have no evidence that KIG 284 has abnormal X-ray properties for an S0 galaxy. Therefore, further CHANDRA observations are required to test the merger hypothesis for these systems.  While there is now significant evidence that many or all ellipticals formed by merging of disk galaxies, the origins of S0 galaxies are more obscure. Because most S0 studies have concentrated on cluster S0s, the proposed origins of these systems mostly involve processes which remove gas from disk galaxies and thus truncate star formation \\citep{mihos95b,quilis00}. However, some of these processes may not be relevant to the histories of very isolated S0 galaxies. The existence of isolated S0s at luminosities comparable to group and cluster S0s suggests that there are at least two ways in which S0s can form, one method which is operable where a dense intra-group or intra-cluster gas is present and one which does not require external gas. However, the dearth of S0s compared to ellipticals (1:1.75) in the KIG sample is in great contrast to their relative abundance in the general population (S0s outnumber Es 4:1 in the CfA survey). This argues that gas stripping or other removal processes that involve a dense external medium are the most efficient method for forming S0 galaxies."
    },
    "0311/astro-ph0311058_arXiv.txt": {
        "abstract": "We introduce \\textsc{ann}\\emph{z}, a freely available software package for photometric redshift estimation using Artificial Neural Networks. \\textsc{ann}\\emph{z} learns the relation between photometry and redshift from an appropriate training set of galaxies for which the redshift is already known. Where a large and representative training set is available \\textsc{ann}\\emph{z} is a highly competitive tool when compared with traditional template-fitting methods.  The \\textsc{ann}\\emph{z} package is demonstrated on the Sloan Digital Sky Survey Data Release 1, and for this particular data set the r.m.s. redshift error in the range $0\\la{z}\\la{0.7}$ is $\\sigma_\\mathrm{rms} = 0.023$.  Non-ideal conditions (spectroscopic sets which are small, or which are brighter than the photometric set for which redshifts are required) are simulated and the impact on the photometric redshift accuracy assessed.  The package may be freely downloaded from \\verb+http://www.ast.cam.ac.uk/~aac+. ",
        "introduction": "\\label{sec.intro}  In its most general sense, the term \\emph{photometric redshift} refers to a redshift estimated using only medium- or broad-band photometry or imaging. Most commonly, photometric redshifts are determined on the basis of galaxies' colours in three or more filters (thus giving a very coarse approximation to the spectral energy distribution, hereafter SED), but they could also be based on other properties which can be derived from images, such as the angular size or concentration index. The method has found successful application to deep-field and wide-field surveys, notably the Hubble Deep Field \\citep*[e.g.][]{1999ApJ...513...34F}, and the Sloan Digital Sky Survey \\citep{2003AJ....125..580C}.  The most commonly used approach to photometric redshift estimation is the \\emph{template-matching} technique. This requires a set of `template' SEDs covering a range of galaxy types, luminosities and redshifts appropriate to the population for which photometric redshifts are required. For a particular target galaxy, the photometric redshift is chosen to be the redshift of the most closely matching template spectrum; this is usually defined as the template which minimizes the $\\chi^2$ between the template and actual magnitudes.  The template spectra are usually derived from a small set of SEDs representing different classes of galaxy at redshift $z=0$, which are then manually redshifted to give a discrete sampling along the redshift axis (note that this method does not account for \\emph{evolution} with redshift). Commonly used template sets are the \\citet*[CWW;][]{1980ApJS...43..393C} SEDs which are derived observationally, or those of \\citet{BruzChar93}, derived from population synthesis models. The template-matching technique owes its popularity to the very few resources required for a basic implementation (i.e. a handful of template SEDs), but the accuracy of the technique strongly depends on the extent to which the template spectra are representative of the target populations: for example, template SEDs derived from observations of low-redshift galaxy populations may be a poor match for populations at higher redshifts.  The chances of success can be improved by increasing the number of templates, or by more carefully matching the templates to the populations being studied. For example, the spectroscopic catalogue of the Sloan Digital Sky Survey (SDSS; \\citealt{2000AJ....120.1579Y}) could be used to produce a set of templates which are very well representative of the SDSS photometric catalogue \\citep{2003AJ....125..580C}. However, in situations with such a large amount of prior redshift information about the sample, the template-matching technique is not the best approach: so-called \\emph{empirical} methods usually offer greater accuracy, as well as being far more efficient.  In essence, empirical photometric redshift methods aim to derive a parametrization for the redshift as a function of the photometric parameters. The form of this parametrization is deduced through use of a suitably large and representative training set of galaxies for which we have both photometry and a precisely known redshift. A simple example is to express the redshift as a polynomial in the galaxy colours (e.g. \\citealt{1995AJ....110.2655C}; \\citealt{2000AJ....119.2598S}). The coefficients in the polynomial are varied to optimize the fit between the predicted and measured redshift. The photometric redshift for the galaxies for which we have no spectroscopy can then be estimated by applying the optimized function to the colours of the target galaxy.  Ideally the training set would be a representative subset of the actual photometric target sample (this has the attractive side-effect of nullifying any systematics in the photometry). However, the training set could also be derived from a set of template spectra or from simulated catalogues (e.g. \\citealt{vanzella}). The photometry for the training set must be for the same filter set and should have the same noise characteristics as that for the target sample. The trained method can usually only be reliably applied to target galaxies within the ranges of redshift and spectral type adequately sampled by the training set.  In this paper we introduce \\textsc{ann}\\emph{z}, a software package for photometric redshift estimation using Artificial Neural Networks (hereafter ANNs) to parametrize the redshift-photometry relation. It can be shown (e.g. \\citealt{jones90}; \\citealt{blum91}) that a sufficiently complex ANN is capable of approximating to arbitrary accuracy any continuous functional mapping. ANNs have previously found a number of applications in astronomy, including morphological classification of galaxies (e.g. \\citealt{1996MNRAS.283..207L}; \\citealt{ball}) star/galaxy separation \\citep{1996A&AS..117..393B} and object detection \\citep[e.g.][]{2000MNRAS.319..700A}. \\citet*{firth03} previously demonstrated the feasibility of using ANNs for photometric redshift estimation, and more recently \\citet{vanzella} have applied the method to the Hubble Deep Fields.  The layout of this paper is as follows. In \\S\\ref{sec.anns} Artificial Neural Networks are introduced, and the particular methods used by \\textsc{ann}\\emph{z} are explained. In \\S\\ref{sec.sdss} \\textsc{ann}\\emph{z} is applied to the Sloan Digital Sky Survey. The results are compared with rival photometric redshift estimators and various extensions to the basic technique are explained and illustrated. Finally, less ideal conditions are simulated to assess the impact on the accuracy of photometric redshift estimation. In \\S\\ref{sec.conc} the results are summarised, and prospects for the application of \\textsc{ann}\\emph{z} discussed. ",
        "conclusions": "\\label{sec.conc}  In appropriate circumstances, \\textsc{ann}\\emph{z} is a highly competitive tool for photometric redshift estimation. However, it does rely on the existence of a sufficiently large training set which is representative of the particular populations being studied. The package's utility therefore lies particularly with large photometric surveys such as the SDSS, GOODS \\citep{2001AAS...198.2501D} or the VIRMOS-VLT Deep Survey \\citep{2003SPIE.4834..173L}, some of which include spectroscopic surveys for subsets of the photometric catalogues (for example, of the eventual 100 million photometric objects which the SDSS expect to catalogue, 1 million will also have spectroscopy, and hence accurate redshifts).  A major problem for empirical photometric redshift estimators is the difficulty in extrapolating to regions of the input parameter space which are not well sampled by the training data. Care should be taken to match the training data to the target sample as closely as possible in terms of the magnitude and colour distributions of each. Use of an evaluation set is essential when applying \\textsc{ann}\\emph{z} to a new data set: the good performance demonstrated here on the SDSS data cannot be guaranteed on different data sets.  A potential solution to the problem of obtaining training sets when spectroscopy is difficult to obtain is to use simulated catalogues as training data (e.g. \\citealt{vanzella}). Since this requires the use of theoretical SEDs it introduces the disadvantages of the template-based methods, such as the need for precise calibration. However, the ANN approach has advantages over standard template-matching: simulated catalogues can contain galaxies representing a large range of complex star formation histories, dust extinction models and metallicities etc., giving fully Bayesian statistics, and ANNs allow much more flexible weighting to be applied to the filters than is possible with the simple $\\chi^2$-weighting of standard template-matching."
    },
    "0311/astro-ph0311602_arXiv.txt": {
        "abstract": "The Surface Brightness Fluctuations (SBF) Method for distance determinations of elliptical galaxies is been modeled in order to investigate the effect of the Point Spread Function (PSF). We developed a method to simulate observations of SBF of galaxies having various properties and located at different distances. We will use this method in order to test the accuracy on the estimates of the extra-galactic distances for PSFs representing typical seeing conditions, AO systems and for future observations with ELTs close to the diffraction limit. ",
        "introduction": "The SBF Method (Tonry \\& Schneider 1988) measures the irreducible mottling in an early-type galaxy image due to the Poisson fluctuations in the finite number of stars per pixel. The method has important applications for both extra-galactic distance (e.g. Jensen et al. 1998) and stellar population (e.g. Blakeslee et al. 2001) studies. Its advantage is based on the fact that if we observe the same region of two identical elliptical galaxies at different distances, the one being twice as far away as the other, the more distant galaxy would have four times more stars contributing the same average flux into the pixel as the stars of the closer. The pixel-to-pixel variation in the flux due to fluctuations in the number of stars then scales inversely to the distance and so the fluctuations in surface brightness can be used as a distance indicator.  \\begin{figure}[t] \\begin{center} \\includegraphics[height=\\textwidth,angle=270]{gouliermisF1.eps} \\end{center} \\caption[]{Grayscale contours (top panel) and 3D brightness distributions (bottom panel) of a simulated galaxy.  The example shows an E5 galaxy with $R_{\\rm eff} \\simeq 7.5$ kpc located at 20 Mpc following a de Vaucouleurs profile. The three steps of the simulations are shown: The model of galaxy's brightness profile ({\\it right}), its SBF ({\\it middle})  and the image convolved with a Gaussian PSF of FWHM = 5 px ({\\it left}).  A 1024 $\\times$ 1024 pixel CCD array with a pixel scale of 0\\farcs 146 was assumed.} \\label{fig1} \\end{figure} ",
        "conclusions": "We present a model of the SBF method for distance estimations. Our aim is to use it for the investigation of hypothetical SBF observations with ELTs. Specifically, the measurements of the SBF amplitude of simulated observations of elliptical galaxies with various properties can lead to estimations of the distances of the galaxies with the use of available SBF calibrations (e.g. Tonry et al.  2001). The comparison between the input distance that was initialy selected for every galaxy and the estimated value will allow us to define the accuracy of the SBF method with the use of various telescopes and under various observational conditions. Thus we will be able to study the accuracy and the limitations of SBF observations with future ELTs."
    },
    "0311/astro-ph0311328_arXiv.txt": {
        "abstract": "We present a measurement of cosmic shear on scales ranging from $ 10\\arcsec $ to $ 2\\arcmin $ in 347 WFPC2 images of random fields.  Our result is based on shapes measured via image fitting and on a simple statistical technique; careful calibration of each step allows us to quantify our systematic uncertainties and to measure the cosmic shear down to very small angular scales.  The WFPC2 images provide a robust measurement of the cosmic shear signal decreasing from $5.2\\% $ at $ 10\\arcsec $ to $2.2\\% $ at $ 130\\arcsec $. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311434_arXiv.txt": {
        "abstract": "We describe an observation of the X-ray luminous SN1978K in NGC 1313 using the ACIS detector on {\\it Chandra} and an archival {\\it Newton} EPIC observation.  The models that provided good fits to the {\\it ASCA} SIS and GIS and the {\\it ROSAT} PSPC spectra no longer do so for the ACIS and EPIC spectra.  The best-fit model to the ACIS and EPIC spectra are dual hot plasma models (vmekal); one component is soft (T = 0.61$^{+0.04}_{-0.05}$ keV, 90\\% errors) and the other is harder (T = 3.16$^{+0.44}_{-0.40}$ keV).  For the varying abundances permitted within the model, only the Si abundance of the soft component differs from solar, with a value n$^{\\rm Si}_{\\rm SN78K}$~/~n$^{\\rm Si}_{\\rm solar}$ = 3.20$^{+1.80}_{-1.90}$ (90\\% errors).  From a ratio of the low- and high-T model fits to the {\\it Chandra} and {\\it XMM-Newton} spectra, we infer an exponent ${\\rm n}$ of the ejecta density distribution, ${\\rho}_{\\rm ejecta}$ $\\propto$ r$^{\\rm -n}$, of $\\sim$5.2 adopting a circumstellar matter distribution exponent of s = 2 (${\\rho}_{\\rm cs}$ $\\propto$ r$^{\\rm -s}$).  The 0.5-2 keV light curve shows essentially no decline; the 2-10 keV light curve, constructed only of the {\\it ASCA}, {\\it XMM-Newton}, and {\\it Chandra} observations, shows a drop of 1.5 from the {\\it ASCA} epoch. The hard band decline, together with the apparently enhanced Si emission, signal the start of the X-ray decline of SN1978K. ",
        "introduction": "X-ray emission may emerge at two stages in the evolution of a supernova: near the time of the shock breakout and via circumstellar interaction at a later date.  Three different processes generate X-rays: from the prompt thermal burst; via Compton-scattered ${\\gamma}$-rays from the synthesized matter's radioactive decay (from which X-rays are emitted while the debris is optically thick to ${\\gamma}$-rays); and from the interaction of the shock with circumstellar material in the vicinity of the progenitor.  Only SN1987A has been seen to emit X-rays definitely attributed to Compton-scattered ${\\gamma}$-rays (Dotani et al. 1987, Nature, 330, 230).  Two mechanisms exist for the circumstellar production of X-rays: X-rays from the outgoing shock or from the reverse shock \\citep{Fran96}.  Below 10 keV, the reverse shock should be the dominant emitter because the density behind the shock is higher, by a factor of $\\sim$4, than the swept-up circumstellar density behind the forward shock.  If a cool shell forms because the reverse shock is radiative, that shell absorbs X-rays below 10 keV and the observed 1-10 keV emission must be from the forward shock as discussed for SN1993J \\citep{Fran96}.  Other than SN1987A, few supernova have been followed in the X-ray band.  The total current pool consists of $<$20 objects; of those, only four supernovae have data sets consisting of more than four observations (SN1978K, SN1993J, SN1998S, SN1999em; \\citealt{S95,IL02}). SN1978K is one of a small number of extremely luminous, X-ray emitting supernovae \\citep{S95,IL02}.  As X-ray observations are difficult to obtain, each bright supernova contributes valuable knowledge to our understanding of the X-ray emission from supernovae.  The X-ray emission of SN1978K was first uncovered in an observation obtained with the {\\it ROSAT} PSPC in 1992 \\citep{Ryder93} (hereafter, R93).  It had been detected optically during a spectrophotometric survey of H II regions two years earlier and reported as a nova \\citep{DopRyd90} at that time.  It had also been detected as a powerful radio source in 1982 but never followed up (R93).  As part of the investigation spurred by the X-ray detection, examination of archival optical plates uncovered a light curve that while crude was sufficient to assign an explosion date near 1 June 1978 (R93).  To further study SN19798K, observations were obtained during the mid- and late-1990s with {\\it ASCA}, {\\it Hubble Space Telescope}, {\\it ROSAT} (High Resolution Imager), {\\it Australia Telescope Compact Array} (ATCA), and the {\\it Anglo-Australian Observatory} (AAO) covering the X-ray, ultraviolet, optical, and radio bands \\citep{Schl99} (hereafter, S99).  An early {\\it ASCA} spectrum and {\\it ROSAT} light curve were reported in \\cite{Petre94} and \\cite{Schl96}, respectively.  The {\\it ATCA} radio light curve showed behavior typical of a circumstellar interaction, with a rapid turn-on followed by a slow decline at all radio wavelengths (e.g., \\citealt{Weiler96}; S99).  Optically, the {\\it AAO} spectra of SN1978K showed strong emission lines of the Balmer series as well as He I, [O III], [Ne III], and [Fe II] (R93).  In the ultraviolet, the Mg II doublet and [Ne IV] lines were detected using the {\\it HST} Faint Object Spectrograph (S99).  The X-ray light curve was consistent with no decline throughout the decade (S99).  The {\\it ASCA} spectra were best-fit with absorbed, single-component continuum models; no evidence for line emission existed in the data (\\citealt{Petre94}, S99).  SN1978K is important because it is among the first known X-ray emitting supernovae with luminosities above 10$^{39}$ erg s$^{-1}$. It is also a nearby object and may be studied in greater detail than more distant supernovae.  In this paper, we present the {\\it Chandra} ACIS spectrum of SN1978K and contrast it with the recent {\\it XMM-Newton} and earlier {\\it ASCA} spectra.  Throughout we adopt a distance to NGC 1313 of 4.13$\\pm$0.11 Mpc as determined by \\cite{Mendez02}. ",
        "conclusions": "We conclude that the {\\it Chandra} and {\\it XMM-Newton} observations detected the first hints of increasing transparency of the matter surrounding SN1978K.  The increasing transparency signals the start of its decline in the X-ray band.  The low-temperature spectral components remain essentially unchanged from the {\\it ASCA} epochs except for the addition of the detected Si. Within the errors, the high-temperature components have also not changed; if we restrict our attention solely to the {\\it XMM-Newton} and {\\it Chandra} observations of 2000 and 2002, then we see a decline of a factor of $\\sim$1.5, consistent with an interpretation of increasing transparency.  Additional X-ray observations will confirm or refute our interpretation.  If X-ray line emission is becoming increasingly visible, line emission in other bands may be expected to change depending upon the atomic cascades.  An observation with the {\\it HST} STI Spectrograph is clearly warranted; such an observation would also aid in sorting out the emission region and clarifying the nature of the line emission uncovered with the FOS observation (S99).  Additional X-ray observations over the coming years are important to investigate possible variations in the wind as well as detecting the expected developing line emission.  If our interpretation is correct, then the decline phase commenced just within the past few years and provides an opportunity to observe the decline.  Comparison with SN1987A will reveal differences likely attributable to the differing densities of the circumstellar media.  At least two radio-emitting supernovae have shown hesitations and increases on the decline, interpreted as enhanced mass loss during the late stages of stellar evolution prior to the progenitor's explosion (e.g., SN2001ig, \\cite{Ryd04}; SN1993J, \\cite{Zim03}; SN1979C, \\cite{Montes00} and references therein).  With the shock moving at about 10$^3$ times the wind velocity, continued X-ray and radio observations are expected to provide in a human lifetime an extended view of several thousand years of late wind and mass loss behavior of a massive star."
    },
    "0311/astro-ph0311420_arXiv.txt": {
        "abstract": "Various methods have been proposed to search for extrasolar planets. Compared to the other methods, microlensing has unique applicabilities to the detections of Earth-mass and free-floating planets.  However, the microlensing method is seriously flawed by the fact that the masses of the detected planets cannot be uniquely determined. Recently, Gould, Gaudi, \\& Han introduced an observational setup that enables one to resolve the mass degeneracy of the Earth-mass planets.  The setup requires a modest adjustment to the orbit of an already proposed microlensing planet-finder satellite combined with ground-based observations.  In this paper, we show that a similar observational setup can also be used for the mass determinations of free-floating planets with masses ranging from $\\sim 0.1\\ M_{\\rm J}$ to several Jupiter masses.  If the proposed observational setup is realized, the future lensing surveys will play important roles in the studies of Earth-mass and free-floating planets, which are the populations of planets that have not been previously probed. ",
        "introduction": "High-precision radial velocity measurements of nearby stars resulted in the detections of more than 100 planetary systems (http:$\\dslash$cfa-www.harvard.edu/planets).  Besides the radial velocity method, planets can also be detected and characterized by using various other methods.  These methods include the pulsar timing analysis, direct imaging, accurate measurement of astrometric displacement, planetary transit, and microlensing (see the review of Perryman 2000).   However, all these methods except the microlensing method can be used to detect only planets which are gravitationally bound to their parent stars.  Theories about the planet formation indicate that planets often do not stay in the same orbit where they were formed. Dynamical interactions with either a massive planetesimal disk \\citep{murray1998, trilling2002} or the protoplanetary nebular gas \\citep{ward1997} can result in inward migration of planets.  On the other hand, gravitational interactions with other planet members \\citep{levison1998} or neighboring stars \\citep{sigurdsson1992} can result in the disruption of the planetary system and thus ejection of planets.  The latter process could lead to a population of free-floating planets.   Microlensing can be occurred by a low-mass object and thus free-floating planets can be, in principle, searched for by detecting short-duration lensing-induced light variations of background stars caused by the planets.  However, this method has not been used due to technical difficulties. These difficulties arise because of the rarity of lensing events combined with the short durations of planetary events. At any given time, only a few out of $10^6$ Galactic bulge stars undergo microlensing with their fluxes magnified enough for detection and a small fraction of those microlensed stars are expected to be lensed by planetary masses. In addition, events produced by planetary lenses last only a short duration of $\\sim 1$ day even for an event caused by a giant planet (with a mass $M\\sim 10^{-3} \\ M_\\odot$) and the duration decreases in proportional to the square root of the lens mass.  Therefore, implementing the microlensing method to search for free-floating planets requires a massive survey for a large number of stars with a very high monitoring frequency.  Several groups have been previously carried out and are currently performing lensing surveys \\citep{alcock1993, aubourg1993, udalski1993, alard1997, bond2002}, but the typical monitoring frequency of these experiments is at most several times per day, which is far too low to detect events caused free-floating planets.  To increase the monitoring frequency, the experiments are employing early warning systems \\citep{alcock1996, afonso2001, bond2001, udalski1994} to issue alerts of ongoing events detected in the early stage of lensing magnification and follow-up observation programs \\citep{alcock1997, rhie1999, albrow1998, yoo2003} to intensively monitor the alerted events hoping to detect short-duration planetary signals on the top of the lensing light curve produced by the parent star.  However, planets detectable from this type of strategy are also bound planets like the ones detectable from other planet search methods.   However, the situation will greatly change with the advent of future lensing surveys with very high monitoring frequencies.  Actually, several such experiments in space and on the ground were proposed. The space-based GEST mission, proposed to NASA by \\citet{bennett2002a}, is designed to continuously monitor $\\sim 10^8$ Galactic bulge main-sequence stars with a frequency of several times per hour by using a 1--2 m aperture space telescope.  Studies about the scientific outcome from a ground-based high-frequency lensing experiment using four 2 m-class telescopes located at good seeing sites around the world and equipped with large format CCD cameras are in progress \\citep{gould2003}. From detailed simulations of Galactic bulge events detectable from the space observation by using the GEST, \\citet{bennett2002a} estimated that the total number of detections of events caused by free-floating planets during the life time of the GEST mission ($\\sim 3.7$ yrs) would be $\\sim ({\\cal O)}10$ for Earth-mass planets and $\\sim ({\\cal O})10^2$ for giant planets under the assumption that one planet was ejected for each star.   However, the microlensing method is seriously flawed by the fact that the detected planet's mass cannot be uniquely determined.  This defect arises because among the three lensing observables related to the physical parameters of the lens, the Einstein time scale $\\te$, the angular Einstein ring radius $\\thetae$, and the projected Einstein radius onto the plane of the observer $\\retilde$, only $\\te$ is routinely measured from the lensing light curve.  These three observables are related to the mass of the lens ($M$), relative source-lens parallax ($\\pi_{rel}= {\\rm AU}/[D_l^{-1}-D_s^{-1}]$, where $D_l$ and $D_s$ are the distances to the lens and source, respectively), and proper motion ($\\mu_{rel}$) by \\citep{gould2000} \\begin{equation} \\te = {\\thetae\\over \\mu_{rel}},\\ \\ \\ \\thetae=\\sqrt{4GM\\pi_{rel} \\over c^2\\ {\\rm AU}},\\ \\ \\ \\retilde = \\sqrt{4GM\\ {\\rm AU} \\over c^2 \\pi_{rel}}. \\label{eq1} \\end{equation} Due to the fact that $\\te$ results from the combination of $M$, $\\pi_{rel}$, and $\\mu_{rel}$, short durations of events do not guarantee the events are caused by planets because they can be produced by low-mass stars or brown dwarfs with exceptionally high proper motions or very small Einstein ring radii.  As a result, even if a large number of short-duration events were detected by future lensing surveys, they could not be definitely identified as planetary events.  To isolate events caused by free-floating planets and determine their masses, it is required to measure the other observables of $\\thetae$ and $\\retilde$ (see \\S\\ 2 for the detailed description about the methods to measure these quantities). Once $\\thetae$ and $\\retilde$ are measured, the lens mass is uniquely determined by \\begin{equation} M = {c^2\\over 4G} \\retilde \\thetae. \\label{eq2} \\end{equation}   Recently, \\citet{gould2003} pointed out that the masses of {\\it bound} Earth-mass planets could be measured by combining a GEST-type satellite with ground-based observations.  They showed that the proper motions of these events can be measured by analyzing the planet-induced perturbations in lensing light curves observed from space.  They also showed that if the satellite is placed in an L2 orbit, the Earth-satellite baseline is sufficient enough to measure lens parallaxes.  In this paper, we show that a similar observational setup proposed by \\citet{gould2003} can also be used for the mass determinations of free-floating planets, which are the population of planets that have not been previously probed.   The paper is organized as follows.  In \\S\\ 2, we describe the methods of $\\retilde$ and $\\thetae$ measurements and explain the reasons for the difficulties in measuring these quantities for events caused by stellar mass lenses.  In \\S\\ 3, we illustrate the relative easiness of $\\retilde$ and $\\thetae$ measurements for events caused by low-mass lenses such as free-floating planets.  In \\S\\ 4, we estimate the mass range over which the proposed mass determination method is sensitive by estimating the uncertainties of the recovered lens masses under the observational setup proposed by \\citet{gould2003}.  In \\S\\ 5, we briefly discuss several problems of the proposed observational setup. We conclude in \\S\\ 6. ",
        "conclusions": "We propose an observational setup that enables one to determine the masses of free-floating planets via gravitational microlensing. The setup is similar to the one of \\citet{gould2003} proposed to resolve the mass degeneracy of the Earth-mass planets and requires a modest adjustment to the orbit of an already proposed microlensing planet-finder satellite combined with ground-based observations. The proposed method is most sensitive to planets with masses ranging from $\\sim 0.1\\ M_{\\rm J}$ to several Jupiter masses. If realized, the proposed high-frequency lensing surveys will be able to play important roles in characterizing Galactic free-floating planets that have not been probed before."
    },
    "0311/astro-ph0311616_arXiv.txt": {
        "abstract": "{We present the kinematical data obtained for a sample of active (Seyfert) and non active isolated spiral galaxies, based on long slit spectra along several position angles in the H$\\alpha$ line region and, in some cases, in the Ca triplet region as well. Gas velocity distributions are presented, together with a simple circular rotation model that allows to determine the kinematical major axes. Stellar velocity distributions are also shown. The main result is that active and control galaxies seem to be equivalent in all kinematical aspects. For both subsamples, the departure from pure circular rotation in some galaxies can be explained by the presence of a bar and/or of a spiral arm. They also present the same kind of peculiarities, in particular, S-shape structures are quite common near the nuclear regions.  They define very similar Tully-Fisher relations. Emission line ratios are given for all the detected HII regions; the analysis of the [NII]/H$\\alpha$ metallicity indicator shows that active and non-active galaxies have indistinguishable disk metallicities.  These results argue in favour of active and non-active isolated spiral galaxies having essentially the same properties, in agreement with our previous results based on the analysis of near infrared images. It appears now necessary to confirm these results on a larger sample. ",
        "introduction": "The DEGAS collaboration was started a few years ago to search for a sufficient condition for a galaxy to host an active nucleus. The active galaxies have been chosen with the following criteria: (a) catalogued as Seyfert 1 or 2 in V\\'eron-Cetty \\& V\\'eron's catalogue (we have checked that they were all classified as Seyferts in the NED database and give their AGN type in Table~\\ref{pargal}); (b) with morphological information in the RC3 catalogue (de Vaucouleurs et al. 1991); (c) isolated, in the sense of not having a companion within 0.4 Mpc (H$_0$ = 75\\kms~Mpc$^{-1}$) and cz$<$ 500 \\kms, or companions catalogued by Nilson without known redshift; (d) nearby, cz $<$ 6000 \\kms; and (e) intermediate inclination (30 to 65 \\degr). Note however that the sample of active galaxies thus defined may not be representative of the whole AGN population, since the selection criteria of the V\\'eron-Cetty \\& V\\'eron catalogue are not necessarily complete. The non-active galaxies have been selected among spirals verifying the same conditions (b), (c), (d) and (e), and with morphologies (given by the complete de Vaucouleurs coding, not just the Hubble type) similar to those of the active spirals. A total of 17 active galaxies and 16 non active galaxies (used as control sample), with similar distribution of morphological types (including the bar type), were selected. Since all galaxies were chosen to be isolated under the criteria defined above, side effects that could be produced by the gravitational interaction with nearby galaxies are avoided. The details of the program and sample definition are given in M\\'arquez et al. (1999, 2000), where it has been shown that the non-active galaxies are well suited to be used as a control sample. Basic properties of the galaxies of our sample are given in Table~\\ref{pargal}.  The observations we carried out included infrared and optical imaging and long slit spectroscopy to derive the kinematical properties of the ionized gas and, for some of them, of the stellar component. Infrared HST images of our selected galaxies were also retrieved and analyzed when available.  The data already presented have shown some interesting results: the analysis of the infrared images indicated that there are no morphological differences between active and non active galaxies in our sample. We also found that most galaxies have bars and sometimes nuclear bars within the large bars. Some of these inner bars were reported for the first time (M\\'arquez et al. 1999, 2000). The full analysis of NGC~6951, the first galaxy of our sample for which all the required data became available (P\\'erez et al. 2000), led us to discover the decoupling between the stellar and gas kinematics in the central regions, and to hint the presence of a younger stellar population there. Velocity drops were detected in the nuclear regions of five of our sample galaxies. The analysis of the Calcium triplet (CaT) lines equivalent widths indicated the presence of a different stellar population spatially coincident with the kinematical drops. Further analysis of the HST images of a total of 14 galaxies with reported velocity dispersion drops (including our five), show evidence for central elongations in most of them, which could be related with the mechanism invoked in numerical models to reproduce such drops (see M\\'arquez et al. 2003 and references therein).  The present work includes the kinematical data on the gas and stars obtained for the galaxies of our sample, together with the emission line ratios measured in the nuclear and extra-nuclear HII regions of the same galaxies. The metallicities of these HII regions, estimated from the observed line ratios, are briefly discussed. A general discussion based on the DEGAS data will be presented in a forthcoming paper.  The paper is organized as follows. The observations and data reduction procedures are presented in Sect.~2. Characteristics of the rotation curves for each galaxy, together with a discussion of the positions of our sample galaxies in the Tully-Fisher relation are presented in Sect.~3.  The properties of nuclear and extra-nuclear HII regions are described in Sect.~4. The summary and conclusions are presented in Sect.~5.  \\begin{table}[t!] \\caption[]{Main properties of the galaxies} \\label{pargal} \\begin{tabular}{llll } \\hline Galaxy & type  & m$_B$ & AGN \\\\ & (RC3) &(mag)  & type \\\\ & (1)   & (2)   & (3) \\\\ \\hline UGC  1395          & .SAT3.. & 14.18       & 1.9 \\\\ IC    184          & .SBR1*. & 14.66       & 2 \\\\ IC   1816          & .SBR2P? & 13.83       & 1 \\\\ UGC  3223          & .SB.1.. & 13.83       & 1 \\\\ NGC  2639          & RSAR1*  & 12.59       & 1.9 \\\\ IC   2510          & .SBT2*. & 14.00$^{*}$ & 2 \\\\ NGC  3660          & .SBR4.. & 12.29       & 2\\\\ NGC  4507          & PSXT3.. & 12.78       & 2 \\\\ NGC  4785          & PSBR3*. & 13.21       & 2 \\\\ NGC  5347          & PSBT2.. & 13.16       & 2 \\\\ NGC  5728          & .SXR1*. & 12.37       & 2 \\\\ ESO139-12          & PSAT4P* & 13.59       & 2 \\\\ NGC  6814          & .SXT4.. & 11.85       & 1.5 \\\\ NGC  6860          & PSBR3.. & 13.68       & 1 \\\\ NGC  6890          & .SAT3.. & 13.05       & 2 \\\\ NGC  6951          & .SXT4.. & 11.91       & 2 \\\\ \\hline NGC   151          & .SBR4..  & 12.23       & \\\\ NGC  1357          & .SAS2..  & 12.46       & \\\\ IC    454          & .SB.2..  & 14.50$^{*}$ & \\\\ NGC  2712          & .SBR3*.  & 12.38       & \\\\ NGC  2811          & .SBT1..  & 12.66       & \\\\ NGC  3571          & PSXT1*.  & 12.99       & \\\\ NGC  3835          & .S..2*/  & 13.20       & \\\\ NGC  4162          & RSAT4..  & 12.55       & \\\\ NGC  4290          & .SBT2*.  & 12.66       & \\\\ NGC  6012          & RSBR2*.  & 12.69       & \\\\ NGC  6155          & .S?....  & 13.20       & \\\\ NGC  7328          & .S..2..  & 13.98       & \\\\ \\hline \\end{tabular}  \\noindent (1) Morphological type taken from the RC3 catalogue;\\\\ (2) B magnitude as given in the RC3 catalogue;\\\\ (3) AGN Seyfert type taken from the NED.\\\\ $^{*}$ Taken from NED\\\\ \\end{table} ",
        "conclusions": "We have presented and discussed the properties derived from the analysis of long slit spectra along several position angles of a sample of isolated active and non-active galaxies. A simple model has been used to trace overall rotation curve properties and obtain the parameters for Tully-Fisher (TF) analysis. The active and non-active spirals define very similar TF relations, populating the region occupied by isolated spirals of early types. In general, the rotation curves broadly agree with circular rotation, but the pure rotation model fails to account for large scale velocity distribution peculiarities, that can be attributed to the presence of a large scale bar, spiral-arm components, or small-scale central deviations indicative of the presence of non-circular motions and/or decoupled disks.  Note that the same kind of small scale central S-shaped deviations from the rotation curve model are found in both active (whether they have a single or double barred host, or are of Sy1 or Sy2 type) and non-active galaxies. When enough position angles are available, a better indication can be obtained of whether an inner bar (most cases) or a nuclear disk (see IC~184, NGC~3660, NGC~5728, ESO~139 and NGC~6012) are at the origin of such features.  Disk HII region metallicities, estimated from the [NII]/H$\\alpha$ ratios, are also very similar for the active and non-active sub-samples.  The analysis of the photometric properties of the DEGAS sample obtained from NIR imaging has already shown that active and non-active isolated spiral galaxies share the same disk, bulge and primary bar properties, and that they are indistinguishable at least down to spatial scales of about 300 pc (M\\'arquez et al. 1999, 2000).  These results are reinforced here, in the sense that both kinematical properties and disk metallicities lead to the same conclusion. Thus, whether or not they have an AGN, the spiral galaxies in the DEGAS sample are equivalent in all respects (a general discussion based on the DEGAS data will be presented in a forthcoming paper).  Therefore, the differences claimed to be related to the mechanisms driving AGN activity should be found in details related to circumnuclear regions still unresolved in our observations. But, interestingly, no significant morphological differences were found, even in the innermost few arcseconds, for a sample of 123 active and non-active galaxies imaged with the HST (Martini et al. 2003). Besides, the AGN activity can involve such different time scales from those of the dynamical processes under study, that any direct connection would be out of scope (see e.g. Hunt \\& Malkan 1999). Finally, and considering the growing evidence that all galaxies above a certain luminosity host central, massive compact objects (McLure \\& Dunlop 2002), our results are fully compatible with the idea that the AGN activity could take place recurrently in a given galaxy; in this way, the AGN itself would be seen when fuel is made available in the most internal regions. Note however that our sample is small and that differences between the active and non-active samples may be lost due to the poor statistics (in the recent works by the SLOAN project, for example, several thousand galaxies were required to statistically identify clear differences between active and non active galaxies (Kauffmann et al. 2003)."
    },
    "0311/astro-ph0311085_arXiv.txt": {
        "abstract": "Contact binaries can be used for distance determinations of stellar systems. They are easy to discover and identify and are very abundant among solar-type stars, particularly for $M_V > +3$. The period--luminosity--colour (PLC) relations have similar properties to those for pulsating stars and can currently predict individual values of $M_V$ to about $\\pm 0.25$ mag. ",
        "introduction": "\\label{intro}  Contact binary stars of the W~UMa-type (also known as EW) are unique objects \\citep{ruc93,web03}. The luminosity, produced almost exclusively in the more massive component is efficiently distributed (by poorly understood processes) through the common envelope so that the surface brightness is practically identical over the whole visible surface of the binary (the gravity variations over the surface introduce a very minor modification). Thus, the effective temperature is everywhere -- over the whole common surface -- the same, in spite of usually very different masses hiding inside the common envelope. Mass ratios are known to span the whole wide range, from almost unity (V753~Mon, $q = 0.97$, \\citealp{ddo3}), to very small values, as small as $q=0.066$ (SX~Crv, \\citealp{ddo5}). The more common  -- but more difficult to detect -- small mass-ratio systems are the best illustration to why the contact binaries are so unique. Simplifying: the primary component provides the luminosity, while both components provide the radiating area. Thus, it is practical to consider the primary mass, $M_1$, and the mass ratio, $q = M_2/M_1 \\le 1$, as the two main parameters. The third parameter, the orbital period, $P$, is related to the amount of the angular momentum in the system. The range of the primary masses is moderate and corresponds to Main Sequence spectral types from early A to early K and roughly maps into the orbital-period range of about 1.5 days to 0.22 days. Since the period distribution appears to continue the $1/P$ shape (as for other binaries) to about 0.45 days, and shows a mild bending down between 0.45 and 0.25 days, the most common are short period systems with periods within 0.25 -- 0.5 days, down to the sharp and currently unexplained cutoff at 0.22 days \\citep{ruc92,ste01}.  The contact binaries are known to exist only within the Main Sequence. Somewhat similar light curves of some giant or other more evolved binaries can be explained by semi-detached binaries or spotted stars. Very rare early-type contact systems on the upper Main Sequence, including a few currently known O-type contact systems, do not appear to be very different from the more typical solar-type W~UMa binaries, but will not be considered here.  With their special properties, the contact binaries of the W~UMa type form a distinct group of objects which are very easy to detect and identify due to: (1)~rather large amplitudes of light variations reaching one magnitude, (2)~short orbital periods, so that limited--duration monitoring is sufficient. Massive, systematic searches of the sky for stellar variability, such as the micro-lensing projects, have led to many detections of EW systems; e.g.\\ in the first catalogue of OGLE-I, as much as two thirds of the one thousand detected eclipsing binaries were EW systems \\citep{ruc97}. Many EW systems have been detected during the galactic cluster searches, both in open and in globular clusters (see \\citealp{ruc98,ruc00} for the references to the original works). ",
        "conclusions": "\\label{concl}  W~UMa-type (EW) contact binaries are very easy to find in massive stellar variability programs, even in short-duration ones, thanks to the short orbital periods and large amplitudes of photometric variations. The $M_V$ calibrations utilizing de-reddened colour indices $B-V$ or $V-I$ and the periods can predict individual values to about $\\pm 0.25$ mag, so that about a thousand systems are needed to reduce the group uncertainty to the level of $\\pm 0.01$. Such large numbers will be discovered in the nearby galaxies once the surveys pass the threshold of $M_V \\simeq 3 - 5$ which, for the Local Group typically corresponds to $V > 23 - 25$. The current calibrations involve de-reddened colour indices and thus remain sensitive to the reddening corrections. The widespread availability of the $K$-band data suggests development of calibrations based on the $V-K$ or $I-K$ indices; this is a matter of the current work. The existing $M_V$ calibrations do not seem to require a correction for metallicity, but existence of a small $[Fe/H]$ term is not excluded.  \\medskip  Thanks are due to Dr.\\ G.\\ Handler for illuminating e-mail discussion on the density of the Main Sequence pulsating stars."
    },
    "0311/hep-ph0311115_arXiv.txt": {
        "abstract": "s{ We consider the scenario of direct leptogenesis from the decays of heavy sneutrino which also plays a role of inflaton. We calculate its predictions for flavour-violating decays of charged leptons. We find that $\\mu \\to e \\gamma$ should occur close to the present experimental upper limits, as might also $\\tau \\to \\mu \\gamma$. }   \\vspace*{1cm} \\noindent ",
        "introduction": "Leptogenesis \\cite{fy} is currently the favorite scenario for explaining the observed baryon asymmetry  of the Universe. It is related to the light neutrino masses and mixings via the seesaw mechanism~\\cite{seesaw}. In supersymmetric models all the neutrino observables can also be related to the lepton flavour violating (LFV) decays of charged leptons which are induced by the heavy singlet neutrino Yukawa couplings via the renormalization effects. Implications of light neutrino masses on the heavy neutrino masses and couplings~\\cite{thomas}, and the relations between thermal leptogenesis and LFV~\\cite{sasha} have been extensively studied. The appeal and popularity of the thermal leptogenesis comes form its predictive power. On the other hand, the second completely predictive leptogenesis scenario - direct leptogenesis form the scalar field which dominates the Universe~\\cite{sn2,us2,us} - has got somewhat less attention.  The purpose of this talk is to review the most predictive scenario of direct leptogenesis in which the lightest heavy singlet sneutrino plays the role of inflaton~\\cite{sn1}. Inflation~\\cite{inf} has become the paradigm for early cosmology, particularly following the recent spectacular CMB data from the WMAP satellite~\\cite{wmap}, which strengthen the case made for inflation by earlier data, by measuring an almost scale-free spectrum of Gaussian adiabatic density fluctuations exhibiting power and polarization on super-horizon scales, just as predicted by simple field-theoretical models of inflation. Ever since inflation was proposed, it has been a puzzle how to integrate it with ideas in particle physics. For example, a naive GUT Higgs field would give excessive density perturbations, and no convincing concrete string-theoretical model has yet emerged. In this conceptual vacuum, models based on simple singlet scalar fields have held sway~\\cite{inf}. Here we assume the simplest scenario in which the lightest heavy singlet sneutrino drives inflation. This scenario constrains in interesting ways many of the 18 parameters of the minimal seesaw model for generating three non-zero light neutrino masses.   This minimal sneutrino inflationary scenario (i) yields a simple ${1 \\over 2} m^2 \\phi^2$ potential with no quartic terms, with (ii) masses $m$ lying naturally in the inflationary ballpark. The resulting (iii) spectral index $n_s$, (iv) the running of $n_s$ and (v)  the relative tensor strength $r$ are all compatible with the data from WMAP and other experiments~\\cite{wmap}. Moreover, fixing $m \\sim 2 \\times 10^{13}$~GeV as required by the observed density perturbations (vi) is compatible with a low reheating temperature of the Universe that evades the gravitino problem~\\cite{gr}, (vii) realizes leptogenesis~\\cite{us2,us} in a calculable and viable way, (viii) constrains neutrino model parameters, and (ix) makes testable predictions for the flavour-violating decays of charged leptons~\\cite{us}. ",
        "conclusions": "The main results of our scenario are the following. {\\it First}, reheating of the Universe is due to the neutrino Yukawa couplings, and therefore can be related to light neutrino masses and mixings. {\\it Secondly}, the lepton asymmetry is created in direct sneutrino-inflaton decays. There is only one parameter describing the efficiency of leptogenesis in this minimal sneutrino inflationary scenario in all leptogenesis regimes - the reheating temperature of the Universe - to which the other relevant parameters can be related. This should be compared with the general thermal leptogenesis case which has two additional independent parameters, namely the lightest heavy neutrino mass and width. {\\it Thirdly}, imposing the requirement of successful leptogenesis and correct amount of cold dark matter, we can predict branching ratios for $\\mu\\to e\\gamma$ and $\\tau\\to\\mu\\gamma$ in a very narrow band within about one order of magnitude of the present experimental bound."
    },
    "0311/astro-ph0311036_arXiv.txt": {
        "abstract": "We present a brief summary of observations of the transiting extrasolar planet, HD\\,209458\\,b, designed to detect the secondary eclipse.  We employ the method of `occultation spectroscopy', which searches in combined light (star and planet) for the disappearance and reappearance of weak infrared spectral features due to the planet as it passes behind the star and reappears. We have searched for a continuum peak near 2.2 \\micron\\ (defined by CO and H$_2$O absorption bands), as predicted by some models of the planetary atmosphere to be $\\sim 6 \\times 10^{-4}$ of the stellar flux, but no such peak is detected at a level of $\\sim 3 \\times 10^{-4}$ of the stellar flux. Our results represent the strongest limits on the infrared spectrum of the planet to date and carry significant implications for understanding the planetary atmosphere. ",
        "introduction": "Our observations cover two predicted secondary eclipse events, and we obtained 1036 individual spectra of the HD\\,209458 system using the SpeX instrument at the NASA IRTF in September 2001.  Our spectra extend from 1.9 to 4.2 \\micron\\ with a resolution ($\\lambda/\\Delta \\lambda$) of 1500.  A summary of the method of occultation spectrosocpy, as well as the details of the data analysis, can be found in Richardson, Deming, \\& Seager (2003). ",
        "conclusions": "The final difference spectrum, shown in Figure~1, represents the average candidate planetary spectrum, as calculated from the `in-eclipse' minus the `out-of-eclipse' spectra from the two nights during which a secondary eclipse was predicted to occur.  The average spectrum represents data from 550 individual spectra of HD\\,209458, as well as an equal number of spectra of the comparison star HD\\,210483. The comparison star was used in our analysis to remove variability due to changes in the terrestrial atmsophere and to normalize the data to the stellar flux density.  Also shown in Figure~1 is the baseline model for HD\\,204958\\,b calculated by Sudarsky, Burrows, \\& Hubeny (2003), which exhibits a peak near 2.2 \\micron.   A least-squares analysis indicates this peak is not present in the candidate planetary spectrum as derived from the data.  We believe this result has significant implications for the structure of the planetary atmosphere. In particular, some models that assume the stellar irradiation is re-radiated entirely on the sub-stellar hemisphere predict this flux peak, which is inconsistent with our observations. Several physical mechanisms can improve agreement with our observations, including the re-distribution of heat by global circulation, a nearly isothermal atmosphere, and/or the presence of a high cloud.  \\begin{figure}[!t] \\epsfysize=3.7in \\plotone{richardson_fig1.eps} \\caption{Final candidate planetary spectrum (symbols).  Solid line represents the baseline model for HD\\,209458\\,b as caluclated by Sudarsky, Burrows, \\& Hubeny (2003). The offset between the model and the data is not important, since we are comparing the shapes of the two spectra; thus, the mean has been subtracted from the data and model, respectively, for plotting purposes. \\label{fig:sw}} \\end{figure}"
    },
    "0311/astro-ph0311170_arXiv.txt": {
        "abstract": "Increasing samples of pulsating variable stars populating the classical instability strip from the horizontal branch to a few magnitudes brighter are being found in several Local Group galaxies, irrespective of the galaxy morphological type. We will review the observational scenario focusing in particular on the Anomalous Cepheids and related objects. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311493_arXiv.txt": {
        "abstract": "{Based on a large sample of disk and halo giant stars, for which accurate effective temperatures derived through the InfraRed Flux Method (IRFM) exist, a calibration of the temperature scale in the Vilnius, Geneva, $RI_\\mathrm{(C)}$ and DDO photometric systems is performed. We provide calibration formulae for the metallicity dependent $\\teff$ vs color relations as well as grids of intrinsic colors and compare them with other calibrations. Photometry, atmospheric parameters and reddening corrections for the stars of the sample have been updated with respect to the original sources in order to reduce the dispersion of the fits. Application of our results to Arcturus leads to an effective temperature in excellent agreement with the value derived from its angular diameter and integrated flux. The effects of gravity on these $\\teff$ vs color relations are also explored by taking into account our previous results for dwarf stars. ",
        "introduction": "In a previous paper (Mel\\'endez \\& Ram\\'{\\i}rez \\cite{melendez03}, Paper~I) we derived $\\teff:$ color $:\\feh$ relations for dwarf stars in the Vilnius, Geneva, $RI_\\mathrm{(C)}$ and DDO systems from the Alonso et~al. (\\cite{alonso96a}) sample. This time, we have performed a similar extension for giants by using the corresponding Alonso et~al. (\\cite{alonso99a}) sample. We have chosen this sample in order to maintain the homogeneity of the calibrations, the effective temperatures of all of the stars in both samples (dwarfs and giants) have been obtained in a single implementation of the InfraRed Flux Method (IRFM) by Alonso and coworkers.  In addition to their primary importance as fundamental relations, these empirical $\\teff$ calibrations are extremely useful for several purposes. Applications include chemical abundance studies (e.g. Smith et~al. \\cite{smith02}, Kraft \\& Ivans \\cite{kraft03}), the transformation of theoretical HR diagrams into their observational counterparts, i.e. color-magnitude planes (e.g. Girardi et~al. \\cite{girardi02}) and the test of synthetic spectra and colors (e.g. Bell \\cite{bell97}). When used along with other studies, combined results may have implications in galactic chemical evolution, population synthesis and cosmology.  A noteworthy feature of our work and that of Alonso et~al. (\\cite{alonso96b}, \\cite{alonso99b}) is the inclusion of Population II stars in the calibrations. They allow to extend the ranges of applicability of the formulae, which in turn become useful also for metal-poor stars. In order to achieve this goal, a large number of cluster giants has been included in the sample adopted for the present work. Even though for a given photometric system observations are usually available for only 2 or 3 clusters, the number of stars contributed by each of them is considerable and thus deserved to be included.  When comparing temperature scales of main sequence stars with those corresponding to giants, small but systematic differences arise. Theoretical calculations can easily take into account these variations with the $\\logg$ value. In empirical studies, however, it is only after a considerable number of stars has been studied and their properties properly averaged that gravity effects become clear. Detailed, careful inspection of gravity effects on colors may improve our understanding of the physics behind stellar spectra formation.  The present work is distributed as follows: Sect. \\ref{secc:...adopted} describes the data adopted, in Sect. \\ref{secc:...cal} we discuss some properties of the general calibration formula employed, perform the calibrations and apply them to Arcturus. Comparison of our work with previously published calibrations is presented in Sect. \\ref{secc:...comp} while intrinsic colors of giant stars and the effects of gravity on the temperature scale are discussed in Sect. \\ref{secc:...escala}. We finally summarize our results in Sect.~\\ref{secc:...conclusion}. ",
        "conclusions": "\\label{secc:...conclusion}  We have calibrated the effective temperature versus color relations for several color indices in 4 important photometric systems using reliable and recent $\\feh$ and $\\logg$ measurements to explore with improved accuracy the effects of chemical composition and surface gravity on the temperature scale.  In general, the present calibrations span the following ranges: 3800 K$<\\teff<$8000 K, $-3.0<\\feh<+0.5$. Ranges of applicability, however, are different and specific for each color calibration. For $(R-I)_\\mathrm{(C)}$ and $C(42-48)$, for example, these ranges are not so wide, specially in $\\teff$. We also provide specific ranges of applicability after every formula.  The standard deviation of the fits amount from 46 K for $(V-I)_\\mathrm{(C)}$ to 99 K for $(Y-V)$ with more than 130 stars in almost every calibration. Residuals of every fit were iteratively checked to reduce the dispersion and undesirable systematic effects.  Finally, from our formulae we have calculated the intrinsic colors of giant stars and have probed their consistency with empirical color-color diagrams, gravity effects on stellar spectra, $\\teff$ versus color relations for stars in clusters and evolutionary calculations. Our results for the main sequence were also explored in this last case."
    },
    "0311/astro-ph0311346_arXiv.txt": {
        "abstract": "{We present preliminary results from a set of near-IR integral field spectroscopic observations of the central, star-burst, regions of the barred spiral galaxy M83, obtained with CIRPASS on Gemini-S. We present maps in the \\PaB and \\Feii\\,1.257\\,$\\mu$m emission lines which appear surprisingly different.  We outline the procedure in which we will use \\PaB emission line strengths and measures of CO absorption to determine the relative and absolute ages of individual star-forming knots in the central kpc region of M83. ",
        "introduction": "The central star-burst in M83 (NGC 5236) has been studied photometrically by Harris \\etal (2001) using \\emph{HST/WFPC} images in the broad-band near-UV and optical, as well as narrow-band \\HA and \\HB to derive colours and line equivalent widths for 45 clusters. Despite the excellent spatial resolution of these observations, optical photometric analyses such as these suffer from patchy dust extinction, a reddening vector which parallels the evolutionary tracks in the two-colour diagram, and selection effects which tend to exclude the very youngest (\\emph{t}$<$5\\,Myr) clusters which have strong emission lines, but only a weak stellar continuum.  Additionally, it is not possible, on the basis of broad-band colours alone, to distinguish an instantaneous burst of star formation from a constant star formation rate.  To help overcome these effects, we use the Integral Field Unit (IFU) of the Cambridge Infra-Red Panoramic Survey Spectrograph (CIRPASS; Parry \\etal 2000)\\footnote{http://www.ast.cam.ac.uk/$\\sim$~optics/} to measure both the \\PaB equivalent width, as well as the CO(6,3) spectroscopic index over the entire star-forming arc and nucleus of M83. As demonstrated by Ryder \\etal (2001) for M100, the combination of two such diagnostics allows one to constrain not just the age, but also the burst duration, for each cluster.  \\begin{figure} \\centering \\includegraphics[width=8cm]{M83_ifumap_small_new.ps} \\caption{\\label{M83 IFU map} \\emph{HST} colour composite image from Harris \\etal 2001 (F300W, F547M and F814W) overlaid with the outline of the CIRPASS IFU indicating the three telescope pointings targeted. The optical galaxy nucleus is the bright spheroidal knot to the North-East of the centre of the image, at the end of the central IFU pointing.  A naive interpretation of this three colour image suggests that much of the star formation activity has taken place in the southern arc of the nucleus, with isolated pockets of activity elsewhere.  Using IFU spectroscopy observations one can survey the entire nuclear region, looking for the signatures of star-formation, past and present, without the bias inherent in previous long slit studies which have, by necessity, only targeted regions highlighted by broad band imaging techniques.} \\end{figure} ",
        "conclusions": "\\label{summary} Quantitative analysis is still in progress, but already we see evidence for a much more complex star formation history than that hinted at by the single long-slit near-IR spectrum of Thatte \\etal (2000).  Figure \\ref{spectra} shows two example spectra. The left hand spectrum is taken from a region visible in narrow band \\emph{HST}/NICMOS \\PaA imaging with extended emission but which does not show signs of active star formation loci in broad band WFPC2/NICMOS imaging.  This highlights one of the key goals of these observations, to perform an accurate census of the star formation in the nuclear region of M83, without bias towards bright continuum regions, such as those targeted by Harris \\etal (2001).  In contrast to the strong \\PaB emission, little CO absorption is seen in the \\emph{H} band spectrum of this region (not shown in Figure \\ref{spectra}), suggesting a young ($<5$\\,Myr) star formation region.  The right hand spectrum of Figure \\ref{spectra} is taken from the optical nucleus of M83.  Strong CO absorption is visible in this \\emph{H} band spectrum which, when taken in comparison with a very weak \\PaB emission line strength in the \\emph{J} band spectrum (not shown), is suggestive of an old stellar population ($>15$\\,Myr).  By comparing the equivalent width ratio for \\PaB and CO(6,3) we will determine the age of star-formation across the nucleus of M83.  Figure \\ref{PaB FeII} compares the distribution of \\PaB line flux, with that of the \\Feii\\,1.644\\,$\\mu$m.  The former provides a measure of the massive star formation rate within the last $\\sim$5\\,Myr, while the latter is thought to be a tracer of shocks from supernova explosions within the past 10$^4$\\,years (Alonso-Hererro \\etal 2003). The \\Feii emission is rather clumpy, and confined mainly to the periphery of the more diffuse \\PaB emission, suggesting that massive star formation is being propagated into regions of undisturbed gas by the passage of the supernova blasts themselves.  \\begin{figure*} \\centering \\includegraphics[width=5cm]{PaB_Fint_BW.ps} \\hspace{1cm} \\includegraphics[width=5cm]{FeII-H_Fint_BW.ps} \\caption{\\label{PaB FeII}A comparison of the \\PaB line flux (left), with that of the \\Feii\\,1.257\\,$\\mu$m line (right) is made above. See discussion in Section \\ref{summary}.} \\end{figure*}"
    },
    "0311/astro-ph0311301_arXiv.txt": {
        "abstract": "The relativistic outflows from gamma-ray bursts are now thought to be narrowly collimated into jets. After correcting for this jet geometry there is a remarkable constancy of both the energy radiated by the burst and the kinetic energy carried by the outflow. Gamma-ray bursts are still the most luminous explosions in the Universe, but they release energies that are comparable to supernovae. The diversity of cosmic explosions appears to be governed by the fraction of energy that is coupled to ultra-relativistic ejecta. ",
        "introduction": "\\label{sec:jets}  In hindsight, since jets are a natural outcome of most high energy phenomena, they should have been expected in gamma-ray bursts (GRBs). While there had been some early indications that relativistic outflows from GRBs might not be isotropic \\cite{wkf98,rho99b}, the real impetus for invoking jets came from the ``energy crisis'' brought on by the spectacular GRB\\,990123 \\cite{kdo+99}. On a timescale of order 80 s, the isotropic energy released in gamma-rays E$_{iso}(\\gamma)$ from this burst approached the rest mass energy of a neutron star! If GRB outflows were not isotropic but instead were collimated into jets with an opening angle $\\theta_j$ then they would only radiate into a fraction $f_b =(1-\\cos \\theta_j) \\cong \\theta_j^2/2$ of the celestial sphere \\cite{rho99,sph99}. Thus the true gamma-ray energy released $E_\\gamma$ would be smaller than $E_{\\rm iso}(\\gamma)$ by the same factor, i.e.~$E_\\gamma=f_b\\times E_{\\rm iso}(\\gamma)$.  The characteristic signature of a jet-like outflow at optical and X-ray wavelengths is an achromatic break at a time $t_j$ in the power-law decay of the light curves (Fig.~\\ref{fig:jet}), in which the exponent $\\alpha$ (defined by $F_\\nu\\propto t^{\\alpha}$) steepens by $\\Delta\\alpha\\sim{1}$ \\cite{sph99}. At radio wavelengths, in which the emission is initially produced by electrons radiating below the peak of synchrotron spectrum $\\nu_m$, a jet break is expected initially to produce only a shallow power-law decay (e.g. t$^{-1/3}$ to t$^0$) of the light curve until a time when $\\nu_m$ passes through the radio band. Thereafter, if the expansion remains relativistic, the power-law decay index will be the same as the optical and X-ray (i.e.~$\\alpha\\sim -2$). GRB afterglow measurements of the degree of polarization and the polarization angle can provide a powerful diagnostic for jets \\cite{gl99,sar99}, but the observations to date have been ambiguous.  \\begin{figure}[t] \\centering \\includegraphics[scale=0.35]{jetbreak.eps} \\vspace{-35pt} \\caption{\\it The UBVRIJK band lightcurves of GRB\\,010222. The fits to the data are made with a broken temporal power law function added to a constant host galaxy component. The break in the light curves occurs at $t_j$=0.93 d. From \\cite{grb+03}. \\label{fig:jet}} \\vspace{-25pt} \\end{figure}  The origin of this sharp break at $t_j$ is due to two effects. The first is a purely geometric transition that occurs when the Lorentz factor $\\Gamma$ drops below $\\theta_j^{-1}$, and the observer begins to see the edge of the jet \\cite{rho99}. A second effect that may become important after $t_j$ is a dynamical transition in which the jet outflow switches from a pure radial outflow to a laterally spreading jet component. With some basic understanding of the jet dynamics, a measurement of $t_j$ and redshift $z$ yields an estimate of $\\theta_j$ and hence $E_\\gamma$. Readers who wish to learn more about the practical difficulties in estimating $t_j$ and possible systematic effects should review \\cite{fks+01,bfs01,bfk03} and references therein. ",
        "conclusions": "For the last decade or more, progress in our understanding of GRB progenitors has marched in lock step with the convergence of the GRB energy scale.  By 1992 we knew from their peak flux and sky distribution that GRBs were consistent with either a cosmological population or a hitherto unseen population of sources in the halo of our Galaxy. The resulting eight orders of uncertainty in the energy scale led to a plethora of possible progenitor models \\cite{nem94}. With the demonstration that GRBs were cosmological \\cite{mdk+97}, this initial uncertainty shrank but the isotropic gamma-ray energy still spanned three orders of magnitude -- reaching $\\sim{10^{54}}$ erg -- creating a real strain on plausible burst models.  The solution, which seems obvious only in hindsight, was to recognize that the relativistic outflows from GRBs are jet-like, not isotropic (\\S\\ref{sec:jets}).  Consequently it appears that there is a near standard energy yield of $\\sim$10$^{51}$ erg in both the radiated and the kinetic energy of the GRB explosion (\\S\\ref{sec:standard}-\\ref{sec:kinetic}). This result is all the more remarkable when it is combined with strong evidence supporting the view that long-duration GRBs are the result of the core collapse of a massive star, aka ``collapsar'' \\cite{smg+03,hsm+03}.  It is convenient to think of a collapsar explosion as having two distinct equal-energy components: quasi-isotropic ejecta expanding with velocity $\\Gamma\\simeq 1.005$ producing the familiar radioactively powered supernova light curves, and a highly collimated flow with $\\Gamma\\simeq 100$ powered by synchrotron emission from electrons accelerated in the relativistic shock. This picture, however, appears to be too simple, since there is now growing evidence that the Lorentz factor of the relativistic ejecta can vary as a function of viewing angle and from one event to the next (\\S\\ref{sec:universe}). Energy appears to be the one quantity that is (roughly) conserved throughout (\\S\\ref{sec:diverse}).  By understanding the GRB energy budget we are gaining fundamental insight into the inner workings of the central engine. As we move forward studying new sub-classes of cosmic explosions, such as X-ray flashes, events with on-going energy injection \\cite{fyk+03}, and bursts with distinct high-energy components \\cite{gdk+03}, we would be well-served to keep calorimetry as an essential tool.  \\vspace{1pc} {\\bf Acknowledgements.} I would like to thank Ed van den Heuvel, who made it possible for me to attend this meeting and give a final farewell to {\\it BeppoSAX} - The Little Satellite That Could. Much of the work summarized in this paper was done by my collaborators, including PhD students Edo Berger, Paul Price, and Sarah Yost."
    },
    "0311/astro-ph0311137_arXiv.txt": {
        "abstract": "VLBI observations have shown that the parsec-jet of 3C 345 is formed by several components, ejected from the core with superluminal velocities and travelling along bent trajectories on the plane of the sky. We interpret the differences in velocity and position angle among the different features at formation time as the result of parsec-scale precession of the relativistic jet and calculate the aperture angle of the precession cone, the angle between the cone axis and the line of sight and the Lorentz factor associated with the jet bulk motion. We assumed a precession period of 10.1 yr, which is one of the {\\it B}-band light curve long-term periods reported in the literature. We propose that boosting of the underlying jet emission, which is time-dependent due to precession, is responsible for this long-term optical variability. Jet precession with periods of several years can be produced in super-massive black hole binary systems, when the secondary black hole is in an orbit non-coplanar with the primary accretion disk, inducing torques in the inner parts of the disk. Assuming that this mechanism is responsible for the jet precession in 3C\\,345, we estimate upper and lower limits for the masses of the two black holes, as well as their mean separation. We found a correlation between the formation of jet components and the occurrence of strong optical flares, as well as a very strong anti-correlation between the intensity of these flares and the time required for the components to reach the maximum flux density at radio frequencies. ",
        "introduction": "The quasar 3C\\,345 (z=0.5928; \\citealt{mar96}), also known as 1641+399 or 4C\\,39.48, was optically identified by \\citet{goki65}, and it is now known to be the nucleus of an elliptical E3 galaxy \\citep{kir99}. Strong variability (typical of an OVV object) is found in its optical light curve, and periodicities of about 5 and 11 yr have been reported \\citep{web88,kidg89,zha98,zha00}, although it has also been suggested that this variability may originate from non-linear and non-stationary stochastic processes \\citep{vio91}.  The radio flux density has been monitored at several frequencies \\citep{wal91,all96,ter98}. The continuum spectrum is flat up to 10 GHz and gets steeper towards higher frequencies, with spectral index ranging from -0.9 (between 10$^{11.9}$ and 10$^{13.2}$ Hz) to -1.4 (between 10$^{13.2}$ and 10$^{15.4}$ Hz) \\citep{bre86}. Outbursts have been detected from centimeter to infrared wavelengths \\citep{bre86,ste96}, besides time variable linear polarization at radio and optical frequencies \\citep{most81,bre86,bro94}. In the optical regime, the degree of polarization decreases monotonically toward shorter wavelengths \\citep{sit85,smi86,mea88,deDie94} and in the X-ray range, the source is weak and possibly variable \\citep{halp82,maki89,wowi90}.  VLA observations showed a faint halo around a bright core and an extended kiloparsec-jet \\citep{kol89}. At parsec-scales, 3C\\,345 exhibits a stationary core \\citep{bar86,tan90} and a jet with superluminal components, which travel apparently along curved paths with variable velocities \\citep{unw83,bir86,zen95,ros00}. These characteristics had been interpreted as a result of helical motion of the components along the jet, either in a pure phenomenological model \\citep{qia96} or as a consequence of HD or MHD instabilities (e.g., \\citealt{koch85,came86,hard87,qia91,qia92,cakr92,ste95,hard00}).  An alternative approach is to assume that the helicoidal appearance of the jet is a consequence of precession. In this case, and considering that the bulk velocity is very large, each plasma element in the jet would move in an almost straight trajectory, defined by the jet direction when this element was ejected. From the observational point of view, superluminal components can be interpreted as these plasma elements, even if they are the results of shock waves propagating along the jet; their velocity should not be very different from the bulk velocity (e.g., \\citealt{mage85}) and, at least close to their formation epoch, they should reflect the precessing jet direction. Curved trajectories, as found by \\citet{zen95} and \\citet{loba96} in 3C\\,345, can be due to several reasons: jet instabilities (e.g., \\citealt{hard00}), influence of the external environment (one of the mechanisms proposed by \\citealt{wal87} to explain the bending of the jet in 3C\\,120), or even the superposition with other components formed at different epochs and moving with different superluminal velocities. In any case, the parameters of the precession model can be completely determined by the superluminal velocities and position angles of the jet components close to their formation epoch.  The parsec-jet precession scenario has already been applied to 3C\\,279, 3C\\,273 and OJ\\,287 \\citep{abca98,abro99,abra00}, where the velocity variations among components were attributed to differences in the angle between their trajectories and the line of sight. Besides, it was found that the periodic outbursts seen in the optical light curves of 3C\\,279 and OJ\\,287 could be due to changes in the boosting parameter \\citep{abca98,abra00}. Although this model explained well the behavior of these objects, the precession parameters were obtained from a limited set of data and should probably be revised to include new and more precise data \\citep{hom01,weh01}. We will show in this paper that the variations in the ejection angles and velocities of the jet components of 3C\\,345 can also be interpreted as the result of parsec-jet precession. Furthermore, we show that the presence of long-term optical variability can be attributed to variable boosting of the non-thermal radiation from the underlying jet.  In Sect.2, we describe the precession model used in this work. In Sect.3, we present the precession model parameters for 3C\\,345, as well as the influence of precession in the absolute core position shifts due to opacity effects. In Sect.4, we discuss, in the framework of the precession model, other observed characteristics, such as the relation between optical flares and the origin of superluminal components, {\\it B}-band long-term periodic variability, polarimetric observations and dispersion in the apparent angular size of the components. Furthermore, orbital parameters of a possible super-massive black hole binary system in the nucleus of 3C\\,345 are estimated based on the precession model, with additional constrains from optical continuum and spectroscopic data. Finally, conclusions are presented in Sect.5. ",
        "conclusions": "In this work, we used data available in the literature involving the parsec-scale radio structure, which comprise almost twenty years of monitoring of 3C\\,345. We assumed that all superluminal features have quasi-ballistic motions in the inner regions ($r<1$ mas).  We showed that, close to the core, the trajectories of the superluminal components can be explained quite well if the jet is precessing with a period of 10.1 yr, one of the {\\it B}-band periodicities reported by \\citet{zha98}. As a consequence of this, we found a correlation between time variation of the Doppler factor and the optical light curve, indicating that probably the boosted optical emission from the underlying jet is responsible for the {\\it B}-band long-term periodic variability. An upper limit of 7 nJy for the intrinsic flux density of the underlying jet has been estimated in this work. Lower limits for the Doppler factor calculated previously from X-ray observations \\citep{unw83,unw97} were used as constrains in our model.  In this work, we also analyzed the influence of jet precession in the observed position of the optically thick core at different frequencies \\citep{blko79,loba98}. We found that precession introduces a time modulation in the shift corrections, as well as in the direction in which the corrections should be applied. However, core position shifts calculated in this work did not exceed 0.31 mas.  Considering that jet precession in 3C\\,345 is driven by a secondary super-massive black hole in a non-coplanar orbit around the primary accretion disk and using the total mass of the two black holes derived from the H$\\beta$ line-width and optical continuum luminosity \\citep{gu01}, we estimated the masses of the primary and secondary black holes as $4\\times 10^9 \\;{\\mathrm M}_{\\sun}\\leq M_{\\mathrm{p}} \\leq 5\\times 10^9$ M$_{\\sun}$ and $3\\times 10^9 \\;{\\mathrm M}_{\\sun}\\leq M_{\\mathrm{s}} \\leq 4\\times 10^9$ M$_{\\sun}$, respectively. In addition, we found that the distance  $r_{\\mathrm ps}$ between them should be $5.5\\times 10^{16}\\; {\\mathrm {cm}} \\leq r_{\\mathrm {ps}}  \\leq 7.3\\times 10^{16}$ cm and their orbital period $2.5\\; {\\mathrm {yr}} \\leq P_{\\mathrm {ps}} \\leq 3.8$ yr.  The observed dispersion in the apparent sizes of jet components could be attributed to jet precession and interactions between jet and environment matter. It is corroborated by sudden variations in the radio flux density of the knots when there are changes in their propagation direction. Furthermore, those interactions seem to be responsible for gentle bending seen at distances larger than 1 mas, which could explain also the difference between jet orientation at parsec- and kiloparsec-scales.  We were able to associate the emergence of jet components to the occurrence of flares in the optical band stronger than about 1.5 mJy. We also found an unexpected relation between the flare flux density at optical wavelengths and the elapsed time between the epochs of superluminal component ejection and occurrence of maximum flux density at radio frequencies, which led us to the conclusion that these optical flares are originated in the accretion disk and not in the shock waves seen as superluminal features.  Finally, the complex polarimetric structure of 3C\\,345 in parsec-scales \\citep{bro94,lep95,tayl98,ros00} may be explained by our precession model, when the superposition of jet components with different position angles is considered."
    },
    "0311/gr-qc0311036_arXiv.txt": {
        "abstract": "Spacetimes with horizons show a resemblance to thermodynamic systems and it is possible to associate  the notions of temperature and entropy with them. Several aspects of this connection are reviewed in a manner appropriate for  broad readership. The approach uses two essential principles: (a) the physical theories  must be formulated for each observer entirely in terms of variables any given observer can access and (b) consistent formulation of quantum field theory requires analytic continuation to the complex plane. These two principles, when used together in spacetimes with horizons, are powerful enough to provide several results in a unified manner.  Since  spacetimes with horizons have a generic behaviour under  analytic continuation, standard results of quantum field theory in curved spacetimes with horizons can be obtained directly (Sections III to VII). The requirements (a) and (b) also  put  strong constraints on the action principle describing the gravity and, in fact, one can obtain the Einstein-Hilbert action from the thermodynamic considerations (Section VIII). The review emphasises the thermodynamic aspects of horizons, which could be obtained from general principles and is expected to remain valid,  independent of the microscopic description (`statistical mechanics') of horizons.    \\bigskip  \\emph{We combine probabilities by multiplying, but we combine the actions ... by adding;  ...... since the logarithm of the probability is necessarily negative, we may identify action  provisionally with minus the logarithm of the statistical probability of the state...}  \\hfill Eddington (1920)\\ \\cite{Eddington:1920}  \\smallskip  \\noindent\\emph{The mathematicians can go beyond this Schwarzschild radius and get inside, but I would maintain that this inside region is not physical space, .... and should not be taken into account in any physical theory.}  \\hfill Dirac (1962)\\ \\cite{Dirac:1962} ",
        "introduction": "The analytic properties of spacetime manifold in the complex plane directly lead to the association of a temperature with a generic class of horizons. In Section \\ref{hawkrad} we also saw that there exist quantum states  in which a flux of thermal radiation will flow away from the horizon if the horizon forms asymptotically. Given these results, it is natural to enquire whether one can attribute other thermodynamic variables, in particular entropy, to the horizons. We shall now discuss several aspects of this important --- and not yet completely resolved --- issue.  The simplest and  best understood situation arises in the case of a Schwarzschild black hole formed due to gravitational collapse of matter. In this case, one can rigorously demonstrate the flow of thermal flux of radiation to asymptotic infinity at late times, which  can be collected by observers located in (near) flat spacetimes at $r\\to\\infty$. Given a temperature and a change in energy, one can invoke classical thermodynamics to define the {\\it change in} the entropy  via $dS=dE/T(E)$. Integrating this equation will lead to the function $S(E)$ except for an additive constant which needs to be determined from additional considerations. In the Schwarzschild spacetime, which is asymptotically flat, it is  possible to associate an energy $E=M$ with the black-hole. Though the  calculation  was done in a metric with a fixed value of energy $E=M$, it seems reasonable to assume that --- as the energy flows to infinity at late times ---  the mass of the black hole will decrease.  {\\it If} we make this assumption,  {\\it then} one can integrate the equation $dS=dM/T(M)$ to obtain the entropy of the black-hole to be \\begin{equation} S=4\\pi M^2=\\frac{1}{4}\\left(\\frac{\\mathcal{A}_H}{L_P^2}\\right) \\label{entropyarea} \\end{equation} where $\\mathcal{A}_H=4\\pi (2M)^2$ is the area of the event horizon and $L_P=(G\\hbar/c^3)^{1/2}$ is the Planck length. This integration  constant is   fixed by  the additional assumption that $S$ should vanish when $M=0$.\\footnote{One may think that this assumption is eminently reasonable since the Schwarzschild metric reduces to the Lorentzian metric when $M\\to 0$.  But note that in the same limit of $M\\to 0$, the temperature of the black-hole diverges. Treated as a limit of Schwarzschild spacetime, normal flat spacetime has infinite --- rather than zero --- temperature.}   The fact that entropy of the Schwarzschild black hole is \\emph{proportional} to the horizon area was conjectured \\cite{Bekenstein:1972tm,Bekenstein:1973ur,Bekenstein:1974ax} even before it was known that black holes have a temperature. The above analysis fixes the proportionality constant between area and entropy to be (1/4) in Planck units. It is  also obvious that the entropy is purely a quantum mechanical effect and diverges in the limit of $\\hbar \\to 0$. Nevertheless, even in the {\\it classical} processes involving black holes, the horizon area does act in a manner similar to entropy. For example, when two black holes coalesce and settles down to a final steady state (if they do),  the sum of the areas of horizons does not decrease. Similarly, in some simple processes in which energy is dumped into the black hole, one can prove an analogue for first law of thermodynamics involving the combination $TdS$. While both $T$ and $S$ depend on $\\hbar$ the combination $TdS$ is independent of $\\hbar$ and can be described in terms of classical physics.  The next natural question is whether the entropy defined by Eq.~(\\ref{entropyarea}) is the same as ``usual entropy\". If so, one should be able to show that for any processes involving matter and black holes, we must have $d(S_{BH}+S_{matter})/dt\\geq 0$ which goes under the name generalised second law (GSL).  One simple example in which the area (and thus the entropy) of the black hole \\emph{decreases} is the Hawking evaporation; but the GSL holds since the thermal radiation produced in the process has entropy. It is generally believed that GSL always holds though a completely general proof is difficult to obtain. Several thought experiments, when analysed properly, uphold this law \\cite{Unruh:1982ic} and a proof is possible under certain restricted assumptions regarding the initial state \\cite{Frolov:1993fy}. All these suggest that the area of the black hole corresponds to an entropy which is  same  as the ``usual entropy''.  In the case of normal matter, entropy can be provided a statistical interpretation  as the logarithm of the number of available microstates that are consistent with the macroscopic parameters which are held fixed. That is,  $S(E)$ is related to the degrees of freedom (or phase volume) $g(E)$ by $S(E)=\\ln g(E)$. Maximisation of the phase volume for systems which can exchange energy will then lead to equality of the quantity $T(E)\\equiv (\\partial S/\\partial E)^{-1}$ for the systems. It is usual to identify this variable as the thermodynamic temperature. (This definition works even for self-gravitating systems in microcanonical ensemble; see eg., \\cite{Padmanabhan:1990gm}.)     Assuming that the entropy of the black hole should have a similar interpretation, one is led to the conclusion that the density of states for a black hole of energy $E$ should vary as \\begin{equation} g(E) \\propto \\exp \\left(\\frac{1}{4}\\frac{\\mathcal{A}_H}{L_P^2}\\right) = \\exp\\left[ 4\\pi \\left(\\frac{E}{E_P}\\right)^2\\right] \\label{gbhofe} \\end{equation} Such a growth implies, among other things, that the Laplace transform of $g(E)$ does not exist so that no canonical partition function can be defined (without some regularization).   This brings us to the next question: What are the microscopic states by counting which one can obtain the result in Eq.~(\\ref{gbhofe}) ? That is, what are the degrees of freedom (or the missing information content)  which lead to this entropy ?  There are two features that need to be stressed regarding  these questions. First, classically, the black hole is determined by its charge, mass and angular momentum and hence has ``no hair\" (for a review, see e.g., \\cite{Bekenstein:1998aw}). Therefore,  the degrees of freedom which could presumably account for all the information contained in the initial (pre-collapse) configuration cannot be classical.  Second, the question is intimately related to what happens to the matter that collapses to form the black hole. If the matter is ``disappearing\" in a singularity then the information content of the matter can also ``disappear\". But since singularities are unacceptable in physically correct theories, we expect the classical singularity to be replaced by some more sophisticated description in the correct theory. Until we know what this description is, it is impossible to answer in a convincing manner what happens to the information and entropy which is thrown into the black hole or  was contained in the initial pre collapse state.  In spite of this fact, several attempts have been made in the literature to understand features related to entropy of black holes. A statistical mechanics derivation of entropy was originally attempted in \\cite{Gerlach:1976hbh}; the entropy  has been interpreted as the logarithm of: (a)  the number of ways in which black hole might have been formed \\cite{Bekenstein:1973ur,Hawking:1976de}; (b) the number of internal black hole states consistent with a single black hole  exterior  \\cite{Bekenstein:1973ur,Bekenstein:1975tw,Hawking:1976de} and (c)  the number of horizon quantum states \\cite{Wheelerbk:1990,'tHooft:1990fr,Susskind:1993if}. There are also  other approaches which are more mathematical ---  like the ones based on Noether charge \\cite{Wald:1993nt,Jacobson:1994vj,Visser:1993nu}, deficit angle related to conical singularity \\cite{Banados:1994qp,Susskind:1993ws}, entanglement entropy  \\cite{Dowker:1994fi,Israel:1976ur,Callan:1994py} and thermo field theory and related approaches \\cite{Frolov:1998vs,Frolov:1997up,Frolov:1997aj}. Analog models for black holes which might have some relevance to this question are discussed in \\cite{Unruh:1995je,Jacobson:1993hn,Brout:1995wp}. There are also attempts to compute the entropy using the Euclidean gravitational  action and canonical partition function \\cite{Gibbons:1977ue}. However, since we know that canonical partition function does not exist for this system these calculations require a non trivial procedure for their interpretation. In fact, once the answer is known, it seems fairly easy to come up with very imaginative derivations of the result. We shall comment on a few of them.   To begin with, the thermal radiation surrounding the black hole has an entropy which one can attempt to compute. It is fairly easy to see that this entropy will proportional to the horizon area but will diverge quadratically. We saw in Section \\ref{ftdim} that, near the horizon,  the field becomes free and solutions are simple plane waves. It is the  existence of such a continuum of wave modes which leads to infinite phase volume for the system. More formally, the number of modes  $n(E)$ for a scalar field $\\phi$ with vanishing boundary conditions at two radii $r=R$ and $r=L$ is given by \\begin{equation} n(E) = \\frac{2}{ 3\\pi} \\int^L_R \\frac{r^2 dr}{ \\left( 1 - 2M/r\\right)^2}\\left[ E^2 - \\left( 1 -\\frac{2M}{ r}\\right)  m^2\\right]^{3/2} \\label{eqn:qnofe} \\end{equation} in the WKB limit. [This result is essentially the same as the one contained in Eq.(\\ref{hjmodes}); see \\cite{'tHooft:1985re,Padmanabhan:1986rs}.] This expression diverges as $R\\to 2M$ showing that a scalar field propagating  in a black hole spacetime has infinite phase volume. The corresponding entropy computed using the standard relations: \\begin{equation} S = \\beta \\left[ \\frac{\\partial}{ \\partial \\beta} - 1 \\right] F; \\qquad F = - \\int_0^\\infty dE \\frac{n(E)}{e^{\\beta E} - 1 }, \\end{equation} is quadratically divergent:  $S = (\\mathcal{A}_H/l^2)$ with $l \\to 0$. The divergences described above occur around any infinite redshift surface and is a geometric (covariant) phenomenon.  The same result can also be obtained from what is known as ``entanglement entropy'' arising from the quantum correlations which exist across the horizon. We saw in Section \\ref{thermalden} that if the field configuration inside the horizon is  traced over in the vacuum functional of the theory, then one obtains a density matrix $\\rho$ for the field configuration outside [and vice versa]. The entropy $ S=-Tr(\\rho \\ln \\rho)$ is usually called the entanglement entropy. This is essentially the same as the previous calculation and, of course, $S$ diverges quadratically on the horizon \\cite{Frolov:1993ym,Zurek:1985gd,Dowker:1994fi,Israel:1976ur,Callan:1994py}. Much of this can be done without actually bringing in gravity anywhere; all that is required is a spherical region inside which the field configurations are traced out \\cite{Bombelli:1986rw,Srednicki:1993im}. Physically, however,  it does not seem reasonable to integrate over all modes without any cut off in these calculations. By cutting off the mode at $l \\approx L_P$ one can obtain the ``correct'' result but in the absence of a more fundamental argument for regularising the modes, this result is not of much significance.  The cut off can be introduced in a more sophisticated manner by changing the dispersion relation near Planck energy scales but again there are different prescriptions that are available \\cite{Padmanabhan:1998jp,Padmanabhan:1998vr,Unruh:1994zw,Corley:1996ar} and none of them are really convincing.   The entropy computed using any non gravitational degrees of freedom will scale in proportion with the number, $g_s$, of the species of fields which exist in nature. This  does not cause a (separate) problem since one can re-absorb it in the renormalisation of gravitational constant $G$. In any calculation of effective action for a quantum field  in curved spacetime, one will obtain a term proportional to $R$ with a quadratically divergent coefficient. This coefficient is absorbed by renormalising the gravitational constant and this procedure will also take care of $g_s$.   In conventional description, entropy is also associated with the amount of missing information and one is tempted to claim that information is missing inside the horizon of black hole thereby leading to the existence of non zero entropy.  It is important to distinguish carefully the separate roles played by the horizon and singularity in this case; let us, for a moment, ignore the black hole singularity inside the horizon. Then the fact that a horizon hides information is no different from the fact that the information contained in a room is missing to those who refuse to enter the room. The observers at $(r, \\theta, \\phi) =$ constant in the Schwarzschild metric do not venture into the horizon and hence cannot access the information at $r<2M$. Observers who are comoving with the collapsing matter, or even those who plunge into the horizon later on, can access (at least part of) the  information which is not available to the standard Schwarzschild observers at $r>2M$. In this respect, there is no difference between a Rindler observer in flat spacetime and a $(r, \\theta, \\phi) =$ constant observer in the Schwarzschild spacetime (see figure \\ref{hyperfig}) and it is irrelevant what happens to the information content of matter which has collapsed inside the event horizon. The information missing due to a horizon is observer dependent since --- as we have stressed before --- the horizon is defined with respect to a congruence of timelike curves (``family of observers''). If one links the black hole entropy with the missing information then the entropy  too will become observer dependent.  In the examples which we have discussed in the previous Sections, the  thermal density matrix and temperature of the horizon indeed arose from the integration of modes which are hidden by the horizon. In the case of a black hole formed by collapse, there is a well defined, non singular, description of physics in the asymptotic past. As the system evolves, the asymptotic future is made of two parts. One part is outside the horizon and the other part (classically) hits a singularity  inside the horizon. The initial quantum state has now evolved to a correlated state with one component inside the horizon and one outside. If we trace over the states inside the horizon, the outside will be described by a density matrix. None of this is more mystifying than the usual phenomenon in quantum theory of starting with a correlated quantum state of a system with two parts (say, two electrons each having two spin states), spatially separating the two components and tracing over one of them in describing the (spatially) localised measurements made on the other. There is no real information loss paradox in such systems.  In the case of the black hole there is an additional complication that the matter collapses to a singularity classically taking the information along with it. In this description, some of the information will be missing even to those observers who dare to plunge inside the event horizon. But, as we said before, this issue cannot be addressed until the problem of final singularity is solved. We have no idea what happens to the matter (or the wave modes of the quantum field) near the singularity and as such it is not possible to do a book keeping on the entropy content of matter inside the black hole. As the black hole evaporates, its mass will decrease but such a semi-classical calculation cannot be trusted at late stages.  There is considerable discussion in the literature on the ``information loss problem'' related to this issue. Broadly speaking, this problem arises because the evolution seems to take a pure quantum state to a state with significant amount of thermal radiation. It is, however, difficult even to attempt to tackle this problem properly since physics loses its predictive power at a singularity. One cannot meaningfully ask what happens to the information encoded in the matter variables which collapses to a singularity. So to tackle this question, one needs to know the correct theory which replaces the singularity. If for example, a Planck size remnant is formed inside the event horizon then one needs to ask whether a freely falling observer can retrieve most of the information at late stages from this remnant. [Some of the discussions in the literature also mixes up results obtained in different domains with qualitative arguments for the concurrent validity. For example, one key assumption in the information loss paradox is that the initial state is pure. It is far from obvious that in a fully quantum gravitational context a pure state will collapse to form a black hole \\cite{Myers:1997qi}].   One immediate consequence, of linking entropy of horizons to the information hidden by them, is that all horizons must be attributed an entropy proportional to its area, with respect to the observers who perceive this horizon. More precisely,  given  a congruence of timelike curves in a spacetime we define the horizon to be the boundary of the union of the causal  pasts of the congruence. Assuming this is non trivial surface, observers moving on this congruence will attribute a constant entropy per unit area $(1/4L_P^2)$ to this horizon. (We shall  say more about this in Section \\ref{thermoroute}.) The analysis given in Section \\ref{complexplane} [see Eq.~(\\ref{pofefromsofe})] shows that whenever a system crosses the horizon with energy $E$, the probability picks up a Boltzmann factor related to the entropy. In the case of a spherically symmetric horizon, one can imagine thin shells of matter carrying some amount of energy being emitted by the horizon. This will lead to the correct identification of entropy for the horizon. It is conceivable that similar effect occurs whenever a packet of energy crosses the horizon even though it will be difficult to estimate its effect on the surface gravity of the horizon. Naive attempts to  compute the corresponding results for other geometries will not work and a careful formalism using the entropy density of horizons --- which is currently not available --- will be required.     \\subsection{Black hole entropy in quantum gravity models}\\label{qgbh}  The above discussion highlights the fact that any model for quantum gravity, which has something to say about the black hole singularity, will also make definite predictions about the entropy of the black hole. There has been considerable amount of work in this direction based on different candidate models for quantum gravity. We will summarise some aspects of this briefly. [More extensive discussions as well as references to original literature can be found in the reviews \\cite{Rovelli:1998yv,Das:2000su}]. The central idea in any of these approaches is to introduce microscopic degrees of freedom so that one can attribute  large number of microscopic states to solutions that could be taken to represent a classical black hole configuration. By counting these microscopic states, if one can show that $g(E)\\propto\\exp(\\alpha E^2)$, then it is usually accepted as an explanation of black hole entropy.  In standard string theory this is done as follows: There are certain special states in string theory, called BPS states \\cite{Peet:2000hn}, that contain electric and magnetic charges which are equal to their mass. Classical supergravity has these states as classical solutions, among which are the extremal black holes with electric charge equal to the mass (in geometric units). These solutions can be expressed as a Reissner-Nordstrom metric with both the roots of $g_{00}=0$ coinciding: obviously, the surface gravity at the horizon, proportional to $g_{00}'(r_H)$ vanishes though the horizon has finite area. Thus these black holes, classically, have zero temperature but finite entropy. Now, for certain compactification schemes in string theory (with $d=3,4,5$ flat directions), in the limit of $G\\to 0$, there exist BPS states which have the same mass, charge and angular momentum of an extremal black hole in $d$ dimensions. One can explicitly count the number of such states in the appropriate limit and one finds that the result gives the exponential of black hole entropy with correct numerical factors \\cite{Strominger:1996sh,Das:2000su,Breckenridge:1996sn}. This is done in the weak coupling limit and a duality between strong coupling and weak coupling limits \\cite{Sen:1994yi,Vafa:1994tf,Sen:1999mg} is used to argue that the same result will arise in the strong gravity regime. Further, if one perturbs the state slightly away from the BPS limit, to get a {\\it near} extremal black hole and construct a thermal ensemble, one obtains the standard Hawking radiation from the corresponding near extremal black hole \\cite{Breckenridge:1996sn}.  While these results are intriguing, there are several issues which are still open: First, the extremality or near extremality was used crucially in obtaining these results. We do not know how to address the entropy of a normal Schwarzschild black hole which is far away from the extremality condition. Second, in spite of significant effort, we do not still have a clear idea of how to handle the classical singularity or issues related to the information loss paradox. This is disappointing since one might have hoped that these problems are closely related. Finally, the result is very specific to black holes. One does not get any insight into the structure of other horizons, especially De Sitter horizon, which does not fit the string theory structure in a natural manner.  The second approach  in which some success related to black hole entropy is claimed, is in the  loop quantum gravity (LQG). While string theory tries to incorporate all interactions in a unified manner, loop quantum gravity \\cite{Rovelli:1998gg,Thiemann:2001yy} has the limited goal of providing a canonically quantised version of Einstein gravity. One key result which emerges from this programme is a quantisation law for the areas. The variables used in this approach are like a gauge field $A^i_a$ and the Wilson lines associated with them. The open Wilson lines carry a quantum number $J_i$ with them and the area quantisation law can be expressed in the form: $\\mathcal{A}_H=8\\pi G\\gamma\\sum \\sqrt{J_i(J_i+1)}$ where $J_i$ are spins defined on the links $i$ of a spin network and $\\gamma$ is free parameter called Barbero-Immirizi parameter. The $J_i$ take half-integral values if the gauge group used in the theory is SU(2) and take integral values if the gauge group is SO(3). These quantum numbers, $J_i$, which live on the links that intersect a given area, become undetermined if the area refers to a horizon. Using this, one can count the number of microscopic configurations contributing to a given horizon area and estimate the entropy. One gets the correct numerical factor (only) if $\\gamma=\\ln m/2\\pi \\sqrt{2}$ where $m=2$ or $m=3$ depending on whether the gauge group SU(2) or SO(3) is  used in the theory \\cite{Krasnov:1998wc,Rovelli:1996dv,Ashtekar:2000eq,Ashtekar:1998yu}.  Again there are several unresolved issues. To begin with, it is not clear how exactly the black hole solution arises in this approach since it has been never easy to arrive at the low energy limit of gravity in LQG. Second, the answer depends on the Immirizi parameter $\\gamma$ which needs to be adjusted to get the correct answer, if we know the correct answer from elsewhere. Even then, there is an ambiguity as to whether one should have SU(2) with  $\\gamma=\\ln 2/2\\pi \\sqrt{2}$ or SO(3) with $\\gamma=\\ln 3/2\\pi \\sqrt{2}$. The SU(2) was the preferred choice for a long time, based on its close association with fermions which one would like to incorporate in the theory. However, recently there has been some rethinking on this issue due to the following consideration: For a classical black hole, one can define a class of solutions to wave equations called quasi normal modes [see e.g.,\\cite{Corichi:2002ty,Kokkotas:1999bd,Berti:2003jh,Cardoso:2003vt}]. These modes have discrete frequencies which are complex, given by \\begin{equation} \\omega_n=i\\frac{n+(1/2)}{4M}+\\frac{\\ln(3)}{8\\pi M}+\\mathcal{O}(n^{-1/2}) \\end{equation} The $\\ln(3)$  in the above equation is not negotiable \\cite{Padmanabhan:2003fx,Choudhury:2003wd,Motl:2003cd,Motl:2002hd}. If one chooses SO(3) as the gauge group, then one can connect up the frequency of quanta emitted by a black hole when the area changes by one quantum in LQG with the quasi normal mode frequency \\cite{Dreyer:2002vy,Hod:1998vk}. It is not clear whether this is a coincidence or of some significance. Third, most of the details of the LQG are probably not relevant to the computation of the entropy. Suppose we have {\\it any} formalism of quantum gravity in which there is a minimum quantum for length or area, of the order of $L_P^2$. Then, the horizon area $\\mathcal{A}_H$ can be divided into $n=(\\mathcal{A}_H/c_1L_P^2)$ patches where $c_1$ is a numerical factor. If each patch has $k$ degrees of freedom (due to the existence of a surface field), then the total number of microscopic states are $k^n$ and the resulting entropy is $S=n\\ln k=(4\\ln k/c_1)(\\mathcal{A}_H/4L_P^2)$ which will give the standard result if we choose $(4\\ln k/c_1)=1$. The essential ingredients are only discreteness of the area and existence of certain degrees of freedom in each one of the patches.  Another key issue in counting the degrees of freedom is related to the effective dimensionality. If we repeat the above argument with the volume $V\\propto M^3$ of the black hole then one will get an entropy proportional to the volume rather than area. It is clear that, near a horizon, only a region of length $L_P$ across the horizon contributes the microstates so that in the expression $(V/L_P^3)$, the relevant $V$ is $M^2 L_P$ rather than $M^3$. It is possible to interpret this as due to the entanglements of modes across the horizon over a length scale of $L_P$, which --- in turn --- induces a nonlocal coupling between the modes on the surface of the horizon. Such a field will have one particle excitations, which have the same density of states as black hole \\cite{Padmanabhan:1998jp,Padmanabhan:1998vr}. While this is suggestive of why we get the area scaling rather than volume scaling, a complete understanding is lacking. The area scaling of entropy has also led to different proposals of holographic bounds [see, e.g. \\cite{Bousso:2002ju}] which is beyond the scope of this review. ",
        "conclusions": "We shall now take stock  of the results discussed in this review from a broader perspective and will attempt to provide an overall picture.   \\begin{figure}[htbp] \\begin{center} \\includegraphics[scale=0.6]{TP_fig04.eps} \\caption{Summary of the logical structure of the approach adopted in this review} \\label{flowchart} \\end{center} \\end{figure}     Combining the principles of  quantum theory with special relativity (and Lorentz invariance) required a fairly drastic change in the description of physical systems.  Similarly, it is natural for new issues to arise when we take the next step of combining quantum theory with the concept of general covariance or when we attempt to do quantum field theory in a curved background spacetime. However, one would have naively expected these issues to be \\emph{kinematical} in the sense that they are independent of the field equations or the action for gravity. Our discussion shows that there is a strong link between the kinematical aspects and the dynamics of gravity because of the structure of  classical general relativity. While it may be convenient to distinguish between the kinematical aspects (discussed in Sections \\ref{horizonsr} to \\ref{exptab}) and the dynamical aspects (discussed in Sections \\ref{entropyhorizon} -- \\ref{thermoroute}), each by itself can only give a partial picture. The overall structure of the theory should allow a seamless transition across these two aspects.   In this review this was attempted by  (i) noting that one needs to use the Euclidean sector  to incorporate the new ingredients which arise when special relativity is combined with quantum mechanics and (ii) using the fact that when quantum theory is formulated in the Euclidean sector, a  unique structure emerges in the presence of horizons.  Using a congruence of  timelike curves  to define a horizon, one finds that it is possible to incorporate the kinematical effects of (at least static) horizons in a general manner and associate the notion of temperature with the horizons. This is achieved by using a coordinate system in which the spacetime region hidden by a horizon is mapped to a single point in the Euclidean sector and constructing an effective manifold  for a family of   observers by removing this point.  The resulting non trivial topology leads to the standard results of quantum field theory in curved spacetimes with horizons.     The importance of the above point of view lies in its ability to provide a deeper relationship between gravity and thermodynamics, as shown in Section \\ref{thermoroute}.  If one accepts the idea --- that the physical theory for a class of observers should be formulated in an effective manifold in which the region inaccessible to those observers is removed --- then one is led to enquire what it implies for the dynamics of gravity. Using the fact that the horizon is the common element between the inaccessible and accessible regions, it is possible to argue that  the action functional for gravity must contain (i) a well defined surface term and (ii) a bulk term which is related to the surface term in a specific manner. Hence,  this point of view allows one to determine the action functional for  gravity from thermodynamic considerations. What is more, it links the kinematical and dynamical  aspects   of the theory in an interesting manner.  This approach is very similar in spirit to that of renormalisation group theory (RGT) in particle physics.  When  an experimenter does not have information about the model at scales $k>\\Lambda$, say, in \\emph{momentum space}, the RGT allows one to use an effective low energy theory with the coupling constants readjusted to incorporate the missing information. This, in turn, puts restrictions on the nature of the theory as well as the ``running\" of the coupling constants. Similarly,  when a given family of observers has limited information because they are blocked by a horizon (in \\emph{real} space rather than momentum space) it is necessary to add certain boundary terms in the action functional in order to provide a consistent description.   Just as the RGT contains nontrivial information about the low energy sector of the theory, our approach allows us to determine the form of the action in the long wavelength limit of gravity. As far as the loss of information due to a horizon is concerned, there is no need to distinguish between the uniformly accelerated observers in flat spacetime and, say,  the observers located permanently at $r>2M$ in the Schwarzschild spacetime.    There are some new insights that arise in this approach which are worth exploring further. \\begin{itemize} \\item Einstein's  equations for  gravity can be obtained from a  variety of action functionals, any two of which differ by a surface term. In the case of Einstein-Hilbert action, the surface term is related in a very specific manner to the bulk term. (See e.g., Eq.~(\\ref{lageh}); it is rather intriguing that this relation has  not been explored in the literature before.)  \\emph{This relation is so striking that it demands an explanation} which is indeed provided by the  thermodynamic paradigm described in Section \\ref{thermoroute}.  \\item The approach makes gravity ``holographic'' in a specific sense of the word. The Einstein-Hilbert Lagrangian has the structure $L_{EH} = L_1 + L_2 $ where $L_1 \\sim (\\partial g)^2, L_2  \\sim \\partial^2 g$. Along any world line, one can choose a coordinate system such that  $(\\partial g)^2 \\to 0$ suggesting that the  dynamics of the theory is actually contained in the $L_2 \\sim   \\partial^2 g$ term which leads to the surface term in the action. We saw in Section \\ref{grfresults} that one could determine the bulk term from the surface term under certain assumptions. This fact, that the structure of the surface term in an action determines the theory, provides a possible interpretation of holographic principle (which is somewhat different from the conventional interpretation of the term).  \\item The approach supports the paradigm that the spacetime is similar to the continuum limit of a solid that is obtained when one averages over the underlying   microscopic degrees of freedom \\cite{Sakharov:1968pk}. As described in Section \\ref{thermoid}, this strongly indicates the possibility that gravity is intrinsically quantum mechanical at all scales just as solids cannot exist in the strict $\\hbar \\to 0 $ limit. Just as  the bulk properties of solids can be described without reference to the underlying atomic structure, much of classical and semi classical gravity (including the entropy of black holes) will be independent of the underlying description of the microscopic degrees of freedom. Clear signs of this independence emerges from the study of Einstein-Hilbert action which contains sufficient structure to lead to many of the results involving the horizon thermodynamics. Hence any microscopic description of gravity which leads to Einstein-Hilbert action as the long wavelength limit will also incorporate much of horizon physics.  \\end{itemize}    \\bigskip"
    },
    "0311/astro-ph0311071_arXiv.txt": {
        "abstract": "We model the spectral energy distribution (SED) from the mid-infrared to submillimeter of the ring-like disk of $\\hra$, the dustiest A-type star. We consider dust made either of coagulated but otherwise unaltered protostellar interstellar grains, or grains that are highly-processed in a protostellar/protoplanetary nebula with silicate dust annealed and carbon dust oxidized. Both types of dust are successful in reproducing the observed SED, provided that the grains are highly fluffy, with a vacuum volume fraction of $\\sim 90\\%$. We find no evidence for the existence of a hot ``zodiacal dust'' component a few AU from the star, which was suggested by previous workers to account for the 10$\\mum$-wavelength emission. ",
        "introduction": "} $\\hra$ is a nearby (distance to the Earth $d\\approx 67\\pm 3\\pc$) young main-sequence (MS) star (age $\\approx 8\\pm 3\\myr$) of spectral type A0\\,V (effective temperature $\\Teff\\approx 9500\\K$). The $\\hra$ dust disk has recently aroused considerable interest because (1) it has the largest fractional infrared (IR) luminosity relative to the total stellar luminosity ($L_{\\rm IR}/L_{\\star}\\approx 5\\times 10^{-3}$; $L_\\star \\approx 21\\,L_\\odot$) among the $\\sim 1500$ A-type MS stars in the {\\it Bright Star Catalogue} (Jura 1991); (2) unlike the majority of circumstellar disks, the $\\hra$ disk displays a ring-like structure peaking at $\\sim 70\\AU$ from the central star and abruptly truncated both interior and outside with a width of $\\simlt 17\\AU$ (Schneider et al.\\ 1999); (3) its young-age nature places it at a somewhat transitional stage between massive gaseous protostellar disks around young pre-MS T-Tauri and Herbig Ae/Be stars ($\\sim 1\\myr$) and much evolved and tenuous debris disks around MS ``vega-type'' stars ($\\sim 100\\myr$).  The physical and chemical properties of the dust in the $\\hra$ disk, which play a significant role in understanding the creation, dynamical evolution and structural properties of the disk, are poorly constrained: (1) Based on the lifetime of dust against loss by Poynting-Robertson drag, Jura et al.\\ (1993) suggested that the $\\hra$ dust is larger than $10\\mum$ in radius and the dust is a remnant of the protostellar nebula; (2) From an analysis of radiation pressure vs. gravitational attraction, Jura et al.\\ (1995) found that the minimum radius of grains orbiting $\\hra$ is about $3\\mum$ and argued that these grains have undergone measurable coalescence; (3) In order to explain the $\\simgt 40\\AU$ inner hole (Jura et al.\\ 1995) and the $\\approx 110\\K$ black-body approximation for the 10--100$\\mum$ dust emission, Jura et al.\\ (1998) proposed that the $\\hra$ grains are largely composed of ice particles with a typical radius of $\\sim 100\\mum$ which originate from a protocometary cloud; (4) In order to explain the entire 12.5$\\mum$ emission and the 20.8$\\mum$ residual emission unaccounted for by their simple model, Koerner et al.\\ (1998) proposed the existence of a tenuous dust component at a distance of 3--6$\\AU$ from the star which is heated to about 200--300$\\K$, similar to the zodiacal dust in our own solar system; (5) From the red reflectance of the disk, Schneider et al.\\ (1999) argued that the mean dust size must be larger than several microns and the dust is circumstellar debris rather than interstellar in origin; (6) Based on a detailed modelling of the spectral energy distribution (SED) of the $\\hra$ disk, Augereau et al.\\ (1999) proposed a two-component model consisting of a cold annulus peaking at $70\\AU$ from the star (made of interstellar dust-type grains ranging from $10\\mum$ to a few meters) and a hot population at 9--10$\\AU$ from the star (made of comet-like grains of radii $\\approx 450\\mum$); (7) Using the flux ratio at 10.8$\\mum$ and 18.2$\\mum$ and assuming ``astronomical silicates'', Telesco et al.\\ (2000) inferred a ``characteristic'' diameter of $\\approx 2-3\\mum$ and argued that these grains are unlikely to be primordial; instead, they are probably products of recent collisions of large bodies.  It is the purpose of this work to constrain the dust properties (size, composition, and morphology) of the $\\hra$ disk. In general, in a dusty system, dust spatial distribution and dust sizes cannot be uniquely determined by the SED {\\it alone}. Given that the distribution of the dust in the $\\hra$ disk is well constrained by the near-IR imaging of scattered starlight (Schneider et al.\\ 1999) and, to a less degree, by the mid-IR imaging of dust thermal emission (Jayawardhana et al.\\ 1998; Koerner et al.\\ 1998; Telesco et al.\\ 2000), an attempt is therefore made in this work to infer the $\\hra$ dust properties by modelling the full spectral energy distribution from the mid-IR to the submillimeter wavelengths.  Lacking a priori knowledge of the composition of the dust in the $\\hra$ disk, we consider two extreme dust types: (1) ``cold-coagulation'' dust -- dust in protoplanetary disks is formed through cold aggregation of unaltered interstellar materials -- this is the case in the outer parts of the disk where interstellar dust originating from the parent molecular cloud (out of which the disk forms) is not likely to have significantly changed its composition during passage through the weak shock front when the disk is first created, except volatile ice mantles around the refractory cores may partly sublimate, but the recondensation of volatiles occurs efficiently behind the shock (see Beckwith, Henning, \\& Nakagawa 2000 for a review); (2) ``hot-nebula'' dust -- dust has undergone significant destruction and modification in protoplanetary accretion disks through annealing of amorphous silicates and oxidation of carbonaceous dust in the warm inner regions of the accretion disk (Gail 2001); turbulent radial mixing (Gail 2001) and/or outflows driven by ``X-winds'' (Shu, Shang, \\& Lee 1996) then carries these heavily processed dust into cold regions of the disk where it is mixed with freshly accreted material from the parent molecular cloud; as an extreme, we consider a dust model in which the carbonaceous dust component has been fully destroyed by oxidation so that only crystalline silicate dust remains. In reality, the dust in protoplanetary disks would undoubtedly be intermediate between heavily processed dust and aggregates of unaltered interstellar dust. Since planetesimals and cometesimals are formed by coagulation of such dust aggregates, the dust generated by collisions of planetesimals and cometesimals should resemble the original dust from which they are built up.  We first discuss in \\S\\ref{sec:method} the general constraints on the $\\hra$ dust and disk properties. We then model the observed SED in \\S\\ref{sec:cold} and \\S\\ref{sec:hot} in terms of the ``cold-aggregation'' model of unaltered interstellar dust and the ``hot-nebula'' model with all amorphous silicate dust annealed and all carbonaceous dust oxidized, respectively. We find that, for both models, (1) highly porous grains with a porosity $P\\approx 0.90$ provide excellent fits to the entire SED; (2) there appears to be no need for a hot ``zodiacal cloud'' dust component (\\S\\ref{sec:zodi}) which was previously suggested by Koerner et al.\\ (1998) and Augereau et al.\\ (1999). Therefore, the presence of either type of particles in the disk is possible. We show in \\S\\ref{sec:pwl} that, although good fits to the SED can be achieved by models with a single power-law dust spatial distribution, they are apparently in conflict with the imaging observations of scattered starlight and dust thermal emission. In \\S\\ref{sec:sirtf} we calculate the dust IR intensities integrated over the SIRTF/IRAC and MIPS bands predicted for our best-fitting models. The major conclusions are summarized in \\S\\ref{sec:summary}. ",
        "conclusions": "} We have modelled the mid-infrared to submillimeter spectral energy distribution of the ring-like dust disk around the dustiest $\\sim 8\\myr$-old A-type star $\\hra$. We start with two extreme dust types: one formed through cold-coagulation of unaltered protostellar interstellar grains, and one formed through aggregation of grains highly-processed in the protostellar nebula with silicate dust annealed and carbon dust oxidized. We adopt a Gaussian-type dust spatial distribution with a peak at $70\\AU$ from the central star and a FWHM of $15\\AU$ as inferred from images of scattered light and dust thermal emission. We take a simple power-law for the dust size distribution ($dn/da\\propto a^{-\\alpha}$ with $\\alpha \\approx 2.8-2.9$) in the range of $1\\mum \\simlt a \\simlt 1\\cm$. Our principal results are: \\begin{enumerate} \\item It is shown that both types of dust are successful in reproducing the observed SED, provided that the dust generated by collisions of planetesimals and cometesimals is highly fluffy, with a vacuum volume fraction of $\\sim 90\\%$ (\\S\\ref{sec:model}). The fact that the dust in the $\\hra$ disk must be somewhere intermediate between these two types implies that our models are robust. Future high resolution spectroscopy of the $\\hra$ disk would allow us to infer the degree to which the dust has been processed (\\S\\ref{sec:robust}).  \\item Although models with a single power-law spatial distribution are also able to reproduce the observed SED, the derived dust distribution is both unphysical and inconsistent with the imaging observations of scattered light and dust mid-IR thermal emission (\\S\\ref{sec:pwl}).  \\item Our models show no evidence for the existence of a hot ``zodiacal dust'' component suggested to lie in a radial distance of a few AU from the star by previous workers to account for the emission at $\\lambda \\simlt 12\\mum$ (\\S\\ref{sec:zodi}). Our upper limit on the total mass of this component is at least 10 times smaller than what would be obtained if the entire $\\lambda \\simlt 12\\mum$ emission is attributed to this hot dust.  \\item Grains smaller than $\\sim 10-20\\mum$ will be radiatively expelled from the disk; grains at $r=70\\AU$ in the range of $\\sim 10-20\\mum \\simlt a \\simlt 100-200\\mum$ will also be removed from the disk due to the Poynting-Robertson inward spiralling drag (at a {\\it closer} distance from the star, {\\it larger} grains will be removed). Collisions of planetesimals/cometesimals must continuously replenish the dust in the disk at a rate of $\\approx 8.3\\times 10^{-7}\\,m_\\oplus/{\\rm yr}$ (\\S\\ref{sec:rppr}). \\end{enumerate}"
    },
    "0311/astro-ph0311592_arXiv.txt": {
        "abstract": "{We consider the simultaneous deblurring of a set of noisy images whose point spread functions are different but known and spatially invariant, and the noise is Gaussian. Currently available iterative algorithms that are typically used for this type of problem are computationally expensive, which makes their application for very large images impractical. We present a simple extension of a classical least-squares (LS) method where the multi-image deblurring is efficiently reduced to a computationally efficient single-image deblurring. In particular, we show that it is possible to remarkably improve the ill-conditioning of the LS problem by means of stable operations on the corresponding normal equations, which in turn speed up the convergence rate of the iterative algorithms. The performance  and limitations of the method are analyzed through numerical simulations. Its connection with a column weighted least-squares approach is also considered in an appendix. ",
        "introduction": "An important problem in image processing is the simultaneous deblurring of a set of observed images (of a fixed object), where each image is degraded by a different spatially invariant point spread function (PSF). For example, data obtained by the Large Binocular Telescope (LBT). This instrument consists of two $8.4 {\\rm m}$ mirrors on a common mount with a spacing of $14.4 {\\rm m}$ between the centers \\citep{ang98}. For a given orientation of the telescope, the diffraction-limited resolution along the center-to-center baseline is equivalent to that of a $22.8 {\\rm m}$ mirror, while the resolution along the perpendicular direction is that of a single $8 {\\rm m}$ mirror. A possible way to obtain an image with an improved and uniform spatial resolution is to simultaneously deconvolve the images taken with different orientations of the telescope. Another example is the multi-frequency observations of the Cosmic Microwave Background (CMB) obtained from satellites such as {\\it PLANCK}. In fact, although some other physical components are present at the microwave frequencies, the images taken at high galactic latitudes are expected to be almost entirely  dominated by the CMB contribution in a large range of observing frequencies \\citep[e.g., see][]{sto02}. Since the PSFs corresponding to the images obtained at the various frequencies can be quite dissimilar, simultaneous deblurring can be useful for improving the extraction of CMB information from the data.  Although a vast literature on the classical problem of single-image deblurring is available, much less has been published on multiple-image deblurring. An excellent general review is \\citet{ber00b}.  In addition, \\citet{teg99} provides some material more directly related to CMB studies. But one of the more serious limitations of the methods currently available is that, although able to provide satisfactory results, they are quite expensive in terms of computational requirements. Thus, these methods may be quite difficult to use for the very large images ($\\approx 10^6 - 10^8$ pixels) expected from LBT and {\\it PLANCK} observations. It is clear that more efficient methods must be developed; the aim of this paper is to provide one such approach.  In Sect.~\\ref{sec:formalization} we formalize the problem and propose an efficient solution in Sect.~\\ref{sec:efficient}. Some of its possible limitations are considered in Sect.~\\ref{sec:problems}. In Sect.~\\ref{sec:experiments} we study its performance through numerical experiments, and also compare it to that of other methods currently available in the literature. Finally, Sect.~\\ref{sec:conclusions} closes with some final comments and conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We have considered the problem of simultaneously deblurring a set of images of a fixed object blurred by different space-invariant PSFs. Although currently available methods seem to provide good results, they are comptutationally expensive. We have developed a method, based on a LS approach, to efficiently transform a multi-image deblurring problem into a single-image one. This approach provides substantial savings in computational requirements and can be implemented using standard currently available numerical algorithms. These conclusions are confirmed by our numerical experiments.  But despite these encouraging results, some questions are still open. In particular, regularization parameter selection methods (such as GCV for the Tikhonov approach) have to be extended to deal with the correlated noise component of the mean image. This is especially important for the deblurring of random fields where the statistical properties of the field are more important than a particular realization.  Another question is the generalization of the proposed method to non-Gaussian noise. For example, for Poisson noise good results have been obtained with the {maximum-likelihood} estimate \\begin{align} \\fb &= \\arg\\max \\,\\sum_{j=1}^p \\sum_{m,n=0}^{N-1} \\nonumber \\\\ & [ \\,g_j(m,n) \\ln (\\Ab_j \\fb)(m,n) - (\\Ab_j \\, \\fb)(m,n) ] , \\end{align} \\citep{cor00, ber00a}."
    },
    "0311/hep-ph0311014_arXiv.txt": {
        "abstract": "We discuss the impact of different solar neutrino data on the spin-flavor-precession (SFP) mechanism of neutrino conversion. We find that, although detailed solar rates and spectra allow the SFP solution as a sub-leading effect, the recent KamLAND constraint on the solar antineutrino flux places stronger constraints to this mechanism. Moreover, we show that for the case of random magnetic fields inside the Sun, one obtains a more stringent constraint on the neutrino magnetic moment down to the level of $\\mu_\\nu \\lsim {\\rm few} \\times 10^{-12}\\mu_B$, similar to bounds obtained from star cooling. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311467_arXiv.txt": {
        "abstract": "Aberration and retardation of the radiation within a pulsar's magnetosphere enable a single corotating absorbing region in the outer magnetosphere to give rise to a `double notch' in the pulsar's observed radio profile. This effect requires that a pulsar has not only absorbing but also emitting regions at heights which are a significant proportion of the light cylinder radius.  The known presence of double notches in the weak secondary emission of three nearby pulsars suggests that this emission comes from open field lines well above their poles and that an obscuring region extending over no more than 5 percent of the light cylinder radius is located at or near the light cylinder. ",
        "introduction": "\\subsection{Background} It is known that at least three pulsars have mysterious double `notches' in their integrated profiles (McLaughlin $\\&$ Rankin 2003, henceforth MR). All are also known to be sources of pulsed X-rays (Wang $\\&$ Halpern 1997, Zavlin et al 2002), and PSR 1929+10 and PSR 0950+08 also have optical identifications (De Luca et al 2003, Mignani et al 2002, Zharikov et al 2003). Their high-energy spectra have variously been interpreted as polar cap thermal emission sustained by particle backflow or as non-thermal cyclotron emission, with the latter being a particularly strong indicator of outer magnetospheric activity. Since one of the three is a millisecond pulsar and the other two are slow normal pulsars they would seem to have little in common - even their magnetic fields computed at the light cylinder differ by several orders of magnitude. The only intrinsic factor all three have in common is a similar \"acceleration parameter\", a measure of the potential difference available to accelerate particles within the magnetosphere.  The three pulsars are nearby and, possibly because of this, have detectable weak emission though a large fraction of their periods. The double notches are embedded in the weak emission and MR report that the phase of their centroid is frequency-independent, although their width and separation seems to be frequency-dependent and the feature may not be seen at all at some frequencies. These are the salient points on which an explanation must be built.  One piece of obvious and valuable information to be gleaned from the observations is the physical scale of the phenomenon. This can be inferred from the notch separation, which MR report to be around $10^{o}$ for the normal pulsars and $3.3^{o}$ for the millisecond pulsars. One might seek an explanation in terms of unusual surface composition or magnetic field structure, which would then imply a lateral scale very small compared to the polar cap dimensions. But we would then require that such 'freak' conditions existed on a similar scale in three very different pulsars. In this paper we offer an explanation based on the effects of aberration and retardation, leading to physical scales based on light-travel times. The notches may thus corespond to lateral or longitudinal features on a scale up to one sixth of a light cylinder radius (or ~2000km in the case of PSR 0950+08 and PSR 1929+10 and just 20km for PSR J0437-4715), a scale far greater than inferred height variations for polar cap emission of a single frequency (Gupta $\\&$ Gangadhara 2003, Dyks et al 2003a). This suggests that the emission in and around the notches comes from regions higher in the magnetosphere - an idea which is important since, if true, it would demonstrate support for a long-held suspicion (e.g. Arons 1979, Hankins $\\&$ Cordes 1981, Gil 1985) that not only new-born but also older normal pulsars are capable of emitting at radio frequencies in the outer magnetosphere.  To account for notches in such emission it is natural to turn to the possibility of an obscuring region lying between the emitting region and the observer. Luo $\\&$ Melrose (2001) and Fussell et al (2003) have suggested that regions of emission suppression can arise in a pulsar magnetosphere as a result of cyclotron absorption. Clearly any obscuring feature must lie within the magnetosphere and corotate with it or the notches would not appear at a fixed phase of the pulse window. The size of the obscuration estimated from the width of the individual notches comes out as 600km in the case of PSR 0950+08, 400km in the case of PSR 1929+10 and a tiny 5km for PSR J0437-4715. Such scales are not impossible to envisage: perhaps a region of locally-enhanced particle density at high energies located either at the light cylinder or on the boundary of the corotating zone defined by the last fieldlines to close within the light cylinder. But how do we explain $\\it{two}$ notches?  It is most unlikely that $\\it{two}$ obscuring regions lie adjacent to one another at any location in the magnetosphere at a fixed or frequency-dependent separation from one another. It would not correspond to any known feature from theory or observation. So instead we turn to one of the unique properties of a pulsar magnetosphere to provide us with a possible explanation. Pulsar magnetospheres differ from conventional magnetospheres in that the particles move tightly bound to magnetic fieldlines and emit their radiation in a extraordinarily narrow beam.  This implies that most regions of the magnetosphere are invisible to the observer at any given instant. We will show in this paper that under these circumstances it is possible for a single obscuring region to produce a double notch effect, and in such a way that observed features of the notches enable us to deduce a simple relation between the location of the emitting and obscuring regions.  Section 2 summarises the basic principles employed in the model, and Sections 3 and 4 give the detailed geometric calculations for the cases of radial and dipolar fieldlines respectively. Section 5 applies the results to MR's observations of the three candidate pulsars, and Section 6 summarises the implications and conclusions for future high-energy observations. ",
        "conclusions": "The principal conclusions of this paper can be summarised as follows: \\subsection{Geometric} (1) Through the combined effects of aberration and retardation a single corotating absorber can cause a system of rigidly rotating and unidirectionally emitting sources to exhibit a double notch feature to an external observer. This effect occurs whether the sources emit radially or at a moderate angle to the radial direction.  (2) The absorbed emission comes from a height approximately half that of the absorbing region. The model predicts that the width and appearance of the double notches will be frequency-dependent but their centroid invariant.  (3) The emitting region in the magnetosphere within which the double notch is found must be elongated and thin relative to the scale of the absorbing region, and be quasi-radial in structure.  \\subsection{Physical} (4) The presence of double notches in three very different nearby pulsars demonstrates that old pulsars are capable of emitting at radio frequencies high in their magnetospheres. There is evidence in all three that the emission is on open fieldlines close to the dipole axis, and may thus possibly be linked to conal or core behaviour at the polar caps (Gil 1985). This interactive feature could be tested in future observations.  (5) Future high-energy observations of the magnetospheres of these pulsars may correlate with intensity-variations of their weak radio emission. The interpretations given here suggest that a pulsar's polar cap and its outer magnetosphere interact, and this may guide future theoretical models (eg Harding $\\&$ Muslimov 2003, Rankin $\\&$ Wright 2003, Wright 2003).  (6)The emission in which notches are found is characteristically weak, highly linearly polarised and extended over long stretches of pulse longitude, suggesting that modelling of a $\\it{relativistic}$ single vector sweep might fix the inclination of these pulsars (eg Dyks et al 2003b).  (7) Some pulsars - maybe all - have absorbing sites located near the light cylinder no bigger than $5\\%$ of a light cylinder radius, possibly created by a relatively high density of particles achieving high Lorentz factors at the light cylinder and causing cyclotron absorption (Luo $\\&$ Melrose 2001, Fussell et al 2003)."
    },
    "0311/astro-ph0311184_arXiv.txt": {
        "abstract": "The NOAO Survey for Ionization in Neutral Gas Galaxies (SINGG) is the largest star formation survey of an \\HI\\ selected sample.  Since the selection is made without regard to optical morphology, it is not biased toward or against ``interesting'' types of galaxies; thus SINGG is an ideal sample for studying galaxy demographics.  Of a sample of 90 extra-galactic sources observed in photometric conditions, all are detected in \\Halpha.  This indicates that dormant galaxies, those containing an appreciable ISM but no star formation, are at best rare. We have made first pass morphological surveys for starbursts, as judged by \\Halpha\\ surface brightness, and outflows as judged by extra-planar \\Halpha.  We find that about 15\\%\\ of the sources contain starbursts, with little dependence on the neutral hydrogen mass \\MHI.  Nearly one half of a sample $\\sim 35$ edge-on galaxies show evidence for extra-planar \\Halpha\\ having a scale size of 0.5 Kpc or larger, while nearly one quarter have extra-planar \\Halpha\\ features 1.0 Kpc in size or larger. There is a hint that high \\MHI\\ systems preferentially have displaced outflows (chimneys, or fountains) while central outflows (galactic winds) preferentially occur in low \\MHI\\ systems.  However, a larger sample (e.g.\\ the full SINGG survey) is needed to confirm this trend. ",
        "introduction": "Here I present preliminary morphological results from {\\em SINGG\\/} - the Survey for Ionization in Neutral Gas Galaxies.  SINGG is an approved and ongoing survey project of the National Optical Astronomy Observatoryical (NOAO).  It is a follow-up survey to HIPASS - the \\HI\\ Parkes All Sky Survey (Staveley-Smith \\etal\\ 1997).  HIPASS has surveyed the entire southern sky for neutral hydrogen emission out to nearly 13000 \\kms. SINGG is intended to look for star formation, as traced by \\Halpha\\ emission, in this \\HI\\ selected sample.  The primary motivation for SINGG is to measure the star formation rate density of the local universe, \\SFRD.  The selection by 21cm emission bypasses all the common {\\em optical\\/} selection biases (e.g.\\ dust content and surface brightness), which should allow a more complete and representative survey of local star formation.  The planned incorporation of ultraviolet data from {\\em GALEX\\/} and infrared data from {\\em IRAS\\/} and {\\em SIRTF\\/} will allow a multi-wavelength determination of \\SFRD, addressing issues of dust extinction and the initial mass function.  SINGG is also intended to examine the \\HII\\ region luminosity function, determining how it varies as a function of morphology and \\HI\\ mass, \\MHI.  In this paper I address some of the morphological questions that motivate SINGG.  In particular (1) What fraction of gas rich galaxies are dormant?  That is, what fraction contains no detectable star formation?  (2) How common are starbursts?  And, (3) how common are galactic winds?  Addressing these questions will give us a better understanding of the life-cycle of star formation, and how star forming galaxies enrich the inter-galactic medium.  After presenting a few more details of the SINGG survey in \\S{2}, I discuss the \\Halpha\\ detectability of HIPASS sources in \\S{3}.  In \\S{4} I present plans for making an objectively defined starburst sample and estimate the starburst incident rate. In \\S{5} I summarize a morphological survey for extra-planar \\Halpha.  Finally \\S{6} summarizes these preliminary results. ",
        "conclusions": "The strategy of SINGG is to survey for star formation in the galaxies known to contain its essential ingredient, interstellar hydrogen, i.e.\\ we know where the fuel is; that's where to look for the fire.  This has proven to be a very effective approach, since we find that there are no dormant galaxies: when a galaxy has $\\MHI > 3\\times 10^7 \\, \\Msun$ it is observed to be undergoing some star formation.  The amount of \\HI\\ does not prove to be a good predictor of star formation morphology which varies greatly at fixed \\MHI.  About 15\\%\\ of \\HI\\ selected galaxies are starbursts.  Extra-planar \\Halpha\\ also appears to be a common phenomenon with half of edge-on galaxies having features with a radius $\\geq 0.5$ Kpc, and about a quarter having a radius $\\geq 1$ Kpc."
    },
    "0311/astro-ph0311190_arXiv.txt": {
        "abstract": "An observed correlation $E_p \\propto (E_{iso})^{1/2}$ extending from Gamma-ray Bursts (GRBs) to X-ray flashes (XRFs) poses problems both for a power-law universal jet model where the energy per solid angle decreases as the inverse square of the angle respect to the jet axis, and for a conical jet model with a uniform energy density within the jet beam and a sharp energy cut-off at the jet edge. Here we show that the current GRB-XRF prompt emission/afterglow data can be understood in terms of a picture in which the GRB-XRF jets are quasi-universal and structured, with a Gaussian-like or similar structure, i.e., one where the jet has a characteristic angle, with a mild variation of energy inside and a rapid (e.g. exponential) decrease of energy outside of it. A Monte Carlo simulation shows that the current data is compatible with such a quasi-universal Gaussian jet with a typical opening angle of $5.7^{+3.4}_{-2.1}$ degrees, and with a standard jet energy of about $\\log (E_j/\\mbox{1 erg})= 51.1\\pm 0.3$. According to this model, the true-to-observed number ratio of the whole GRB-XRF population is about 14 with the current instrumental sensitivity. ",
        "introduction": "The achromatic steepening in the afterglow lightcurves of some GRBs as well as energy budget arguments suggest that GRBs are produced by collimated jets (Rhoads 1999; Kulkarni et al. 1999; Harrison et al. 2001). This suggestion receives indirect support from the intriguing fact that a geometry-corrected jet energy appears to be standard (Frail et al. 2001; Panaitescu \\& Kumar 2001; Bloom, Frail \\& Kulkarni 2003; Berger, Kulkarni \\& Frail 2003a). There are two straightforward interpretations of this fact. One is that different GRBs collimate the same total energy into different angular openings (ranging from 1 to 30 degrees) (Rhoads 1999; Frail et al. 2001). Another is that all GRBs have a universal jet shape with a varying energy per solid angle which is of the form $\\epsilon(\\theta) \\propto \\theta^{-2}$, where $\\theta$ is the polar angle respect to the jet axis (Rossi, Lazzati \\& Rees 2002; Zhang \\& \\Mesz~2002a). In principle, an individual GRB could also have other jet structure, e.g. a Gaussian or even arbitrary function (Zhang \\& \\Mesz~2002a). When the jet parameters are allowed to have some dispersion around the mean values (which is demanded by the data), one has a ``quasi-universal'' jet structure, and a power-law jet structure is no longer a pre-requisite (Lloyd-Ronning, Dai \\& Zhang 2003).  Recently the so-called X-ray flashes (XRFs, Heise et al. 2003; Kippen et al. 2003) are identified and proposed as a natural extension of GRBs into a softer and fainter regime. Another intriguing empirical fact is that the GRB-XRF spectral break energy ($E_p$) in the cosmic rest frame appears to be correlated with the ``isotropic-equivalent'' energy (or luminosity) of the explosion (Amati et al. 2002; Sakamoto et al. 2003; Lamb, Donaghy \\& Graziani 2003; Lloyd et al. 2000) according to \\be E_p \\sim 100~{\\rm keV} \\left(\\frac{E_{iso}}{10^{52} ~{\\rm erg}}\\right)^{1/2}. \\label{Ep} \\ee Although the data sample in the XRF regime is currently very small (because the faintness of XRFs hinders the detection of their possible afterglows and the measurements of their redshifts), this intriguing result, if confirmed by further future data, strongly suggests that GRBs and XRFs are related events. A successful model should therefore be able to interpret both GRBs and XRFs within a unified framework. It is worth noticing that GRB 980425 has a very low $E_{iso}$ ($\\sim \\sim 10^{48}$ ergs) but a relative large $E_p$ ($\\sim 70$ keV), which is apparently inconsistent with eq.(\\ref{Ep}). Because it is atypical in other aspects as well, we do not interpret it within the current framework. Since there is no unified terminology, in this paper we call the combined population of GRBs and XRFs as GRB-XRF.  In this Letter, we do not discuss the nature of the $E_p \\propto (E_{iso})^{1/2}$ empirical law\\footnote{We note however that this information potentially contains important information to constrain the nature of the GRB fireball (whether or not it is magnetic-energy-dominated) and the location of the GRB prompt emission (whether internal or external). See Zhang \\& \\Mesz~(2002b) for more detailed discussions.}. Rather we conjecture that it is a universal law for the majority of GRB-XRFs, and evaluate its implications for various GRB-XRF jet models. We will show that a quasi-universal Gaussian-like structured jet model is compatible with the current prompt emission and afterglow data of GRBs and XRFs. This is consistent with earlier suggestions that GRBs would be seen as X-ray transients when viewed at large angles (e.g. \\Mesz~\\& Rees 1997; Woosley, Eastman \\& Schmidt 1999 and subsequent work). ",
        "conclusions": "We have argued that in order to incorporate both the prompt emission and afterglow data for GRBs and XRFs within a unified theoretical framework, both the simple $k=2$ power-law universal jet model and the on-beam uniform jet model encounter difficulties. With a Monte Carlo simulation, we show that the current data are compatible with a quasi-universal Gaussian-like structured jet model.  Other possible models may still interpret the data. First, a uniform jet also produces an equivalent isotropic energy outside the jet cone due to the off-beam Doppler effect. So an off-beam model for XRFs (Yamazaki et al. 2003a) does not over-generate XRFs. However, the decline of the isotropic energy at larger viewing angles is even steeper than exponential (e.g. Yamazaki, Yonetoku \\& Nakamura 2003b). So one can only generate XRFs at viewing angles slightly larger than the jet opening angle. This tends to {\\em under-generate} XRFs. The afterglow issue as discussed in \\S2 also poses some constraints on this model. In any case, a population simulation like the one presented here is needed to validate that model. Second, it is possible that some or even all GRB-XRF jets include two (or more) components (Li \\& Chevalier 1999; Frail et al. 2000; Lipunov, Postnov \\& Prokhorov 2001; Berger et al. 2003b; Sheth et al. 2003; Huang et al. 2003), as reproduced from the collapsar models (Zhang et al. 2003a, 2003b; \\Mesz \\& Rees 2001; Ramirez-Ruiz, Celotti \\& Rees 2003). In such a model, the jet structure may be approximated as a superposition of two (or more) uniform or Gaussian components. Such a model may be also able to reproduce the GRB-XRF data. However, the afterglow lightcurves could show distinct signatures under certain conditions (Berger et al. 2003b; Huang et al. 2003). The lack of such signatures in most of the afterglow lightcurves poses constraints on the parameters for such two-component structured jet models. Third, we do not exclude the possibility that a $k=2$ power law structure extends only up to a certain angle, above which the jet energy has a much steeper decline (Zhang \\& \\Mesz~2002a). However, the current Gaussian model is the simplest among these, and can accommodate the widest range of prompt emission and afterglow data for the majority of GRBs and XRFs. More detailed Monte Carlo simulations for this as well as other models and their comparison to a wider set of detailed prompt emission and afterglow data will be presented elsewhere (Dai \\& Zhang 2003, in preparation, see also Liang, Wu \\& Dai 2003)."
    },
    "0311/astro-ph0311473_arXiv.txt": {
        "abstract": "The bright Seyfert 1 galaxy \\mcg\\ has provided some of the best evidence to date for the existence of supermassive black holes in active galactic nuclei.  Observations with \\asca\\ revealed an X-ray iron line profile shaped by strong Doppler and gravitational effects. In this paper the shape of the iron line, its variability characteristics and the  robustness of this spectral interpretation are examined using the long \\xmm\\ observation taken in 2001.  A variety of spectral models, both including and excluding the effects of strong gravity, are compared to the data in a uniform fashion. The results strongly favour models in which the spectrum is shaped by emission from a relativistic accretion disc.  It is far more difficult to explain the $3-10 \\keV$ spectrum using models dominated by absorption (either by warm or partially covering cold matter), emission line blends, curved continua or additional continuum components.   These provide a substantially worse fit to the data and fail to explain other observations (such as the simultaneous \\sax\\ spectrum).  This reaffirms the veracity of the relativistic `disc line' interpretation.  The short term variability in the shape of the energy spectrum is investigated and explained in terms of a two-component emission model. Using a combination of spectral variability analyses the spectrum is successfully decomposed into a variable  power-law component (PLC) and a reflection dominated component (RDC). The former is highly variable while the latter is approximately constant throughout the observation, leading to the well-known spectral variability patterns. Consideration of the long term X-ray monitoring of \\mcg\\ by \\xte\\ demonstrates that the long \\xmm\\ observation sampled the `typical' state of the source.  These results and those of other analyses of the large \\xmm\\ dataset are summarised and their implications for understanding this enigmatic source are discussed. ",
        "introduction": "\\label{sect:intro}  X-ray spectroscopy holds great promise for probing physics close to the accreting black holes thought to power Active Galactic Nuclei (AGN) and Galactic Black Hole Candidates (GBHCs). The high X-ray luminosities observed from these  systems are thought to be generated close to the black hole ($\\ls 100 \\rg$; where $\\rg \\equiv GM/c^2$ is the gravitational radius) where the effects of strong gravity, such as gravitational redshift and light bending, will imprint characteristic signatures on the emerging X-ray spectrum.  As it escapes towards the observer, the primary X-ray emission will interact with material in the innermost regions, modifying the observed X-ray spectrum.  Reprocessing of the X-rays by an accretion disc will lead to a particular set of spectral features. Absorption, fluorescence, recombination, and Compton scattering in the surface layers of the X-ray illuminated disc produce a `reflection' spectrum. Under a wide range of physical conditions the most prominent observables will be an emission line at $\\approx 6.4 \\keV$ (the iron \\ka\\ fluorescence line) and a broad `hump' in the spectrum which peaks in the range $20-40 \\keV$ (e.g. Guilbert \\& Rees 1988; Lightman \\& White 1988; George \\& Fabian 1991; Matt, Perola \\& Piro 1991).  Observations of Seyfert 1 galaxies with \\ginga\\ (Pounds \\et 1990; Nandra \\& Pounds 1994) demonstrated the simultaneous presence of a $\\approx 6.4 \\keV$ emission line and an up-turn in the spectrum above $\\sim 10 \\keV$, indicative of the reflection hump.  The presence of the reflection hump has since been confirmed by \\cgro\\ (e.g. Zdziarski \\et 1995) and \\sax\\ (e.g. Perola \\et 2002).  The bright Seyfert 1 galaxy \\mcg\\ ($z=0.007749$) has received particular attention since \\asca\\ observations revealed the presence of a broad, asymmetric emission feature at energies between $4 - 7 \\keV$ identified with a highly broadened Fe \\ka\\ emission line (Tanaka \\et 1995).  Similar features were subsequently observed in other Seyfert 1 galaxies (Mushotzky \\et 1995; Nandra \\et 1997).  The profile of the line can be explained in terms of fluorescent emission from the surface of  an accretion disc extending down to $\\ls 6 \\rg$ about a supermassive black hole (SMBC): the relativistic `disc line' model (Fabian \\et 1989; Stella 1990; Laor 1991). For comparison, the emission lines seen  in other wavebands (e.g. broad optical lines) are thought to originate at distances $\\gs 10^{3} \\rg$. The broad Fe \\ka\\ line therefore potentially offers a powerful diagnostic of the physical conditions in the immediate environment of the black hole (see Fabian \\et 2000; Nandra 2001 and Reynolds \\& Nowak 2003  for recent reviews).  Repeated observations of \\mcg\\ with \\asca\\ (Iwasawa \\et 1996, 1999; Shih \\et 2002), \\sax\\ (Guainazzi \\et 1999), \\xte\\ (Lee \\et 1999, 2000; Vaughan \\& Edelson 2001), \\chandra\\ (Lee \\et 2002) and \\xmm\\ (Wilms \\et 2001, hereafter W01, and Fabian \\et 2002, hereafter paper~I) have confirmed the presence of the broad, asymmetric emission feature. However, while this is clearly not an artifact of any particular instrument or observation, there are  uncertainties associated with the detailed modelling of the spectrum.  The most challenging of these is that \\mcg, along with many other Seyfert 1s, shows complex absorption in its X-ray spectrum. This complicates the process of identifying the correct underlying continuum and thereby measuring the superposed line emission.  Indeed, the presence of relativistically broadened  emission lines in the X-ray spectra of other Seyfert 1 galaxies has recently become something of a {\\it cause celebre} (see e.g. Lubinski \\& Zdziarski 2001; Inoue \\& Matsumoto 2001; Branduardi-Raymont \\et 2001; Pounds \\& Reeves 2002; Page, Davis \\& Salvi 2003; Pounds \\et 2003a,b).  The long ($\\sim 320\\ks$) \\xmm\\ observation of \\mcg\\ is  of great importance to studies of the broad iron line, since  \\mcg\\ offers the best established example of a relativistic line profile.  In this paper the X-ray spectrum and spectral variability of \\mcg\\ are examined using the long \\xmm\\ observation.  The interpretation of the spectrum in terms of relativistic Fe \\ka\\ emission is  investigated by fitting a range of models including both simple phenomenological models and models accounting for the detailed physics expected in photoionised accretion discs. These are compared to models that do not include relativistic effects in order to test the strength of the argument for strong gravity in \\mcg.  A variety of techniques are employed to characterise the spectral variability of the source and thereby further constrain the range of  viable spectral models.  The two \\xmm\\ observations are considered in the context of the long term behaviour of \\mcg\\ using \\xte\\ monitoring observations spanning more than six years.  The rest of this paper is organised as follows. Section~\\ref{sect:data} gives details of the \\xmm\\ and \\xte\\ observations and their data reduction.   Section~\\ref{sect:history} briefly reviews the \\xte\\ monitoring of \\mcg\\ and shows how the two \\xmm\\ observations fit in with the long timescale properties of the source.   Section~\\ref{sect:fit} discusses the spectral fitting results using models that include relativistic effects (Section \\ref{sect:fit_with}) and exclude relativistic effects (Section \\ref{sect:fit_without}).   In section~\\ref{sect:spec_var} the spectral variability of the source is examined using flux-flux plots (section~\\ref{sect:ff}), flux-resolved spectral analysis (section~\\ref{sect:fluxes}), rms spectra (section~\\ref{sect:rms}), time-resolved spectral analysis (section~\\ref{sect:time-res}) and Principal Component Analysis (PCA; section~\\ref{sect:pca}). Section~\\ref{sect:big_diff} discusses differences between the long \\xmm\\ observation of 2001 and the previous observation taken in 2000. Section~\\ref{sect:review} summarises the new results from the present work as well as the other analyses of  these \\xmm\\ data. Section~\\ref{sect:disco} presents a discussion of the main results and the main conclusions are given in section~\\ref{sect:conc}. Some details of the EPIC spectral calibration are investigated in Appendix~\\ref{app:cal}. ",
        "conclusions": "\\label{sect:conc}  The long \\xmm\\ observation of \\mcg\\ has confirmed the existence of a highly broadened iron line. The main results of the present investigation are as follows:  \\begin{enumerate}  \\item It is not yet possible, given the present state of knowledge of the detailed properties of the warm absorber, to unambiguously dis-entangle the red wing of the line from the absorbed continuum.  \\item Nevertheless, the best-fitting model requires the reflection/line emitting region to extend inwards to $\\sim 2\\rg$. This result was robust to various possible biases in the continuum determination.  \\item Alternative models for the spectral region around the iron line, not allowing for the effects of strong gravity, gave unsatisfactory results.  \\item The spectral variability of \\mcg, including the correlation between $2-10\\keV$ slope and flux and the reduced variability of the iron line, has been explained using a two-component model.  \\item The relative constancy of the reflection dominated component (RDC), in the presence of large variations in the power-law component (PLC), is the primary cause of the spectral variations. The reasons for the lack of RDC variations are unclear. A model invoking gravitational light bending near the SMBH can qualitatively explain the suppressed variations, the relative strength of the RDC and the small inner disc radius.  \\end{enumerate}"
    },
    "0311/astro-ph0311159_arXiv.txt": {
        "abstract": "In this paper I present evidence suggesting that the absence or presence of Hidden Broad Line Regions (HBLRs) in Seyfert 2 galaxies is regulated by the rate at which matter accretes onto a central supermassive black hole, in units of Eddington. I discuss the above findings in the context of a model proposed by Nicastro (2000), in which the existence of the BLRs is regulated by the disk accretion rate. For low enough accretion rates, the disk is stable, and the BLRs do not form. ",
        "introduction": "In the framework of orientation-based unificaton schemes (Antonucci, 1993), type-2 AGN are believed to be intrinsically the same as type-1 AGN, but seen at different angles (i.e. edge-on {\\em versus} pole-on). This scenario was first proposed by Antonucci \\& Miller (1985) to explain the presence of polarized broad lines in the archetypical Seyfert 2 (Sy2), NGC 1068, and is now supported other than by spectropolarimetric observations of Hidden Broad Emission Line Regions (HBLRs) in several other sources, also by X--ray observations, which demonstrate that Sy2s have usually absorption columns largely exceeding the Galactic ones.  Despite observations do generally support orientation-based unification models for AGNs, important exceptions do exist. Purely orientation based unification schemes, for example, fail to account for all of the observed differences between two sub-classes of type 1 AGN: normal Seyfert 1s (Sy1s), and the so called Narrow Line Seyfert 1s (NLSy1s). NLSy1s have [OIII]/H$\\beta$ ratio typical of type-1 sources, but also BEL FWHM narrower than 2000 km s$^{-1}$. They also differ significantly from Sy1s in the X-rays, where they show much steeper soft (0.1-2 keV: Boller et al., 1996) and hard (2-10 keV: Brandt et al. 1997; Leighly, 1999a) continua, as well as faster variability (Leighly, 1999b). It has been suggested that the black hole mass, or the accretion rate, or both (e.g. Pounds et al. 1995; Laor et al. 1997; Nicastro, 2000: hereinafter N00) are responsible for the observed differences between Sy1s and NLSy1s.  \\noindent More importantly, orientation-based models fail in their fundamental prediction: only about 50 \\% of the brightest Seyfert 2s show the presence of HBLRs in their polarized optical spectra, while the remaining half do not (Tran, 2001; Gu \\& Huang, 2002).  Here I present evidence that suggests that the absence or presence of HBLRs is regulated by the ratio between the X-ray luminosity and the Eddington luminosity which, in the accretion--powered scenario, is a measure of the rate at which matter accretes onto the central supermassive black hole. This evidence is consistent with the model proposed by N00, in which the BLRs are formed due to accretion disk instabilities occurring in proximity of the critical radius at which the disk changes from gas pressure dominated to radiation pressure dominated. This radius diminishes with decreasing $\\dot{m}$; for low enough accretion rates (and therefore luminosities), the critical radius becomes smaller than the innermost stable orbit, the disk is not longer unstable, and so BLRs do not form. Under the Keplerian assumption, this model is able to reproduce the whole range of observed BEL FWHM in AGN, as a function of the accretion rate in units of Eddington, and so also to explain the observed difference in BEL-FWHM between Sy1s and NLSy1s.  \\subsection{The Model} The model is based on two basic assumptions: (a) a vertical disk wind, originating at a critical distance in the accretion disk, is the origin of the BEL-Clouds (BELCs), and (b) the widths of the BELs are the Keplerian velocities of the accretion disk at the radius where the wind arises. The wind forms because of the existence of Lightman-Eardley (Lightman \\& Eardley, 1974) instabilities in the radiation pressure dominated region of a Shakura-Sunyaev (SS) disk (Shakura \\& Sunyaev, 1973), and so for external accretion rates higher than a minimum value below which a standard SS-disk is stable and extends down to the last stable orbit. The region of the disk where the wind forms is bound internally by the transition radius between the radiation pressure and gas pressure dominated parts of the disk, \\begin{equation} \\label{rtran} r_{tran} f^{-16/21} \\simeq 15.2 (\\alpha m)^{2/21} \\left({1 \\over \\eta} \\dot{m} \\right)^{16/21}, \\end{equation} and externally by the radius at which, in the context of the dynamical disk-corona model proposed by Witt, Czerny and Zycki in 1997 (WCZ97), the fraction of energy dissipated in the corona ($(1-\\beta) \\simeq 0.034 (\\alpha f {1 \\over \\eta} \\dot{m})^{-1/4} r^{3/8}$) is maximum (see N00 for additional details) \\footnote{In the WCZ97 model a transonic vertical outflow from the disk develops where the fraction of total energy dissipated in the corona is close to unity} : \\begin{equation} \\label{rmax} r_{max} f^{-2/3} \\simeq 8,000 \\left( \\alpha {1 \\over \\eta} \\dot{m} \\right)^{2/3}. \\end{equation} In the above equations all quantities are dimensionless: $m = M/M_{\\odot}$, $\\dot{m} = \\dot{M} / \\dot{M}_{EDD}$, $r = R/R_0$, with $\\dot{M}_{EDD} = 1.5 \\times 10^{17} \\eta^{-1} m \\hbox{ }$ g s$^{-1}$, and $R_0 = 6GM/c^2$, for a non-rotating black hole; here $\\eta$ is the maximum efficiency. Finally, $f$ gives the boundary conditions at the marginally stable orbit: $f = f(r) = (1 - r^{-0.5})$.  The critical radii $r_{tran}$ and $r_{max}$ (equations 1 and 2) depend both on the accretion rate and, interestingly, with similar powers. This results in a quasi-rigid radial shifting of the region delimited by these two distances as the accretion rate (in critical units) varies. Finally, for the transonic disk outflow (and so for the BELRs), the model adopts an intermediated radius $r_{wind}$ computed weighting the radial distance between $r_{tran}$ and $r_{max}$ by $(1-\\beta)$.  Equation 1 allows us to define the minimum external accretion rate needed for a thermally unstable radiation pressure dominated region to exist. From the condition $r > 1.36$ (the limit of validity of the SS-disk solution) we have: $\\dot{m} \\gs \\dot{m}_{min}(m) \\simeq 0.3 \\eta (\\alpha m)^{-1/8}$. Throughout this paper I assume $\\eta = 0.06$, and a viscosity coefficient of $\\alpha = 0.1$, which give a minimum external accretion rate of $\\dot{m}_{min} \\sim (1-4) \\times 10^{-3}$, for $m$ in the range $10^6 - 10^9$ (see Fig. 1 of N00). At lower accretion rates a SS-disk (SS73) is stable down to the last stable orbit: we propose that all the available energy is dissipated in the disk and no radiation pressure supported and driven wind is generated. AGN accreting at these low external rates should show no BELs in their optical spectra. ",
        "conclusions": "In this contribution I presented a model that links the physical mechanism responsible for the AGN activity (i.e. the accretion) with the existence and width of the BELs. This model seems to naturally (a) reproduce the Sy1-{\\em versus}-NLSs1 observed differences, and (b) account for those 50 \\% of Sy2s for which no HBLRs has been observed. Additional testing of this model is needed, and requires larger samples both in the optical (e.g. SDSS photometry and spectra), and in the X-rays (e.g. {\\em Chandra} and {\\em Newton}-XMM spectra)."
    },
    "0311/astro-ph0311535_arXiv.txt": {
        "abstract": "We present a dark energy model with a double exponential potential for a minimally coupled scalar field, which allows general exact integration of the cosmological equations. The solution can perfectly emulate a $\\Lambda$-term model at low-medium redshift, exhibits tracking behavior, and does not show eternal acceleration. ",
        "introduction": "Among the many problems posed by the discovery of positive acceleration of the Universe, one of the most puzzling is that of eternal acceleration. In the case of a true cosmological constant, as well as for most of the models proposed in the literature, the present time accelerating era seems doomed to last forever. This creates problems in string theory, where asymptotic flatness of spacetime is required \\cite{1}. Several possible solutions are based on cycles in the evolution of the Universe \\cite{2}, or on the introduction of negative potentials for the scalar field \\cite{3}. A much simpler approach is obtained with a negative cosmological constant, together with a suitable potential \\cite{4}. However, this last approach poses the same type of problems which make unsatisfactory the simple use of a standard positive $\\Lambda $. A common feature of many of these models is that they foresee a recollapsing universe in the context of a spatially flat metric. But a recollapsing universe can also be obtained simply by setting $k=+1$, an option which has not been ruled out by the BOOMERANG and MAXIMA missions, as instead people often say. On the contrary, the most recent data by Wilkinson Microwave Anisotropy Probe (WMAP) have reopened the way to nonflat cosmological models \\cite{5}, \\cite{6}.  In this paper, we remain in the context of a flat space $\\left( k=0\\right) $ and propose a very simple and attractive model with a minimally coupled scalar field plus dust. We shall find a general exact solution of the equations, and show that it meets all the requests which are usually made for a dark energy model, including tracking behavior and switch off of the acceleration. Moreover, we shall not get a recollapsing universe.  Let us consider the following potential (in units such that $8\\pi G=1$): \\begin{eqnarray} V(\\varphi ) &=&\\left( A\\exp \\left( \\frac{1}{2}\\sqrt{\\frac{3}{2}}\\varphi \\right) -B\\exp \\left( -\\frac{1}{2}\\sqrt{\\frac{3}{2}}\\varphi \\right) \\right) ^{2}  \\nonumber \\\\ &=&A^{2}\\exp \\left( \\sqrt{\\frac{3}{2}}\\varphi \\right) +B^{2}\\exp \\left( -% \\sqrt{\\frac{3}{2}}\\varphi \\right) -2AB.  \\label{eq1} \\end{eqnarray}  \\noindent  This potential has already been discussed in the literature (see Refs. \\cite{7}, \\cite{8}, \\cite{9} for instance). In Ref. \\cite{8}, which is a paper closely related to this one, the potential is not equal but only similar to that in Eq. (\\ref{eq1}). Comments on the attractor properties of the (generic) solutions due to such a kind of potential in connection with initial conditions and the values of model parameters have also been given, without trying to find any exact solution \\cite{7}. On the other hand, Ref. \\cite{9} inverts our point of view, since it postulates a special scale factor behavior, so deducing a potential form similar to the one in Eq. (\\ref{eq1}% with some relevant differences with respect to it. (The main point is the different relation among the values of the potential parameters.) The potential (\\ref{eq1}) also turns out to be a particular case of the wider class discussed in Ref. \\cite{4}. For a discussion on the possibility of finding a general exact solution for the kind of potentials to which the one in Eq. ( \\ref{eq1}) belongs, however, see Ref. \\cite{10}.  In any case, here\\ we do not present a new technique for finding solutions; instead, we select the particular class of potentials in Eq. (\\ref{eq1}) and a treatment of the related solutions, which allows us to obtain the desired results and puts them in greatest evidence.  Some comments are now needed on Eq. (\\ref{eq1}).  (1) The second form of the potential can also be interpreted as containing a negative cosmological constant, fine tuned with the parameters of the potential itself. The first form shows that this is in fact only a matter of interpretation. Such a form as a whole, of course, still retains the right to be used as a potential in itself.  (2) In particular, if we set $A=B$, the expansion of $V(\\varphi )$ around its minimum is \\begin{equation} V(\\varphi )\\approx \\frac{1}{2}(3B^{2})\\varphi ^{2}+\\frac{1}{16} (3B^{2})\\varphi ^{4}+...,  \\label{eq2} \\end{equation} showing that it can be interpreted as a massive self-interacting potential, which is quite reasonable in the situation we are studying. We shall see that posing $A = 0$ only affects the asymptotic value of $\\varphi $. As a matter of fact, we consider $A\\neq B$ just for sake of generality.  (3) The real problems with potential (\\ref{eq1}) are the particular value of the exponent, \\ and the link between the mass term and the self-interaction term in Eq. (\\ref{eq2}). We are not really able to explain them in terms of fundamental theory\\footnote{But, about this, see Refs. \\cite{11},\\cite{12} for some dimensional considerations on exponential potential, once one wants power law solutions for the scale factor.}. It is however intriguing that, as we shall see, the dynamics entailed by this potential on the minisuperspace $\\{a,\\varphi \\}$ admits a Noether symmetry.  After presenting the cosmological solution in Sec. II, we analyze it in Sec.\\ III. Finally, Sec. IV is dedicated to some conclusive remarks. ",
        "conclusions": "The model presented here, in our opinion, shows a possible solution for all the main problems posed by the dark energy theory, at least from a mathematical point of view.  The physical attitude should be however more careful. Though very simple and reasonable, the potential proposed here has at the moment not satisfactory justification, independent of this model.  We have seen that, were this the real situation, it could be probably impossible to discriminate from a $\\Lambda $-term model with any realistic plan of observations.\\footnote{ Note that this tendency also results from the very recent WMAP data \\cite{18} \\par {}}. As for the switch of the acceleration in the future, it is of course unobservable in principle!  The main result of this paper seems to be an indication of a possible way to solve the problem of switching off the acceleration by means of minimal tools. It also suggests that oscillations of the scalar field could be an interesting field of study also in a context very far from inflation and reheating. (For some considerations about this, for example, see Refs. \\cite{19}, \\cite{20}.)  A last remark has to be addressed to the fact that the radiation dominated era is not recovered at early times. As a matter of fact, this is due to the model itself and the way it is built. Adding a radiation component is not simple on an analytical ground, since in fact it is not possible to find a Noether symmetry and, therefore, an analytical exact general solution. This means that including radiation, so extending the model backwards, at least needs numerical techniques, if one also wants to retain the other two components, dust and scalar field."
    },
    "0311/astro-ph0311229_arXiv.txt": {
        "abstract": "In the last few years a significant population of ultracool L and T dwarfs has been discovered.  With effective temperatures ranging from $\\sim$ 2200 to 700 K, these objects emit most of their radiation in the near infrared and their spectral energy distributions are dominated by strong molecular absorption bands. These highly structured energy distributions lead to $JHK$ magnitudes that are extremely sensitive to the exact filter bandpass used.  In the case of the T dwarfs, the differences between commonly used photometric systems can be as large as 0.4 mag at $J$ and 0.5 mag at $J-K$.  Near--infrared magnitudes have been published for L and T dwarfs using a variety of photometric systems.  Currently, the data obtained with these systems cannot be accurately compared or combined as transformations based on the colors of hotter stars are not valid for L and T dwarfs.  To address this problem, we have synthesized $J, H$, and $K$ magnitudes for some of the common photometric systems and present transformation equations with respect to the most atmospheric--independent system, the Mauna Kea Observatory (MKO) filter set.  If the spectral type of the dwarf is known, our transformations allow data to be converted between systems to 0.01 mag, which is better than the typical measurement uncertainty.  Transforming on the basis of color alone is more difficult because of the degeneracy and intrinsic scatter in the near--infrared colors of L and T dwarfs; in this case $J$ magnitudes can only be transformed to $\\lesssim$0.05 mag and $H$ and $K$ to $\\lesssim$0.02 mag. ",
        "introduction": "L and T dwarfs have unusual spectral energy distributions (SEDs), with most of their flux emitted through windows in the near--infrared.  Normalized spectra of an L5 and T4.5 dwarf \\citep{geb02} are shown in Figure 1 to illustrate how absorption bands of $\\rm H_{2}O$, CO, and $\\rm CH_{4}$ regulate the near--infrared emission and create the flux windows. The $\\rm H_{2}O$ bands are the same features responsible for the telluric absorption that defines the conventional $J$ (1.1---1.4 $\\micron$), $H$ (1.5---1.8 $\\micron$), and $K$ (2.0---2.4 $\\micron$) bandpasses.  Thus, the presence of $\\rm H_{2}O$ in the atmospheres of the L and T dwarfs forces much of their flux to be emitted within these bands, resulting in the extreme far--optical and near--infrared colors that are used to identify L and T dwarf candidates from photometric surveys like the 2 Micron All Sky Survey (2MASS, \\citealt{2m}), and the Sloan Digital Sky Survey (SDSS, \\citealt{sdss}).  There are now hundreds of L and T dwarfs with near--infrared magnitudes published in various photometric systems.  To maximize the science potential of these observations and their impact on brown dwarf theory, transformation equations are desirable to convert these and any future magnitudes to a common photometric system.  Because the $J$, $H$ and $K$ bandpasses include  $\\rm H_{2}O$, CO, and $\\rm CH_{4}$ absorption features, any variation in the width of the filters will lead to system--dependent magnitudes.  Therefore, transformation equations need to be derived as a function of near--infrared spectral type or, less desirably, color (see later discussion in \\S 4) to correctly account for the presence of the molecular bands.  Figure 2 shows $J-K$ color as a function of spectral type for typical main sequence stars (A0--M5) in the Bruzual--Persson--Gunn--Stryker Atlas\\footnote{http://www.stsci.edu/instruments/observatory/cdbs/astronomical\\_catalogs.html}, and late M, L and T dwarfs reported in \\citet{le02}.  Important features to note, all of which are explained more fully in other works (e.g. \\citealt{le02}) are: the intrinsic spread in $J-K$ for the L dwarfs that may be produced by variations in the extent and location of the condensed grain layer in the photosphere; the increasingly bluer $J-K$ color for the late L dwarfs and T dwarfs that is mostly due to the appearance of the $\\rm CH_{4}$ band at 2.2\\micron; and the scatter in $J-K$ for the latest T dwarfs likely due to gravity--dependent $\\rm H_{2}$ opacity.  Although T dwarfs can have the same value of $J-K$ as A through M stars, convolving a T dwarf spectrum with its strong molecular absorption bands (Figure 1) with any $JHK$ bandpass will produce a very different result from the convolution of e.g. the Rayleigh Jeans curve of an A0 star with the bandpass. Consequently, transformations based on the colors of hotter stars are not valid for L and T dwarfs and photometric transformations must be derived directly from observations of these ultracool dwarfs.  In this paper we present synthesized $J, H$ and $K$ magnitudes using near--infrared spectroscopic observations of L and T dwarfs from each spectral subtype, for the photometric systems in which L and T dwarf photometry has most frequently been published and for established systems in which future observations may occur.  The photometric systems presented are: 2MASS \\citep{c01}, Caltech  (CIT, \\citealt{el82}), the DEep Near-Infrared Survey (DENIS, \\citealt{fou00}), Las Campanas Observatory (LCO, \\citealt{per98}), Mauna Kea Observatory (MKO, \\citealt{sim02,tok02}), the United States Naval Observatory Flagstaff Station (NOFS, \\citealt{d02,gue03}) and the United Kingdom Infrared Telescope (UKIRT, \\citealt{haw01}).  We also generate equations  to transform $J, H$ and $K$ magnitudes between the other systems and the MKO system. The MKO photometric system was chosen as the reference point because MKO filters are narrower than classical $J$, $H$ and $K$ filters, thus avoiding the telluric absorption bands that can vary with time and observing location (see discussion in \\citealt{sim02}, \\citealt{tok02}).  As a result MKO magnitudes have little dependence on local observing conditions, and  their use produces transformation equations with less uncertainty than would be obtained using another photometric system.  In addition, the MKO filters have been widely adopted and the system is endorsed by the IAU Working Group on Infrared Photometry as the preferred photometric system for ground-based near--infrared observations.  In \\S 2 we present the observed differences in the 2MASS and MKO magnitudes that have been measured for several L and T dwarfs, showing that system transformations cannot be reliably derived empirically due to significant uncertainty in the observational data. \\S 3 discusses the inputs for synthesizing magnitudes: filter transmission profiles, telluric absorption bands, instrument optics and observed spectra. Our results are presented in \\S 4 and our conclusions given in \\S 5. ",
        "conclusions": "To obtain accurate and stable photometry, filter bandpasses should not go into poor regions of the terrestrial atmosphere  --- most (classical) $J$  filters are poorly defined from this point of view.  Variations in the water vapor content change the effective bandpass of such filters and, for objects with extremely structured spectral energy distributions  such as T dwarfs, these changes produce photometric deviations of  $\\sim$0.05---0.10~magnitudes.  To measure magnitudes and colors to better than this requires use of a filter set that is well matched to the atmosphere, such as the MKO filter set.  $JHK$ magnitudes for L and T dwarfs are highly dependent on the photometric system used for the observation; for T dwarfs differences between systems can be several tenths of a  magnitude.  However, we have shown that $JHK$ magnitudes for L and T dwarfs can be transformed between the 2MASS, CIT--$H$\\&$K$, DENIS, LCO, NOFS and UKIRT systems and the MKO system to  $\\sim$0.01~mag if the spectral type of the dwarf is known. This is significantly better than the typical measurement uncertainty, i.e. the original uncertainty in the measurement will determine the accuracy of the transformed value. For the CIT system the uncertainty in the $J$ bandpass effects the derived magnitudes by 0.05--0.10 mag for L and T dwarfs. Variations between the optical elements of common infrared instrumentation are expected to impact the measured magnitudes of the late L and T dwarfs at the 0.01~mag level. If spectral type is not known, then $J-K$ color can be used to transform $H$ and $K$ magnitudes measured in different systems with an accuracy of about 0.02~mag but the $J$ value can only be derived to $\\sim$ 0.05~mag on the basis of color alone.  The results presented here will be valuable for researchers in the very active field of ultracool dwarf studies, where imaging data are plentiful, and where the data have unfortunately been obtained with a variety of photometric systems.  Transformations based on the colors of hotter stars, even if the stars have the same color, cannot be applied to objects with strong molecular absorption bands such as those seen in L and T dwarfs.  For these ultracool objects intercomparison of photometric data requires knowledge of the filter profiles at instrument temperature, and knowledge of the local atmospheric transmission.  If $JHK$ photometry is obtained with a well understood photometric system we have shown that such datasets can be accurately combined or compared."
    },
    "0311/gr-qc0311022_arXiv.txt": {
        "abstract": "We argue that an infinite universe should not necessarily be avoided on philosophical grounds. Paradoxes of repeating behaviour in the infinite, or eternal inflationary, universe can be alleviated by a realistic definition of differing lives: not simply permutations of various quantum states. The super-exponential growth in the rules of a  cellular automata is used as an example of surpassing the holography bound.  We also critically question the notion that our universe could simply be a simulation in somebody else's computer. \\\\    PACS numbers: 04.20.Gz, 98.80.-k\\\\ Keywords: Cosmology, Inflation, Cellular Automata ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311586_arXiv.txt": {
        "abstract": "{ We compare the tau neutrino flux arising from the galaxy and the earth atmosphere for $10^{3} \\leq E/\\mbox{GeV} \\leq 10^{11}$. The intrinsic and oscillated tau neutrino fluxes from both sources are considered.  We find that, for $E\\geq 10^3$ GeV, the oscillated $\\nu_{\\tau}$ flux along the galactic plane dominates over the maximal intrinsic atmospheric $\\nu_{\\tau}$ flux, i.e., the flux along the horizontal direction. We also briefly comment on the prospects for observing these high-energy tau neutrinos. \\PACS{ {95.85.Ry}{Neutrino, muon, pion, and other elementary particles; cosmic rays}   \\and {13.85.Tp}{Cosmic-ray interactions} } % } % ",
        "introduction": "\\label{intro} The Milky way is one of the nearby astrophysical sources producing high energy neutrinos, besides the familiar earth atmosphere \\cite{Athar:2003gw}. Measurements of galactic neutrino and photon fluxes could provide information about the distribution of matter and cosmic rays in the galaxy. Furthermore, the above flux is also a background for the search of more distant high energy neutrino sources such as the AGNs and the GRBs. In this talk, we shall focus on the flux of $\\nu_{\\tau}$. It is clear that the observation of astrophysical $\\nu_{\\tau}$, with a flux comparable to the flux of $\\nu_{e}$ and $\\nu_{\\mu}$, directly confirms the neutrino oscillations. In order to observe galactic $\\nu_{\\tau}$ flux, it is essential to study the atmospheric background. We shall focus on the energy range $E_{\\nu}\\geq 10^3$ GeV.  We first discuss the distinction between intrinsic and oscillated neutrino fluxes arising from the Milky way and the earth atmosphere. We then present results for galactic and atmospheric tau neutrino fluxes. Finally we comment on the prospects for observing galactic tau neutrinos. ",
        "conclusions": ""
    },
    "0311/astro-ph0311065_arXiv.txt": {
        "abstract": "Earlier works, which we review, have shown that if the Fe core in a presupernova star is to be sufficiently massive to collapse into a black hole, earlier in the evolution of the star the He core must be covered (clothed) by a hydrogen envelope during He core burning and removed only following this, in, e.g. common envelope evolution. This is classified as Case C mass transfer. These previous arguments were based chiefly on stellar evolution, especially depending on the way in which $^{12}$C burned.  In this work we argue for Case C mass transfer on the basis of binary evolution. The giant progenitor of the black hole will have a large radius $\\sim 1000\\rsun$ at the end of its supergiant stage. Its lifetime at that point will be short, $\\sim 1000$ yrs, so it will not expand much further. Thus, the initial giant radius for Case C mass transfer will be constrained to a narrow band about $\\sim 1000\\rsun$. This has the consequence that the final separation $a_f$ following common envelope evolution will depend nearly linearly on the mass of the companion $m_d$ which becomes the donor after the He core of the giant has collapsed into the black hole. The separation at which this collapse takes place is essentially $a_f$, because of the rapid evolution of the giant. (In at least two binaries the black hole donor separation has been substantially increased because of mass loss in the black hole formation. These can be reconstructed from the amount of mass deposited on the donor in this mass loss.)  We show that the reconstructed preexplosion separations of the black hole binaries fit well the linear relationship. ",
        "introduction": "}  The black hole Soft X-ray Transient (SXT) \\astrobj{A0620$-$00} consisting of an $\\sim 11\\msun$ black hole and $\\sim 0.7\\msun$ K-star companion was evolved theoretically by De Kool et al. (1987) with Case C mass transfer; i.e., the common envelope evolution in which the companion removed the hydrogen envelope of the giant black hole progenitor through spiral-in took place after He core burning was completed in the giant. In Table~1 of Lee, Brown, \\& Wijers (2002) (denoted as LBW), the  seven SXTs with shortest periods had K- or M-star companions and the unclassified companion in \\astrobj{XTE 1859$+$226} may also well be K or M because of its short period $P=0.380$ days. The progenitor binaries of these would have involved $\\sim 25\\msun$ giants and $\\sim 1-2\\msun$ companions, the latter having had some mass stripped off by the black hole. In other words, all of the shortest period SXTs are successfully evolved with the {\\it same} Case C mass transfer.  We emphasized in Lee \\& Brown (2002) that in Case C mass transfer the orbital separations for the above progenitor binaries in Roche Lobe contact are at $\\sim 1700\\rsun (\\pm \\sim 10\\%)$, $\\sim 8$ AU, for ZAMS $20\\msun$ black hole progenitor and $1\\msun$ companion star. Progenitors with more massive companions and the larger initial separation necessary for Case C mass transfer could have removed the H-envelope of the giant with spiral-in to larger final separations $a_f$, since their drop in gravitational energy of the more massive companion is then sufficient to remove the envelope, and we shall see that this is indeed what happens.   In LBW we listed \\astrobj{Nova Scorpii}  and \\astrobj{IL Lupi} as undergoing mass transfer while in main sequence. Beer \\& Podsiadlowski (2002) have carried out a detailed, convincing numerical evolution of \\astrobj{Nova Scorpii}, showing that the orbit has widened substantially under nearly conservative mass transfer. Podsiadlowski et al. (2002) (denoted as PRH) have recently extended such calculations to the other binaries with evolved companions, showing that they all began mass transfer in main sequence, although \\astrobj{V404 Cyg}, \\astrobj{J1550$-$564} and probably \\astrobj{GRS 1915$+$105} will have progressed beyond main sequence. (As noted later, we differ with PRH in our suggested evolution of \\astrobj{Cyg X-1}.) The PRH calculations generally support the schematic LBW calculations of mass transfer, but have the added advantage that by beginning the transfer in main sequence, sufficient mass can be transferred in the traditional sub Eddington limit. Whereas we do not believe this to be necessary in the case of black holes, seeing no reason why the accretion across the event horizon could not be substantially hyper Eddington (and PRH also covers this case, as a possibility) the standard PRH scenario allays the fear of the greatly hyper Eddington scenario which may go against \"accepted wisdom\".  In Case C mass transfer there is a great regularity expressed in the roughly linear dependence of companion mass on orbital separation of the giant black hole progenitor and companion on the companion mass \\be a_f \\propto \\frac{M_d}{\\msun} \\left(\\frac{M_{\\rm giant}}{\\msun} \\right)^{-0.55} R \\label{eq1} \\ee following the spiral-in stage which removes the envelope of the giant (Lee, Brown \\& Wijers 2002). We derive Eq.~(\\ref{eq1}) in the Appendix. Except for the roughly square root dependence on giant mass, this relation is linear. Here the companion (donor) mass is labelled $M_d$, $a_f$ is the separation of the He-star, companion binary following spiral-in in common envelope evolution, and $R$ is the initial radius of the giant at the start of common envelope evolution. The dependence on $M_{\\rm giant}$ is weak, the interval \\be 20\\msun < M_{\\rm giant} < 30\\msun \\ee being used by Lee, Brown \\& Wijers (2002). The term depending on giant mass originates from the term $M_{\\rm He}/M_{\\rm giant}^2$ in common envelope evolution. The relation Eq.~(\\ref{eq1}) is particularly useful because, as we shall argue, $R$ is nearly constant, $\\sim 1000\\rsun$, to within $\\sim 10\\%$. Because the giant evolutionary time is so short, $a_f$ is essentially the preexplosion separation of black hole and donor.   We thus have three classes:  \\begin{enumerate}  \\item[(i)] The 8 AML (angular momentum loss) SXTs with K or M-star main sequence companions come from binaries which overfill their Roche Lobe during spiral in, as discussed in LBW. Their periods are decreased as they transfer mass to the companion black hole, as they lose angular momentum by magnetic braking and gravitational waves.  \\item[(ii)] The next six SXTs which established Roche contact while in main sequence, some of them having evolved beyond.  \\item[(iii)] The special case of the continuously shining \\astrobj{Cyg X-1} which we place just before its Roche Lobe, the companion now undergoing unstable mass transfer to its lower mass companion black hole.  \\end{enumerate}  Interestingly, we find that the division between the unevolved main sequence class (i) and evolved companion (ii) is given accurately by Fig.~2 of de Kool et al., who plot the mass of the companion which undergoes angular momentum loss by gravitational waves and magnetic braking, both as functions of time. They obtain the companion mass of $2\\msun$ as giving the division. We find this to be true for $M_{\\rm giant} =20\\msun$ in Eq.~(\\ref{eq1}).  Note that the binaries with late main sequence companions \\astrobj{Nova Scorpii} and \\astrobj{IL Lupi} are special in that these binaries have experienced large mass loss, which can be explained as in LBW by magnetohydrodynamic effects, not included in the evolution discussed here. One may wonder why just there two binaries, with ZAMS companion masses, have lost a sizable fraction of their progenitor He star masses, whereas there is no sign of a kick outwards in separation from copious mass loss in the Class (i) SXTs.   The above regularity shows immediately that at least the first class with K and M companions must have a very large $a_i$ (LBW find $a_i\\sim 1700\\rsun$ for a $20\\msun$ black hole progenitor with $1\\msun$ companion) corresponding to a giant radius of $\\sim 1000\\rsun$, so that the binding energy of the giant envelope, which decreases inversely with its radius, is small enough to be furnished by the drop in gravitational binding energy of the low-mass companion as it spirals in to its Roche Lobe.  As developed in many papers by Brown and collaborators, most lately in Lee \\& Brown (2002) and LBW, and backed by evolutionary calculations by Brown et al. (2001) and PRH, the above delineation into three classes, depending upon companion mass, can be understood if the giant is required to finish (or nearly finish) He core burning before common envelope evolution takes place; i.e., if the mass transfer is essentially Case C.   In Sec.~\\ref{caseC}, we discuss the stellar evolution necessary to produce models which allow Case C mass transfer for ZAMS $20-30\\msun$ stars. In Sec.~\\ref{sec_dev}, we review the role of carbon burning. In Sec.~\\ref{popul}, we discuss that Case B mass transfer would not only allow too much of the He envelope to blow away and leave too much $^{12}$C after He core burns, but also is disfavored by the population of SXTs. In Sec.~\\ref{sec_pop}, we compare our approach with other works, especially those by PRH, and discuss population synthesis of Case B mass transfer. We summarize our conclusion in Sec.~\\ref{disc}. ",
        "conclusions": "\\label{disc}  The work of LBW (Lee, Brown, \\& Wijers 2002) has been amalgamated with that of PRH (Podsiadlowski et al. 2002). Both papers agree that common envelope evolution must come following helium core burning; i.e., be Case C. Our work is based on the nearly linear relationship Eq.~(\\ref{eq1}) between separation following common envelope evolution and companion (donor) mass, which is spelled out in our Appendix. By choosing the common envelope parameter $\\lambda\\alpha_{ce}=0.2$, we are able to evolve those binaries with K and M-star companions, which we believe to be the success of LBW and our earlier works.  The PRH evolutions clarify that all of the binaries with the possible exception of \\astrobj{Cyg X-1} could have made Roche contact in main sequence evolution of the companion, as in the earlier work by Beer \\& Podsiadlowski (2002). This makes it possible to evolve all of the binaries with sub-Eddington rate of mass transfer (although we do not believe this to be necessary in the case of black holes).  LBW and PRH agree that the present evolutionary tracks of supergiants from ZAMS masses $\\sim 20-30\\msun$ and possibly greater, must be changed so as to allow Case C mass transfer. An example of how to do this was constructed (by hand) in LBW.   We would like to point out that there are uncertainties due to the radius evolution of massive stars and the parameterization of mixing, etc. In particular, stars that use the LeDoux criterion or a small amount of semiconvection burn helium as red supergiants, while those with Schwarzschild or a lot of semiconvection stay blue much of the time. Rotationally induced mixing also plays a role as in Langer and Maeder (1995)."
    },
    "0311/astro-ph0311253_arXiv.txt": {
        "abstract": "The results of a numerical study of the growth of solid particles, ranging from 1 micron to 1 millimeter in size, in the vicinity of an azimuthally symmetric density enhancement of a protostellar disk are presented. It is shown that the combined effect of gas drag and pressure gradients, which causes solid objects to rapidly migrate toward the location of the local maximum density, can also enhance the rate of growth of dust particles to larger objects. The results of numerical simulations of such growth processes are presented and the effects of the changes of the physical properties of particles on the rates of their growth are also discussed. ",
        "introduction": "The focus of this paper is on the study of the growth rate of dust grains in the vicinity of a local maximum density of a nebula. It has recently been shown that in a nebula with local gas density enhancements, the combined effect of gas drag and pressure gradients causes solid objects to rapidly migrate toward the location of maximum density (Haghighipour $\\&$ Boss, 2003a$\\&$b, hereafter HB03a$\\&$b). In those papers, the authors have studied the dynamics of a single object in a non-uniform gaseous disk with no particulate background material. The only interaction considered in those papers between the object and the disk was through the gas drag. If the disk contains a mixture of gas and small solid particles, however, the motion of an object will also be affected by its interaction with this particulate material. For instance, an object may sweep up smaller particles and grow larger during its gas-drag induced migration. In this paper, I consider particulate material coupled to a gaseous disk with a local density enhancement, and study the rate of growth of small particles while sweeping up the particles of the background material.  For the past four decades, the growth of dust particles in a solar nebula was studied by many authors. In 1969, Safronov studied the growth of dust particles during their vertical descent in a quiescent nebula and showed that, considering their sticking growth, settling of dust particles could be quite rapid. Safronov's formulation indicated that dust particles could grow to 1 centimeter in size at a 1 AU radial distance from the Sun, assuming a 100$\\%$ sticking probability \\citep{Safronov}. Growth and sedimentation of dust particles have also been studied extensively by \\citet{Weid80,Weid84} and \\citet[][1986]{Nakagawa81}. \\citet{Weid80,Weid84} considered coagulation of dust grains due to van der Waals force in their settling toward the midplane and showed that, while such particles grow to 1 centimeter in size, they create a turbulent layer in the vicinity of the midplane which increases the rates of their collision and their coalescence to larger objects. \\citet{Nakagawa81} studied the growth and settling of dust grains at the location of Earth in a nebula and showed that, similar to the results presented by Safronov, the settling of particles and the formation of a layer of centimeter-sized objects in the vicinity of the midplane can occur in a relatively short time ($\\simeq$300 years). This study was later extended by \\citet{Nakagawa86} to the regions of Jupiter and Neptune in which the authors studied the settling and growth of dust particles in laminar phase of a nebula and presented analytical expressions for the rates of the growth and also vertical and radial motions of solids. By treating a dust layer as a two-dimensional fluid, \\citet{Nakagawa86} showed that the total settling time of dust grains in the two stages of gas- and dust-dominant phases are smaller than those estimated by \\citet{Weid80} and \\citet{Nakagawa81}. More recently, \\citet{Sup00} have studied the coagulation of small solids in a turbulent protostellar disk with an application to the formation of icy planetesimals in the outer region of the nebula.  In this paper, I extend the previous studies presented in HB03a$\\&$b by considering the growth of a dust particle due to interaction with the particulate background of a nebula in the vicinity of its local density enhancement. The main purpose of this paper is to study how the combined effect of gas drag and pressure gradients can speed up the coagulation of dust particles and the formation of centimeter-sized objects.  The model nebula considered in this paper is presented in $\\S$ 2. Section 3 presents the equations of motion of a solid particle and  $\\S$ 4 deals with the analysis of the numerical results. Section 5 concludes this study by reviewing the results and discussing their applications. ",
        "conclusions": "The results of the numerical study of the growth rates of small particles in the vicinity of a local density enhancement of a solar nebula were presented. It was shown that, as a consequence of migration toward the location of a local density enhancement, solid objects can rapidly grow in size by sweeping up the smaller particles of the background material. Particles of interest in this study had sizes between 1 to 1000 micron and their growth-rates were studied until they reached 1 cm in radius. The results presented here indicate that, compared to the time of settling and growth of dust particles in a nebula without local gas density enhancements \\citep{Safronov,Weid80,Nakagawa81,Weid88}, in the vicinity of a local maximum gas density, the combined effect of gas drag and pressure gradients can increase the growth rates of dust grains by one order of magnitude.  In this study the focus was on the growth of small grains to objects of 1 cm in size. Although centimeter-sized objects can still acquire mass and grow larger by sweeping up smaller particles, their dynamics is not solely driven by gas drag and the sweeping process. The velocities of such particles are large enough to cause them to fragment and produce dust when they collide. On the other hand, as noted by \\citet{Safronov}, \\citet{Weid80}, \\citet{Nakagawa81}, \\citet{Nakagawa86}, and \\citet{Weid88}, the largest that dust particles can grow to while settling toward the midplane is centimeter-sized. The concentration of centimeter-sized particles in the vicinity of the midplane creates a turbulent layer \\citep{Weid80} which requires a more detailed analysis. All these introduce complexities in treating centimeter-sized objects that are unnecessary for the purpose of this study. Studies are, however, underway in which the mutual interactions of centimeter-sized particles and the effect of turbulence on their collision and coalescence have been considered.  The spatial density of the particulate background of the nebula was chosen to be proportional to the gas distribution function. The solid/gas ratio $f$ was chosen to be constant and equal to 0.0034. This is a minimal value for $f$ that corresponds to a dust composition of metals and silicates on the midplane of a standard nebula at 1 AU \\citep{Pod74,Weid80,Weid88}. In a more realistic scenario, the value of $f$ is a function of position and time. For instance, in the model nebula considered here, as time passes, the background material of the nebula settles toward the midplane and also migrates toward $r=1$ AU, where the gas distribution function maximizes. This increases the dust/gas ratio around $r=1$ AU on the midplane and decreases $f$ at other locations. In such cases, one has to take the time variation of $f$ into consideration. However, because in this study, the rates of radial and vertical migration of the background material are smaller than those of the objects of interest, and also because it has been assumed that the sticking probability is 100$\\%$, which results in faster growth of the particle, growing to 1 cm in size occurs so rapidly that the particles of the background will not have enough time to undergo appreciable radial and vertical migrations.  As mentioned earlier, the focus of this study has been on the growth of particles to 1 cm in size. For such small objects, the assumption of growth merely due to sweeping up smaller particles is quite plausible since these objects are coupled to the gas and they approach one another slowly. Once the sizes of objects increase 1 cm, their collisions may result in adhesion and/or fragmentation. As a result, their growth will not be entirely due to sweeping up of smaller particles and will be affected by their individual velocities. Replacing the relative velocities of such objects with their velocities relative to the gas is no longer a valid approximation. In such cases, one has to follow individual particles and consider their velocities relative to one another. Such an extension of this work is currently in preparation for publication."
    },
    "0311/astro-ph0311409_arXiv.txt": {
        "abstract": "We present year-scale morphological variations of the Crab Nebula revealed by {\\it Chandra} X-ray observatory. Observations have been performed about every 1.7 years over the 3 years from its launch.  The variations are clearly recognized at two sites: the torus and the southern jet. The torus, which had been steadily expanding until 1.7 years ago, now appears to have shrunk in the latest observation. Additionally, the circular structures seen to the northeast of the torus have decayed into several arcs. On the other hand, the southern jet shows the growth of its overall kinked-structure. We discuss the nature of these variations in terms of the pulsar wind nebula (PWN) mechanism. ",
        "introduction": "The 5-month monitoring observations of {\\it Chandra} and HST showed short-term (days to weeks) morphological variations of the Crab Nebula (Mori et al. 2002; Hester et al. 2002). A variable inner ring and wisps emerging from it visually revealed the existence of the termination shock of the pulsar wind and its downstream flow, respectively. In addition to such short-term, therefore more noticeable, variations, expansion of the outer torus on longer time scale (months) was also discovered. However, this seems to contradict with the fact that the angular extent of the torus is almost constant over 25 years (Mori et al. 2002); if the torus kept expanding at the observed rate, it would have become almost twice larger during this period.  On the other hand, the southern jet did not show a strong variation of its overall structure. This differs from the highly variable jet of Vela (Pavlov et al. 2003), which is an another famous example of a variable PWN.  Here we present long-term (years) morphological variations of the Crab Nebula which are newly discovered comparing {\\it Chandra} observations over 3 years. ",
        "conclusions": "In contrast with short-term variations are seen at inner regions, the long-term variations are prominent at outer regions: the torus and the southern jet. In the following sections, we will discuss these variations separately.  \\subsection{The northeast of the torus}  \\begin{figure}[htb] \\centerline{\\psfig{file=morik1_1.eps,width=11cm}} \\caption{Expanded views around the northeast of the torus of 1st (top left), 2nd (top right), and 3rd (bottom left) epoch observations. Black lines seen in the 3rd epoch image are instrumental effects.} \\label{figure:torus} \\end{figure}  Figure~\\ref{figure:torus} shows expanded views of the northeast of the torus. The brightness scales of these images are normalized to the bright region of the torus, not by exposure time, because of the differences of the observational configurations.  Images of the 1st and 2nd epoch observations are quite similar. The bright ``thick'' torus is encircled by circular structures at its northeastern end. As reported previously, these two structures appear to expand between the 1st and 2nd epoch observations (Mori et al. 2000; Mori 2002). In contrast with the similarity between the 1st and 2nd epoch images, the morphological transition from the 2nd to the 3rd epoch image is remarkable.  The bright region of the torus appears to have became shrunken and ``thin'', and the circular structures have decayed into several arcs.  Now it is clear that the torus is not steadily expanding; the overall extent is almost constant over decades, but the boundary of the torus varies with an angular scale of a few arcseconds and a time scale of years. It is as if we were seeing a top of a fountain. Greiveldinger \\& Aschenbach (1999) reported surface brightness variations of the Crab Nebula using $ROSAT$ HRI observations spanning 6 years. Although $ROSAT$ could not detect a morphological variation, they showed a monotonic increase of the surface brightness at the northeastern region of the torus. The brightness variation discovered by $ROSAT$ is most likely related to the morphological variation discovered by \\cxo, suggesting that the time scale of the variation of the torus might be about a decade rather than a few years.   \\subsection{The southern jet}  \\begin{figure}[htb] \\centerline{\\psfig{file=morik1_2.eps,width=11cm}} \\caption{Expanded views of the southern jet of 1st (top left), 2nd (top right), and 3rd (bottom left) epoch observations. The line in each image traces the axis of the jet. All three lines are superimposed in the bottom right image.} \\label{figure:jet} \\end{figure}  Figure~\\ref{figure:jet} shows expanded views of the southern jet of the 1st, 2nd, and 3rd epoch observations. Although the displacement of the jet is quite small, the series of images clearly shows the growth of its overall kinked-structure. The variability of the jet is reminiscent of northern variable jet of Vela (Pavlov et al. 2003). Table~\\ref{table:jet} compares the time scale of the variability and the width between the Crab and Vela Jets. The Crab jet is 10 times larger and varies 10 times slower than the Vela jet. If these variations are due to MHD instability, the time scale is proportional to the Alven crossing time, $\\tau \\sim r/\\nu$, where $r$ is the width of the jet and $\\nu$ is the Alven velocity (Begelman 1998). Considering that the Alven velocity in an ultrarelativistic plasma like the Crab and Vela PWN is more or less the same, above equation applies to both the Crab and Vela jet. Therefore, it is suggested that a common mechanism is responsible for the variability of the Crab and Vela jet.  \\vspace{-0.5cm} \\begin{table}[htb] \\begin{center} \\caption{Comparison of variable jets of the Crab Nebula and Vela} \\label{table:jet} \\begin{tabular}{lcc} \\tableline \\tableline & Time scale (day) & Width (cm)  \\\\ \\tableline Crab & 150-500 & 2.9 $\\times$ 10$^{17}$ \\\\ Vela & 10-30 & 3 $\\times$ 10$^{16}$ \\\\ \\tableline \\end{tabular} \\end{center} \\end{table}  \\vspace{-0.8cm}"
    },
    "0311/astro-ph0311123_arXiv.txt": {
        "abstract": "We investigate multicolour imaging data of a complete sample of 19 low redshift ($z<0.2$) quasar host galaxies. The sample was imaged in four optical (\\textsl{BVRi}) and three near-infrared bands (\\textsl{JHK}$_s$). Galaxy types, structural parameters and robust host galaxy luminosities are extracted for all bands by means of two-di\\-men\\-sio\\-nal deblending of galaxy and nucleus. For the disc dominated fraction of host galaxies (Sa and later) the optical and optical-to-NIR colours agree well with the average colours of inactive galaxies of same type. The bulge dominated galaxies (E/S0) on the other hand appear a significant 0.35~mag bluer in $(V-K)$ than their inactive counterparts, being as blue as the discs in the sample. This trend is confirmed by fitting population synthesis models to the extracted broad band SEDs: the stellar population age of the bulge dominated hosts lies around a few Gyr, much younger than expected for old evolved ellipticals. Comparison to other studies suggests a strong trend for stellar age in elliptical host galaxies with luminosity. Intermediately luminous elliptical hosts have comparably young populations, either intrinsically or from an enhanced star formation rate potentially due to interaction, the most luminous and massive ellipticals on the contrary show old populations. The correspondence between the nuclear activity and the blue colours suggests a connection between galaxy interaction, induced star formation and the triggering of nuclear activity. However, the existence of very symmetric and undisturbed disks and elliptical host galaxies emphasised that other mechanisms like minor merging or gas accretion must exist. ",
        "introduction": "Since the discovery of massive black holes in the centres of normal galaxies and an identical relation between bulge and black hole mass for active and inactive, quasar host galaxies start to be accepted as not-so-different members of the galaxy population than previously thought. The question remains how different quasar hosts are from `normal' galaxies. Is a quasar phase part of every host galaxies life? The connection of the masses of galactic bulges and central black holes suggests a common evolutionary path. Nuclear activity might thus be a special phase of black hole fuelling. Obvious candidates for triggering active galactic nuclei (AGN) include major and minor merger in the scenario of hierarchical clustering, as well as internal instabilities. The effects of these events in the shape of nuclear activity and, potentially, enhanced induced star formation thus might give an opportunity to study a generic phase in the evolutionary cycle of galaxies.  Single optical or near infrared (NIR) bands are generally sufficient to characterise morphological properties like galaxy types, host and nuclear luminosities, or to conduct environment studies \\citep[e.g.][]{mcle95b,mclu99}. It was also attempted to detect and quantify events of merging or interaction from pure morphology of the visible host galaxy \\citep[e.g.][]{hutc92}.  To assess the dominant stellar populations in the host, spectral information is needed. While one optical colour already permits a characterisation of young stellar populations, the restframe NIR on the other hand holds information about the old stars, dominating the host in mass.  Ideally, spectra of the host galaxy would be needed, spanning all the optical and NIR domains. However, due to the unavoidable contamination of the final host galaxy spectrum with a nuclear component, increasing with decreasing angular resolution, spectroscopy is so far of limited use for luminous quasars (QSOs), already at low redshifts.  With this study we attempt to assemble spectral information on quasar hosts by combining optical and NIR broad band data \\citep[as first performed by][]{koti94}. At the price of a coarse sampling of the SED compared to spectroscopy, the deblending of nuclear and stellar components in broad band images is by now a robust and reliable technique.  In the present paper we present multicolour data for a complete sample of low redshift quasars. The sample selection and observations are described in Section~2. The deblending into nuclear and host contributions is presented in Section~3, followed by a description of the photometry (Sect.~4). In Section~5 we discuss the results, including notes on individual objects. The derived host galaxy colours are compared to those of inactive galaxies in Section~6, while Section~7 contains results from simple stellar population model fits. In Sections~8 and 9 we present an overall discussion and some conclusions.  Throughout this article we use a cosmology with $H_0=50$~km~s$^{-1}$~Mpc$^{-1}$, $q_0=0.5$ and $\\Lambda=0$. ",
        "conclusions": "Are QSO host galaxies just normal galaxies with an active nucleus? We find strong evidence that at least a subpopulation has significantly different colour properties than their inactive counterparts. While the disc dominated or spiral galaxies in our sample show statistically no difference in colour and age to inactive discs, the ellipticals exhibit abnormally blue colours and SEDs consistent not with old, evolved stellar populations.  Even though elliptical galaxies can be significantly more luminous than disc galaxies, the abnormally blue colours found are not a function of only luminosity. Both for our sample as well as that by \\citeauthor{scha00}, the two classes of bulge vs.\\ disc dominated hosts contain objects of similar luminosity (see Figs.~\\ref{fig:vh_disk_ellipse} and \\ref{schade}). This suggests that the main difference lies in the {\\em morphological type} of the host galaxy. However, at some point a luminosity -- or mass -- dependence has to exist, since the more massive elliptical host galaxies studied by \\citeauthor{dunl03} show old evolved stellar populations and signs of only small recent star formation.  The correspondence between the QSO activity and the blue colours suggests a connection between galaxy interaction, induced star formation, and the onset of nuclear activity. However, we find that it is likely that mechanisms other than major mergers, such as minor merger or gas accretion, are responsible for the activity of both the high luminosity ellipticals as well as the disc-dominated host galaxies.  At this point the picture of how activity and interaction are connected is not clear. However, the abnormally blue colours of a sub-population are a strong indicator that for spheroid-dominated host galaxies, merging is an important factor in triggering nuclear activity."
    },
    "0311/astro-ph0311315_arXiv.txt": {
        "abstract": "{ The quantitative spectroscopy of stellar objects in complex environments is mainly limited by the ability of separating the object from the background. Standard slit spectroscopy, restricting the field of view to one dimension, is obviously not the proper technique in general. The emerging Integral Field (3D) technique with spatially resolved spectra of a two-dimensional field of view provides a great potential for applying advanced subtraction methods. In this paper an image reconstruction algorithm to separate point sources and a smooth background is applied to 3D data. Several performance tests demonstrate the photometric quality of the method. The algorithm is applied to real 3D observations of a sample Planetary Nebula in M31, whose spectrum is contaminated by the bright and complex galaxy background. The ability of separating sources is also studied in a crowded field in M33. ",
        "introduction": "\\vspace{-0.3cm} 3D spectroscopy becomes more and more an important tool not only for spatially extended objects, but also for point sources. This may sound controversial at first. But the more the development of modern telescopes enable access to spectra of individual objects in nearby galaxies, the more the problem of crowding of stars and subtraction of the bright and spatially varying background become limiting factors. Standard slit spectroscopy has a strong disadvantage, as it restricts the field of view to one dimension. Additionaly adverse slit effects are intrinsic limitations for this technique, e.g.:\\\\ The slit width affects the spectral resolution and the amount of background, which is combined with the object spectrum. On the other hand, when minimizing the slit width, light losses limit the photometric quality of the object spectrum. Atmospheric refraction effects, causing a shift of the object position with wavelength, a wavelength dependant point spread function (PSF), as well as pointing inacurracies are additional factors causing potential light losses differrentially with wavelength.  The 3D technique can circumvent the instrumental limitations imposed by the slit spectroscopy. It provides the full spectral and spatial information of an observed object and its environment. Interpreting a 3D data cube as a collection of monochromatic images of the identical object rather than a spatially distributed collection of spectra allows the usage of image restoration algorithms to separate the object from its crowded field environment in each wavelength bin. The images are connected in the sense that the shape of the PSF (Fried 1966) as well as the object position (Filippenko 1982) vary as a function of wavelength. Provided the algorithm works with high photometric accuracy, the resulting object intensities of all wavelength bins can be combined to yield a background cleaned object spectrum.\\\\ The first popular image restoration algorithm was made available by Richardson (1972) and Lucy (1974). With an a priori known PSF, this iterative algorithm sharpens blurred images. However, like similar maximum likelihood algorithms, it does not preserve the photometry, but typically shows large residuals, especially around point sources.\\\\ In the 1990s much work to improve the RL-algorithm was done at ST-ECF (e.g. Lucy \\& Hook 1991, Hook \\& Lucy 1993). Lucy (1994) introduced the {\\em cplucy} or {\\em 2 channel} algorithm which allows the reconstruction of point sources with high photometric accuracy. Instead of a blind deconvolution as performed by the RL-algorithm, the procedure simultaneously models the intensity distribution of a discret number of point sources with {\\em known} positions and the background with an assumed smoothness controlled by an additional entropy term. For a detailed description see Lucy (1994), Hook et al. (1994). In the following sections {\\em cplucy} and its application to 3D data will be studied. The algorithm will be applied to real 3D data of a point source on a complex but smooth background as well as to a star in a crowded field, surrounded by an emission nebula. ",
        "conclusions": "\\vspace{-0.3cm} We have demonstrated that the cplucy algorithm can be used to improve the image quality of 3D data. Applied to datacubes of a faint, background limited planetary nebula in M31, it is shown that systematic errors of emission line intensities, which seem to be unavoidable with the limited geometry of slit spectroscopy, can be eliminated by using the knowledge of the PSF. The absence of a measurable continuum of the PN is a favourable case, since it allows one to model the background continuum in between the PN emission lines. Subtracting the background model prior to iterating cplucy helps to make sure that the algorithm is converging with the same speed in each wavelength bin. With an example of the luminous star B416 in M33, located in a gaseous emission region, we have also shown that the cplucy  method is even useful in cases where the image quality does not seem to warrant any deconvolution effort, helping nevertheless to disentangle the stellar PSF from the gaseous background component. \\vspace{-0.3cm}"
    },
    "0311/astro-ph0311178_arXiv.txt": {
        "abstract": "The Sunyaev-Zel'dovich (SZ) effect expected to be associated with massive star-formation activity produced during the formation of the most luminous bulges of normal galaxies is discussed.  Using 1-D chemohydrodynamical models for spheroidal galaxy evolution we  show that during the early epochs of galaxy evolution the gas in massive events of star formation may reach temperatures and densities high enough to produce values of the comptonisation parameter \\y comparable to those present in galaxy clusters.  In this scenario, we discuss the possibility of detection of the SZ signature in high redshift starforming galaxies with the next generation of mm telescopes capable of arcsecond resolution and equipped with high sensitivity detectors.  We show how millimeter colour--colour diagrams or diagnostic diagrams could be used to distinguish between the dust emission and the SZ effect and suggest the use of simultaneous  multifrequency observations to improve the chances of detecting the SZ effect. ",
        "introduction": "The SZ effect (Sunyaev \\& Zel'dovich 1972) is produced when the cosmic microwave background (CMB) photons from the Rayleigh-Jeans region, due to the interaction with the fast electrons, move to the Wien tail of the CMB spectrum thus producing a unique spectral signature. This signature shows an increase with respect to the mean CMB brightness of the observed intensity for wavelengths shorter than 1.34 mm and a decrease for wavelengths longer than 1.34 mm. This is known as the thermal SZ effect. Another component is the kinetic SZ effect, due to the interaction with CMB photons of plasma coherently moving with respect to the reference frame where the CMB is isotropic (e.g. Rephaeli 1995, Birkinshaw~1999, Church, Jaffe and Knox~2001).  The SZ effect constitutes a powerful tool to study the physical properties of the intercluster gas, and when combined with  X-ray  observations, also to  estimate  cosmological constants such as the Hubble constant H$_0$, the current mass density of the universe $\\Omega_{\\rm M}$ or the cosmological constant $\\Lambda$ (e.g. Majumdar 2001, Diego et al. 2002, Levine, Schulz \\& White 2002, Reese et al. 2002, Majumdar \\& Mohr 2003).   It is understood that  in clusters the thermal SZ effect due to the interaction of CMB photons with electrons with temperatures around 10$^8$K dominates over the kinetic SZ effect. This fact is evidenced by combining observations at 1 mm where the thermal effect is close to the maximum, at 1.4 mm where it is almost zero and the kinetic SZ effect can be marginally detected and at 2 mm where the thermal SZ has its minimum and is negative (e.g. Lamarre et al. 1998, Mauskopf et al. 2000).  The advent of a new generation of mm telescopes (LMT/GTM or ALMA, described in section~\\ref{DetTel}) that combine relatively high angular resolution with high sensitivity will allow to explore the possibility of detecting the SZ effect in volumes much smaller than that of a cluster of galaxies. Would it be possible for these new instruments to detect the SZ signature in individual galaxies?  Natarajan \\& Sigurdsson (1999), Rosa-Gonz\\'alez et al.~(2000, 2001), Majumdar, Nath and Chiba (2001) and Rosa-Gonz\\'alez (2002) have indeed explored the detactability of the SZ effect in individual galaxies.  Natarajan \\& Sigurdsson (1999) studied the SZ effect due to galactic outflows powered by the mechanical energy provided by the accretion of matter to the central supermassive black hole in QSOs. Majumdar et al. (2001) propose the evolution of  supernova driven galactic winds during the early stages of evolution of a normal galaxy, as the cause of the distortion of the CMB radiation due to the kinetic SZ effect. A similar model was originally proposed by Tegmark, Silk and Evrard (1993) to explain the absence of absorption lines in the observed spectra of high redshift quasars (Gunn-Peterson effect, Gunn \\& Peterson~1965). The explosion model assumes an initial energy input which  is equal to the 2\\% of the total luminosity generated in the supernova explosions during about 5$\\times 10^7$ years. The range of galaxy masses covered by this model goes from 5$\\times 10^7$\\Msolar\\ to  $10^{11}$\\Msolar.  Rosa-Gonz\\'alez et al.~(2000, 2001) and Rosa-Gonz\\'alez (2002) discussed the possibility of observing the  SZ effect in the centre  of young giant spheroidal galaxies. They include in the calculations both the thermal and the kinetic SZ effect.  In  the {\\it accepted} galaxy formation scenario described by the hierarchical model (e.g. Navarro, Frenk \\& White 1995), massive galaxies form after merging with smaller system that have been formed in a previous epoch. However, there is observational evidence that a significant number of L$^\\ast$ galaxies  exist at redshift of about 5 and beyond (Hu, Cowie \\& McMahon 1998, Stanway, Bunker \\& McMahon 2003). Jimenez et al. (1999) discussed the possibility of the existence of massive elliptical galaxies at high redshift, contrary to early predictions of the hierarchical models of galaxy formation. The study of massive galaxies at high redshift that could be detected through the unique SZ signature could be used to constrain the  different galaxy formation scenarios.  The SZ effect due to young galaxies together with radio sources (Holder 2002) and dusty galaxies (Blain, Ivison \\& Smail~1998, Blain~1998) are potential sources of noise in sensitive studies of SZ clusters and upcoming measurements of the CMB fluctuations (e.g. {\\it FIRST}, {\\it Plank}). Therefore, their characteristic emission and physical scale must be quantified.  In this paper we explore in detail the possibility of detection of the SZ effect in young galaxies undergoing a massive starformation event. We discuss the confusion due to dust emission and suggest a diagnostic to segregate between dust emission and SZ effect distortions.  \\subsection{Thermal SZ effect} \\label{SZbasic}   The SZ effect describes the brightness change of the CMB due to the inverse Compton scattering by hot electrons~\\cite{1972Sunyaev}.  The change in intensity between the unaffected CMB radiation  and the radiation which go through the hot gas can be expressed by the integral, \\begin{equation}\\label{eq:Int} \\Delta I = \\tau_e \\frac{2h}{c^2}\\int_{-\\infty}^{+\\infty} ds P(s) \\left[ \\frac{\\nu_0^3}{\\exp{\\left(\\frac{h\\nu_0}{kT_{rad}}\\right)}}  - \\frac{\\nu^3}{\\exp{\\left(\\frac{h\\nu}{kT_{rad}}\\right) }} \\right] \\end{equation} where the energy of the CMB photons is given by $h\\nu_o$, the energy of the photons which come out from the cluster is $h\\nu$, $s$ is defined by $s=log(\\nu/\\nu_0)$, $\\tau_e$ is the optical depth for scattering and P(s) is the probability that an incident photon with energy $h\\nu_o$ becomes a photon with an energy $h\\nu$ \\cite{1999Birkinshaw}.  The integral in the Equation~\\ref{eq:Int} can be calculated by using the non relativistic expression given by the Kompaneets approximation \\cite{1957Kompaneets}, \\begin{equation} \\label{eq_SZII} \\Delta I(x) = h(x) y I_0 \\end{equation} \\noindent where the distortion of the photon spectrum, $h(x)$ is defined by \\begin{equation}\\label{eq_SZ} h(x) = x^4 \\frac{e^x}{(e^x-1)^2} \\left( x\\ \\hbox{coth}\\,\\frac{x}{2}-4\\right) \\end{equation} $x$ is the dimensionless frequency, $x=h\\nu/(k T_{rad})$ with $T_{rad}=2.728$ the temperature of the CMB, \\begin{equation}\\label{I0} I_0 = \\frac{2h}{c^2}\\left(\\frac{k_B T_{rad}}{h} \\right)^3 \\end{equation} and the Comptonization parameter ``\\y '' is  defined by the integral of the electron pressure along the line of sight \\begin{equation} \\label{y_eq} y=\\sigma_T\\int dl\\ n_e\\frac{k_B T_e}{m_e c^2} \\end{equation} In this equation $\\sigma_T$ is the Thomson scattering cross section, n$_e$ and $T_e$ are  the electron number density and the electron temperature respectively, $k_B$ is the Boltzmann constant and $m_e c^2$ is the rest energy of an electron; (see Rybicki \\& Lightman~1980 for a more detailed derivation of the different expressions described above and Birkinshaw 1999 for a recent review). The effect of the hot gas on the CMB spectrum is to produce a fractional change in the brightness of the CMB radiation. We remark that the amplitude of the SZ effect is  redshift-independent and depends only on the integrated  properties of all the hot gas along the line of sight. The SZ effect can be considered as a tool to detect high pressure electron gas over cosmological distances.   \\begin{figure} \\setlength{\\unitlength}{1cm} \\begin{picture}(7,6) \\put(-1.2,-6.7){\\special{psfile=TherKin.ps hoffset=0 voffset=0 hscale=50.0 vscale=50.0 angle=0}} \\end{picture} \\caption{\\label{TerKin} The kinetic SZ effect (solid and dot-dashed lines) signature is compared to the thermal SZ effect (dashed line) for the case where the kinetic effect is  half of the thermal one. The two cases of the kinetic SZ effect correspond to gas moving away from (dot-dashed line) and towards (solid line) the observer.} \\end{figure}   \\subsection{Kinetic SZ effect}   \\begin{figure*} \\setlength{\\unitlength}{1cm} \\begin{picture}(7,6) \\put(-3.5,-11.5){\\special{psfile=DyTprof.ps hoffset=0 voffset=0 hscale=65.0 vscale=65.0 angle=0}} \\end{picture} \\caption{\\label{Profi} Electron density and temperature profiles for the model of baryonic mass 2x10$^{12}$\\Msolar. The different lines correspond to different times: 20 Myr for the solid line, 600 Myr for the dotted line, 1 Gyr for the dashed line, 4 Gyr for the dot--dashed line and 11 Gyr for the dot-dot-dashed line.} \\end{figure*}   The kinetic SZ effect, as already mentioned, is due to the movement of the plasma with respect to the CMB rest frame. For a shell of matter which is moving with a radial velocity $v$ with respect to the Hubble flow, the change of the intensity due to the kinetic SZ effect ($\\Delta I_{kin}$) is given by, \\begin{equation} \\label{KinEq} \\Delta I_{kin} = - \\frac{v}{c} \\tau_e I_0 \\frac{x^4e^x}{(e^x - 1)^2} \\end{equation} where $\\tau_e$ is the scattering optical depth,  $I_0$  is given by equation~\\ref{I0} and $x$ is the dimensionless frequency.  At low frequencies, the ratio between the kinetic SZ effect ($\\Delta I_{kin}$) to the thermal SZ effect ($\\Delta I_{th}$) is independent of the scattering optical depth and can be calculated by, \\begin{equation} \\label{KTratio} \\frac{\\Delta I_{kin}}{\\Delta I_{th}} = \\frac{1}{2}\\frac{v}{c}\\frac{m_e c^2}{kT_e} \\end{equation} where $v$ is the velocity of the flow, $m_e c^2$ is the electron rest energy and $T_e$ is the gas temperature~\\cite{1999Birkinshaw}. The signatures of  the thermal and kinetic SZ effects are plotted in Figure~\\ref{TerKin}. ",
        "conclusions": "In this paper we have explored the possibility that the central comptonization \\y parameter in young spheroidal galaxies or galaxies undergoing a massive starforming event may reach values comparable to those  observed in galaxy clusters.  However due to the poor angular resolution-sensitivity combination  of present day mm facilities, the detection of the SZ effect in individual galaxies may only be  possible with the next generation of high angular resolution - high sensitivity mm telescopes (e.g. GTM/LMT or ALMA) and then probably only in the most massive systems capable of producing a 800~$\\mu$m flux of more than a few $\\mu$Jy for periods of time of several hundred megayears.    We have included instrumental noise in our calculations, but neglected the noise contribution caused by high redshift millimeter background point sources (e.g. Blain 1998, Hughes et al. 1998).  The confusion due to unresolved sources depends on the actual details (i.e. resolution) of the observation, the (unknown) emission radial profile of these sources and the extrapolation of the number counts to fainter fluxes. Blain et al. (2000)  give an estimation of the confusion limit of about  10~$\\mu$Jy for the case of  observations at 1 mm and a beam size of 5\\arcsec. This value is comparable to the expected SZ flux of galaxies with baryonic masses of 2$\\times 10^{12}$\\Msolar\\ and an age of $\\sim 3\\times 10^{8}$ years (Figure~\\ref{TotFlux}). Confusion will be a major problem when combining independent observations at different frequencies but its effect will be greatly reduced using  instruments that simultaneously observe at three selected frequencies.    A major difficulty in the detection of the SZ effect in an individual galaxy will be its confusion with the dust emission originated in the same star formation event. To solve the dust contamination problem we suggest to exploit the particular spectral signature of the  SZ spectrum and maximize its segregation from dust emission by using simultaneous multifrequency observations or differential detection at 1, 2 and 3 mm. In the diagnostic diagram presented in Figure~\\ref{ColCol} there is a wide region centred at I(2 mm)/I(1 mm)=--0.3 and I(3 mm)/I(1 mm)=--0.3 where the thermal SZ emission can be clearly separated from the dust emission over a wide range of redshifts.  Due to its particular spectrum the detection of the SZ signature in individual galaxies and the discrimination between sources dominated by the SZ effect from those sources dominated by dust emission would be enhanced  by the use of differential mapping systems that simultaneously observe at three selected frequencies (Section~\\ref{DetTel}). Additional advantages of such systems will be the lowering of the minimum detected flux due to a better sampling of the time correlated atmospheric fluctuations (i.e. reducing the sky noise) and the relative absence of source confusion generated by dust emitting sources.  New generation of mm cameras such as SCUBA-2 \\footnote{http: $\\backslash\\backslash$www.roe.ac.uk$\\backslash$atc$\\backslash$projects$\\backslash$scuba\\_two} or SPEED \\footnote{http: $\\backslash\\backslash$www.astro.umass.edu$\\backslash\\sim$fcrao$\\backslash$instrumentation} -- which will have the capability of observing simultaneously in more than a band -- would be ideal to detect compact and faint SZ sources."
    },
    "0311/astro-ph0311208_arXiv.txt": {
        "abstract": "{We describe a project to study the state of the ISM in $\\sim$20 low redshift ($z<0.3$) QSO host galaxies observed with the PMAS integral field spectrograph. We describe method developement to access the stellar and gas component of the spectrum without the strong nuclear emission to access the host galaxy properties also in the central region. It shows that integral field spectroscopy promises to be very efficient to study the gas distribution and its velocity field, and also spatially resolved stellar population in the host galaxies also of luminous AGN. ",
        "introduction": "\\label{intro}  The vast majority of QSO host galaxy studies uses imaging techniques to investigate photometry, the morphological composition, signs for tidal interaction and star formation, companions or clustering. With imaging it is possible to constrain the types of the host galaxies which turned out to be dependent on the luminosity of the nucleus \\citep[e.g.][]{mcle95a}, or vice versa. Using imaging photometry we recently found signs for bluer colours in elliptical host galaxies \\citep{jahn03e} and obtained the first luminosity function of host galaxies \\citep{kuhl01,kuhl03}. Imaging studies showed also e.g.\\ that interaction seems to play a major role in a significant fraction of intermediately luminous host galaxies, but clearly not in some parts of the high luminosity regime \\citep{dunl03}.  Imaging was and is successful from the closest AGN up to high redshifts and by now some standard analysis methods have emerged, like two-dimensional modelling of the surface brightness distribution of nucleus and host galaxy. This allows to determine and remove the nuclear contribution to access the light of the host alone -- being the fundamental problem of host galaxy studies.  Spectroscopy in comparison is used only in a few studies. The application of spectroscopy is even stronger constrained than imaging by the presence of the bright nucleus. Studies are in most cases restricted to the outer part of the host galaxy, away from the light emitted by the nucleus -- so the inner parts of host galaxies were unaccessible. Placing a slit across the nucleus, the obvious choice for an investigation of the inner part of the host, would as a disadvantage have integrated most of the point like nuclear flux, but only very little of the host galaxies and thus leading to an unfavorable contrast. Only very recently, newly developed techniques involving deconvolution \\citep{cour02b} and again spatial modelling of the light distribution \\citep{jahn02} made on-nuclear spectroscopy feasible, with some success.  We want to describe here an approach using integral field units for efficient spectroscopy of quasar host galaxies, combining imaging and spectroscopy in one observation. While the analysis is just fully beginning, we want to present our goals, the data obtained and, in Section~\\ref{methods}, a technique and an example for the removal of the nuclear light, required for studying the inner parts of the host galaxies. ",
        "conclusions": "\\label{summary} It is clear that 3d spectroscopy requires new methods for analysis also for the application to the study of AGN host galaxies. As we described it is possible to remove the light of the active nucleus and to extract information on the host galaxy alone, allowing the use of techniques as for inactive galaxies. In this way 3d spectroscopy becomes extremely powerful for the study of QSO host galaxies. Spatially resolved spectral information can be the key in the search for traces in the dynamic and ionisation state of the ISM in host galaxies, to yield information on the role of merger and interaction in the fuelling process."
    },
    "0311/astro-ph0311514_arXiv.txt": {
        "abstract": "The relatively large Thomson optical depth, $\\tau_{es}$, inferred recently from the WMAP observations suggests that the Universe was reionized in a more complex manner than previously believed.  However, the value of $\\tau_{es}$ provides only an integral constraint on the history of reionization and, by itself, cannot be used to determine the nature of the sources responsible for this transition.  Here, we show that the evolution of the ionization state of the intergalactic medium at high redshifts can be measured statistically using fluctuations in 21 centimeter radiation from neutral hydrogen.  By analogy with the mathematical description of anisotropies in the cosmic microwave background, we develop a formalism to quantify the variations in 21 cm emission as a function of both frequency and angular scale.  Prior to and following reionization, fluctuations in the 21 cm signal are mediated by density perturbations in the distribution of matter.  Between these epochs, pockets of gas surrounding luminous objects become ionized, producing large HII regions.  These ``bubbles'' of ionized material imprint features into the 21 cm power spectrum that make it possible to distinguish them from fluctuations produced by the density perturbations. The variation of the power spectrum with frequency can be used to infer the evolution of this process.  As has been emphasized previously by others, the absolute 21 cm signal from neutral gas at high redshifts is difficult to detect owing to contamination by foreground sources.  However, we argue that this source of noise can be suppressed by comparing maps closely spaced in frequency, i.e. redshift, so that 21 cm fluctuations from the IGM can be measured against a much brighter, but smoothly varying (in frequency) background. ",
        "introduction": "\\label{intro}  One of the long-standing goals of cosmology is to understand how structures have grown through time.  In the usual paradigm, weak density perturbations were imprinted on the Universe during the inflationary era.  These grew through gravitational instability, eventually forming bound halos as well as the cosmic web of sheets and filaments.  Precise measurements of the cosmic microwave background (CMB) anisotropies have fixed the initial conditions of the picture (e.g., \\citealt{spergel03}).  The challenge now is to take structure formation beyond the well-understood linear regime in order to understand how baryons collapsed into the bound objects that we observe today, such as galaxies and galaxy clusters, and to understand how these objects affect their surroundings.  At low or moderate redshifts ($z \\la 6$), galaxies and quasars can be studied in detail with existing technology.  Unfortunately, the first generations of protogalaxies are not yet accessible observationally.  Their properties are nevertheless crucial to understanding both later generations of galaxies, which form out of these early protogalaxies in any hierarchical picture of structure formation, and the gross evolution of baryons in the Universe, because these objects exhibit strong feedback on their surroundings.  Perhaps the most important such channel is the reionization of the intergalactic medium (IGM). When the first protogalaxies or quasars form, they ionize pockets of surrounding gas.  These \\ion{H}{2} regions grow with time and eventually overlap.  The timing, morphology, and duration of this event contain a wealth of information about both the ionizing sources and the IGM (e.g., \\citealt{wyithe03,cen03,haiman03,mackey03,yoshida03-semian,ybh}).  A great deal of effort has gone into constraining the transition from a neutral to ionized IGM.  Unfortunately, existing observational techniques are not optimized to the needed measurements; they have provided tantalizing constraints on reionization but cannot be used to map the event in detail.  The most straightforward method is to extend the ``Ly$\\alpha$ forest'' to high redshifts: regions with relatively large \\ion{H}{1} densities appear as absorption troughs in quasar spectra, which presumably deepen and come to dominate the spectra as we approach the reionization epoch.  Indeed, spectra of $z \\sim 6$ quasars selected from the Sloan Digital Sky Survey\\footnote{See http://www.sdss.org/.}  (SDSS) show at least one extended region of zero transmission \\citep{becker}, indicating that the ionizing background is rising at this time \\citep{fan}. However, the optical depth of the IGM to Ly$\\alpha$ absorption is $\\tau_{Ly\\alpha} \\approx 6.45 \\times 10^5 \\xh [(1+z)/10]^{3/2}$ \\citep{gp}, where $\\xh$ is the neutral fraction.  A neutral fraction $\\xh \\ga 10^{-3}$ will therefore render the absorption trough completely black; quasar absorption spectra can clearly probe only the latest stages of reionization.  A second constraint comes from the effects of the ionized gas on the CMB.  The free electrons Thomson scatter the CMB photons, washing out the intrinsic anisotropies but generating a polarization signal.  The total scattering optical depth $\\tau_{\\rm es}$ is proportional to the column density of ionized hydrogen, so it provides an integral constraint on the reionization history.  Recently, the \\emph{Wilkinson Microwave Anisotropy Probe}\\footnote{See http://map.gsfc.nasa.gov/.} (\\emph{WMAP}) used the polarization signal to measure a large $\\tau_{\\rm es}$, indicating that reionization began at $z_r \\ga 14$ \\citep{kogut03,spergel03}.  More detailed information on the reionization history could be obtained by measuring the (large) angular scales over which CMB polarization is generated \\citep{zal97,kaplinghat03,hu03} or the (small) scales over which secondary anisotropies are generated by the patchiness of reionization \\citep{gruzinov98,knox98}, but these signals promise to be difficult to extract \\citep{holder03,santos03}.  A third constraint comes from measurements of the temperature of the Ly$\\alpha$ forest at $z \\sim 2$--$4$, which suggest an order unity change in the ionized fraction at $z_r \\la 10$ \\citep{theuns02-reion,hui03}, although this argument depends on the timing and history of \\ion{He}{2} reionization (e.g., \\citealt{sokasian02}).  Taken together, these three sets of observations imply a complex reionization history extending over a large redshift interval ($\\Delta z \\sim 10$).  This is inconsistent with a ``generic'' picture of fast reionization (e.g., \\citealt{barkana01}, and references therein).  The results seem to indicate strong evolution in the sources responsible for reionization, and a detailed measurement of the reionization history would contain a rich set of information about early structure formation \\citep{sokasian03a,wyithe03,cen03,haiman03}. The optimal reionization experiment would: (1) be sensitive to order unity changes in $\\xh$ (to probe the crucial middle stages of reionization), (2) provide measurements that are well-localized along the line of sight (rather than a single integral constraint), and (3) not require the presence of bright background sources, which may be rare at high redshifts.  The most promising candidate proposed to date is the 21 cm hyperfine transition of neutral hydrogen in the IGM \\citep{field58,field59a}, which fulfills all three of these criteria. So long as the excitation temperature $T_S$ of the 21 cm transition in a region of the IGM differs from the CMB temperature, that region will appear in either emission (if $T_S > T_{\\rm CMB}$) or absorption (if $T_S < T_{\\rm CMB}$) when viewed against the CMB.  Variations in the density of neutral hydrogen (due either to large-scale structure or to \\ion{H}{2} regions) would appear as fluctuations in the sky brightness of this transition.  Because it is a line transition, the fluctuations can also be well-localized in redshift space.  Thus, in principle, high resolution observations of the 21 cm transition in both frequency and angle can provide a three-dimensional map of reionization.  Together with radio absorption spectra of bright background sources (which can probe much smaller physical scales in the IGM; \\citealt{carilli,furl-21cm}), these observations promise to shed light both on the early growth of structure and on reionization.  The physics of this transition has been well-studied in the cosmological context.  Early work focused on fluctuations due to large-scale structure \\citep{scott,kumar,mmr,tozzi,iliev}, because the signals could be estimated through linear cosmological perturbation theory.  \\citet{shaver} were the first to explicitly consider the signal at reionization, although they focused on the ``all-sky'' signature rather than the fluctuations.  Recently, \\citet{ciardi03} and \\citet{furl-21cmsim} used numerical simulations of reionization to estimate how the fluctuations would behave during that epoch.  We show in Figure \\ref{fig:sim} three time slices from the simulation analysis described by \\citet{furl-21cmsim}, corresponding to the early, middle, and late stages of reionization (from left to right).  It is clear that both the mean signal and the fluctuations drop abruptly.  Interestingly, the fluctuations during reionization have a very different morphology than those due to large-scale structure; the spectrum of fluctuations thus has the potential to constrain the process of reionization.  \\begin{figure*} \\plotone{sim.eps} \\caption{ The brightness temperature of the 21 cm transition at several redshifts, as predicted by the ``late reionization'' simulation analyzed in \\citet{furl-21cmsim}.  Each panel corresponds to the same slice of the simulation box (with width $10 h^{-1}$ comoving Mpc and depth $\\Delta \\nu=0.1 \\Mhz$), at $z=12.1$, 9.2, and 7.6, from left to right. The three epochs shown correspond to the early, middle, and late stages of reionization in this simulation.  (For details about the simulations, see \\citealt{sok01,sh03a,sh03b}.)} \\label{fig:sim} \\end{figure*}  In this paper, we present a new approach to 21 cm fluctuations.  We draw an analogy between these measurements and those of the CMB: in both cases we wish to measure the level of inhomogeneity as a function of scale on the sky.  Previous treatments of the 21 cm signal have focused on measuring fluctuations on a particular patch of the sky, implicitly referring to imaging observations.  Here we show that a statistical treatment of the fluctuation power spectrum contains a great deal of information about reionization.  Furthermore, a large set of tools for CMB predictions and data analysis has already been developed (see \\citealt{hu02} for a review), so there is much to be gained by connecting the two.  Indeed, some steps in this direction have already been taken by \\citet{pen03} and \\citet{cooray03}, both of whom considered the effects of lensing on the 21 cm signal.  The analogy with the CMB is not perfect, however, because the 21 cm signal can be separated in redshift space; in other words, we can make maps at a series of frequencies, each of which samples an independent volume.  In this sense the analogy is closer to redshift surveys (e.g., \\citealt{peebles80}) or weak lensing tomography (e.g., \\citealt{hu99}).  We therefore develop our formalism with explicit consideration of how multifrequency information can be used with power spectrum statistics, in effect generalizing the methodology used to analyze CMB anisotropies.  After reviewing the physics of the 21 cm transition in \\S 2, we show how to compute the angular power spectrum of 21 cm fluctuations in \\S 3.  In \\S 4 we show some simple applications of our approach.  We give predictions for the angular power spectrum of 21 cm fluctuations from a fully neutral medium and for a simple toy model of reionization.  In the former case, fluctuations in the signal are due only to large-scale structure.  We show that in this regime 21 cm measurements essentially yield the power spectrum of density fluctuations (see also \\citealt{pen03}).  We then show that variations in the neutral fraction during reionization distort the power spectrum.  Another advantage of our approach is that the angular power spectra are closely related to the physically observed quantities.  This is especially true for the interferometers that will most likely be used to measure the redshifted 21 cm signal.  Our results thus connect theoretical predictions to the potential observations.  For example, inhomogeneities in the 21 cm signal must be separated from fluctuations in any foreground sources.  This is particularly important because the absolute foreground signal will swamp the 21 cm signal by many orders of magnitude.  While Galactic foregrounds are expected to be fairly smooth on the relevant angular scales, faint radio galaxies, starbursts, and even the galaxies responsible for reionization fluctuate strongly on arcminute scales and dominate those of the 21 cm signal by at least an order of magnitude \\citep{dimatteo,oh03}.  These results have been used to argue that the prospects for large angular-scale measurements of the reionization epoch are dim.  However, both \\citet{dimatteo} and \\citet{oh03} also pointed out that all of the (known) foreground sources have featureless power-law spectra.  Both suggested that the foregrounds could therefore be removed in frequency space.  As an example, consider the simple case in which every foreground source has the \\emph{same} spectral index.  Then the foregrounds between two maps at nearby frequencies would be exactly correlated, while the 21 cm fluctuations will be uncorrelated because each frequency samples an independent volume of the IGM.  Comparing two maps closely spaced in redshift therefore allows one to remove the foreground component.  We show in \\S 5 how the foreground sources can be modeled with our multi-frequency formalism.  In \\S 6 we quantify how well their contamination can be removed.  We find that foregrounds are much less important than previously assumed so long as the range of allowed spectral indices for faint sources is similar to that already measured for brighter sources.  Finally, we estimate in \\S 7 how well the power spectrum can be measured with the next generation of low-frequency radio telescopes, and we conclude in \\S 8.  When necessary, we assume a $\\Lambda$CDM cosmology with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $\\Omega_b=0.04$, $H_0=100h \\hunits$ (with $h=0.7$), and a scale-invariant primordial power spectrum with $n=1$ normalized to $\\sigma_8=0.85$ at the present day. ",
        "conclusions": "\\label{sec:conclusion}  The recent successful efforts to map anisotropies in the cosmic microwave background have determined the global properties of the Universe to high precision (e.g. Spergel et al. 2003).  When combined with the highly successful paradigm for the hierarchical growth of structure (e.g. White \\& Rees 1978), these results imply that the evolution of the Universe on large scales is now relatively well-understood.  However, on the smaller scales characteristic of individual objects, many uncertainties remain owing to our ignorance of the nature of dark matter, the lack of a full physical model for the origin of primordial density fluctuations, and our poor understanding of galaxy formation.  Determining the properties and consequences of the first luminous objects at $z\\sim 15-30$ may help to clarify these issues.  The evolution of the ionized part of the IGM at high redshifts provides a fossil record of the Universe at these times.  In principle, the physical state of this diffuse gas constrains when and where the first luminous objects formed as well as the nature of the sources responsible for reionization.  The relatively large electron scattering optical depth obtained from the WMAP observations is indicative of a complex reionization history but, by itself, does not provide unambiguous answers to the remaining questions.  In this paper, we have argued that fluctuations in 21 cm emission from the IGM can be used to measure the three-dimensional distribution of neutral hydrogen at high redshifts.  Measurements of the angular power spectrum at different frequencies can be used to mitigate contamination by foreground sources that would otherwise overwhelm the 21 cm signal.  Our approach is similar to that employed in analyses of CMB anisotropies, but is more general because of the frequency-dependent nature of redshifted 21 cm emission, making it possible to use 21 cm fluctuations study the {\\it evolution} of reionization. We note here that the most basic kind of experiment is to seek an ``all-sky'' signature corresponding to a global phase transition in the neutral hydrogen (which could be reionization, reheating, or the onset of Ly$\\alpha$ coupling; \\citealt{shaver}).  While valuable (especially because they have much less stringent sensitivity requirements), these measurements only provide the most basic information.  Moreover, they are subject to severe foreground contamination \\citep{gnedin03} because the large scale signal varies relatively smoothly with frequency.  Thus frequency differencing will not be as effective as with small-scale fluctuations.  The small-scale fluctuations we have described will therefore ultimately be a better route to pursue.  Using a simple conceptual model for the morphological evolution of ionized regions, we have demonstrated how reionization imprints characteristic features into the angular power spectrum of 21 cm fluctuations.  Our formulation is general, but for illustrative purposes we have adopted simplifying assumptions to emphasize salient features of our methodology.  For example, in our toy model of reionization, we have ignored redshift space distortions in the 21 cm signal and neglected correlations between the density and neutral fraction fields.  In future work, we will investigate the impact of these effects explicitly using semi-analytic methods and numerical simulations to show how the formalism can be adapted to account for these complications.  In principle, measurements of the type we propose can be used to distinguish between various evolutionary histories.  Models with multiple epochs of reionization (e.g. Cen 2003; Wyithe \\& Loeb 2003) lead to a behavior in which the volume fraction of ionized gas in the universe shows a complex dependence on redshift (e.g. Figures 8 \\& 9 of Sokasian et al. 2003b), unlike single episodes of reionization where the trend is simpler (e.g. Figure 5 of Sokasian et al. 2003a).  These differences will be imprinted onto the frequency dependence of 21 cm fluctuations of the IGM and can be discerned using the approach described here. In fact, the simple model described in \\S 4 is only part of the story available to us through 21 cm observations.  For example, the fluctuations constrain the thermal history of the IGM as well as the ionization history (see \\S 2).  If the earliest ionizing sources have hard spectra (such as quasars), we would expect the IGM to be heated rapidly, while if cool, low-mass stars are responsible for reionization, the heating would occur much closer to overlap.  Another way to look at this is that the reionization history determines how much information is available to us through 21 cm fluctuations.  For instance, in a scenario with rapid heating, we will have a long epoch where density variations dominate the signal. The power spectrum of the 21 cm fluctuations is a direct tracer of the matter power spectrum. Thus, detailed measurements of this signal could provide invaluable constraints on the primordial spectrum on small scales which could tightly constrain the physics of the early Universe. In particular, because one gets many independent maps by varying the frequency, the constraints on the power spectrum obtained in this way could be significantly better than those coming from the CMB primary anisotropies. Also, 21 cm fluctuations allow one to probe the power spectrum at $z \\sim 10$--$20$, an epoch inaccessible to both the CMB and galaxy surveys and a relatively large lever with which to study the growth of fluctuations.  On the other hand, the 21 cm signal depends on complex physical processes (see \\S 2) and interpreting the results will require careful modeling.  Differences in the evolutionary state of the ionized IGM also contain information about the matter power spectrum on small scales.  For example, in cosmological models with reduced small-scale power, the halos hosting star-forming regions form late, delaying the ionizing effects of the first luminous objects (e.g. Somerville et al. 2003; Yoshida et al. 2003b,c).  Thus, universes with a large component of warm dark matter or those in which the matter power spectrum has a running spectral index should exhibit lower amplitude fluctuations in 21 cm emission from higher redshifts than in conventional $\\Lambda$CDM models.  A detailed study of 21 cm fluctuations can also constrain the nature of the sources responsible for reionization.  In some scenarios, it is conjectured that reionization occurs ``outside-in,'' affecting voids first and high density regions later (e.g. Miralda-Escud\\'e et al. 2000).  This progression would be expected if the sources are rare, but bright.  Alternatively, reionization could proceed in the opposite sense, ``inside-out,'' particularly if the sources are numerous, but faint (e.g. Gnedin 2000; Sokasian et al. 2003a).  The former possibility would apply if quasars were the primary sources of ionizing radiation, while stars in low-mass galaxies would be more relevant to the ``inside-out'' scenario.  The morphological appearance of the ionized gas is different in these two cases, and these differences would be reflected in the angular power spectrum of 21 cm fluctuations and how this quantity varies with frequency. In particular, the size of the \\ion{H}{2} regions at a given $x_H$ constrains the number density of ionizing sources.  Also, if voids are ionized first the density and $x_H$ fields will be correlated (i.e., the neutral fraction falls first in regions that are already underdense), while the opposite is true in an ``inside-out'' scenario. Thus the amplitude of the brightness temperature fluctuations contains information about the process of reionization.  As we have argued in \\S 7, the technological requirements for detecting 21 cm fluctuations from cosmic gas at high redshifts, while demanding, are within reach.  In the near future, it is likely that measurements of the power spectrum of 21 cm fluctuations will reveal the physical state of the Universe at an epoch that is currently inaccessible to other observational probes."
    },
    "0311/astro-ph0311044_arXiv.txt": {
        "abstract": "We present preliminary results from spectroscopic observations of a sample of 11 low surface brightness galaxies (LSB). We measured the stellar and gaseous kinematics along their major and minor axes. Such information will allow us to accurately investigate the dark matter (DM) content within their optical regions, providing further constraints on the predictions of standard CDM models. Although dynamical modeling is still in progress, our observations already show that the derived stellar kinematics is more regular than the ionized gas one, which often shows evidence for non-circular and asymmetric motions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311334_arXiv.txt": {
        "abstract": "Big Bang Nucleosynthesis can provide, via constraints on the expansion rate at that time, limits on possible variations in Newton's Constant, $G$.  The original analyses were performed before an independent measurement of the baryon-to-photon ratio from the cosmic microwave background was available.  Combining this with recent measurements of the primordial deuterium abundance in quasar absorption systems now allows one to derive a new tighter constraint on $G$ without recourse to considerations of helium or lithium abundances.  We find that, compared to todays value, $G_0$, $G_{BBN}/G_0=1.01^{+0.20}_{-0.16}$ at the $68.3$\\% confidence level.  If we assume a monotonic power law time dependence, $G\\propto t^{-\\alpha}$, then the constraint on the index is $-0.004 < \\alpha < 0.005$. This would translate into $-3\\times10^{-13}\\;\\textrm{yr}^{-1} < \\left(\\dot G/G\\right)_{\\textrm{today}} < 4 \\times 10^{-13}\\;\\textrm{yr}^{-1}$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311102_arXiv.txt": {
        "abstract": "In the curvaton model of inflation, where a second scalar field, the `curvaton', is responsible for the observed inhomogeneity, a non-zero neutrino degeneracy may lead to a characteristic pattern of isocurvature perturbations in the neutrino, cold dark matter and baryon components. We find the current data can only place upper limits on the level of isocurvature perturbations. These can be translated into upper limits on the neutrino degeneracy parameter. In the case that lepton number is created before curvaton decay, we find that the limit on the neutrino degeneracy parameter is comparable with that obtained from Big-bang nucleosynthesis. For the case that lepton number is created by curvaton decay we find that the absolute value of the non-Gaussianity parameter,$|\\fnl|$, must be less than 10 (95\\% confidence interval). ",
        "introduction": "In single field models  of inflation, the inflaton gives rise to an era of accelerated expansion in the early universe during which quantum fluctuations in the inflaton field are expanded beyond the horizon size \\cite{LLBook}. At the end of inflation, the inflaton decays into the standard matter fields.  The primordial fluctuations, imprinted onto the gravitational potential in form of small perturbations, become the ``seeds'' that start the formation of large scale structure. As there is only a single degree of freedom, the inflaton's value, the resulting perturbations are always {\\em adiabatic,} i.e.\\ the ratio of the number densities of different particle species is spatially homogeneous, there is no isocurvature mode.  Recently a different model has been proposed, the curvaton scenario \\cite{curvaton, LUW}, which has already received considerable attention \\cite{allcurvaton,MWU,DHLDW,gordonlewis}. In this model the accelerated expansion is still provided by the inflaton, but the primordial perturbations are generated by another scalar field, the {\\em curvaton}. The curvaton energy density is small compared to the inflaton energy density and being a light field, it acquires an almost scale-invariant spectrum of perturbations during inflation. After the decay of the inflaton, the isocurvature perturbation in the curvaton fluid is transformed into an adiabatic perturbation \\cite{Mollerach,LM,MWU} and the curvaton decays into radiation and other particles. A similar model to the curvaton scenario, in which a second scalar field determines the inflaton decay rate, has also been  proposed recently \\cite{vardecay}.   The curvaton scenario requires different constraints on the inflationary potential to standard inflation and might therefore be better suited for inflation at a lower energy scale \\cite{lythdimopolos,lythsmallscale}. It also allows for the possibility of an isocurvature mode in the standard matter fields \\cite{curvaton,LUW}, i.e.\\ the ratio of particle number densities of different particle species is not necessarily spatially homogeneous. This arises since the final constituents of the standard big bang model (photons, neutrinos, baryons and cold dark matter (CDM)) do not all necessarily originate entirely from the spatially inhomogeneous curvaton, but may originate instead from the spatially homogeneous inflaton or some other source.   In Ref.~\\cite{gordonlewis} the isocurvature components in the CDM and baryons that could arise in the curvaton scenario were constrained using WMAP and other observational data. It was found that the current data rule out either CDM or baryons being created before curvaton decay. In this article we examine what the current observational constraints are on neutrino isocurvature perturbations created by curvaton decay. These can only arise if there is a non-zero neutrino degeneracy parameter $\\xi = \\mu_\\nu / T_\\nu$, where $\\mu_\\nu$ is the neutrino chemical potential and $T_\\nu$ is the neutrino temperature \\cite{LUW}. There may be a separate degeneracy parameter for each of the three species of neutrinos, but the recent atmospheric and solar neutrino data indicate that approximate flavor equilibrium between all active neutrino species is established well before the BBN epoch \\cite{solar} and hence we take the degeneracy parameters to be equal.  Various studies have constrained the value of $\\xi$ using CMB and other large scale structure data in the adiabatic case \\cite{xibg}. However, the tightest constraints still come from Big-bang nucleosynthesis (BBN) \\cite{bbnxi}. In this article we examine how the constraints on $\\xi$ change in the curvaton scenario where neutrino isocurvature perturbations can be induced by a non-zero $\\xi$. It turns out that the neutrino isocurvature perturbations then also induce baryon and CDM isocurvature perturbations \\cite{LUW}. It is therefore not possible to use constraints that have simply considered the addition of a neutrino isocurvature mode such as Ref.~\\cite{garcia-bellido} or those that have considered arbitrary combinations of different isocurvature perturbations \\cite{multiiso}.  In the next section we evaluate the possible magnitudes of the neutrino isocurvature perturbations. Then in \\sec{section3} we estimate analytically the effects of the isocurvature perturbations on the temperature fluctuations of the CMB and their dependence on $\\xi$. A full likelihood analysis of the allowed ranges of $\\xi$ is given in \\sec{data}. Finally, the conclusions and the discussion are given in \\sec{conclusions}. ",
        "conclusions": "\\slabel{conclusions}  We have examined the current cosmological constraints on the neutrino degeneracy $\\xi$, in the curvaton scenario, using WMAP, ACBAR, CBI, 2dF, HST and BBN data. The expressions for the isocurvature perturbations were evaluated to all orders in $\\xi$ and the effects of the induced baryon and CDM isocurvature perturbations were accounted for. Analytical formulas were given for the modification of the Sachs-Wolfe effect on large scales.  Comparisons were made with the BBN constraint which results from the effect of the neutrino degeneracy on the pre-BBN proton to neutron ratio and using the large mixing angle solution in which the neutrino degeneracies of the different species take on the same value.  It was found that the constraint on $\\xi$ was comparable to that from BBN when the lepton number was created before curvaton decay. As seen in \\fig{cls}, a non-zero $\\xi$ affects the acoustic peaks whose data will dramatically improve in the future. So the CMB constraints on $\\xi$ may also improve dramatically. Without the BBN constraint, the neutrino isocurvature perturbation was found to be smaller than 10\\% the size of the adiabatic mode (95\\% CI), while including the BBN constraint reduced this by a factor of a few. In the case where the lepton number is created by curvaton decay, there is a degeneracy between the value of the curvaton energy density at the time of curvaton decay and $\\xi$. However, this does not stop the BBN data strengthening the constraint on the isocurvature perturbation as the value of the curvaton energy density at the time of curvaton decay is constrained by the limits on the non-Gaussianity parameter from WMAP. So in this case, the neutrino isocurvature mode is less than about 5\\% (95\\% CI) of the adiabatic mode. Unless the BBN constraints on $\\xi$ become tighter around zero, the case of the lepton number being created by curvaton decay would be unlikely if future observations found $\\fnl < -10$, as this would imply a larger isocurvature perturbation than is observed."
    },
    "0311/astro-ph0311428_arXiv.txt": {
        "abstract": "A satisfactory understanding of the origin of the dependence of galaxy properties on their environment has remained, so far, out of reach. In the light of numerous observational results and substantial theoretical progress obtained for clusters of galaxies in the last years, a primary goal is to understand how the star formation activity depends on cluster substructure, i.e. on the merging/accretion history of a cluster.  In this contribution we present a case in which it is possible to identify the cluster environment, and in particular the intracluster medium and the recent infall history of galaxies onto the cluster, as the cause for an abrupt change in the star formation histories of a subset of galaxies in the Coma cluster. ",
        "introduction": "Fundamental galaxy properties such as star formation history, morphology and gas content vary in a systematic way from dense environments like clusters of galaxies to low density regions in the general field.  Similarly, it has long been known that there are systematic variations of {\\it galaxy} properties with {\\it cluster} properties.  Already in 1974, on a sample of 15 clusters, Oemler tentatively divided clusters into three types. He noted that clusters with cD galaxies and rich in ellipticals have a smooth spherical distribution and are centrally concentrated, while clusters that are dominated by S0s are not as smooth, dense or centrally concentrated as cD clusters, and clusters of irregular appearance are dominated by spiral, star-forming galaxies.  A few years later, in 1977, Bahcall first showed that the fraction of spiral galaxies in a cluster decreases rapidly with increasing X-ray luminosity (see also Edge \\& Stewart 1991).  Subsequent studies on the correlation between galaxy properties and global properties of clusters have been innumerable, especially searching for correlations between the star formation activity in galaxies and the cluster X-ray luminosity (Smail et al. 1998, Ellingson et al. 2001, Balogh et al. 2002a), the cluster richness (Margoniner et al. 2001) and the presence of substructure (Wang \\& Ulmer 1997, Metevier et al. 2000, Pimbblet et al. 2002).  At some level the global cluster properties (X-ray and optical luminosities, richness, concentration etc.) are all correlated with each other, though with a large scatter, and it is hard to pin down what exactly is the primary correlation with galaxy properties that drives all the others.  A central question is surely how the star formation activity depends on cluster substructure, i.e. on the merging/accretion history of the cluster.  There are some very good examples of cluster-cluster mergers where the star formation activity in galaxies, as measured by radio continuum emission, is exceptionally high, such as Abell 2125 (Dwarakanath \\& Owen 1999) and Abell 2255 (Miller \\& Owen 2003). In principle it is not justified to inequivocally attribute this activity to an ``enhancement'' of star formation due to the cluster-cluster merging: the activity observed could simply reflect an higher average activity in the substructures that are present in the cluster, given that the average star formation is known to be higher in lower density/less massive environment. Thus, the cause of the observed activity needs to be investigated on a one-by-one case. In Abell 2255, for example, the alignment of the star-forming galaxies along a preferential direction, perpendicular to the merger axis, strongly supports the hypothesis that the bursts of star formation were indeed triggered by the cluster-cluster merger, induced by ram pressure variations at the shock front (Miller \\& Owen 2003).  On the other hand, there are also merging clusters where only a low star formation activity is detected, suggesting that the gas reservoir in the galaxy populations of the two merging clusters must also play a role (Miller 2003). The dynamical state of the cluster and the type of galaxies that happen to be falling in at the time we observe it are the two crucial factors also for understanding the gas (HI) content of galaxies in clusters (van Gorkom 2003).  Galaxy properties are also known to be closely linked with their {\\it local} environment, regardless of the ``global'' environment.  The most striking evidence for this is the morphology-density relation (Dressler 1980, Dressler et al. 1997), paralleled by a star-formation/density relation (Lewis et al. 2002, Gomez et al. 2003).  The higher incidence of ellipticals in regions of high local density, and the corresponding preference for spirals to reside in low density regions is known to be valid across a wide range of ``global'' environments, for both concentrated and irregular clusters, for rich and poor systems (Dressler 1980, Postman \\& Geller 1984), though an additional second-order effect related to the {\\it global} environment seems to be present, in the sense that at a given local density more massive and relaxed clusters show a small excess of early-type galaxies compared to less massive and irregular clusters (Dressler 1980, Balogh et al. 2002b).  One of the open questions is therefore to what extent are the differences in galaxy properties in various types of clusters due to the fact that the {\\it average local} environment varies as a function of the cluster type.  The findings cited above, together with many others in the literature, prove that both {\\it local} % and {\\it global} effects are at work. In this talk we present the case of a cluster-related phenomenon, where it is possible to identify the cluster environment, and in particular the intracluster medium and the recent infall history of galaxies onto the cluster, as the cause for an abrupt change in the star formation histories of a subset of the Coma galaxies. ",
        "conclusions": "The position of faint post-starburst galaxies in Coma relative to X-ray substructure in the cluster strongly suggests that the interruption of the star formation activity in these galaxies is a cluster-related phenomenon, most likely due to the impact with the intracluster medium. Post-starburst galaxies in Coma are much fainter than those observed in distant clusters, a difference probably reflecting a change in the galaxy populations infalling in clusters.  The dynamical state of the cluster and its recent history of accretion are clearly important factors in determining the galaxy properties observed. Identifying substructure, however, it is often hard, and no easily measurable ``global'' cluster property (such as for example X-ray luminosity or richness) is sufficient to characterize the cluster dynamical status.  We have shown here an example in which only a detailed analysis of the X-ray residuals has revealed the link between substructure and rapid changes in the star formation histories of some of the cluster galaxies."
    },
    "0311/astro-ph0311116_arXiv.txt": {
        "abstract": "We derive self-similar solutions including cooling and heating in an Einstein de-Sitter universe, and investigate the effects of cooling and heating on the gas density and temperature distributions. We assume that the cooling rate has a power-law dependence on the gas density and temperature, $\\Lambda$$\\propto$$\\rho^{A}T^{B}$, and the heating rate is $\\Gamma$$\\propto$$\\rho T$. The values of $A$ and $B$ are chosen by requiring that the cooling time is proportional to the Hubble time in order to obtain similarity solutions. In the region where the cooling rate is greater than the heating rate, a cooling inflow is established, and the gas is compressed and heats up. Because the compression is greater in the inner region than in the outer region, the temperature becomes an increasing profile toward the center. In particular, when a large infall velocity is produced due to an enormous energy loss, the slope of the density approaches a value that depends on $A$, $B$, and the velocity slope, and the slope of the temperature approaches $-$1. On the other hand, in the region where the heating rate is greater than the cooling rate, the infall velocity is suppressed, compression of the gas is weakened, and the gas cools down. The slope of the density becomes shallow due to suppression of the contraction, and the temperature is lower than that without heating. The self-similar collapse presented here gives insights to the effects of cooling and heating on the gas distributions in galaxies and clusters of galaxies. ",
        "introduction": "\\subsection{Spherical collapse in an expanding universe} We consider here that an over-dense region has a density contrast ($\\delta_{i}$) at an initial Hubble time of $t_{H,i}$, and that the background universe is an Einstein de-Sitter universe. The gravitational collapse in the flat universe becomes a similarity evolution because bounded shells extend to infinity. Fillmore \\& Goldreich\\shortcite{Fillmore} introduced a scale-free initial density contrast, \\begin{equation} \\delta_{i}=\\frac{\\delta M_{i}}{M_{i}}=(\\frac{M_{i}}{M_{0}})^{-\\epsilon} =(\\frac{r_{i}}{R_{0}})^{-3\\epsilon}, \\label{eq:equation1} \\end{equation} where $M_{i}$=$(4/3)\\pi\\rho_{H,i} r_{i}^{3}$ is the mass contained within the proper radius $r_{i}$, $\\rho_{H,i}$=$1/(6 \\pi G t^{2}_{H,i})$ is the background density at $t_{H,i}$, and $M_{0}$=$(4/3)\\pi\\rho_{H,i} R_{0}^{3}$ is a reference mass. Since $\\delta M_{i}$ must be constant or increase with increasing radius, and $\\delta_{i}$ has to decrease with increasing radius, $\\epsilon$ is restricted to $0< \\epsilon \\leq 1$. The variance of the density fluctuation is \\begin{equation} |\\delta(x,t_{H})|^{2}\\simeq k^{3}|\\delta_{k}|^{2} \\simeq a(t_{H})^{2} M^{-\\frac{n+3}{3}}, \\label{eq:equation2} \\end{equation} where $x$ is the comoving radius, and $|\\delta_{k}|^{2}$ $\\propto$ $a(t_{H})^{2} k^{n}$ is the power spectrum, where $a(t_{H})$ is an expansion factor. The relation between the index $n$ and the parameter $\\epsilon$ is \\begin{equation} n=3(2\\epsilon-1). \\label{eq:equation3} \\end{equation} We assume that the initial velocity of matter is pure Hubble flow, and that the initial gas pressure is zero. The expanding matter slows down with the growth of the perturbation. Eventually, it reaches a maximum radius and decouples from the Hubble expansion. Since $a(t_{H})\\propto t_{H}^{2/3}$, the perturbation grows as $|\\delta(r,t_{H})|^{2}$ $\\propto$ $a(t_{H})^{n+5}r^{-(n+3)}$ $\\propto$ $t_{H}^{2(n+5)/3}r^{-(n+3)}$, based on equation (\\ref{eq:equation2}). The density contrast at the turn-around radius is $\\delta(r_{ta},t_{H})$=(3$\\pi$/4)$^{2}-$1$\\sim$4.55(Peebles 1980). Thus, the time evolution of the turn-around radius and the enclosed mass are \\begin{equation} r_{ta}(t_{H})\\propto t_{H}^{\\xi}, \\label{eq:equation4} \\end{equation} \\begin{equation} m_{ta}(t_{H})\\propto t_{H}^\\frac{2}{3\\epsilon}, \\label{eq:equation5} \\end{equation} where \\begin{equation} \\xi=\\frac{2}{3}(1+\\frac{1}{3\\epsilon}). \\label{eq:equation6} \\end{equation} After the turn-around, particles fall toward the center. The collapse of matter causes an increase of the central density. Since the gas particles are decelerated, a shock wave is formed and propagates outward. The system undergoes self-similar evolution at a later time.  \\subsection{Cooling and heating laws in similarity evolutions} In the case of the collapse of adiabatic gas, the physical scale of the system is only determined by the gravitational dynamics. When radiative cooling and heating are included, the similarity evolution is broken due to the existence of many physical scales. In order to maintain the similarity evolution, it is required that all of the physical scales are proportional to the dynamical scale.  The scale-free condition was constructed by ensuring that the cooling time is proportional to the dynamical time (Owen et al. 1998; Abadi et al. 2000).  We assume that the cooling rate per unit volume is $\\Lambda$=$\\Lambda_{0} \\rho^{A} c_{s}^{2B}$. Here, $A$ and $B$ are free parameters, $\\Lambda_{0}$ is the cooling coefficient, $\\rho$ is the gas density, and $c_{s}$ is the sound speed. The cooling time becomes \\begin{equation} t_{cool}=\\frac{E}{\\frac{DE}{Dt}} =E(\\frac{\\Lambda}{\\rho})^{-1} =\\frac{\\rho^{1-A}c_{s}^{2(1-B)}}{\\gamma(\\gamma-1)\\Lambda_{0}}, \\label{eq:equation7} \\end{equation} where $E$ is the specific thermal energy of the gas, and $\\gamma=5/3$ is the adiabatic index. The Hubble time evolves as \\begin{equation} t_{H}=\\frac{2}{3}H_{0}^{-1}(\\frac{a}{a_{0}})^{\\frac{3}{2}}, \\label{eq:equation8} \\end{equation} where $a_{0}$ and $H_{0}$ are the expansion factor and Hubble constant at the present epoch, respectively. The sound speed and density evolve as \\begin{equation} c_{s}=(\\frac{a}{a_{0}})^{\\frac{1-n}{2(3+n)}} c_{s,0}, \\label{eq:equation9} \\end{equation} \\begin{equation} \\rho=(\\frac{a}{a_{0}})^{-3} \\rho_{0}, \\label{eq:equation10} \\end{equation} where $c_{s,0}$ and $\\rho_{0}$ are the sound speed and the gas density at the present epoch, respectively. We assume that the Hubble time is proportional to the cooling time by the constant $\\hat{t_{c}}$, in order to maintain scale-free conditions (i.e., $t_{H}$=$\\hat{t_{c}}t_{cool}$), \\begin{equation} t_{H}=\\hat{t_{c}}\\frac{\\rho^{1-A}c_{s}^{2(1-B)}}{\\gamma(\\gamma-1)\\Lambda_{0}}. \\label{eq:equation10-1} \\end{equation} Substituting equations (\\ref{eq:equation8})-(\\ref{eq:equation10}), the above equation becomes \\begin{equation} \\frac{2}{3}H_{0}^{-1} (\\frac{a}{a_{0}})^{\\frac{3}{2}-3(A-1)-\\frac{1-n}{3+n}(1-B)} =\\hat{t_{c}} \\frac{\\rho_{0}^{1-A}c_{s,0}^{2(1-B)}} {\\gamma(\\gamma-1)\\Lambda_{0}}. \\label{eq:equation11} \\end{equation} The $a$-dependence term must vanish in order to keep $\\hat{t_{c}}$ constant. Using equation (\\ref{eq:equation3}), we find \\begin{equation} \\frac{3}{2}-3(A-1)-\\frac{2-3\\epsilon}{3\\epsilon} (1-B)=0. \\label{eq:equation12} \\end{equation}  Figure 1 displays the parameter space of $A$ and $B$ in relation to $\\epsilon$. The shaded areas correspond to $0< \\epsilon \\leq1$. The solid lines display the lines of $\\epsilon$=1, 2/3, 1/3, and 0. Abadi et al.(2000) selected at the point ($A$,$B$)=(3/2,1), which is independent of $\\epsilon$. The points of free-free emission, ($A$,$B$)=(2,1/2), and line cooling, ($A$,$B$)=(2,$-$1/2), are indicated by crosses.  Next, we assume that the heating rate per unit volume is given by \\begin{equation} \\Gamma=\\alpha \\rho c^{2}_{s}, \\label{eq:equation12-1} \\end{equation} where $\\alpha$=$(t_{H,0}/t_{H})\\alpha_{0}$, $t_{H,0}$ is the current Hubble time, and $\\alpha_{0}$ is the model parameter. The heating time becomes \\begin{equation} t_{heat}=\\frac{E}{\\frac{DE}{Dt}}=E(\\frac{\\Gamma}{\\rho})^{-1} =\\frac{1}{\\gamma(\\gamma-1)}\\frac{1}{\\alpha}. \\label{eq:equation12-2} \\end{equation} We assume that the heating time is proportional to the Hubble time by the constant $\\hat{t_{h}}$(i.e., $t_{H}$=$\\hat{t_{h}}t_{heat}$), \\begin{equation} \\frac{2}{3}H_{0}^{-1}=\\frac{\\hat{t_{h}}}{\\gamma(\\gamma-1)\\alpha_{0}}. \\label{eq:equation12-3} \\end{equation}  \\begin{figure} \\begin{center} \\includegraphics[width=8cm,height=10.0cm,clip]{fig1.eps} \\caption{Relation between ($A$, $B$) and $\\epsilon$. The shaded areas indicate the inside of the range $0< \\epsilon \\leq1$. The solid lines are constant values of $\\epsilon$=1, 2/3, 1/3, and 0. The closed circles represent the parameter positions that we calculated in this paper. The points of free-free emission and line emission are indicated by crosses. } \\end{center} \\label{fig:figure1} \\end{figure}%  \\subsection{Basic equations of collisional gas} Neglecting the effect of dark matter, the basic equations with additional terms of cooling and heating are as follows: \\begin{equation} \\frac{\\partial{m}}{\\partial r}=4\\pi r^{2}\\rho, \\label{eq:equation14} \\end{equation} \\begin{equation} \\frac{\\partial\\rho}{\\partial t}+\\frac{1}{r^{2}}\\frac{\\partial} {\\partial r}(r^{2}\\rho v)=0, \\label{eq:equation15} \\end{equation} \\begin{equation} \\frac{\\partial v}{\\partial t}+v\\frac{\\partial v}{\\partial r}= -\\frac{\\partial}{\\gamma \\rho \\partial r} (\\rho c_{s}^{2\\setlength{\\oddsidemargin}{-0.1in}})-\\frac{Gm}{r^{2}}, \\label{eq:equation16} \\end{equation} \\begin{equation} \\frac{1}{\\gamma-1}\\frac{dc_{s}^{2}}{dt}-\\frac{c_{s}^{2}}{\\rho}\\frac{d\\rho}{dt} =- \\gamma \\Lambda_{0} \\rho^{A-1} c_{s}^{2B}+\\gamma \\alpha c_{s}^{2}, \\label{eq:equation17} \\end{equation} where $v$ is the velocity, $G$ is the gravitational constant, and $m$ is the mass of the gas.  To derive self-similar solutions, we introduce the similarity variable $X$ and similarity functions, which are given by the following forms: \\begin{equation} X=\\frac{r}{r_{ta}}, \\label{eq:equation18} \\end{equation} \\begin{equation} \\rho=\\rho_{H}D, \\label{eq:equation20} \\end{equation} \\begin{equation} v=\\frac{r_{ta}}{t_{H}}V, \\label{eq:equation19} \\end{equation} \\begin{equation} c_{s}=\\frac{r_{ta}}{t_{H}}C, \\label{eq:equation21} \\end{equation} \\begin{equation} m=\\frac{4}{3}\\pi\\rho_{H}r^{3}_{ta}M, \\label{eq:equation22} \\end{equation} \\begin{equation} \\alpha=\\frac{\\mathcal{H}}{t_{H}}, \\label{eq:equation24} \\end{equation} where $\\rho_{H}$=$1/(6 \\pi G t^{2}_{H})$ is the mean density of the universe, and ${\\mathcal{H}}$ is the same as $\\alpha_{0}t_{H,0}$ or $3H_{0}{\\hat{t_{h}}}t_{H,0}/2\\gamma(\\gamma-1)$ from equation (\\ref{eq:equation12-3}). The density and the sound speed are expressed as \\begin{equation} \\rho=\\rho_{0}\\frac{\\rho_{H}}{\\rho_{H,0}}, \\label{eq:equation24-1} \\end{equation} \\begin{equation} c_{s}=c_{s,0}(\\frac{r_{ta}}{t_{H}})/(\\frac{r_{ta,0}}{t_{H,0}}), \\label{eq:equation24-2} \\end{equation} where $\\rho_{H,0}$ and $r_{ta,0}$ are the background density and the turn-around radius at the present epoch, respectively. We substitute equations (\\ref{eq:equation24-1}) and (\\ref{eq:equation24-2}) for (\\ref{eq:equation10-1}), and define $\\Lambda_{0}$ as \\begin{eqnarray} \\Lambda_{0}&=&\\frac{\\hat{t_{c}}(\\frac{\\rho_{0}}{\\rho_{H,0}})^{1-A} (\\frac{c_{s,0}t_{H,0}}{r_{ta,0}})^{2(1-B)}}{\\gamma(\\gamma-1)} \\rho_{H}^{1-A} r_{ta}^{2(1-B)}t_{H}^{2B-3},\\nonumber \\\\ &=&\\frac{K_{0}}{\\gamma(\\gamma-1)}\\rho_{H}^{1-A} r_{ta}^{2(1-B)}t_{H}^{2B-3}. \\label{eq:equation25} \\end{eqnarray} Substituting equations (\\ref{eq:equation18})-(\\ref{eq:equation24}), and (\\ref{eq:equation25}) for (\\ref{eq:equation14})-(\\ref{eq:equation17}), the fluid equations become, \\begin{equation} \\frac{dM}{dX}=3X^{2}D, \\label{eq:equation26} \\end{equation} \\begin{eqnarray} (V-\\xi X)\\frac{dD}{dX}+D\\frac{dV}{dX}+2D(\\frac{V}{X}-1)=0, \\label{eq:equation27} \\end{eqnarray} \\begin{eqnarray} (V-\\xi X)\\frac{dV}{dX}-(1-\\xi)V =-\\frac{1}{\\gamma D}\\frac{d}{dX}(DC^{2})-\\frac{2M}{9X^{2}}, \\label{eq:equation28} \\end{eqnarray} \\begin{eqnarray} \\frac{2}{\\gamma-1}[{(\\xi-1)+\\frac{(V-\\xi X)}{C}\\frac{dC}{dX}}]+2 -\\frac{(V-\\xi X)}{D}\\frac{dD}{dX} \\nonumber \\\\ =-\\frac{K_{0}}{(\\gamma-1)}D^{A-1}C^{2B-2} +\\gamma {\\mathcal{H}}. \\label{eq:equation29} \\end{eqnarray} The non-dimensional mass $M$ is derived by another integration(Bertschinger 1985). The equation of continuity is rewritten as follows: \\begin{eqnarray} \\frac{\\partial{m}}{\\partial t}+4\\pi r^{2}\\rho v=0. \\label{eq:equation29-1} \\end{eqnarray} Substituting equations (\\ref{eq:equation18})-(\\ref{eq:equation22}) for (\\ref{eq:equation14}) and (\\ref{eq:equation29-1}), the dimensionless mass is \\begin{eqnarray} M=\\frac{-3}{3\\xi-2}X^{2}D(V-\\xi X). \\label{eq:equation29-2} \\end{eqnarray}  \\subsection{Numerical integration and boundary conditions} \\label{subsec:numerical} \\begin{figure*} \\resizebox{8cm}{!}{\\includegraphics{fig2D.eps}}% \\resizebox{8cm}{!}{\\includegraphics{fig2V.eps}}\\\\ \\resizebox{8cm}{!}{\\includegraphics{fig2C2.eps}}% \\resizebox{8cm}{!}{\\includegraphics{fig2tc.eps}}\\\\ \\ \\\\ \\caption{Three types of solutions with $\\epsilon$=1 ($n$=3), ${\\mathcal{H}}$=0, $A$=2, $B$=$-$7/2, and $K_{0}$=0.01. Each line shows the resultant profiles of the 'stagnation' solution (dashed line), 'eigensolution' (solid line), and 'adiabatic' solution (dotted line), respectively. The shock radii are $X_{s}$=0.338976, 0.326595, and 0.3, respectively. Each panel represents radial profiles of the density (top left), the proper velocity (top right), the square of the sound speed corresponding to the temperature (bottom left), and the ratio of the Hubble time to the cooling time (bottom right), respectively.} \\label{fig:figure2} \\end{figure*}% Assuming that the gas pressure is zero, the physical variables of a gas shell are given by the parameter $\\theta$(Peebles 1980): \\begin{equation} X=\\frac{r}{r_{ta}}=\\rm{sin}^{2}(\\frac{\\theta}{2}) (\\frac{\\theta-\\rm{sin}\\theta}{\\pi})^{-\\xi}, \\label{eq:equation30} \\end{equation} \\begin{equation} D=\\frac{9}{2}\\frac{(\\theta-\\rm{sin}\\theta)^{2}}{(1-\\rm{cos}\\theta)^{3} (1+3\\epsilon\\chi)}, \\label{eq:equation31} \\end{equation} \\begin{equation} V=X\\frac{\\rm{sin}\\theta(\\theta-\\rm{sin}\\theta)} {(1-\\rm{cos}\\theta)^{2}}, \\label{eq:equation32} \\end{equation} \\begin{equation} C=0, \\label{eq:equation32-1} \\end{equation} \\begin{equation} M=\\frac{9}{2}\\frac{(\\theta-\\sin \\theta)^{2}}{(1-\\cos \\theta)^{3}}X^{3}, \\label{eq:equation33} \\end{equation} with $\\chi$=$1-3V/2X$. The outer boundary condition is defined at the turn-around radius, $\\theta$=$\\pi$. The collisional gas shell falls toward the center after the turn around, and the central gas density increases. Eventually, adiabatic compression occurs, and a shock propagates outward. The shock radius appears at a fixed fraction of the turn-around radius. The shock jump conditions are written as \\begin{equation} D_{2}=\\frac{\\gamma+1}{\\gamma-1}D_{1}, \\label{eq:equation34} \\end{equation} \\begin{equation} V_{2}=\\xi X_{s}+(V_{1}-\\xi X_{s})\\frac{\\gamma-1}{\\gamma+1}, \\label{eq:equation35} \\end{equation} \\begin{equation} C_{2}^{2}=\\frac{2\\gamma(\\gamma-1)}{(\\gamma+1)^{2}}(V_{1}-\\xi X_{s})^{2}, \\label{eq:equation36} \\end{equation} where subscripts 1 and 2 represent the pre- and post-shock quantities, and $X_{s}$ is the non-dimensional shock radius. We need to determine the shock position. Bertschinger(1989) and Abadi et al.(2000) showed three types of solutions by changing the shock radius. Figure 2 represents the three possible solutions of the 'stagnation' solution (dashed line), 'eigensolution' (solid line), and 'adiabatic' solution (dotted line) for $\\epsilon$=1($n$=3), ${\\mathcal{H}}$=0, $A$=2, $B$=$-$7/2, and $K_{0}$=0.01. Each shock radius is $X_{s}$=0.338976, 0.326595, and 0.3, respectively. The ratio of the Hubble time to the cooling time at each radius is shown in Fig. 2, which is defined by \\begin{equation} \\frac{t_{H}}{t_{cool}}=K_{0}D^{A-1}C^{2B-2}. \\label{eq:equation38} \\end{equation}  The first type is the 'stagnation' solution (dashed line), which diverges at a small radius. We found that the similarity solution with $X_{s}$=0.338976, which corresponds to the shock position of Bertschinger's adiabatic solution (Bertschinger 1985), becomes the 'stagnation' solution. We analytically explain that the divergence at the small radius occurs due to cooling. Because the profiles of the fluid variables are not influenced by cooling in the outer region, we assume that the velocity approaches $V$=$V_{0}X$, which is the same as that of the adiabatic solution (Chuzhoy \\& Nusser 2000). The density and temperature can be expressed from equations (\\ref{eq:equation27}) and (\\ref{eq:equation29}) as follows: \\begin{equation} D=D_{0}X^{s},\\nonumber\\\\ C=C_{0}X^{t},\\nonumber\\\\ \\label{eq:equation39} \\end{equation} where \\begin{equation} s=\\frac{1}{V_{0}-\\xi}(2-3V_{0}), \\label{eq:equation40} \\end{equation} \\begin{eqnarray} t=\\frac{1}{V_{0}-\\xi}[-(\\xi-1) -\\frac{K_{0}}{2}D^{A-1}C^{2B-2} -\\frac{3}{2}(\\gamma-1)V_{0} \\nonumber \\\\ +\\frac{\\gamma(\\gamma-1)}{2}{\\mathcal{H}}]. \\label{eq:equation41} \\end{eqnarray} As the gas relaxes to a hydrostatic equilibrium in order to vanish the velocity as $X\\rightarrow0$, 2$t-s$=2 is satisfied by equation (\\ref{eq:equation28}), in which case $V_{0}$ is \\begin{eqnarray} V_{0}=\\frac{1}{3\\gamma-4}[-K_{0}D^{A-1}C^{2B-2} +\\gamma(\\gamma-1)\\mathcal{H}]. \\label{eq:equation42} \\end{eqnarray} The first and second terms on the right-hand side represent the cooling and heating terms, respectively. We consider the behavior of the cooling term in the inner region in the case of $\\mathcal{H}$=0. Assuming that the cooling term converges to zero at the center, $D^{A-1}C^{2B-2}$ becomes \\begin{eqnarray} D^{A-1}C^{2B-2}\\propto X^{\\kappa}, \\label{eq:equation44} \\end{eqnarray} where \\begin{eqnarray} \\kappa=-\\frac{2}{\\xi}(A-1)+\\frac{\\xi-1}{\\xi}(2B-2)>0. \\label{eq:equation45} \\end{eqnarray} However, the parameters $A$ and $B$, which satisfy $\\kappa>$0, are not in the range of $0< \\epsilon \\leq1$. Thus, $\\kappa<0$ is satisfied and the first term on the right-hand side of equation (\\ref{eq:equation42}) diverges, resulting in a steeper drop of the temperature. That is to say, a cooling catastrophe occurs in the inner region.  The second type is an 'adiabatic' solution (dotted line). It is a free-fall solution to a point mass. Thus, the density and velocity asymptotically approach $D$$\\propto$$X$$^{-3/2}$ and $V$$\\propto$$X$$^{-1/2}$, respectively. Because the cooling and heating times are longer than the flow time, they are negligible for $X$$\\rightarrow$0. For the law conservation of energy, the sound speed approaches $C$$\\propto$$X$$^{-1/2}$ as $X\\rightarrow$0.  The third type of solution is the 'eigensolution' (solid line), which is a marginal case of 'stagnation' solutions. That is to say, the 'eigensolution' satisfies the condition $t_{flow}$=$t_{cool}$ at the origin. A divergence of the physical variables does not occur, and the flow extends to $X$=0. The 'eigensolution' is the only solution where the mass is zero at the origin. We only derive the 'eigensolution' below. ",
        "conclusions": "In this paper, we present self-similar solutions including cooling and heating for collisional gas. We assume that the cooling rate has a power-law dependences of the gas density and temperature, $\\Lambda$$\\propto$$\\rho^{A}T^{B}$, and that the heating rate is $\\Gamma$$\\propto$$\\rho T$. In order to obtain a similarity solution, $A$ and $B$ are selected by requiring that the cooling time is proportional to the dynamical time. The main results in this paper are the following:  \\begin{enumerate} \\item In the region where the cooling rate is stronger than the heating rate, a cooling inflow is established, and the gas is compressed and heats up. Because the compression is stronger in the inner region than in the outer region, the temperature increases with decreasing radius. Furthermore, if the large infall velocity, $V$$<<$$-\\xi X$, is produced due to an enormous energy loss, the temperature slope approaches $-$1, and the density slope approaches a value that depends on $A$, $B$, and the velocity slope $w$.  \\item In the region where the heating rate is larger than the cooling rate, the infall velocity is suppressed by heating, the compression is weakened and the gas cools down. The slope of the density becomes shallow due to suppression of the collapse, and the temperature is lower than that without heating.  \\item For $\\epsilon$=1/3 ($n$=$-$1) and ($A$,$B$)=(2,5/2), the 'stagnation' solutions are only derived in the cases of $K_{0}$=0.1 or 0.01. The reason for this is that the cooling is very strong in the inner region, and a cooling catastrophe occurs. We include the heating and a small $K_{0}$ in order to decrease the effect of cooling, and find heating solutions that can balance the cooling with $K_{0}$=10$^{-4}$ and ${\\mathcal{H}}$=0.7-0.8.  \\end{enumerate}  We consider that the self-similarity solution presented here can be used for modeling of structure formations. For example, the value of the slope parameter, $\\beta_{fit}$, obtained from observations of the surface brightness profile indicates the typical value for rich clusters, $\\beta_{fit} \\sim 2/3$. On the other hand, group systems have significantly flatter slopes with $\\beta_{fit} \\sim 0.4$ (Arnaud \\& Evrard 1999; Helsdon \\& Ponman 2000). As an origin of such discrepancies, it has been proposed that the intracluster medium (ICM) extends to the outer regions by non-gravitational heating, leading to flatter density profiles (Metzler \\& Evrard 1997). Furthermore, in the model $A$=2 and $\\epsilon$=1/3($n$=$-$1), which is close to those of real clusters, a cooling catastrophe occurs if there is no heating, and there are the heating solutions that can balance the cooling. This may mean that real clusters oscillate between a heating dominated regime and one with a cooling catastrophe (Kaiser \\& Binney 2003), or that the heating sources, such as AGN (Churazov et al. 2001; Quilis et al. 2001; Reynolds et al. 2002; Br$\\ddot{\\rm{u}}$ggen et al. 2002; Br$\\ddot{\\rm{u}}$ggen \\& Kaiser 2002; Br$\\ddot{\\rm{u}}$ggen 2003) or thermal conduction (Fabian et al. 2002; Ruszkowski \\& Begelman 2002; Voigt et al. 2002; Zakamska \\& Narayan 2003), balance the cooling. Of course, it is difficult to directly apply the solution presented here to a real cluster, because the time evolutions of the cooling and heating differ from the Hubble time. For a realistic treatment of this problem, it is necessary to perform a numerical simulation while assuming regular radiative cooling and heating rates. However, the accuracy of the numerical simulation is limited by the finite resolution. We consider that a comparison between the simulation and the similarity solution gives insights into the thermal evolution of clusters of galaxies."
    },
    "0311/astro-ph0311320_arXiv.txt": {
        "abstract": "An understanding of the oscillations of differentially rotating systems is key to many areas of astrophysics.  It is of particular relevance to the emission of gravitational waves from oscillating neutron stars, which are expected to possess significant differential rotation immediately after birth or binary merger.  In a previous paper we analysed the normal modes of a simple system exhibiting differential rotation. In this complementary paper we address the initial value problem for the same simple model using both analytical methods and numerical time evolutions. We derive a necessary and sufficient condition for dynamical shear instability.  We discuss the dynamical behaviour of the continuous spectrum in response to an initial perturbation, and show that certain singular solutions within the continuous spectrum appear physically indistinguishable from the discrete modes outside the continuous spectrum. ",
        "introduction": "The oscillations of differentially rotating systems are important in many areas of astrophysics.  They are of particular relevance to neutron star astrophysicists, because non-axisymmetric oscillations of compact stars are promising sources of gravitational waves. Detection and interpretation of gravitational waves from oscillating neutron stars could answer many questions about neutron star physics that are difficult to answer using electromagnetic observations, for example regarding the supranuclear equation of state.  Emission of gravitational waves from neutron star oscillations will be strongest if the oscillations can grow via some dynamical or secular instability. The most promising instability scenarios arise for newly-born, accreting or newly-merged neutron stars. But it is in precisely these situations that we expect the stars to be differentially rotating.  The latest simulations of core collapse indicate that the nascent neutron star emerges  rotating differentially \\citep{dm02a,aki03,ott04}.  \\citet{fuj93} has shown that rapid accretion from a companion could drive differential rotation in the surface layers of the neutron star, and a similar effect may occur during rapid accretion of supernova fallback material \\citep{wat02}.  \\citet{shi00a} and \\citet{bau00} have shown that a long-lived massive differentially rotating neutron star may be generated as the result of the merger of two neutron stars. In addition, recent work on r-modes in neutron stars suggests that the modes may drive a uniformly rotating star into differential rotation via non-linear effects \\citep{rez00,lev01,lin01}. Understanding the effect of differential rotation will be critical if we are to model gravitational wave emission from these systems correctly.  Differential rotation introduces two interesting phenomena.  Firstly, it may lead to dynamical shear instabilities \\citep{pap84,pap85,bal85,luy90a}.  Secondly, the dynamical equations of the inviscid problem are complicated by the presence of corotation points.  These are points where the pattern speed of a mode matches the local angular velocity, and at these points the governing equations are formally singular.  The range of frequencies for which modes possess corotation points is termed the corotation band. It gives rise to a continuous spectrum, the dynamical role of which is unclear.  Neither issue has received significant attention in the stellar perturbation literature (see \\citet{wat03a}, hereafter referred to as Paper I, for a summary of previous work).  We have begun to address the problem by modelling a simple system, a differentially rotating incompressible spherical shell. This model is simplistic, but exhibits the main features of the more complex three-dimensional problem.  In Paper I we conducted a normal mode analysis for the shell. In the uniform rotation limit the shell problem admits only the standard r-modes as solutions.  As the degree of differential rotation increases, the frequencies of these modes gradually approach the lower edge of the expanding corotation band.  Most modes approach the boundary tangentially and appear to cease to exist when the degree of differential rotation exceeds a certain threshold value. In only a few cases do modes cross the corotation band boundary.  For real frequencies in the corotation band the shell problem admits a continuous spectrum of solutions with discontinuous first derivatives.  In general, the continuous spectrum solutions have both a logarithmic singularity and a finite step in the first derivative at the corotation point.  For certain special frequencies, however, the step in the first derivative vanishes. We refer to these solutions as zero-step solutions.  Since these solutions appear when regular normal modes cross into the corotation band, they are likely to be of particular importance. Furthermore, in differentially rotating cylinders, these solutions can merge to give rise to dynamical instabilities \\citep{bal85}. Even though we do not observe such behaviour in our problem it is an interesting possibility that should not be overlooked in future studies.  The character of the modes that cross into the corotation band varies. Only in one exceptional case, for the simplest rotation law examined, do we observe a mode crossing the corotation band boundary and continuing to exist as a regular stable mode.  We think that such regular modes within the corotation band are unlikely to exist for more realistic rotation laws. In general, at all points where modes cross the corotation band boundary we observe the development of zero-step solutions.  At some points we also observe the development of dynamically unstable modes.  A necessary (but not sufficient) condition for dynamical instability is that the equilibrium vorticity have a turning point on the shell.  Some of the rotation laws examined in Paper I meet this condition when the degree of differential rotation exceeds a certain threshold value.  All modes crossing the corotation band boundary above this threshold become dynamically unstable.  In addition, we observe for one rotation law the development of dynamical instability via the merger of two stable modes outside the corotation band.  Paper I raised a number of questions.  Is it possible to derive a necessary and sufficient condition for instability? Do the singular solutions of the continuous spectrum have physical meaning?  And do the zero-step solutions appear distinct from the continuous spectrum in time evolutions?  In this paper we resolve these questions by addressing the initial value problem for the shell. By trying to isolate the various contributions to the evolution of smooth initial data, we hope to clarify the physical role of different parts of the spectrum of a differentially rotating system. The results for this simple toy model will guide us in future investigations of the three dimensional problem \\citep{wat03b}. ",
        "conclusions": "The main question arising from the normal mode analysis of Paper I was the nature of the physical perturbations associated with the singular eigenfunctions of both the zero-step corotating solutions, and the continuous spectrum within which they exist.  By studying the initial value problem both analytically and numerically we have been able to resolve this issue.  Firstly, we have confirmed that the physical perturbations associated with these phenomena are not singular.  Their time dependence in response to an initial perturbation, however, is complicated.  For appropriate initial data, the zero-step solutions behave in much the same way as stable normal modes outside corotation. That is, they produce a clear peak in the power spectrum at a fixed frequency, standing out from the rest of the continuous spectrum.  The fact that such oscillatory solutions can exist within the continuous spectrum is of great relevance for future studies of the oscillations of differentially rotating systems.  The behaviour of the rest of the continuous spectrum is difficult to predict analytically, and can only be elucidated by numerical time evolutions.  The analytic calculations suggest that the physical perturbation associated with the continuous spectrum is oscillatory, with a complicated frequency dependence that varies in part according to the position on the shell.  The amplitude of the oscillations could grow, remain at constant amplitude or die away with time (as a an inverse power of time), depending on the initial data and the rotation law.  For the rotation laws and initial data that we have studied we observe only decay, but we cannot rule out either constant amplitudes or growth.  The results of numerical time evolutions have verified to high accuracy the predictions of mode frequencies made by the normal mode analysis.  This is true both for modes outside the corotation band (including dynamically unstable modes), and the zero-step solutions within it.  This is encouraging, as when we move on to study more realistic stellar models, we are likely to rely more on time evolutions than on increasingly complicated mode analysis."
    },
    "0311/astro-ph0311050_arXiv.txt": {
        "abstract": "We report the discovery of optical and radio afterglow emission from the extremely soft X-ray flash, XRF 020903.  Our spectroscopic observations provide the first redshift for an X-ray flash, thereby setting the distance scale for these events.  At $z=0.251$, XRF 020903 is one of the nearest cosmic explosions ever detected, second only to the recent GRB 030329 and the unusual GRB 980425/SN 1998bw.  Moreover, XRF 020903 is the first X-ray flash for which we detect an optical afterglow. The luminosity of the radio afterglow of XRF 020903 is 1000 times greater than that of Ibc supernovae but similar to those of GRB afterglows. From broadband afterglow modeling we show that the explosion energy of XRF 020903 is not dissimilar from values inferred for typical gamma-ray bursts, suggesting that these cosmological explosions may derive from a similar mechanism. ",
        "introduction": "Prior to the detection of afterglows, gamma-ray bursts (GRBs) were enshrouded in mystery for nearly thirty years.  Great progress in our understanding of these energetic events came with the first redshift measurement which placed GRBs at cosmological distances \\citep{mdk+97}.  In a similar fashion, the mystery of X-ray flashes (XRFs) has been fueled by the absence of confirmed redshifts.  These events, identified in the 1990s by {\\em BeppoSAX}, are characterized by a peak energy in $\\nu F_{\\nu}$ of $E_{\\rm peak}\\sim 25$ keV (compared to $E_{\\rm peak}\\sim 250$ keV for GRBs).  With a distribution of durations similar to those observed for GRBs, it has been assumed that XRFs are associated with GRBs and therefore share their extragalactic distance scale \\citep{hzk+01}.  The subsequent discovery of XRF X-ray and radio afterglows with properties similar to those observed in GRB afterglows has further strengthened this association \\citep{hyf+01, acf+02, tfk01}.  Still, the question of whether the difference between GRBs and XRFs is intrinsic or extrinsic remains unanswered.  Assuming XRFs are simply GRBs observed away from the jet collimation axis, they would have less $\\gamma$-ray emission and the difference is extrinsic, based solely on the line-of-sight to the observer.  On the other hand, the difference could be intrinsic, namely XRFs may represent a class of explosions which are similar in energetics to GRBs yet characterized by less relativistic ejecta possibly due to a heavier baryonic load.  It is clear that by setting the distance scale for XRFs (and hence their energy scale) we can begin to distinguish between extrinsic and intrinsic effects.  In this paper we present the first spectroscopic redshift for an X-ray flash, XRF 020903, placing this event among the nearest high energy explosions and offering confirmation that X-ray flashes are cosmological and produce a total energy output similar to that observed in GRBs. ",
        "conclusions": "We present radio and optical observations of the afterglow of XRF 020903 -- the first X-ray flash to have a detected optical afterglow and a spectroscopic distance determination. At a redshift of $z=0.251$, this burst has set the distance scale for XRFs and confirmed the assumption that they are cosmological in origin. The host galaxy of XRF 020903 appears to be a typical star-forming galaxy similar to those of GRB host galaxies.  The isotropic energy release of the prompt (X-ray) emission of XRF 020903 is at least two orders of magnitude smaller than those of GRBs. However, from our broadband modeling of the afterglow we determine that the total kinetic energy is about $10^{50}\\,$erg, not dissimilar to those inferred for GRBs \\citep{pk02,bkf03}.  In comparison with a larger sample of cosmic explosions, XRF 020903 is a clear example where less energy is coupled to high Lorentz factors. This source highlights the diversity in high energy transients and underscores the importance of studying spherical explosions."
    },
    "0311/astro-ph0311099_arXiv.txt": {
        "abstract": "{ While there are many parsec\\,-\\,scale Herbig-Haro (HH) outflows known to be driven by low\\,-\\,mass young stars, few are associated with their intermediate mass counterparts. Here we present the discovery of five such bipolar outflows. Of these, LkH$\\alpha$\\,198, 1548C27\\,IRS\\,1, \\LH and \\Lk were previously known to possess small-scale HH flows, while no such activity was observed before near IRAS\\,19395+2313. The largest of the newly discovered outflows are seen in the vicinity of \\Lk and 1548C27\\,IRS\\,1, and stretch (in projection) 8\\,pc and 7.5\\,pc respectively. \\LH which was previously known to power a spectroscopically detected small-scale ($\\leq$10\\arcsec) jet is now seen to drive a 3\\,pc outflow and LkH$\\alpha$\\,198 is shown here to power a 2\\,pc outflow. Two HH objects in the vicinity of IRAS\\,19395+2313 lead us to suggest that it may also be responsible for a 5\\,pc outflow. In total, 27 new HH objects/complexes were discovered. Examination of these parsec\\,-\\,scale outflows show that they have similar lengths, morphologies, and dynamical timescales as those from low\\,-\\,mass sources. Many appear to have blown out of the parent cloud, suggesting that their total lengths are much greater than optically observed. The degree of collimation of these outflows is similar to those from low\\,-\\,mass sources suggesting that the transition to more poorly\\,-\\,collimated outflows must occur at higher masses than the sources observed here. ",
        "introduction": "\\indent HH objects are the shock\\,-\\,excited nebulous tracers of outflows from pre\\,-\\,main sequence stars. In many cases these outflows are collimated in the form of highly supersonic jets the existence of which appear to be intrinsicly linked to accretion by the underlying young stellar object (YSO) \\citep{Hartigan95}. Most known optical jets have low mass ($\\sim$ 1\\Msun) sources: either the  embedded (IRAS Class I) counterparts of classical T\\,Tauri stars or classical T\\,Tauri stars themselves \\citep{Reipurth97}.  Turning to higher mass YSOs ($\\verb+>+$ 10\\Msun), such as those driving the Orion OMC\\,1 or Cepheus A outflows, one sees a very different picture \\citep{Allen93, ODell97, Hartigan20}. Their outflows, although highly energetic, often appear poorly collimated and more chaotic \\citep{Reipurth01} than their low mass counterparts. This transformation begs two obvious inter-related questions: at what point does the transition occur and is it a smooth function of mass? To answer these questions one must examine outflows from YSOs of mass greater than $\\sim$2\\,\\Msun$\\!\\!$.  Optical outflows have been observed from a number of intermediate\\,-\\,mass (2\\,\\Msun$\\leq$\\,\\Mstar$\\leq$\\,10\\,\\Msun$\\!\\!$) YSOs, for example R\\,Mon, \\Lk$\\!\\!$, and AFGL\\,4029 (Ray et al., 1990 and references therein). Such stars, where optically visible, are known as Herbig Ae/Be stars (HAEBES) although their embedded counterparts have also been seen. However optical outflows from these YSOs are rare and there are a number of reasons for this. The initial mass function favours the production of low mass stars and therefore intermediate mass YSOs tend to be found at relatively large distances. More massive stars also have a faster evolutionary timescale -- i.e.\\ they evolve more quickly than low\\,-\\,mass YSOs and so their outflow phase is shorter. This would also make their outflows more difficult to detect. Another contributing factor could be that massive stars tend to be more obscured : intermediate and massive stars tend to be surrounded by large amounts of circumstellar gas and dust, making it harder to find an outflow at visual wavelengths. Finally there may well have been a historical bias towards studying outflows from low\\,-\\,mass stars. The situation however has changed in recent years as more and more studies focus on higher mass YSOs.  With these ideas in mind, we have investigated the occurrence of large-scale outflows from intermediate mass stars. By large scale we mean those stretching several parsecs e.g.\\ the PV Cephei outflow at 2.6 pc \\citep{Gomez97,Reipurth97}, the HH\\,80/81 5.3 pc outflow \\citep{Marti93} and the HH\\,354 outflow at 2.4 pc \\citep{Reipurth97}. We emphasise that these outflows have vastly longer associated timescales than those of ``traditionally'' observed flows.  A small\\,-\\,scale HH jet close to its source has a dynamical timescale of only a few hundred years, whereas the HH objects in these parsec\\,-\\,scale flows trace mass ejection over tens of thousands of years. They are, in effect, fossil records of the mass\\,-\\,loss histories of their parent star.  Newly detected parsec\\,-\\,scale outflows around five intermediate\\,-\\,mass young stars are discussed here (see also Table\\ \\ref{HHPositions} and Table\\ \\ref{parameters}). Of these, LkH$\\alpha$\\,198, \\LH and \\Lk are of spectral type A and 1548C27 is A7\\,-\\,F0. All of these stars were known to possess small-scale optical outflows. The one optically invisible YSO in our sample, IRAS\\,19395+2313, was not previously known to drive any outflow.  Section 2 describes how we made our observations and in Section 3 we present our newly discovered large-scale flows. The implications of our findings are discussed in Section 4 and our conclusions are drawn in Section 5. ",
        "conclusions": "We have investigated the occurrence of parsec-scale outflows from intermediate\\,-\\,mass YSOs. As is the case with lower mass YSOs, such flows appear to be common and we report the discovery of four here with well defined sources. These include LkH$\\alpha$\\,198, 1548C27\\,IRS\\,1, \\LH\\, and IRAS\\,19395+2313.  All of these YSOs, with the exception of the last, were previously known to have small-scale outflow activity.  The region surrounding the Herbig Ae star \\Lk is cluttered with outflows and candidate sources. Twelve new HH objects are reported on here, with many of them lying along a preferential axis centred near \\Lk and orientated in a northeast-southwest direction. Their morphology suggests that at least some are part of a large-scale flow centred on the core of the NGC\\,7129 cluster.  Parsec\\,-\\,scale outflows from intermediate\\,-\\,mass YSOs show a number of similarities to those from their low\\,-\\,mass counterparts. In particular :\\\\ - their lengths and degree of collimation appear comparable, \\\\ - they share the same morphological trends such as decreasing frequency, increasing dimension and increasing complexity of HH emission with distance from the source.\\\\  The lengths of these large-scale outflows are usually comparable to the clump size of their associated clouds. As their expected lifetimes are much larger than the apparent dynamical timescales, this suggests that many have ``blown-out'' of the cloud complex.  Finally, it is evident that the transition from highly collimated jet\\,-\\,like flows to poorly collimated wide\\,-\\,angle outflows such as OMC\\,1 must lie at higher masses and luminosities than the sources studied here.  \\acknowledgement{ FMcG and TPR acknowledge support from Enterprise Ireland. JBs research was supported by NSF grant AST-9819820, NASA grant ANG5-8108 (LTSA), and NASA grant NCC2-1052 (CU Center for Astrobiology).}"
    },
    "0311/astro-ph0311266_arXiv.txt": {
        "abstract": "{ { Self-gravitating accretion discs have only been studied in a few nearby objects using maser spots at the parsec-scale. We find a new spectral window for observing the self-gravitating accretion disc in super-Eddington accreting Active Galactic Nuclei (AGNs). This window is determined by calculating the outermost radius ($\\rsg$) of a non self-gravitating disc and the corresponding emission wavelength ($\\lsg$) as a function of various disc parameters. We find that $\\lsg$ reaches $\\sim 4000 \\rmAA$ for $\\alpha=0.1$, when $\\Mdot \\gtrsim 70 \\, (\\Mbh / 10^7 \\Msun)^{-1} \\, \\Leddc2$ (where $\\alpha$, $\\Mdot$, $\\Mbh$ and $\\Ledd$ are, respectively, the viscosity parameter, gas accretion rate onto the central black hole (BH), the BH mass and the Eddington luminosity). Moreover, $\\lsg$ is as small as $\\sim 1500 \\rmAA$ for $\\alpha =0.001$, which is the smallest $\\alpha$ case in this study. Therefore, the window for observing the self-gravitating part of an AGN accretion disc is from $\\sim 2 \\mu$m to $\\lsg$. Incidentally, $\\rsg$ can be less than the photon trapping radius for $\\Mdot \\gtrsim 10^{3.3} \\Leddc2$. Namely, a self-gravitating, optically-thick, advection-dominated accretion disc is expected to appear in the extremely high accretion rate regime.  Next, we demonstrate that the Mid-Infrared to X-ray spectrum of a bright, well-studied Narrow-Line Seyfert 1 galaxy, Ton~S~180, is indeed well fitted by the spectrum arising from the following three components: an inner slim disc (with a corona), an outer, self-gravitating non-Keplerian disc and a dusty torus. The total mass, BH mass plus the entire disc mass, is found to be about $(1.4 - 8.0) \\Mbh$. If the surface density varies with radius $r$ in proportion to $r^{-0.6}$, the total mass is consistent with the central mass estimated by H$\\beta$ and [O III] widths. } ",
        "introduction": "The Optical/UV/X-ray emission from Active Galactic Nuclei (AGNs) is thought to arise from a hot % accretion disc around a super-massive Black Hole (BH; e.g., Mushotzky, Done \\& Pounds 1993; { Koratkar} \\& Blaes 1999), while the { Mid- and Near Infrared (IR) radiation comes from a cold}, dusty torus (e.g., Telesco \\ea 1984; Thatte \\ea 1997; Pier \\& Krolik 1993). % Disc self-gravity can generally be ignored in the innermost region where the high energy spectrum is formed. At large distances (typically $\\gtrsim 10^3 - 10^4$ Schwarzschild radii, depending on accretion parameters), the disc mass is expected to play a significant role. Disc self-gravity has been tested only at the parsec-scale via maser spots detected in a few nearby objects (e.g. Nakai, Inoue \\& Miyoshi 1993; Miyoshi \\ea 1995; Hur\\'e 2002; Lodato \\& Bertin 2003).  Narrow Line Seyfert 1 galaxies (NLS1s) and their high luminosity analogue, % Narrow-Line QSOs, are supposed to have high accretion rates among the AGN population (e.g, Brandt \\& Boller 1998; { Mineshige \\ea 2000}). The gas accretion rate $\\Mdot$ for { those} objects { is} expected to be comparable to or larger than $10 \\Leddc2$, where $\\Ledd$ is the Eddington luminosity. For such super-Eddington accretion rates, the outer edge of the non self-gravitating part of the accretion disc { radiates} optical { continuum} emission ($\\sim 4000 \\rmAA$) for a certain % parameter set (Kawaguchi 2003). Thus, they are potentially good candidates to study the self-gravitating part of the disc observationally at longer wavelengths (Collin et al. 2002; Kawaguchi 2003).  In this paper, we qualitatively describe the wavelength corresponding to the emission from the outer edge of a non self-gravitating disc for various accretion parameter sets. { This is done in \\S 2.} We present { in \\S 3 } a Mid-IR to X-ray spectral modeling of a bright, well-studied NLS1, Ton~S~180 (PHL912), for which numerous, multi-waveband observations are available (e.g., Wisotzki \\ea 1995; Turner et al. 2002; Vaughan et al. 2002). In \\S 4, several discussions are presented. The final section is devoted to a summary. ",
        "conclusions": "We have presented the outermost radius of the non self-gravitating accretion disc ($\\rsg$) around a super-massive black hole (BH), i.e. the inner edge of the self-gravitating disc, as a function of the BH mass $\\Mbh$, accretion rate $\\Mdot$ and viscosity parameter $\\alpha$. We then showed the corresponding wavelength ($\\lsg$) that is relevant to the emission from $\\rsg$.  When $\\Mdot \\lesssim 70 \\, \\Leddc2$ (for $\\Mbh = 10^7 \\Msun$), $\\lsg$ is roughly proportional to $\\Mdot^{-0.5}$ for fixed $\\Mbh$ and $\\alpha$. With a higher $\\Mdot$, $\\lsg$ reaches and stays at $\\sim 4000 \\rmAA$ for $\\alpha=0.1$. Interestingly, $\\lsg$ is as small as $\\sim 1500 \\rmAA$ for the lowest $\\alpha$ case in this study (i.e. $\\alpha =0.001$). Therefore, the continuum emission from $\\sim 2 \\mu$m to $\\lsg$ arises from the self-gravitating part of the super-Eddington accreting disc. Thus, we have discovered a spectral window for observing the self-gravitating disc, which has only been studied by maser spots at the parsec-scale for a few nearby objects.  Next, we demonstrated that the mid-IR to X-ray spectrum of a bright, well-studied Narrow-Line Seyfert 1 galaxy, Ton~S~180, is indeed well fitted by the spectrum arising from the following three components: an inner slim disc (with a corona), an outer, self-gravitating non-Keplerian disc and a dusty torus. % { Comparing the model spectra of the slim disc with the observed UV--X-ray one, the following parameters are favored: $\\Mbh \\approx 10^{6.8} \\Msun$, $\\Mdot \\approx 1000 \\Leddc2$, and $\\alpha \\approx 0.002$.} In the model for the outer, self-gravitating disc, { we allow the viscosity parameter $\\alpha$ to increase with radius, and } an outer radius of $3 \\times 10^4 \\rs$ is inferred. { Various profiles for the surface density $\\Sigma$ are acceptable, e.g., from $\\Sigma \\propto r^{0.3}$ to $\\Sigma \\propto r^{-1.5}$, depending on $\\alpha$ at the outermost radius of the disc.}  { The accretion disc is quite massive, and the disc mass can be $(0.4 - 7.0) \\Mbh$.} The total mass, BH mass plus the entire disc mass, is found to be about $2.4 \\, \\Mbh = 10^{7.2} \\Msun$ { if $\\Sigma \\propto r^{-0.6}$}, which is consistent with the central mass estimated by { H$\\beta$ and [O III] line widths. This indicates that $\\Mbh$-estimations using line widths systematically overestimate the BH mass.} Although the disc mass is quite high relative to the BH mass, the Toomre $Q$ parameter remains greater than one in the whole disc, so that the disc is marginally stable. %  Moreover, $\\rsg$ can be less than the photon trapping radius for $\\Mdot \\gtrsim 10^{3.3} \\Leddc2$. Thus, a self-gravitating, optically-thick, advection-dominated accretion disc is expected to appear in the extremely high accretion rate regime."
    },
    "0311/hep-th0311071_arXiv.txt": {
        "abstract": "{ An exponential potential of the form $V\\sim \\exp(-2c\\,\\varphi/M_p)$ arising from the hyperbolic or flux compactification of higher-dimensional theories is of interest for getting short periods of accelerated cosmological expansions. Using a similar potential but derived for the combined case of hyperbolic-flux compactification, we study the four-dimensional flat (and open) FLRW cosmologies and give analytic (and numerical) solutions with exponential behavior of scale factors. We show that, for the M-theory motivated potentials, the cosmic acceleration of the universe can be eternal if the spatial curvature of the 4d spacetime is negative, while the acceleration is only transient for a spatially flat universe. We also briefly comment on the size of the internal space and its associated geometric bounds on massive Kaluza-Klein excitations. }  \\preprint{\\hepth{0311071} \\\\ IC/2003/148}   \\begin{document} ",
        "introduction": "\\label{sec-intro}  A tentative acceptance of cosmic acceleration of the universe~\\cite{astro1} has stimulated many works on string or M theory cosmology~\\cite{Cornalba02a,Townsend03a,CHNW03b, Ohta1,Emparan03a,MNR1,Ohta03b,GKL03a,CHNOW03b,MNR2,Townsend03b,IPN03c,Misc2}. Though it is not explored yet whether string or M-theory predicts the distinctive features of our universe, like a spatially flat (or open) four dimensional expanding universe with small matter fluctuations, most physicists would believe that a four dimensional cosmology should be derived from a fundamental theory of gravity and particle interactions. This fundamental theory could be a theory of supergravity or superstring origin, or even a theory without the need of extra dimensions or supersymmetry (e.g. loop quantum gravity). Nonetheless, it is quite interesting that cosmic acceleration can arise from compactifications of higher-dimensional theories on hyperbolic spaces without or with fluxes~\\cite{Townsend03a,CHNW03b,Ohta1,Emparan03a,CHNOW03b}. These solutions, along with their many generalizations ~\\cite{MNR1,Ohta03b,GKL03a,CHNOW03b,MNR2,Townsend03b,IPN03c} deal with time-dependent scale factors of internal spaces, thereby overcoming a difficulty of ``no-go theorem'' given for a de Sitter type compactification in supergravity theories with static (and warped) extra dimensions~\\cite{nogo1}.  When studying the dynamics of the inflationary universe, in a four-dimensional effective theory, one considers the Lagrangian density \\begin{equation}\\label{action1} {\\cal L} = \\sqrt{-g}\\, \\left( \\frac{M_P^2}{2}\\,{\\cal R}[g]-(\\partial\\varphi)^2-2V(\\varphi)\\right)\\, \\end{equation} with canonically normalized kinetic and potential terms. Of particular importance is the choice of the potential $V(\\varphi)$ whose form depends on particle physics or compactifications used to derive a lower dimensional effective field theory. One would like to derive potentials that allow inflation in the early universe and/or acceleration in the current cosmological epoch. Both cases may be explained in a typical manner by postulating some suitable scalar potentials.  Recently some progress has been made about inflation within string theory~\\cite{KKLT} by keeping the system in the string theory de Sitter vacuum~\\footnote{However, this minimum is only metastable. There are arguments~\\cite{Giddings03a} that the types of potentials one may get in warped de Sitter string theory geometries, with trapped fluxes present in the extra dimensions, can lead to instability of our four-dimensional world.} and by constructing a scalar potential in the $D3/\\bar{D}3$ warped brane background~\\cite{Kachruetal}. These mathematical frameworks assume that the metric of a compact non-singular internal $m$-manifold $\\Sigma_{m}$ is time-independent, and in turn, inflation is non-generic which involves some degree of functional fine tuning. In this paper we consider scalar fields in cosmology by allowing the metric on $\\Sigma_m$ to be time-dependent.  For many classical compactifications of effective supergravity theories, the scalar potentials are of the form $V\\sim \\exp (-2c\\,\\varphi/M_P)$ with the (dilaton) coupling constant $c\\sim {\\cal O}(1)$. The cosmological models with such a potential have been known lead to interesting physics in a variety of context, ranging from existence of accelerated expansions~\\cite{Halliwell87a} to cosmological scaling solutions~\\cite{Ratra88a}, see also, e.g.,~\\cite{scaling,IMS03a} for some related discussions. It is also worth noting that an attractor solution with a simple exponential potential of the above form may lead to cosmic acceleration for natural values of model parameters~\\cite{datafit}.  A number of authors have studied the effects of exponential potentials on M-theory cosmology, including the recent papers concerning hyperbolic or flux compactification and accelerated expansion. So far no general discussion has been given for the combined effects of hyperbolic-flux compactification, see, however, Refs.~\\cite{Emparan03a,MNR2} for some qualitative arguments. In addition, most authors have studied only a spatially flat universe, and this in turn leads to the observation that the accelerated period in these models is typically preceded by a period when the universe was dominated by the kinetic energy of the scalar field. So in this note we allow the spatial curvature ($k$) of the 4d spacetime to take the value $k=0$ or $k=\\pm 1$. We show that the accelerated expansion of our universe may continue forever if $k=-1$. This possibility was pointed out in~\\cite{CHNOW03b,Halliwell87a}; here we make the earlier discussions more precise with further analysis and new results, taking into account the combined effects of hyperbolic extra dimensions and background fluxes. ",
        "conclusions": "In this paper, we have used the M-theory motivated scalar potentials to study the FRW type four dimensional flat (and open) universe cosmologies and shown that a scalar potential of the form $V \\sim  \\Lambda_1 \\exp (-2c\\kappa \\varphi)+ \\Lambda_2 \\exp (-(6/c)\\kappa\\varphi)$ arising from the hyperbolic-flux compactification might be of interest for getting transiently or eternally accelerating expansion of the universe. For a spatially flat universe the accelerated expansion will continue forever only if the coupling $c\\leq 1$, regardless of the background flux. In light of the result that for all known classical compactifications of supergravity theories on some non-trivial curved internal manifolds or toroidal spaces with trapped fluxes, only $c\\gtrsim 1$ arises in practice, we are led to explore other alternatives for cosmic acceleration. We find that an eternally accelerating expansion is possible with $c>1$ if the spatial curvature of the universe is negative. For the case of hyperbolic compactification, this claim was made before in ref.~\\cite{CHNOW03b} based on a result from linearized approximations. Here we have also generalized those results by solving the classical field equations with non-zero fluxes.  It is known, at least, for $m=2$ and $m=3$, that the most compact manifolds are hyperbolic. In the context of hyperbolic compactification, we point out that the existence of a mass gap in the ${\\cal O}(10^2)\\, \\mbox{GeV}$ is not impossible, since the curvature radius of the internal space can be extremely small, like $r_c \\sim {\\cal O}(10) M_P^{-1}$.   One direction of further investigation is to analyse the field equations as a dynamical system, in the manner of Halliwell, using a phase-plane method and see how the solutions behave in phase space asymptotically in the combined case of the hyperbolic-flux compactification. A step in this direction, with the settings $r_c=1$ and ${b}=0$, is recently taken in~\\cite{Pedro} (see also the last section of ref.~\\cite{Miguel03a}), which generalize the earlier phase-plane analysis in ref.~\\cite{Halliwell87a}.  Another direction is to extend our construction by taking some of the extra dimensions to be flat, like $T^4\\times H^3/{\\Gamma|_{free}}$, instead of the $H^7/\\Gamma$ geometry, since the energy spectrum on a CH-$3$ manifold is better known as compared a CHM in four and higher-dimensions. In addition, every hyperbolic manifold $\\Sigma_{m}=\\H^m/\\Gamma$ of finite volume, with $m\\geq 3$, may be decomposed disjointedly into a relatively compact $\\Sigma_0$ and many finite cusps. One may also include the brane (or brane-instanton) effects by allowing a delocalized flat transverse space. In this case, even though the flat part will have zero contribution to a scalar potential, the coupling constant $c$ will change. We hope to return to this and related problems elsewhere."
    },
    "0311/astro-ph0311500_arXiv.txt": {
        "abstract": "We present a detailed analysis of the magnetic field structure found in a set of four general relativistic 3D MHD simulations of accreting tori in the Kerr metric with different black hole spins. Among the properties analyzed are the field strength as a function of position and black hole spin, the shapes of field lines, the degree to which they connect different regions, and their degree of tangling. Strong magnetic field is found toward small radii, and field strength increases with black hole spin.  In the main disk body, inner torus, and corona the field is primarily toroidal. Most field lines passing through a given radius in these regions wander through a narrow radial range, suggesting an overall tightly-wound spiral structure. In the main disk body and inner torus sharp field line bends on small spatial scales are superimposed on the spirals, but the field lines are much smoother in the corona. The magnetic field in the plunging region is also comparatively smooth, being stretched out radially by the infalling gas. The magnetic field in the axial funnel resembles a split monopole, but with evidence of frame dragging of the field lines near the poles of the black hole.  We investigate prior speculations about the structure of the magnetic fields and discuss how frequently certain configurations are seen in the simulations.  For example, coronal loops are very rare and field lines connecting high latitudes on the event horizon to the disk are not found at all. Almost the entire system is matter-dominated; the only force-free regions are in the axial funnel. We also analyze the distribution of current density, with a view toward identifying possible locations for magnetic energy dissipation. Regions of high current density are concentrated toward the inner torus and plunging region. Dissipation inside the marginally stable orbit may provide a new source of energy for radiation, supplementing the dissipation associated with torques in the stably-orbiting disk body. ",
        "introduction": "Accretion onto black holes is potentially the most efficient form of energy generation possible; it is expected that as much as several tens of percent of the rest-mass energy of the accreted matter can become available for radiation (Novikov \\& Thorne 1973; but see also Gammie 1999 and Agol \\& Krolik 2000).  The rate and character of accretion are regulated primarily by angular momentum transport, making the nature of accretion torques a fundamental problem of astrophysics.  These torques result from magnetohydrodynamic (MHD) turbulence, driven by the magneto-rotational instability (MRI) (see the review by Balbus \\& Hawley 1998).  Large-scale numerical simulation is the best tool we have for exploring turbulent systems.  The first attempts to describe the global structure of accretion onto black holes employed codes solving the equations of pseudo-Newtonian physics, i.e. Newtonian dynamics but with a gravitational potential of the Paczy\\'nski-Wiita form (Hawley 2000; Hawley \\& Krolik 2001, 2002; Armitage et al. 2001; Armitage \\& Reynolds 2003; Machida \\& Matsumoto 2003).  For obvious reasons, these could provide only an approximate description of the dynamics, and could not explore the effects of black hole rotation.  With the construction of a fully general-relativistic three dimensional MHD code (De Villiers \\& Hawley 2003, hereafter DH03), it is now possible to achieve a physical description of black hole accretion that is much more complete.  Using this code we have simulated accretion flows onto black holes with four different spin parameters, $a/M = 0$, $0.5$, $0.9$, and $0.998$.  These runs were labeled KD0, KDI, KDP, and KDE in the first paper of this series (De Villiers, Hawley, \\& Krolik 2003, hereafter Paper I), which presented both a detailed description of how these simulations were done and an introduction to the several distinct physical regimes found within the accretion flows.  In this paper we continue our report on these simulations with a discussion of the structure of the magnetic field.  Before presenting this report, we briefly review the nature and conduct of the simulations.  We adopt geometrodynamic units $G = c = 1$ (Misner, Thorne, \\& Wheeler, 1973), and express evolution times in terms of black hole mass $M$.   We also follow the convention that Greek indices represent spacetime vector and tensor components, whereas Roman indices represent purely spatial components. The simulations analyzed in this paper were evolved in Boyer-Lindquist coordinates ($t,r,\\phi,\\theta$) on a $192 \\times 64 \\times 192$ spatial grid ($r \\times \\phi \\times \\theta$). This grid is described as ``high-resolution\" in Paper I. The radial grid runs from a point just outside the event horizon (the event horizon and the inner boundary vary with $a/M$) to $r=120\\,M$ using a hyperbolic cosine distribution to concentrate zones near the inner boundary.  The polar angle $\\theta$ runs from $0.045\\,\\pi$ to $0.955\\,\\pi$ with reflecting boundaries at the polar axis.  Only the quarter plane from $0 \\leq \\phi \\leq \\pi/2$ was included, with periodic boundary conditions in $\\phi$.  The temporal step size, $\\Delta t$, is determined by the extremal light crossing time for a zone on the spatial grid, and remains constant for the entire evolution, as described in DH03. Step sizes are on the order of ${10}^{-2}\\,M$ for each of the four models discussed here.   The state of the relativistic fluid at each point in the spacetime is described by its density, $\\rho$, specific internal energy, $\\epsilon$, $4$-velocity, $U^\\mu$, and isotropic pressure, $P$.  The relativistic enthalpy is $h=1 + \\epsilon + P/\\rho$.  The pressure is related to $\\rho$ and $\\epsilon$ through the equation of state of an ideal gas, $P=\\rho\\,\\epsilon\\,(\\Gamma-1)$, where $\\Gamma$ is the adiabatic exponent.  For these simulations we take $\\Gamma=5/3$.  The magnetic field is described by two sets of variables.  One is the constrained transport (CT) magnetic field, ${\\cal B}^i = [ijk] F_{jk}$, where $[ijk]$ is the completely anti-symmetric symbol and $F_{jk}$ are the spatial components of the electromagnetic field strength tensor. We also define the magnetic field $4$-vector $\\sqrt{4\\pi}\\,b^\\mu = {}^{*}F^{\\mu\\nu}\\,U_\\nu$, where $^{*}F^{\\mu\\nu}$ is the dual of the electromagnetic field tensor.  Because the CT field is guaranteed to be divergence-free, it is the magnetic field description most readily identified with field lines.  On the other hand, the magnetic field four-vector is the one most readily interpreted as a magnetic field for dynamical purposes because it arises naturally in algebraically simple forms for the stress tensor, $T^{\\mu\\nu}$.  The magnetic field scalar is $||b||^2 = b^\\mu\\,b_\\mu$, and appears in the definition of the total four momentum, $S_\\mu = (\\rho\\,h\\ + {\\|b\\|}^2)\\,W\\,U_\\mu$ , where $W$ is the boost factor. Magnetic pressure is given by $P_{mag}={1\\over 2}{\\|b\\|}^2$.  We also define auxiliary density and energy functions $D = \\rho\\,W$ and $E = D\\,\\epsilon$, and transport velocity $V^i = U^i/U^t$.  For each of the four simulations, the initial conditions were basically the same:  a torus with a near-Keplerian initial angular momentum distribution and pressure maximum at $r = 25\\,M$, containing an initial magnetic field consisting of weak poloidal field loops lying along isodensity contours inside the torus.  Each simulation was run to time $8100\\,M$, which is approximately 10 orbits at the initial pressure maximum.  The radial boundaries permit outflow only, so the simulations have only a finite reservoir of matter to accrete.  However, even in the case with the greatest integrated accretion ($a/M = 0$), only 14\\% of the initial mass was accreted in the course of the simulation. After initial transients, the regions of the accretion flow inside the initial pressure maximum evolve into a moderately thick, nearly-Keplerian, highly-turbulent disk.  This structure gives the appearance of rough statistical time-stationarity but also evolves on long timescales.  As outlined in Paper I, the quasi-steady state system is usefully divided into five regions:  the main body of the disk, the coronal envelope, the inner torus and plunging region, the funnel wall jet, and the evacuated axial funnel.  In this paper we make a more extensive examination of the properties and structure of the magnetic fields in these regions.  Note that outside $r \\simeq 25M$ (the location of the initial pressure maximum), we are not simulating an ``accretion disk\" at all, as material in that region moves outward as it absorbs the angular momentum taken from matter at smaller radii.  For this reason, our analysis concentrates on the region $r \\lesssim 25M$, and especially on behavior in the vicinity of the innermost stable circular orbit. In \\S2 we examine the magnetic field strength and distribution throughout the flow.  In \\S3 we quantify properties of the field geometry, and in \\S4 we determine where field dissipation would likely be important.  Section 5 summarizes our findings. ",
        "conclusions": "In Paper I we divided the accretion flow into five distinct regions, each with its own characteristics, as follows: the main disk body, the inner torus and plunging region, the coronal envelope, the funnel, and the funnel wall jet.  In this paper we have examined the magnetic field strengths and topologies in these regions.  \\subsection{Field Characteristics in the Different Flow Regions}  The main disk body is characterized by MHD turbulence due to the MRI. As has been shown in previous studies, the magnetic field is very tightly wound by the differential rotation, but with significant poloidal tangling due to turbulence.  The field wandering index shows that the poloidal field lines extend farther in the radial direction than in the polar.  The field energies are subthermal, with plasma $\\beta \\approx 10$--100 on average.  Strong fluctuations are the rule, however, and $\\beta$ can exceed 1000 in some portions of the main disk body.  In the plunging region, the flow transitions from turbulence-dominated to spiraling inflow.  A consequence of this is that the tight field winding loosens somewhat as the radial component of the velocity grows relative to the azimuthal component.  This behavior is quantified approximately by the field-wandering and field-tangling indices, which show (see Figs.~\\ref{wander} and \\ref{tangle}) that there is a sharp change of behavior near the marginally stable orbit.  Field lines inside that radius travel farther in radius and much less in polar angle than in the disk body.  This is a consequence of the transition from turbulent motions to more regular spiral inflow; that is, this break is yet another way of marking the turbulence edge defined by Krolik \\& Hawley (2002).  Inside the plunging region the field is amplified relative to the fluid pressure.  Near the horizon above and below the equator, $\\beta$ falls from $\\sim 10$ near $r_{ms}$ to less than 1. The equatorial flow itself maintains $\\beta \\sim 10$ or greater throughout the plunging region.  These conclusions apply most clearly to the $a/M=0$, 0.5, and 0.9 simulations; the plunging region in the $a/M=0.998$ simulation is less well-resolved.  The field in the corona is, on average, in equipartition with the thermal energy there, and is dominated by the toroidal component; field line plots look smooth compared to those in the disk body.  There are relatively extended poloidal fields, but these are more associated with outflowing toroidal coils.  When field emerges from the disk body it carries some of the turbulent tangling with it, but this seems to be rapidly smoothed out.  Gas is carried into the corona as well as magnetic field, so the corona does not become a magnetically dominated force-free region.  The magnetic field does have a larger scale-height than the gas pressure, though.  By contrast, the field in the axial funnel is essentially radial outside the ergosphere, with a toroidal component that becomes larger with increasing black hole spin.  In the funnel the magnetic field energy dominates over the matter energy, and the field can be regarded as force-free.  The magnetic energy is not large compared to the magnetic energy in the remainder of the flow, however.  The magnetic pressure in the funnel is in equilibrium with the total pressure in the corona and inner disk. From the field-connectivity data, we see that the magnetic field structure may be divided into two almost wholly independent regions: the accretion flow proper and the axial funnel.  Field lines passing through the funnel almost never enter the other regions.  In the funnel-wall jet, the magnetic field configuration resembles that of the axial funnel, but the field lines are more tightly wound, consistent with the observation that the gas in the jet has specific angular momentum comparable to that at its injection point, above the inner torus. The jet is also less strongly magnetized than the axial funnel, with $\\beta \\simeq 0.3$.  The categorization of field line types listed in Table~\\ref{fieldlinetable} and illustrated in Figure~\\ref{fieldlines} provides a handy summary of the global field structure.  It is important to recognize that only a subset of all possible field line types are found with any frequency in these simulations; the table indicates which these are.  One of the most important global properties of the field is its inter-region connectivity. From this point of view, the different regions fall into two groups: the main disk body is linked to both the corona and the inner torus, while many field lines stretch from the inner torus to the plunging region.  On the other hand, the field lines in the axial funnel have essentially no connection to the other regions.  \\subsection{Astrophysical Implications}  The simulation results provide some insights into potentially important astrophysical processes that might occur in black hole accretion.  For example, Livio et al. (1999) argued that energy extracted electromagnetically from the accretion disk proper would always outweigh that extracted from spinning black holes because the product of magnetic stress and area would be larger for the disks than for black hole event horizons.  As Figure~\\ref{b2polview} shows, the total intensity of the magnetic field increases sharply inward.  This is as true of the poloidal portion as the toroidal.  Increases by factors $\\sim 30$ from $r=5M$ to the plunging region are typical.  Although it is true that, outside the funnel region, there is no large-scale poloidal field in this simulation, such large contrasts in field strength from disk to plunging region raise doubts about the Livio et al. hypothesis.  Moreover, the force-free field in the evacuated funnel seems best configured to permit processes closely related to the Blandford-Znajek mechanism (Blandford \\& Znajek 1977), while the strength of the funnel field is comparable to that in the plunging region.  These results suggest that accretion could well lead to a strong enough field near the event horizon that a genuine Blandford-Znajek process might yield more energy than competing processes associated more directly with the main disk body.  It is often asserted (in fact, frequently in connection with studies of the Blandford-Znajek mechanism) that the magnetospheres of black holes should be, in large part, effectively force-free (e.g., Okamoto 1992, Ghosh \\& Abramowicz 1997, Ghosh 2000, Park 2000, Blandford 2002, Komissarov 2002).  In our simulations we find that the only region where the force-free approximation is applicable is the funnel. In the disk proper and most of the corona, the field structure is thoroughly matter-dominated; in the outer corona and parts of the plunging region, the energy density in the magnetic field begins to approach $\\rho h$, but never exceeds it.  Thus, the relativistic MHD approximation may be better-suited to work in this area than the relativistic electrodynamic approximation.  It is possible, though, that cooling would alter the properties of the corona somewhat by significantly reducing the gas pressure scale height relative to that of the magnetic field.  A common suggestion for the origin of high-energy emission from accretion disks around black holes is that the footpoints of magnetic field lines in the corona are twisted in opposite directions by orbital shear, leading to reconnection (e.g., Blandford 2002).  This process is initiated by establishment of a radial magnetic field link between two fluid elements arranged so that the outer one is advanced in orbital phase relative to the inner one.  As the inner fluid element catches up to and passes the outer, an X-point would be formed.  We see little evidence of this sort of process.  What appears to happen instead is that when radial magnetic field links are created, it is between two fluid elements at nearly the same azimuthal angle (the closer two points are, the more likely they are to be magnetically connected).  The orbital shear then creates a long azimuthal field line as the fluid element at the smaller radius moves ahead of the other fluid element.  If the inner fluid element moves a full $2\\pi$ in azimuth ahead of the outer one before reconnection occurs, the field lines wrap parallel to one another, rather than crossing near their footpoints.   It is generally believed that regions of high current density are likely to promote anomalous resistivity, and therefore be markers for regions of concentrated heat dissipation.  Although we find that, if anything, current sheets are especially rare in the corona, they are very common and very intense in the innermost parts of the accretion flow, particularly the plunging region. The possibility of substantial heat release inside $r_{ms}$ has an interesting consequence.  In the Novikov-Thorne theory of relativistic accretion disks, the energy made available for radiation is the difference between the potential energy lost by accreting matter and the work done by accretion stresses.  Because they assumed that all stresses vanish inside $r_{ms}$, there is no energy released inside that radius.  If the Novikov-Thorne model is generalized to allow non-zero stress at $r_{ms}$ (e.g., Agol \\& Krolik 2000), the amount of energy dissipated inside the disk rises, but it is still determined in essentially the same way.  Dissipation inside $r_{ms}$ is an entirely new and independent mechanism of energy release.  It is not necessarily associated with angular momentum transport; it could result from entirely local processes (e.g., Machida \\& Matsumoto 2003).  Nonetheless, if the time to radiate this heat is shorter than an infall time, the photons released could carry off the energy to infinity, adding to the light seen by distant observers.  One of the main purposes of carrying out these simulations across a broad range of black hole spins is to extract information previously unobtainable from pseudo-Newtonian simulations of accreting tori. The spacetime rotation caused by spinning black holes is a notable example of this sort of effect. Perhaps the most striking spin-dependent property of the magnetic fields is the increase in magnetic pressure near the black hole with increasing spin. This pressure increase is found at the surface of the inner torus and plunging region (i.e., off the equatorial plane), as well as deep in the axial funnel. There is also a corresponding increase in gas pressure in this region, so that the ratio of pressures ($\\beta$) remains effectively unchanged. The growth of magnetic pressure is due in part to the dragging of field lines by the black hole, and in part to the inward displacement of the marginally stable orbit with black hole spin, which tends to allow the MRI to operate closer to the event horizon. There is also a strong spin-dependent growth in the intensity of currents in the vicinity of the black hole. In the $a/M=0$, 0.5, and 0.9 simulations, the high current regions are organized in sheets through the inner torus; these current sheets also have a spiral character, inherited from gas motion in the vicinity of the plunging region. The $a/M=0.998$ simulation shows a much more intense and turbulent current distribution, reflecting the fact that the plunging region is extremely narrow in Boyer-Lindquist coordinates and the the inner torus extends deep into the ergosphere. This strong spin-dependence of the currents suggests that magnetic dissipation may be especially intense in accretion flows near extreme black holes."
    },
    "0311/astro-ph0311446_arXiv.txt": {
        "abstract": "We have analyzed two long BeppoSAX observations of the bright Seyfert galaxy NGC 4151, searching for short timescale ($\\sim$10-200 ksec) X-ray spectral variability. The light curve of a softness ratio, chosen as most sensitive to pinpoint changes of the column density of the absorbing gas along the line of sight, shows significant variations. We try to model these variations by performing a detailed, time resolved, spectral analysis. We find significant, large (factors of 1.5-6) variations of the absorber column densities on time scales of 40-200 ksec. These values are 10-100 times shorter than those found by Risaliti et al. 2002 in a sample of Seyfert 2 galaxies, and provide strong constraints on the geometry of the obscuring medium.  \\vspace{1pc} ",
        "introduction": "The geometry of the obscuring matter in Active Galactic Nuclei, predicted by unified models \\cite{ant}, is still unclear. The study of variability of the absorbing column density (N$_H$) in these objects, provides invaluable informations on the size and location of the absorber.  Recently, Risaliti et al. \\cite{ris} found N$_H$ variability on time scales from $\\sim 2$ month to $\\sim$ 5 years in nearly all the objects belonging to a sample of bright, Compton-thin Seyfert 2 galaxies with multiple X-ray observations. This variability suggests that the absorber could be much closer to the nucleus than usually assumed.  To further investigate this issue, we have analyzed two BeppoSAX long observations of the brightest Seyfert galaxy, NGC 4151, and present here some preliminary results.  \\hspace{0.cm} ",
        "conclusions": "We have performed a detailed time-resolved spectral study of two long BeppoSAX observations of the Seyfert galaxy NGC4151.  We find strong evidences for spectral variability which we model in terms of variations of two absorbing screens.  The timescales of these variations are between 40 ksec and 200 ksec, 10-100 times shorter than the typical timescales found by Risaliti et al.\\cite{ris} in a sample of Compton thin Seyfert 2 galaxies.  If we associate these timescales to a crossing time, the resulting linear size would be of about one light day, which is the size of the innermost part of the Broad Line Region, as inferred from reverberation mapping in NGC 4151. This interpretation is supported by the argument, used by Risaliti et al.\\cite{ris}, which assume that the absorbing medium might be made up of spherical clouds moving with Keplerian velocities around the central black hole.  This requires that the cloud gas density is  of the order of 10$^{9}$ cm$^{-3}$, typical of Broad Line Region clouds \\cite{bl}."
    },
    "0311/astro-ph0311393_arXiv.txt": {
        "abstract": "The optical light curve of the Be/neutron star binary AX~J0049.4$-$7323 has been investigated using data from the MACHO and OGLE-II projects.  This X-ray source, whose neutron star has a very slow rotation rate (P$_{pulse}$= 755.5 sec), shows optical outbursts every $\\sim$394 days. The regularity of these outbursts suggests that their recurrence time is the orbital period of the system.  During the outbursts the system brightens and becomes slightly redder.  A possible interpretation is that a portion of the equatorial disk is excited as the neutron star passes through it during periastron passage.  In the intervals between outbursts the light curve shows $\\sim$11-day quasi-periodic varability which may be associated with the rotation of the Be star's extended disk. ",
        "introduction": "Numerous X-ray pulsars have been found in the Magellanic Clouds, many of which are associated with binary systems containing a Be star and a neutron star (e.g. Yokogawa et al. 2000, Yokogawa et al. 2003, Haberl \\& Sasaki 2000, and references therein).  Corbet (1984) presented evidence that a relationship exists between the neutron star spin period and the orbital period in these systems.  This was further substantiated by Waters \\& van Kerkwijk (1989) who explained the correlation in terms of the equilibrium spin period of the pulsar as a result of the slow equatorial winds from the Be star.  Subsequent data have further strengthened this spin/orbit correlation (e.g. see Fig.\\ 6 by in't Zand et al. 2001).  ASCA observations by Yokogawa et al.\\ (2000) showed that AX~J0049.4$-$7323 has an unusually long pulse period (755.5 s), making it the slowest known rotator among the X-ray pulsars in the SMC.  Hence, one expects that the system to have a long orbital period, based on Corbet's results from other Be/neutron star systems.  Using photometric data from the MACHO and OGLE-II Projects, we have identified what we believe to be the orbital period of $\\sim$394 days for AX~J0049.4$-$7323, based on recurrent optical outbursts.  From optical observations, Edge \\& Coe (2003) identified AX~J0049.4$-$7323 with a Be star which showed a double peaked H$\\alpha$ emission profile at the time of their observations.  They give the magnitude and color of the system as $V$=14.99 and $B-V$=0.05, but they did not address the question of the orbital period or light variability. Using MACHO and OGLE-II data we have investigated the light curve of this Be star and found what we propose to be the orbital period as well as quasi-periodic variability between outbursts. ",
        "conclusions": "We have shown that AX~J0049.4$-$7323 exhibits small, recurrent optical outbursts with a period of $\\sim$394 days.  This period is identified with the orbital period of the system and may be caused by the interaction of the neutron star with the equatorial disk of the Be star during periastron passage.  During the outbursts, the system's color becomes redder, consistent with a brightening in the disk.  In addition, between outbursts the system displays color-independent, quasi-periodic variability, with a timescale of $\\sim$11 days.  We tentatively suggest that this behavior may result from the observer's changing view of the disturbed region in the equatorial disk as it rotates around the Be star."
    },
    "0311/astro-ph0311283_arXiv.txt": {
        "abstract": "{ Accretion of gas onto a central supermassive black hole is generally accepted to be the source of the emitted energy in active galactic nuclei.  The broad emission lines we observe in their optical spectra are probably formed in the wind of an accretion disk at distances of light days to light years from the central black hole. The variable fraction of the emission lines originates at typical distances of only 1 to 50 light days from the central supermassive black hole. We derived a central black hole mass of  $M_{\\rm orbital}=1.8\\pm 0.4\\times 10^{7} M_{\\odot}$  in the Seyfert galaxy Mrk\\,110 assuming the broad emission lines are generated in gas clouds orbiting within an accretion disk. This figure depends on the inclination angle of the accretion disk. Here we report on the detection of gravitational redshifted emission in the variable fraction of the broad emission lines. We derive a central black hole mass of $M_{\\rm grav}=14.0\\pm 3.0\\times 10^{7} M_{\\odot}$. These measurements are independent on the orientation of the accretion disk. The comparison of both black hole mass estimates allows to determine the projection of the central accretion disk angle $i$ to $21\\pm5 \\deg$. in Mrk\\,110 and therefore the orientation of the spin axis of the central  black hole. ",
        "introduction": "Mass and spin of the central black holes in galaxies as well as geometry and kinematics of their immediate stellar and gaseous surroundings are of fundamental interest for understanding the origin of the enormous energies emitted in the nuclear regions of active galaxies (e.g. Blandford \\& Begelman \\cite{blandford99}, Murray \\& Cheng \\cite{murray97}). Furthermore, the central black hole masses and their evolution might be connected to the general properties of their host galaxies. The black hole mass has been estimated in some active galaxies by reverberation mapping (e.g. Wandel et al. \\cite{wandel99}). A basic assumption of this method is that the broad emission line gas is gravitationally dominated by the central black hole. Here a major uncertainty is introduced by the geometry and orientation of the line-emitting gas.  An independent mass determination of the central black hole can be obtained by the gravitational redshift effect in the optical emission lines of active galaxies.  However, in this respect only the detection of a gravitationally redshifted Fe K$\\alpha$ X-ray line has been reported in a few active galaxies (Tanaka et al. \\cite{tanaka95}, Fabian et al. \\cite{fabian00}). ",
        "conclusions": "\\subsection{Mean and rms line profiles}  We determined mean and rms emission line profiles of the most intense lines H$\\alpha$, H$\\beta$, HeI$\\lambda$5876, and  HeII$\\lambda$4686 from our variability campaign. The rms profiles are a measure of the variable component of the line profile. The profiles shown in Fig.~1 are plotted in velocity space and normalized to the same maximum  intensity. \\begin{figure*} \\hbox{\\includegraphics[bb=40 90 380 700,width=55mm,height=85mm,angle=270]{Fj141_f1a.ps}\\hspace*{7mm} \\includegraphics[bb=40 90 380 700,width=55mm,height=85mm,angle=270]{Fj141_f1b.ps}} \\hbox{\\includegraphics[bb=40 90 380 700,width=55mm,height=85mm,angle=270]{Fj141_f1c.ps}\\hspace*{7mm} \\includegraphics[bb=40 90 380 700,width=55mm,height=85mm,angle=270]{Fj141_f1d.ps}} \\vspace*{5mm} \\caption{Normalized mean (dashed lines) and rms (solid lines) Balmer and Helium emission line profiles.} \\end{figure*} The fraction of the variable component in these lines amounts to 9, 12, 14, and 39\\% , respectively. The line widths (FWHM) of the rms profiles are listed in Table~1.  \\begin{table} \\caption{Rms line widths (FWHM) of our strongest emission lines: H$\\alpha$, H$\\beta$, HeI$\\lambda$5876, and  HeII$\\lambda$4686; differential redshift of rms line centers $\\Delta$\\,$v_{\\rm cent}$(rms) ; cross-correlation lags $\\tau$ and central black hole mass estimation $M_{\\rm grav}$ derived from gravitational redshift.} \\begin{tabular}{lcccc} \\hline \\noalign{\\smallskip} Line & FWHM(rms) & $\\Delta$\\,$v_{\\rm cent}$(rms) & $\\tau$ & $M_{\\rm grav}$\\\\ & [km s$^{-1}$] & [km s$^{-1}$] & [days] & [$10^7 M_{\\odot}$]\\\\ (1) & (2) & (3) & (4) & (5)\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} HeII~ & 4444 $\\pm$ 200 & 541 $\\pm$ 60 & 3.9 $\\pm$ 2 & 13 $\\pm$ 3\\\\ HeI~ & 2404 $\\pm$ 100 & 186 $\\pm$ 60 & 10.7 $\\pm$ 6 & 12 $\\pm$ 4\\\\ H$\\beta$ & 1515 $\\pm$ 100 & 118 $\\pm$ 50 & 24.2 $\\pm$ 4 & 17 $\\pm$ 4 \\\\ H$\\alpha$ & 1315 $\\pm$ 100 &  74 $\\pm$ 50 & 32.3 $\\pm$ 5 & 14 $\\pm$ 5 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table} \\subsection{Virial mass of the central black hole}  We calculated the characteristic distances of the line emitting regions by cross-correlating the light curves of the emission lines with the optical continuum light curve. The broader emission lines originate systematically closer to the galaxy center in Mrk\\,110 (see Table~1) (Kollatschny et al. \\cite{kollatschny01}, Kollatschny \\cite{kollatschny03}). The fraction of the variable line component  is correlated with the distance of the line emitting region from the black hole. The distances of the line emitting regions we determined from the cross-correlation function (CCF) analysis correspond to the widths of the rms profiles (Table~1). The non-variable components of the emission lines originate at larger distances from the black hole. We determined the central black hole mass assuming Keplerian orbits of the line emitting clouds: \\[ M_{\\rm orbital} = f v^{2} G^{-1} R .\\] where v is the characteristic emission-line velocity (rms line width), $G$ is the gravitational constant, and $R = \\tau$c is the distance of the line emitting region ($\\tau$ = calculated cross-correlation lag of line emitting region). The derived black hole mass in Mrk\\,110 is \\[ M_{\\rm orbital}= 1.8\\pm0.4\\times 10^{7} M_{\\odot}. \\]  The dimensionless factor $f$ depends amongst other things on the geometry and (unknown) orientation of the BLR. We could show that the velocities in the broad emission line region are not randomly oriented as for instance in a turbulent flow but that the lines are formed in an accretion disk (Kollatschny et al. \\cite{kollatschny02}, Kollatschny \\cite{kollatschny03}). The line widths of the rms spectral lines are interpreted as the line-of-sight Doppler widths. We observe only the projected velocity $v \\sin i$ of the rotating accretion disk. Thus, the derived mass is a lower limit. The systematic errors in the estimation of black hole masses by reverberation mapping have been discussed by Krolik (\\cite{krolik01}). He pointed at a systematic underestimate of the central AGN mass caused by the possible inclination of the line emitting region. \\subsection{Emission line shifts} We determined from the spectra in Fig.~1 the shift of the rms line centers (uppermost 20\\%). It is a known effect that rms line profiles sometimes show narrow forbidden line residuals that are caused by variable seeing and small shifts in wavelength calibration (e.g. Kaspi et al. \\cite{kaspi00}). The degree of this contamination is strongest in the HeI$\\lambda$5876 line which is the weakest one. The redshift ($\\Delta v$) of the rms profiles with respect to the narrow emission lines increases as a function of line width and ionization potential as seen in Fig.~1. The correlation of rms line width with observed central line shift is shown in Fig.~2. \\begin{figure} \\includegraphics[width=100mm,height=98mm,angle=270]{Fj141_f2.ps} \\vspace{-28mm} \\caption{Relative redshift of H$\\alpha$, H$\\beta$, HeI$\\lambda 5876$, and HeII$\\lambda 4686$  rms line centers as a function of rms line width (FWHM)(see Table~1).} \\end{figure}  High-ionization broad emission lines originate closer to the central supermassive black hole than the low-ionization ones (e.g. Peterson \\& Wandel \\cite{peterson99}, Kollatschny et al. \\cite{kollatschny01}). The redshifts derived from different integrated emission lines in AGN often do not agree with each other. High-ionization broad emission lines are sometimes blueshifted with respect to the low-ionization broad emission lines and narrow emission lines (Gaskell \\cite{gaskell82}, Sulentic et al. \\cite{sulentic00}). Therefore, the shift of the integrated broad line profiles cannot be explained by gravitationally redshift effects. But is has been demonstrated recently that the observed blueshift in the highly ionized lines is mimicked by absorption in the red line wings (Richards et al. \\cite{richards02}). This effect is more pronounced in the higher ionized lines. Additionally, one has to consider that the integrated emission lines originate at distances of light weeks to light years from the black hole. A verification of a gravitational redshift in these lines is beyond the detection limit. \\subsection{Central black hole mass derived from gravitational redshift} The variable broad line components originate far more closely to the central black hole than the constant line components. A gravitational redshift effect in this variable component might thus be present. We measured the differential shifts $\\Delta v = \\Delta z~c$ of the strongest rms emission profiles with respect to their mean line profiles (Fig.~1, Table~1). The observed shifts of the rms profiles are identified as gravitational redshifts. The mean profiles of the strongest broad emission lines are not shifted with respect to the forbidden narrow lines on the other hand. Therefore, the differential shifts of the rms profiles with respect to the narrow lines are identical to their shifts with respect of the mean profiles.  We calculated the central black hole mass $M_{\\rm grav}$ in  Mrk\\,110 using the formula (e.g. Zheng \\& Sulentic \\cite{zheng90}) \\[ M_{\\rm grav} = c^{2} G^{-1} R \\Delta z  .\\]  The distances $R$ ($R = c \\tau$) of the individual line emitting regions are known from the cross-correlation analysis (Table 1). We determined a black hole mass $M_{\\rm grav}$ of $12\\pm4$ to $17\\pm4\\times 10^{7} M_{\\odot}$ with a mean value of \\[ M_{\\rm grav} = 14\\pm3\\times 10^{7} M_{\\odot} . \\] from the line shifts. The rms line profile shifts are shown in Fig.~3 as a function of distance to the center.  The internal shifts of all four emission lines (H$\\alpha$, H$\\beta$, HeI$\\lambda$5876, and  HeII$\\lambda$4686) are different. They depend on the distance of the line emitting region from the center. It is intriguing that the black hole masses we derived from the four different line shifts are identical within the error limits. This verifies our interpretation of the observed differential line shifts as gravitaional redshifts. \\begin{figure} \\includegraphics[width=105mm,angle=270]{Fj141_f3.ps} \\vspace{-30mm} \\caption{Redshift of the rms profiles of the HeII$\\lambda 4686$, HeI$\\lambda 5876$, H$\\beta$ and H$\\alpha$ lines as a function of the distance of the rms line emitting region (see Table~1). The dotted and dashed curves show computed lines of gravitational redshifts for central masses of $5.0, 10.0, 13.5, 17.0, 22.0\\times 10^7 M_{\\odot}$ (from bottom to top).} \\end{figure}  The dotted and dashed lines represent theoretical models. They show the correlation between distance (i.e. cross-correlation lag $\\tau$) and gravitational redshift $\\Delta v$ for different central black hole masses $M_{\\rm grav}$ of 5.0, 10.0, 13.5, 17.0, and $22.0\\times 10^7 M_{\\odot}$, respectively.  Other evidence for the detection of a gravitational redshift has been shown before in the spectra of the Seyfert galaxy NGC4593 (Kollatschny et al. \\cite{kollatschny97}). In this galaxy the Balmer lines originate at distances of 2 to 4 light days only. On the other hand, the quality of the spectra in the NGC4593 variability campaign was not as good as in Mrk110.  The derived black hole mass $M_{\\rm grav} = 2.3\\times 10^{7} M_{\\odot}$ in NGC4593 was by a factor of 1.6 higher in comparison to the mass  $M_{\\rm orbital}$ derived from the line widths. An additional hint for a gravitational redshift can be seen in the HST spectra of the NGC5548 variability campaign. The line center of the rms CIV$\\lambda$1550 profile is redshifted with respect to the mean profile by about 10\\AA\\ (see Fig. 4 in Peterson \\cite{peterson01}). \\subsection{Inclination angle of central accretion disk} The determination of black hole masses in AGNs through their gravitational redshift effect has clear advantages. The derived central black hole mass $M_{\\rm grav}$ is not affected by the orientation of the central accretion disk in contrast to the mass $M_{\\rm orbital}$ derived from the emission line widths. It has been mentioned before that only an accretion disk model of the BLR in Mrk\\,110 can reproduce the observed line profile variations (Kollatschny et al. \\cite{kollatschny01}, Kollatschny \\cite{kollatschny03}). The mutual comparison of both mass estimates gives us information on the orientation angle $i$ of the accretion disk: $M_{\\rm orbital}/M_{\\rm grav} = \\sin^{2} i$. For the narrow line Seyfert galaxy Mrk\\,110 we derive a nearly pole-on line-of-sight of the central accretion disk. The inclination angle $i$ amounts to $21\\pm5 \\deg$. A disk inclination angle of  $30\\pm20 \\deg$ has been previously estimated from line profile variations (Kollatschny \\cite{kollatschny03}). This confirms independently the nearly pole-on view of the accretion disk in Mrk\\,110.  The variable HeII line component originates at a mean distance of 3.9 light-days ($\\stackrel{\\scriptscriptstyle\\wedge}{=}9.8\\times 10^{15}$ cm) only. This distance corresponds to 230 Schwarzschild radii $r_{\\rm s}$ ($r_{\\rm s}= 2GM_{\\rm grav}/c^{2}$) for a central black hole mass of $14\\pm3\\times 10^{7} M_{\\odot}$. It is very plausible that the rotation axis of the inner accretion disk is oriented parallel to the spin axis of the central black hole. Altogether, the black hole mass $M_{\\rm grav}$  derived from the gravitational redshift is not only more reliable: a comparison with the black hole mass $M_{\\rm orbital}$ deduced from the line widths opens up the possibility to determine  the orientation of the central black holes in radio-quiet AGN. A good test of the method proposed here would be to check the inclination of optical broad emission line regions with that of radio/optical jets in radio-loud objects."
    },
    "0311/astro-ph0311410_arXiv.txt": {
        "abstract": "We present an analysis of data from both the {\\it R\\\"{o}ntgen Satellite} (\\rosat), and the {\\it Advanced Satellite for Cosmology and Astrophysics} (\\asca) of the supernova remnant (SNR) G18.95$-$1.1. We find that the X-ray emission from G18.95$-$1.1 is predominantly thermal,  heavily absorbed with a column density around 10$^{22}$~atoms~cm$^{-2}$ and can be  best described by an NEI (nonequilibrium ionization) model with a temperature around 0.9~keV, and an ionization timescale of 1.1$\\times$10$^{10}$~cm$^{-3}$~s$^{-1}$.  We find only marginal evidence for non-solar abundances. Comparisons between 21 cm HI absorption data and derived parameters from our spectral analysis strongly suggest a  relatively near-by remnant (a distance of about 2~kpc).  Above 4~keV, we identify a small region of emission located at the tip of the central, flat spectrum bar-like feature in the radio image. We examine two possibilities for this emission region: a temperature variation within the remnant or a pulsar wind nebula (PWN). The current data do not allow us to distinguish between these possible explanations. In the scenario where this high-energy emission region corresponds to a PWN, our analysis suggests a rotational loss rate for the unseen pulsar of  about  7$\\times$10$^{\\rm 35}$ erg~s$^{\\rm -1}$  and a ratio $L_{r}$/$L_{x}$ about  3.6  for the entire PWN, slightly above the maximum ratio (3.4 for Vela) measured in known PWN. ",
        "introduction": "G18.95$-$1.1 was discovered as a non-thermal extended radio source during a continuum survey of the Galactic plane \\citep{rei84}. The source was first suggested to be energized by a binary system containing a compact object \\citep{fur85}. This classification was based on the unusual radio morphology of G18.95$-$1.1, consisting of various arcs pointing toward a central radio peak. Follow-up radio observations \\citep{ode86} were done soon afterwards to confirm this identification. Based on a detailed analysis of the variation of the spectral index across G18.95$-$ 1.1 and on a re-analysis of the morphological characteristics used by \\citet{fur85}, \\citet{ode86} concluded that G18.95$-$1.1 was more likely a low surface brightness Crab-like supernova remnant (SNR) powered by a hidden central compact object, although the binary interpretation could not be completely ruled-out. He suggested further high-resolution low-frequency observations at 327~MHz to resolve the ``bar-like'' structure at the center of the source. Such observations were carried out by \\citet{pat88}, and their analysis yielded the detection of an almost complete shell surrounding a compact, central object. This result established that G18.95$-$1.1 was indeed a supernova remnant, but the origin of the energy sustaining the central emission from G18.95$-$1.1 remained unknown. \\citet{ode86} suggested that the SNR was powered by a central pulsar. However, \\citet{fur89} argued that  G18.95$-$1.1 could be a binary system consisting of the neutron star created in the supernova explosion and a low mass companion unseen because of the known large optical absorption. To solve the mystery of the energy source of the radio emission and search for a possible central object, a pointed \\rosat\\ PSPC observation (12~ks) of the SNR was carried out by \\citet{fur97}. They found that  the remnant could be best described by a thermal emission model \\citep{ray77} at a temperature of about 0.95~keV and a hydrogen column density of 3.4$\\pm$1.5$\\times$10$^{21}$~atoms~cm$^{-2}$.  A non-thermal model (power-law) yielded a photon index of $\\alpha$~=~$-$8.9 which was therefore dismissed as an ``unrealistic\" model. \\setcounter{footnote}{0} In addition to the X-ray data, \\citet{fur97} obtained a detailed radio map (at an angular resolution of 69$^{\\prime\\prime}$) using the Effelsberg 100-m telescope at 10.55 GHz where Faraday depolarization effects are small. The image extracted from these data is shown in Figure~1\\footnote{The data were kindly provided by Professor F\\\"{u}rst at the Max-Planck-Institute f\\\"ur Radioastronomie.}. It shows  a distinct central bar  surrounded by a faint shell of diffuse emission. About 80\\% of the radio emission is in the central diffuse component and has a total flux  of about 20.4$\\pm$0.2~Jy at 10.55~GHz. This result is consistent with existing radio data at 1.4~GHz and 4.75~GHz for a spectral index ($S_\\nu \\propto \\nu^\\alpha$) $\\alpha$~=~$-$0.14$\\pm$0.03 for the diffuse component and $\\alpha$~=~$-$0.22$\\pm$0.07 for the central bar. A prominent arc on the western part of the remnant (see Figure~1) has a steeper spectral index of $\\alpha$~=~$-$0.36$\\pm$0.04. The linear polarization across the remnant is about 6\\% at 10.55~GHz, a little more than twice that at 4.75~GHz. The polarization intensity is highly concentrated on small scales: when the diffuse contribution is subtracted, the polarization increases to about 40\\% in the central bar and in the arc.  As is often the case in SNR studies, the distance to G18.95$-$1.1 is not well known. There exist 21~cm HI absorption data \\citep{bra84} which were re-analyzed by \\citet{fur89}. Based on a radial velocity of 18~km~s$^{\\rm -1}$ associated with G18.95$-$1.1, they deduced a distance of either 2~kpc (within the Sagittarius arm) or 15~kpc (on the far side of the Galaxy). Because of the large luminosity and velocity expansion implied by this latter distance, they considered the closer distance as the most probable and derived all physical quantities using a 2~kpc measure to the SNR. We will derive all our results normalized to this distance unless otherwise specified and we comment at the end of the paper on the possibility that the remnant is as far away as 15~kpc.  Our motivation, in starting the following investigation of the SNR G18.95$-$1.1 was to solve, once and for all, the mystery of the energy source of the emission from the remnant, namely:  was a compact object powering the emission of the remnant? In view of the \\asca\\ data, we have reanalyzed the existing \\rosat\\ PSPC data, an analysis  we present in \\S 2.  In \\S 3 we describe the data extraction methods, and the spatial and spectral analysis of the \\asca\\ data. In \\S 4 we discuss the implications of our analysis and a summary of our principal conclusions. ",
        "conclusions": "As mentioned at the beginning of this paper, the distance to G18.95$-$1.1 is not well known and we use the lower distance estimate (2~kpc) derived from 21~cm HI absorption data \\citep{bra84,fur89}. At this distance, the supernova remnant defines an X-ray emitting volume of $V \\simeq 1.2\\times10^{\\rm 59}\\,f\\,D_{\\rm 2}^{\\rm 3}\\theta_{\\rm 17}^{\\rm 3}$~cm$^{\\rm -3}$, where $f$ is the volume filling factor of the emitting gas within the SNR, and $\\theta_{\\rm 17}$ the angular radius in units of 17$\\myarcmin$. Because of the large variations in the spectral results using both the standard CIE model and our {\\it pshock} NEI analysis, all the following results were computed using error bars twice as large as that found in the analysis. This allows us to gauge more effectively the different evolution scenarios for G18.95$-$1.1 while acknowledging the difficulty in reaching definitive conclusions with these data alone. We deduce a hydrogen number density $n_{\\rm H}$ = (0.13 -- 0.24)$\\,D_{\\rm 2}^{\\rm -1/2}\\,\\theta_{\\rm 17}^{\\rm -3/2}\\,f^{\\rm -1/2}$~cm$^{\\rm -3}$.  The mass of  the X-ray emitting plasma $M_{X}$, is (17 -- 31) $\\,D_{\\rm 2}^{\\rm 5/2}\\,f^{\\rm 1/2}\\theta_{\\rm 17}^{\\rm 3/2}M_\\odot$. At a distance of 15~kpc this would imply a  mass of  3000 $M_\\odot$, for a Taylor-Sedov value of 0.25 for the filling factor $f$ \\citep{tay50,sed59} . With the preceding numerical values we estimate the supernova explosion energy $E$ to be (0.4 -- 0.9)$\\times 10^{\\rm 50}\\,D_{\\rm 2}^{\\rm 5/2}\\,\\theta_{\\rm 17}^{\\rm 3/2}\\,f^{\\rm -1/2}$~ergs. For a nominal value of $f$ and at 2~kpc, E  is within (although on the low side) standard values for SN explosions. We note that this value is derived under  the hypothesis that electrons and ions are in temperature equilibrium (a simplification that may not be correct in this case) and that a larger explosion energy is allowed if the ions are substantially hotter than the electrons. An estimate of the age of the remnant can be made using our spectral analysis results. We found $t$ between 4400 and 6100 yr, where the range reflects the uncertainties in the fit, but does not include any error on the distance (assumed at 2~kpc) -- the age varies linearly with the distance. The preshock ISM number density is $n_{\\rm 0} \\sim \\rho_0/{\\rm m_{\\rm H}}= (0.06\\pm0.02)\\,D_{\\rm 2}^{-1/2}\\,\\theta_{\\rm 17}^{\\rm -3/2}\\,f^{\\rm -1/2}$~cm$^{\\rm -3}$.  As noted in the introduction, 21~cm HI absorption data \\citep{bra84, fur89} do not exclude a distance of 15~kpc  to G18.95$-$1.1.  If we compute the physical characteristics that this distance would imply for the remnant, we find an explosion energy larger than  $E = 2\\times 10^{\\rm 52}$~ergs, one order of magnitude larger than the value derived from the current SN explosion models.  In addition, because of the large swept-up mass derived in that case, we argue that the results from our X-ray analysis actually rule out this distance.  Based on the radio data alone, it is not surprising that G18.95$-$1.1 was suspected to harbor a compact object.  As described in the spectral analysis summarized above (see \\S 3), the X-ray emission does provide some clues to the real nature of the emission. The X-ray emission from the remnant is clearly dominated by thermal emission and we argue that the remnant is in non-equilibrium ionization, a hypothesis  compatible with the age (between 4400 and 6100 yr) derived in the context of this model.   Our data reveal a clear region of high-energy emission (above 4~keV) located at the tip of the radio bar, although our spectral analysis does not lead to a definitive answer as to the origin of this emission. We  have examined two hypotheses. In the first scenario, this emission results from a high-energy nebula powered by a  hidden pulsar. Our X-ray analysis yields a non-thermal X-ray unabsorbed flux of  about 7$\\times$10$^{\\rm 32}$ erg~s$^{\\rm -1}$ between 0.2 and 4 keV  or about  0.7\\% of the flux from the rest of the remnant. This, in turn, suggests a rotational energy loss \\.{E} of about 7$\\times$10$^{\\rm 35}$ erg~s$^{\\rm -1}$ \\citep{sew88,bec97}. By comparison,  this value is about twice that of the pulsar/SNR in W44 whose pulsar's luminosity is about 1\\% that of the total remnant \\citep{har96}. We  computed the radio luminosity associated with the region from which the high-energy emission is coming and we find about 2.5$\\times$10$^{\\rm 33}$ erg~s$^{\\rm -1}$  at a distance of 2~kpc and using the spectral index from the bar $\\alpha$=-0.22 given in \\citet{fur97}. This value leads to a ratio $L_{r}$/$L_{x}$  about 3.6,  a bit larger than measured values  for other PWNe (that ratio varies from 3.4 for Vela to  3.2$\\times$10$^{\\rm -4}$ for MSH 15-52).  We have presented the results of \\rosat\\ and \\asca\\ X-ray spectral and spatial studies of the SNR G18.95$-$1.1, a middle-aged SNR (about 5000 yr old) most probably in its Taylor-Sedov phase of evolution and at a distance of 2~kpc. The remnant is best described by a nonequilibrium ionization model with solar elemental abundances. There are several X-ray bright spots in the remnant that are generally consistent with thermal emission. We find that we cannot make a definitive identification of the optical counterpart to the unresolved \\rosat\\ PSPC X-ray source. The nature of an unresolved hard \\asca\\ X-ray source located near the tip of the central radio bar remains enigmatic. Future Chandra and XMM-Newton observations may shed more light on this question. \\par"
    },
    "0311/astro-ph0311626_arXiv.txt": {
        "abstract": "The evolution of the temperature distribution function (TDF) of X-ray clusters is known to be a powerful cosmological test of the density parameter of the Universe. Recent XMM observations allows us to measure accurately the $L - T$ relation of high redshift X-ray clusters. In order to investigate cosmological implication of this recent results, we have derived theoretical number counts for different X-ray clusters samples, namely the RDCS, EMSS, SHARC, 160 deg$^2$ and MACS at $z > 0.3$ in different flat models. We show that a standard hierarchical modeling of cluster distribution in a flat low density universe , normalized to the local abundance, overproduces cluster abundance at high redshift ($z>0.5$) by an order of magnitude. We conclude that presently existing data on X-ray clusters at high redshift strongly favor a universe with a high density of matter, insensitively to the details of the modeling. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311140_arXiv.txt": {
        "abstract": "We have analyzed the influence of the stellar populations, from which SN progenitors come from, on the observational outcome, including the metal free Pop. III. We use our models to study the evolution of the progenitor, the subsequent explosion and the light curves.  For Type Ia, the variation of the main sequence mass of the progenitor of the exploding WD, produces an offset in the maximum-decline relation of 0.2 mag. This effect is critical for the use of high redshift Type Ia SNe as cosmological standard candles. In contrast, the metallicity does not change the above relation (at maximum, $\\Delta M_V \\le$0.06 mag).  For Type II, we find a dependence of the light curve properties with both, main sequence mass and metallicity of the progenitor, and we identify a rather homogeneous subclass, {\\it Extreme} II-P, that may be used as a quasi-standard candle. Note that, although not as good as Type Ia for distance determinations, Type II are expected to occur since the first stars were formed. ",
        "introduction": "Due to their high brightness, SNe are among the best candidates to measure distances on cosmological scales; in particular, Type Ia SNe (SNe~Ia) are the favorites as they are brighter by about 1 to 2 magnitudes than other class of SNe and show very homogeneous properties. Recently, SNe~Ia observed at high redshifts, z, (see \\cite{S98} and \\cite{P99}) have provided results that are consistent with a Universe composed mainly of  {\\it dark energy} or, equivalent, with a positive cosmological constant, $\\Omega_\\Lambda $  $\\sim 0.7$. Note that this conclusion relies on the application of an empirical relation between light curve (LC) shape and maximum luminosity which is obtained for the well observed nearby SNe Ia occurring in  galaxies at known distance ( \\cite{P99} and \\cite{R96}).  To understand the nature of this {\\it dark energy}, we need to observe standard-candles at high z. The quantity and quality of the data is increasing due to systematic searches, space telescopes and 8m-class ground telescopes. However, SNe or other standard-candles are calibrated locally, and their application to very high z may be limited by evolutionary effects. In fact, there is observational evidence indicating a correlation of SN properties with those of their host galaxies (\\cite{B96}, \\cite{Ca97}, \\cite{Ha00}, \\cite{IHP00} and \\cite{WHW97}).  SNe~Ia are thought to be the  thermonuclear explosion of CO white-dwarfs \\cite{HF60} triggered by  mass accretion from a companion.  The proposed scenarios are merging of WDs by gravitational radiation  (double degenarate) \\cite{We84} and \\cite{IT84}, and accretion of H and/or He from a non-degenerate companion (single degenerate) \\cite{WI73}. Moreover, they occur in elliptical galaxies indicating that some of their progenitors, at least, have evolutionary time-scales comparable to the Hubble time and, thus,  at high z SNe~Ia should be rare.   The other class of SNe, core collapse SNe, are thought to be the explosion of massive stars (greater than $\\sim 10 M_\\odot$) caused by the collapse of its central parts into a neutron star or a black hole. Their evolutionary time scales are short compared to the age of the universe even at high z (\\cite{T82}, \\cite{MB76} and \\cite{WW86}). Core collapse SNe show a wide range of brightness, up to 6 magnitudes, and properties of their light curves (\\cite{Fi00}, \\cite{Pa94} and \\cite{YB89}) which prevents their use as standard-candles. However, our knowledge of the event is improving and it may be possible to derive the absolute magnitude in a similar way as Type Ia if appropriate empirical correlations can be identified (see Hamuy, this volume and \\cite{HP02}). These objects will occur soon after the initial star formation period and, therefore, can be used to probe the structure of the universe at high z. E.g. at $z \\sim 5$, galaxies are expected to be small and dim and core collapse supernovae may be the brightest objects in the Universe \\cite{MR97}.  It is expected that, going back in time, the stellar population from which SN progenitors come out, would be composed of more low metallicity and rapid evolving, more massive, stars. In this work we focus, mainly, on the exploration of the sensitivity of the light curve characteristics on the underlying progenitor properties: initial mass and metallicity (Z). ",
        "conclusions": ""
    },
    "0311/astro-ph0311376_arXiv.txt": {
        "abstract": "We study numerically the collapse of non-rotating, self-gravitating, magnetized, singular isothermal toroids characterized by sound speed, $a$, and level of magnetic to thermal support, $H_0$. In qualitative agreement with previous treatments by Galli \\& Shu and other workers, we find that the infalling material is deflected by the field lines towards the equatorial plane, creating a high-density, flattened structure -- a pseudodisk.  The pseudodisk contracts dynamically in the radial direction, dragging the field lines threading it into a highly pinched configuration that resembles a split monopole. The oppositely directed field lines across the midplane and the large implied stresses may play a role in how magnetic flux is lost in the actual situation in the presence of finite resistivity or ambipolar diffusion. The infall rate into the central regions is given to 5\\% uncertainty by the formula, $\\dot M = (1+H_0)a^3/G$, where $G$ is the universal gravitational constant, anticipated by semi-analytical studies of the self-similar gravitational collapses of the singular isothermal sphere and isopedically magnetized disks.   The introduction of finite initial rotation results in a complex interplay between pseudodisk and true (Keplerian) disk formation that is examined in a companion paper. ",
        "introduction": "\\subsection{Low-Mass Star-Formation Theory}  Stars form in the dense cores of molecular clouds (Myers 1995, Evans 1999). The nearby cores that form low-mass stars have a typical diameter $\\sim$ 0.1 pc, a number density $\\sim 3\\times 10^4$ cm$^{-3}$, a mass ranging from a fraction of a solar mass to about 10 $M_\\odot$, and an axial ratio for flattening of typically 2:1.  Agents other than isotropic thermal or turbulent pressures must help to support cores against their self-gravity, although there is some debate whether the true shapes are oblate, prolate, or triaxial (Myers 1998, personal communication; see also Jones \\& Basu 2002). Observed cloud rotation rates are generally too small to account for the observed flattening (Goodman et al. 1993), and tidal effects from neighboring cores cannot explain the flattening of inner density contours. This leaves magnetic fields, implicated in many other physical processes of importance in contemporary star formation (see the reviews of Shu et al. 1987, 1999).  In one scenario, ambipolar diffusion leads to a continued contraction of a cloud core with ever growing central concentration (Nakano 1979, Lizano \\& Shu 1989, Basu \\& Mouschovias 1994).  In a finite time given by slippage of neutrals past ions and magnetic fields, the molecular cloud core reaches a configuration where the central density formally tries to reach infinite values. Shu (1995) proposed the name ``gravomagneto catastrophe'' for this mechanism by analogy with the ``gravothermal catastrophe'' (Lynden-Bell \\& Wood 1968, Lynden-Bell \\& Eggleton 1980, Cohn 1980) that overtakes globular-cluster core-evolution because of the diffusion of random velocities by stellar encounters.  Myers (1999) has emphasized that inward fluid velocities of magnitude $\\sim 0.1$ km s$^{-1}$ exist at large radii in some starless cores, and that this feature is not predicted by any of the ambipolar-diffusion calculations cited above (see, however, Li 1998a and Ciolek \\& Basu 2000). Myers \\& Lazarian (1998) suggest instead that pre-existing inward motions prior to true gravitational collapse and star formation are better modeled as resulting from the decay of turbulence.  If the sequence of decaying turbulent parameter $K = 12, 10, 6, 2$ in Figure 2 of Lizano \\& Shu (1989) were to occur over $\\sim$ a few Myr, then inward motions of order half the isothermal sound speed, $0.5 a\\sim$ 0.1 km s$^{-1}$, could indeed be generated at distances of order $0.1$ pc from the center of the cloud core.  Large-scale inward motions and generally shorter time scales may distinguish the ``gravoturbo catastrophe'' resulting from the decay of turbulence in a pre-existing, magnetically supercritical, core with dimensionless mass-to-flux ratio $\\lambda > 1$ (see Li \\& Shu 1966 for the definition of $\\lambda$) from a core that reaches supercriticality via ambipolar diffusion.  As we shall see later, inclusion of initial inward motions of order $0.5 \\, a$ increases the net central mass-accretion rate due to subsequent gravitational collapse by about a factor of 2.  \\subsection{The Pivotal State}  Li \\& Shu (1996) termed the cloud configuration when the central isothermal concentration first becomes formally infinite as the ``pivotal'' state ($t \\equiv 0$).  This state separates the nearly quasi-static phase of core evolution ($t < 0$) that leads to gravomagneto or gravoturbo catastrophe ($t=0$) from the fully dynamic phase of protostellar accretion ($t > 0$). (See the first two stages depicted in Fig. 7 of Shu, Adams, \\& Lizano 1987.)  The numerical simulations indicate that the pivotal states have a number of simplifying properties, which motivated Li \\& Shu to approximate them as scale-free {\\it equilibria} with the power-law radial dependences for the density and magnetic flux function: \\begin{equation} \\rho(r,\\theta)={a^2\\over 2\\pi G r^2}R(\\theta), \\qquad \\Phi(r,\\theta)={4\\pi a^2 r\\over G^{1/2}}\\phi(\\theta). \\label{rhoPhi} \\end{equation} In equation (\\ref{rhoPhi}), we have adopted a spherical polar coordinate system $(r,\\theta,\\varphi)$, and $a$ is the isothermal sound speed of the cloud, while $R(\\theta)$ and $\\phi(\\theta)$ are the dimensionless angular distribution functions for the density and magnetic flux given by force balance along and across field lines.  The resulting differential equations and boundary conditions yield a linear sequence of possible solutions, characterized by a single dimensionless free parameter, $H_0$, which represents the fractional over-density supported by the magnetic field above that supported by thermal pressure. Comparison with the typical degree of elongation of observed cores suggests that the over-density factor $H_0\\approx 0.5 - 1$ (corresponding to $\\lambda \\approx 2$).  In Figure ~\\ref{fig:crete2}, we plot isodensity contours and field lines for the cases $H_0$ = 0.125, 0.25, 0.5, 1.0.  The corresponding values of $\\lambda$ from Table 1 of Li \\& Shu (1966) are $\\lambda$ = 8.38, 4.51, 2.66, 1.79. Below each case are the column density contours that would be seen by an observer in the equatorial plane of the configuration. The presence of very low-density regions at the poles of the density toroid has profound implications for the shape and kinematics of bipolar molecular outflows (Li \\& Shu 1996, see also Torrelles {et al.} 1983, 1994).  \\subsection{Self-Similar Collapse of Singular Isothermal Toroids}  It is well-known that the collapse of the singular isothermal sphere (corresponding to the case $H_0 = 0$ in the previous section) occurs in a {\\it self-similar} manner (Shu 1977). The same holds true for any member in the sequence of isopedically magnetized, singular isothermal toroids (SITs), whether we start them from rest, or with a initial, radially directed, velocity field that is a uniform fraction of the isothermal sound speed $a$. For given $H_0$, we have only two dimensional parameters in the problem, $a$ and $G$.  From $a$ and $G$, we cannot form dimensions for all required physical quantities without using the independent variables of the problem themselves: $r$, $\\theta$, and $t$. The quantities $r$ and $at$ both have the dimensions of length, and they can be combined to form the dimensionless similarity variable \\begin{equation} \\label{xdef} x=r/at , \\end{equation} which can then appear, along with $\\theta$, as the argument for nondimensional mathematical functions.  Thus, the time-dependent axisymmetric solution that describes the $t > 0$ evolution of the unstable equilibrium state (\\ref{rhoPhi}) must take the {\\it self-similar} form: \\begin{equation} \\rho (r, \\theta, t) = {\\alpha (x,\\theta)\\over 4\\pi Gt^2},\\;\\; {\\bf u}(r,\\theta, t) = a {\\bf v}(x,\\theta),\\;\\; {\\bf B}(r,\\theta,t) = {a\\over G^{1/2}t}{\\bf b}(x,\\theta), \\label{rhouB} \\end{equation} where $\\bf u$ and $\\bf B$ are the poloidal fluid velocity and magnetic field, respectively. The dimensionless dependent variables $\\alpha (x,\\theta)$, ${\\bf v}(x,\\theta)$, and ${\\bf b}(x,\\theta)$ are called the {\\it reduced} density, velocity, and magnetic field. Solutions with the form (\\ref{rhouB}) are called ``self-similar'' because, apart from a simple scaling of dependent and independent variables, the solution at one instant in time looks the same as the solution at another instant in time.  For the case $H_0$ = 0 (corresponding to ${\\bf b} = 0$), we may assume that the flow occurs spherically symmetrically.   With no $\\theta$ dependence and with ${\\bf v}$ having only a radial component, the similarity solution has a simple description if the SIS is started from rest (Shu 1977).  At $t=0^+$, consider initiating collapse of the densest innermost portions of a (nonrotating and unmagnetized) singular isothermal sphere to form a condensed central object (a protostar, whose physical dimensions are so small as to allow us here to approximate it as a point mass).  An observer located in the outer portions of the cloud does not realize anything has happened gravitationally, because Newton's theorem for the gravitational field of spherical objects misleads the observer into thinking that all the mass interior to the observer's location equivalently lies at the center anyway.  Only when a sound wave traveling outwards at the speed $a$ reaches the observer, does the observer realize, hydrodynamically, that the bottom has dropped out.  This information initiates a wave of falling whose location at time $t$ has a head described by the locus $r = at$. The similarity variable description for the same locus is $x=1$.  For $x$ larger than 1 (or $r>at$), the density profile retains its unperturbed value $\\alpha = 2/x^2$ (or $\\rho = a^2/2\\pi Gr^2$), and the fluid is at rest, ${\\bf v} =0$ (or ${\\bf u}=0$).  The solution interior to the head of the expansion wave is modified from the initial equilibrium. At $x$ = 1, the solutions for $\\alpha$ and $v$ connect without a jump, but with a discontinuous first derivative (because of the presence of the head of the wave) to the unperturbed state corresponding to the singular isothermal sphere.  Near the origin, the solutions take the asymptotic forms, \\begin{equation} \\label{innerasymp} \\alpha \\rightarrow \\left[ {m_0\\over 2x^3}\\right]^{1/2}, \\qquad v(x) \\rightarrow -\\left[ {2m_0\\over x}\\right]^{1/2} , \\qquad {\\rm as} \\qquad x \\rightarrow 0, \\end{equation} that correspond to gravitational free-fall onto a (reduced) central mass found by numerical integration to be $m_0 = 0.975$. In dimensional units, the origin contains a central mass point (star) that grows linearly with time: \\begin{equation} \\label{centralmass} M = m_0 {a^3t\\over G} . \\end{equation} Equation (\\ref{centralmass}) constitutes the most important property of the solution -- the conclusion that the gravitational collapse of a singular isothermal sphere yields, not a characteristic mass, but a characteristic {\\it mass infall rate} given by \\begin{equation} \\label{Mdotin} \\dot M = m_0\\, {a^3\\over G} . \\end{equation}  Inclusion of the effects of finite levels of the magnetic field (i.e., $H_0 \\neq 0)$ does not change the qualitative situation. The collapse still occurs self-similarly from ``inside-out,'' but the departures from spherical symmetry implies that non-zero motions can be induced {\\it gravitationally} ahead of the front of an expansion wave (see also Terebey, Shu, \\& Cassen 1984 for the case when the cause of departure from spherical symmetry is rotation).  If the gravitational precursor motions are sufficiently compressive, the expansion wave, now propagating outward at the fast MHD speed, may even be preceded by a weak shock wave (Li \\& Shu 1997).  But the main new feature introduced by the presence of the magnetic field is a flattening of the flow in the inner regions to form a ``pseudodisk'' (see the singular perturbational treatment by Galli \\& Shu 1993a,b of this problem). We will concentrate on the formation and structure of this pseudodisk in our numerical calculations, and we shall focus for simplicity most of our attention on equilibrium configurations started from rest.  At the end, however, we will discuss the modifications introduced by starting with uniform inward motions present already at the pivotal instant.  Strict field freezing implies that the mass-to-flux ratio is a conserved quantity: $\\lambda$ = const.  Except for the time-dependent behavior that tapers the surface density distribution near the edge, pseudodisks formed from the gravitational collapse of magnetized SITs therefore will resemble nonrotating, magnetized, singular isothermal disks with values of the dimensionless mass-to-flux ratio $\\lambda$ greater than, but close to, unity. Since such disks, started from rest, collapse to yield a central mass-accretion rate $\\dot M = m_0 a^3/G$ with $m_0 = 1.05(1+H_0)$, Li \\& Shu (1997) speculated that magnetized SITs, for arbitrary initial values of $H_0$ from 0 to $\\infty$, would develop central mass-infall rates well-approximated (to 5\\% accuracy) by the simple formula: \\begin{equation} \\label{magMdotin} \\dot M = (1+H_0){a^3\\over G} , \\end{equation} where $H_0$ is the degree of support by magnetic fields provided in the pivotal configuration. One of the major goals of the present numerical simulation is to justify explicitly the above equation.  For $a \\approx 0.2$ km s$^{-1}$ and $H_0 \\approx 0.5$ (Taurus region), equation (\\ref{magMdotin}) yields $\\dot M \\approx 3\\times 10^{-6}$ $M_\\odot$ yr$^{-1}$, in agreement with the spectral energy distribution (SED) and light-scattering models of Kenyon et al. (1993a, b).  The time $t$ required to build the typical T Tauri star of mass 0.5 $M_\\odot$ is then $1.7\\times 10^5$ yr.  The inference of constant accretion rates, with associated mass-accumulation time scales of $\\sim 10^5$ yr and observed luminosities according to standard models of protostellar evolution (e.g. Stahler et al. 1980, Palla \\& Stahler 1990), is consistent with the survey of Hirano et al. (2002) of 17 protostellar sources in isolated dark clouds, Taurus, and Rho Ophiuchus star-forming regions. They are not consistent with claims for coefficients in the relationship, $\\dot M \\propto a^3/G$,  which are either one to two orders of magnitude larger or highly time-variable, that have been claimed in the literature based on numerical simulations (for low-mass star-forming regions) starting from conditions rather different from singular isothermal toroids (see, e.g., Tomisaka 2002, who started the collapse calculation from an infinitely long cylinder, with the collapse induced by a sinusoidal perturbation in density along its length). Myers \\& Fuller (1993; see also Myers, Ladd, \\& Fuller 1991; Myers 1995) made the intriguing suggestion that total mass-accumulation time scales $\\sim 10^5$ yr may apply across the board, from the formation of the lowest-mass to the highest-mass stars. If this suggestion has merit, and if we continue to use {\\it static}, magnetized SITs to model the pivotal states of such objects, then the final mass of the formed star must scale with the factor $a^3(1+H_0)$, with this factor being larger to form more massive stars.  In other words, massive star-forming regions must be hotter, or more highly magnetized, or more turbulent (if we heuristically include the effect of turbulent support into $H_0$), than low-mass star-forming regions (see the reviews of Shu et al. 1987 and Evans 1999).  Non-isothermal equations of state to mimic the effects of turbulence (where the signal speeds are larger at lower densities) can yield pivotal states with density laws that are shallower than the $\\rho \\propto r^{-2}$ dependences that are considered in this paper (see, e.g., Lizano \\& Shu 1989, McLaughlin and Pudritz 1997, Galli et al. 1999). The gravitational collapse of such non-isothermal configurations will yield mass infall rates $\\dot M$ that increase with time. In contrast, pre-existing inward motions in the central region at the pivotal instant $t=0$ tend to yield $\\dot M$ that decreases with time (Basu 1997, Ciolek \\& K\\\"onigl 1998, Li 1998a).  When added to the support provided quasistatically by thermal pressure, all of these mechanisms {\\it enhance} the mass infall rate above the base value given by the inside-out collapse solution for the initially static, singular isothermal sphere: $\\dot M = 0.975 a^3/G$. None of the considerations imply that $\\dot M$ should {\\it cut off} after some definite time and leave us with a star with a well-defined mass. We return to this difficulty in the last section of this paper (see also the discussion in Shu \\& Terebey 1984, Shu et al. 1999).  In any case, whether or not ideally magnetized SITs represent good pivotal states at $t=0$ for star formation models, the self-similar nature of their $t > 0$ behavior greatly simplify both the collapse calculations and data storage requirements. For the latter, we need to store the reduced variables that are only functions of the similarity coordinates $(x,\\theta)$.  Moreover, these reduced variables depend only on a single dimensionless parameter $H_0$ of the initial state (deducible under the assumption of field freezing from a measurement of $\\lambda$ at any instant during the dynamical collapse). Comparisons with the variables of actual observations at any dimensional time $t$ can then be obtained from the single parameter family of solutions by appropriate scalings of dependent and independent variables by various factors of $t$, $a$, and $G$. ",
        "conclusions": "\\subsection{Protostellar Mass Accretion Rate}  An interesting question is how much our computed mass-accretion rates would have changed if we had allowed for the presence of inward motions in the pivotal state.  We have performed a number of numerical simulations starting with the density and magnetic field configuration of equilibrium toroids (characterized non-dimensionally by the over-density parameter $H_0$), and added an initially uniform, radially inwardly directed, velocity field of $0.5 \\, a$. An example is shown in Figure~\\ref{fig:inwardmotion}.  As we commented upon at the beginning, the results retain their exact self-similarity in spacetime.  The qualitative features of the calculation -- the formation of a pseudodisk, the trapping of the magnetic flux by the central cell to yield a split-monopole geometry, etc. -- also remain unchanged from the corresponding case starting the gravitational collapse from rest. The only substantive modification comes from an increase by about a factor of 2 in the mass-accretion rate onto the central object.  Since unstable cloud cores probably have dimensionless mass-to-flux ratios $\\lambda = $ 1--3, whereas T Tauri stars have $\\lambda >$ 5000, the assumption of field freezing must break down at high collapse densities (Nakano \\& Umebayashi 1980, Li \\& Shu 1997). Whether flux loss in forming protostars ultimately occurs by magnetic reconnection (e.g., Mestel \\& Strittmatter 1967, Galli \\& Shu 1993b) or by C-shocks (Li \\& McKee 1996, Li 1998a, Ciolek \\& K\\\"onigl 1998) remains to be seen.  We elaborate on these possibilities in our companion paper after we include the effects of rotation.  Self-similar collapse calculations carried out with the assumption of strict field freezing yield important ``outer solutions'' for more realistic calculations of the flux-loss problem in the inner regions (e.g., Li 1998b).  An interesting numerical instability inevitably develops to complicate all of our actual calculations. A point mass acting as a split magnetic monopole grows linearly with time at the origin. But, accretion into the point mass occurs primarily through the pseudodisk. Gas in the pdesudodisk feels a magnetic barrier from the split monopole field that opposes rapid accretion.  As the split monopole field grows, mass may pile up in the pseudodisk until it becomes too heavy for the field to support. Until this time, there is lower than average accretion into the origin. Then the excess mass rushes into the point mass, dragging more field to be assimilated by the split monopole. At this time, there is above average accretion into the origin. Thus, the magnetic barrier grows, and mass must pile up in the pseudodisk for a longer time interval, before it can join the point mass, and the oscillations in accretion rate grow in amplitude. The instability has a purely numerical origin because the strictly self-similar, ideal problem contains, of course, no dimensional capability for supporting a temporal period. Nevertheless, the interesting question remains whether the situation with more realistic physics, mimicked by the numerical effects described here, might not sustain relaxation oscillations of a type that are responsible for FU Orionis outbursts.  \\subsection{Limited Role for Magnetic Tension}  There is another important lesson to be learned from the self-similar collapse calculations. Contrary to earlier speculations made in the literature, magnetic tension in the envelope never suffices to suspend the envelope against the gravity of the growing protostar. \\footnote{The concept that magnetic tension can stop infall implicitly assumes that field lines at large distances are held laterally apart by external agencies.  Our example avoids this assumption by allowing the matter and field to have dynamics all the way to infinity. Depending on the nature of the applied outer boundary conditions, caveats may need to be added if the outer parts of a cloud are subcritical (see Shu, Li \\& Allen 2003 for a detailed analysis of this problem).} When the wave of infall reaches successively larger portions of the envelope, those portions move inward to fill the hole and becomes part of the infalling material. In the specific case of the magnetized singular isothermal toroid, without some other influence, the mass infall rate into the central regions continues indefinitely at the constant rate given approximately by equation (\\ref{magMdotin}).  In all cases, the inclusion of the effects of magnetic forces (when $H_0 > 0$) {\\it increases} the rate of central mass accumulation in comparison with the nonmagnetized case ($m_0 = 0.975$), and does not decrease it. The main reason for the increase is that the cloud core is denser on average before it begins to collapse when supported by magnetic fields than when this support is absent.  The slowing down of the infall speed by magnetic forces once collapse begins is largely offset by the increase of the signal speed at which the wave of infall propagates outward (Galli \\& Shu 1993a,b).  \\subsection{Defining Stellar Masses}  Because magnetic tension in the envelope is unable to halt the continued accretion, one of three things must happen if stellar masses are to result from the infall process (Shu et al. 1999).  First, the growing star might run out of matter to accumulate.  For example, the wave of infall might run into a vacuum because the material previously there has already fallen into some other star.  This alternative requires the star formation process to be nearly 100\\% efficient, which we know is not the case except, possibly, in the relatively rare cases of bound cluster formation.  Second, the star might move out of its surrounding cloud core.  Large motions relative to the local pit in the gravitational potential are likely only in very crowded regions of star formation, which we have decided to exempt from the present discussion (but see Shu et al. 1999).  Third, the material otherwise destined to fall onto the star might be removed by some hydrodynamic, magnetohydrodynamic, or photo-evaporative process.  Because of the ubiquity of bipolar outflows in star-forming regions, Shu \\& Terebey (1984) implicated YSO winds as the dominant practical mechanism by which forming stars help to define their own masses, particularly in isolated regions of low-mass star-formation.  In a future study, we will introduce a protostellar wind into the collapse simulation, aiming to determine not only its effects on the mass distribution, but also the kinematics of the wind-swept ambient materials.  \\subsection{Flattened Mass Distribution}  Using a simple thin-shell model, Li \\& Shu (1996) computed the kinematic signature of an X-wind propagating into a static, magnetized SIT and showed that the mass-with-velocity distribution was compatible with observations (for a summary, see Shu et al. 2000).  However, a truly self-consistent model would account for the dynamic infall in the inner regions of the molecular cloud core that produced the central source. The anisotropy in the mass distribution of an initially magnetically supported structure is amplified during the collapse, creating a more extended low-density polar region and a denser equatorial pseudodisk. This effect can be see in Fig.~\\ref{fig:column}, where we plot the iso-column density contours viewed perpendicular to the axis, for the cases of $H_0=0, 0.125, 0.25$ and $0.5$ shown in Fig.~\\ref{fig:allss}.   Note that the column density contours become more elongated closer to the protostar, reflecting the formation of pseudodisk in the inner collapse region. It remains to be seen what are the modifications to be introduced in the calculations of Li \\& Shu (1996) by these effects, particularly in young outflow sources before the swept-up lobes have had a chance to penetrate far beyond the wave of infall.  \\subsection{Effect of Rotation}  YSO outflows arise, many people believe, because of a fundamental interaction between the rapid rotation in a Keplerian disk and some strongly magnetized object (either the central star or the disk itself). By ignoring the effects of rotation in this paper, we have chosen to focus on the pure effects of interstellar magnetic fields on the dynamics of gravitational collapse. However, as discussed above, magnetic fields by themselves do not resolve the fundamental question of why stellar masses result as a product of the gravitational collapse of molecular clouds.  Fortunately, rotation is observed ubiquitously in star-forming dense cores (e.g., Goodman et al. 1993). Although typically not fast enough to contribute significantly to the core support, it is presumably responsible for the formation of Keplerian disks around young stars.  In many studies, disk sizes are computed under the assumption that the total angular momentum observed in the cores is entirely carried into the collapsed object, although there may be some viscous redistribution within the Keplerian accretion disk. The following paper in this series, which finds relatively strong magnetic braking to occur even during the collapse process, casts doubt on the validity on this basic assumption when molecular cloud cores are realistically rotating {\\it and} magnetized."
    },
    "0311/astro-ph0311006_arXiv.txt": {
        "abstract": "We use a high-resolution cosmological simulation to predict the distribution of \\ion{H}{1} Ly$\\alpha$ emission from the low-redshift ($z \\la 0.5$) intergalactic medium (IGM).  Our simulation can be used to reliably compute the emission from optically thin regions of the IGM but not that of self-shielded gas.  We therefore consider several models that bracket the expected emission from self-shielded regions. Most galaxies are surrounded by extended ($\\ga 10^2 \\kpc$) ``coronae'' of optically thin gas with Ly$\\alpha$ surface brightness close to the expected background.  Most of these regions contain smaller cores of dense, cool gas.  Unless self-shielded gas is able to cool to $T < 10^{4.1} \\kel$, these cores are much brighter than the background. The Ly$\\alpha$ coronae represent ``cooling flows'' of IGM gas accreting onto galaxies.  We also estimate the number of Ly$\\alpha$ photons produced through the reprocessing of stellar ionizing radiation in the interstellar medium of galaxies; while this mechanism is responsible for the brightest Ly$\\alpha$ emission, it occurs on small physical scales and can be separated using high-resolution observations.  In all cases, we find that Ly$\\alpha$ emitters are numerous (with a space density $\\sim 0.1\\, h^3 \\Mpc^{-3}$) and closely trace the filamentary structure of the IGM, providing a new way to map gas inside the cosmic web. ",
        "introduction": "\\label{intro}  One of the most striking images to emerge from cosmological simulations is the ``cosmic web''.  Matter collapses into a web of moderately overdense sheets and filaments, with galaxies and galaxy clusters forming through continued collapse at their intersections. This picture explains both the qualitative distribution of galaxies in redshift surveys (e.g., \\citealt{delapparent86}) and many of the characteristics of the Ly$\\alpha$ forest (see e.g. \\citealt{rauch98}). However, mapping the cosmic web remains a difficult venture. Ly$\\alpha$ forest spectra probe only one-dimensional lines of sight, while galaxy redshift surveys find only the small fraction of baryons embedded inside galaxies, offering an indirect picture of the gaseous filaments.  In this \\emph{Letter}, we use a cosmological simulation to show that surveys of hydrogen Ly$\\alpha$ ($\\lambda$1216 \\AA) emission can provide a powerful method to map the cosmic web.  Existing analytic estimates suggest that the surface brightness of dense portions of the IGM can be substantial \\citep{hogan87,gould96}.  Such regions generally lie within filaments but outside of galaxies, so Ly$\\alpha$ emission offers a more direct map of the gas distribution in filaments than do galaxy surveys.  In fact, Ly$\\alpha$ emission from galaxies has been used to detect a filament at high redshift \\citep{moller01}; here we focus on intergalactic emission from low redshifts, where the distances and sky background are relatively small and where Ly$\\alpha$ maps of the cosmic web can be compared to existing galaxy surveys.  We also show that Ly$\\alpha$-emitting gas constitutes an intrinsically interesting phase: gas that is cooling around collapsed halos \\citep{haiman01-lya,fardal01}. The distribution of Ly$\\alpha$ emission therefore provides a window onto the growth of bound objects. ",
        "conclusions": "\\label{discussion}  We study Ly$\\alpha$ emission from gas in the IGM using a high-resolution cosmological simulation.  We find that galaxies are typically surrounded by ``coronae'' of Ly$\\alpha$ emission with bright, central cores (of diameter $\\la 40\\, h^{-1} \\kpc$) surrounded by much larger ($\\ga 10^2\\, h^{-1} \\kpc$) regions with surface brightness near the background.  While the number of Ly$\\alpha$ photons due to star formation in the host galaxies can exceed that emitted by the surrounding gas, the spatial extent of the star formation is typically much smaller.  High angular-resolution observations (probing physical scales $\\la 15\\, h^{-1}\\kpc$) should be able to separate the two components.  Detection of the bright cores is feasible with existing technology, although the wide-field ultraviolet spectrographs most useful for these studies have not yet been built. The \\emph{Galaxy Evolution Explorer} (\\emph{GALEX}), with a large field of view but relatively low spectral resolution ($\\sim 10$ \\AA), may detect the brightest cores.  Furthermore, the strong correlation between Ly$\\alpha$ emission and galaxies may enable a statistical detection of the signal in deep exposures.  A project underway to construct a balloon-borne wide-field UV spectrograph will have a limiting sensitivity $\\Phi \\sim 500 \\photflux$ in a single night's observation (D. Schiminovich, private communication).  We emphasize that the emission from self-shielded gas is uncertain, because the simulation does not include all of the relevant physics. We have considered a range of models bracketing these uncertainties (see \\S \\ref{selfshield}): in our fiducial model, dense, cool gas produces Ly$\\alpha$ photons only through collisional processes, but we also consider cases in which it has zero emissivity or in which we allow photoionization.  The choice has little effect on the large-scale, low surface brightness emission, which comes primarily from optically thin gas, but it does strongly affect the luminous cores., where most of the gas is shielded from the ionizing background.  We have also neglected dust, which can efficiently destroy Ly$\\alpha$ photons.  Because the emitting region is often considerably larger than the associated galaxy, much of the emitting gas may be relatively pristine.  In the simulation, we find that the majority of particles with large Ly$\\alpha$ emissivity contain no metals (or dust), though note that the simulation does not include mixing between simulation particles.  At higher redshifts, several large-scale Ly$\\alpha$-emitting ``blobs'' have already been observed (e.g., \\citealt{steidel00}), limiting the amount of dust in these environments.  The structure and relatively large spatial size of Ly$\\alpha$ emitters, together with their locations, suggest that they are ``cooling flows'' of gas onto halos.  Detailed, high resolution Ly$\\alpha$ observations therefore offer a means to study the accretion and cooling of gas onto galaxies.  Galactic winds will also affect the distribution of Ly$\\alpha$ emission.  However, we find that the winds in our simulations have only minor effects on the statistics of Ly$\\alpha$ emitters.  On larger scales, the Ly$\\alpha$ emitters are distributed along sheets and filaments.  Because these objects are numerous (with a space density $\\sim 0.1 h^3$ Mpc$^{-3}$), Ly$\\alpha$ line observations offer a new and powerful way to map the cosmic web.  Figure \\ref{fig:lyamap} shows that such observations do not require high resolution; wide-field observations could efficiently locate (in three dimensions) a large number of Ly$\\alpha$ emitters. In contrast to galaxy surveys, this approach reveals the location of baryons \\emph{outside} of galaxies.  Because Ly$\\alpha$ emission traces the cool gas accreting onto galaxies, it also selects a different population of dark matter halos, including many with few (or even no) associated stars.  Follow-up observations at high resolution would allow us to compare Ly$\\alpha$ emitters to galaxy surveys and learn how galaxies accrete gas at the present day."
    },
    "0311/astro-ph0311230_arXiv.txt": {
        "abstract": "We present a set of mock redshift catalogs, constructed from N-body simulations, designed to mimic the DEEP2 survey.  Galaxies with a range of luminosities are placed within virialized halos in the simulation using a variant of the halo model.  The halo model parameters are chosen to reproduce local clustering and abundance data and assumed to be independent of redshift.  This allows us to predict the luminosity function, two-point correlation function, luminosity- and scale-dependent bias, redshift space distortions etc. of our galaxies at higher redshifts. We show that the low order clustering properties are consistent with  preliminary DEEP2 data. The catalogs can be used to evaluate the selection effects of the survey and to test new algorithms and statistics that are to be used in the analysis of DEEP2 data. ",
        "introduction": "As our theoretical understanding of structure formation has advanced and the questions we ask of modern redshift surveys have sharpened, the use of realistic mock catalogs has become widespread.  Almost all modern galaxy redshift surveys have made extensive use of mock catalogs both in their design and analysis.  Mock catalogs which reproduce the low order clustering statistics and galaxy redshift distributions can serve to test algorithms for biases and quantify the impact of numerous observational selection effects. They are an indispensable tool in interpreting a wide range of observational data. A suitably constructed mock catalog can also allow us to connect the observations directly to modern models of structure formation, which are based exclusively on hierarchical build-up of dark matter halos in an inflationary cold dark matter universe.  In this paper we describe in some detail the construction of mock catalogs currently being used by the DEEP2 redshift survey (Davis et al.~\\cite{DEEP2}). These mock catalogs are created by populating large, high-resolution dark matter simulations of the $\\Lambda$CDM cosmology with galaxies in such a way as to reproduce the DEEP2 observations as closely as possible. The mock catalogs have undergone a continual evolution over the last few years, and we intend to continue to refine these mock catalogs as further observations become available.  In this paper we describe the current status of the programme and the mocks which are being used in analysis of current DEEP2 data. We describe some of the properties of the galaxies in the simulations which will be useful in interpreting future DEEP2 observations and finish by describing directions for future work.  The outline of the paper is as follows.  In \\S\\ref{sec:mock} we place our work in the context of previous attempts to make mock catalogs.  The N-body simulations and halo model are described in \\S\\ref{sec:sim}. We describe the intrinsic properties of our mock galaxies, derived from single time outputs of the simulation in \\S\\ref{sec:galbox} and how we turn these `galaxy boxes' into lightcone outputs in \\S\\ref{sec:cone}. After presenting some basic results in \\S\\ref{sec:basic} we describe the limitations of the current generation of mock catalogs and some directions for future development in \\S\\ref{sec:future}. We conclude in \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have presented mock catalogs based on high resolution N-body simulations which are designed to match the low order clustering statistics of the DEEP2 redshift survey. The mock catalogs are based upon the halo model for clustering and situate galaxies in their correct cosmological context. The simulations used are sufficiently large to give many independent realizations for deep pencil-beam surveys and have sufficiently dense time sampling to model possible evolution information embedded in the survey.  The mock catalogs provide positions, velocities and luminosities for galaxies in multiple fields, for a number of cosmologies within the $\\Lambda$CDM class.  The galaxies have a non-trivial scale and luminosity dependent bias, which is stochastic, but trace the large scale structure showing voids, filaments, sheets, walls and clusters. In each case the $z$ distribution, luminosity function, correlation function and projected correlation function match well the preliminary results from the DEEP2 survey.  In the process of building these mock catalogs we have noticed several interesting results. \\begin{itemize} \\item[(i)] The halo model and its extensions (such as the CLF) provide a powerful tool for building mock catalogs for galaxy redshift surveys. Galaxies are put into the simulations in a physically motivated way. This has  two advantages.  First, the HOD and CLF can be constrained from observations so that it almost guarantees the mocks will mimic reality. Second, tuning of the model parameters teaches us about galaxy evolution. The HOD and its extensions provide an ideal meeting ground between theories of cosmology and galaxy formation. \\item[(ii)] To get the correct $z$ distribution of galaxies is not trivial, indicating that the observed $z$ distribution contains a lot of information. Within our formalism, the $z$ distribution involves a $K$-correction, $L_*$ evolution, halo mass function evolution and CLF evolution. The first two factors are measurable, the third one depends on cosmology, and the last one on the details of galaxy formation. If the $K$-correction and $L_*(z)$ can be well measured, the remaining degeneracy between cosmology and CLF evolution can be broken by other measurements, e.g., the fluctuations in the $z$ histogram in a narrow beam survey, the LF, counts of groups and their richness, the correlation function etc. \\item[(iii)] Our simple assumption of a HOD which does not evolve from $z=0$ to $z>1$ provides a reasonable fit to the early DEEP2 data.  This is a surprising result.  Additionally there may be hints that a non-evolving HOD will fail to fit improved data at higher redshift, which could provide valuable information on the formation and evolution of galaxies. \\end{itemize}  We anticipate that these mock catalogs will be crucial in understanding the data coming from the DEEP2 survey and will help us to understand the formation and evolution of galaxies."
    },
    "0311/astro-ph0311224_arXiv.txt": {
        "abstract": "Many globular cluster systems have a distinct bimodal metallicity distribution function (MDF) which has strikingly similar features in many large galaxies of all types. By using the Milky Way cluster system as a typical example, we show that bimodal MDFs can be very well matched with a double ``accreting-box'' chemical enrichment model in which both the halo (metal-poor) and bulge (metal-rich) clusters form during an early phase of gas inflow simultaneously with star formation.  However, differences in effective yield between the two phases are not enough by themselves to reproduce the observed MDF shape:  gas infall is required for both phases, and either the initial gas or the infalling gas must have very different metallicities in the two separate phases. ",
        "introduction": "The old globular clusters found in all large galaxies have presented long and persistent challenges to our understanding of the early formation processes of galaxy halos and bulges.  A feature of globular cluster systems (GCSs) which has drawn considerable attention is their metallicity distribution (MDF), which consistently takes on a roughly bimodal form in many large elliptical and large spiral galaxies \\citep[e.g.][among many others]{whi95,geb99,zep95,gei96,for97, nei99,kun01,lar01,lfb01,bea03,har03}. Most clusters fall either in the ``metal-poor'' mode centered at [Fe/H] $\\simeq -1.5$ or ``metal-rich'' mode centered near [Fe/H] $\\simeq -0.5$, with each mode having an intrinsic dispersion $\\sigma$[Fe/H] $\\simeq 0.3$ dex. Furthermore, the mean [Fe/H] values and dispersions of each mode are surprisingly similar from one galaxy to the next despite the vast range of galaxy types \\citep[e.g.][]{for97,h01,for01}.\\footnote{To date, most of the MDF data for the many galaxies in the studies cited above are from broadband photometric indices such as $(V-I)$, which are intrinsically not very sensitive to metallicity. In addition, the two standard ``modes'' have Gaussian dispersions large enough to allow a significant degree of intrinsic overlap, so that if they are blurred out further by low-precision data, the validity of a bimodal fit becomes debatable. For examples of MDFs where the two modes may not have clearly separated peaks, see, e.g., NGC 524, 1399, or 4473 in \\citet{lar01}.  It is, however, important to note that when data are taken with higher-precision photometry or spectroscopy, bimodal distributions tend to emerge more clearly.  That is, the galaxies with the best available data (such as M87, NGC 4472, NGC 1404, NGC 4594, NGC 5128, or the Milky Way itself; see the references cited above) are very well matched by a bimodal MDF in which the two modes display clear separation.  It is these ``best case'' examples which have driven the bimodality paradigm in the literature and which our present discussion is aimed at.}  The Milky Way GCS is no exception to this general trend.  Because it is a completely uncontaminated sample of clusters with well established metallicities, it provides what is still perhaps the best classic example of a bimodal MDF (see Figure 1).  The shape of the MDF and its close connection with the cluster kinematics was first clearly established by \\citet{zin85} and has held up through many subsequent analyses \\citep[see, e.g.][for a review]{h01}. The MDF, in its usual plotted form as number per unit [Fe/H] (left panel of Figure 1), is often approximated by a pair of Gaussian functions (cf. the references cited above). These numerical fits have no physical significance (and have not usually been claimed to), and give little insight into the evolutionary history of the system except for the strong hint that a distinct two-phase formation history is required.  In this brief paper, we take one step further into interpreting these bimodal MDFs using some simple and physically plausible assumptions. Standard first-order chemical evolution models, such as the one-zone closed-box or ``simple model'' \\citep{pag75,bin98} work reasonably well at matching the broad, unimodal MDFs of the halo field-star populations in the Milky Way \\citep[e.g.][]{pag75,har76,rya91,pra03} and in M31 \\citep{dur01} with appropriate choices of the effective nucleosynthetic yield $y_{\\rm eff}$. \\footnote{For a normal IMF and a closed-box system, the nucleosynthetic yield is typically $y_0 \\simeq 0.008$. In ``leaky-box'' models, the value of $y_{\\rm eff}$ can be much smaller and is strongly driven by the fraction of gas ejected from the zone by supernovae and stellar winds at each round of star formation; see Hartwick (1976).} Extensions to the simple model to incorporate some smoothly varying gas inflow (the ``accreting-box'' model) have recently been used to model the unimodal, metal-rich MDF of the halo stars in NGC 5128 \\citep{har02}.  An early, brief phase of primordial gas inflow has even been applied to explain the relatively small number of the {\\sl very} lowest-metallicity stars ([Fe/H] $< -3$) in the Milky Way halo \\citep{pra03}. Accreting-box models can be thought of as approximate versions of more extensive hierarchical-merging models of galaxy formation, in which a large population of small pregalactic gas clouds merges successively into bigger and bigger pieces while star formation within the clouds goes on simultaneously \\citep[see, e.g.][for an application of the GALFORM code to the case of NGC 5128]{bea03}.  By contrast, the distinctly bimodal MDFs of globular cluster systems do not easily fit into the continuous sequence of starforming events in normal hierarchical-merging or accreting-box models. Some new feature of the history needs to be invoked, either to boost the formation of metal-poor clusters at early times, or inhibit the slightly later formation of intermediate-metallicity clusters.  \\citet{har02} and \\cite{bea02} suggest that the metal-poor clusters may have formed preferentially early in the first rounds of star formation, but that their formation was then truncated while field-star formation continued onward.  In this picture, the metal-rich clusters would have formed a bit later, during the major starbursts accompanying the merging of the last, biggest gas clouds.  \\citet{san02}, for example, suggests that the metal-poor cluster formation epoch was truncated at the reionization era.  An alternate approach \\citep{cot98,cot02} is to assume that the metal-poor clusters all formed separately in dwarf galaxies and were dissipationlessly accreted later by the giants.  In this case the bimodality of the combined MDF is a stochastic accident which results only for a restricted mass range and mass spectrum of the accreted satellites. \\citet{ash92} proposed instead that the metal-poor clusters formed {\\sl in situ} while the metal-richer clusters formed later from gas brought in by major mergers (though this model was primarily intended to explain E-galaxy formation rather than big spirals).  Both of these routes can produce bimodal MDFs, though many other interpretive problems exist \\citep[see, e.g.,][for extensive discussions]{for97,h01,hh01}. ",
        "conclusions": "The primary results for our four starting assumptions are summarized in Table 1.  In the last line of the table, mean uncertainties (corresponding to one standard deviation) are listed.  Calculations of extensive model grids showed that the initial masses $M_0$ and $M_{\\rm add}(0)$ were were closely coupled and very weakly constrained.  (In some cases it was possible to start with a virtually ``empty'' box and build the system entirely from infalling material.)  The most interesting constraints were on $y_{\\rm eff}$, $Z_{\\rm add}$, and to a lesser extent, the infall decay time $\\tau$.  However, even these are partially correlated:  for example, an increase in $Z_{\\rm add}$ can be compensated by a shorter time $\\tau$, leading to almost the same final MDF. In Table 1 we also list the typical (one-sigma) uncertainties of the fitted parameters.  All four initial assumptions are capable of leading to good fits to the observed MDF; these are displayed in Figures 2(a-d).  Nevertheless, a few clear trends emerge from inspection of the results. For the ``halo'' stage, {\\sl both the initial gas and infalling gas need to be nearly unenriched, primordial material}, and the steep rise in $dn/dZ$ at very low metallicity ($Z < 0.02 Z_{\\odot}$) is the characteristic mark of early heavy infall.  The effective yield for the halo in all the models is always near $y_{\\rm eff} \\sim 0.0015$, five times smaller than the basic nucleosynthetic yield $y_0$. This $y-$level suggests that $\\sim$80\\% or more of the protocluster gas in the halo was expelled before it could convert into stars, but the presence of infall allows $y_{\\rm eff}$ to be larger than in models without inflow \\citep{har76,bin98}.  For the ``bulge'' formation stage, $y_{\\rm eff}$ is typically three times larger than for the halo phase.  If, as we expect, the mean difference is due mainly to different ratios of ejected gas, then the bulge region, sitting much deeper in the Galaxy's potential well, would have held onto more than half of its gas rather than only 20\\%.  The exponential decay time for gas infall is consistently near $20 \\delta t$ in phase 1, and 25 to 40 $\\delta t$ in phase 2. We cannot convert this to an actual time interval without additional information on merging and star formation rates which only more advanced models can provide. But if, following the GC ages mentioned above, we assume that each stage takes $\\sim 2$ Gy over $\\sim 100 \\delta t$ steps, then in {\\sl very} rough terms the main period of heaviest gas infall should last about half a Gigayear in each stage.  We emphasize here again that the assumption of {\\sl no} infalling gas ($\\tau \\rightarrow 0$, corresponding to the closed-box model) does not reproduce the data adequately; such models produce too many clusters in either phase at the low-metallicity end.  Said differently, a two-phase {\\sl closed-}box model, where the only free parameters are the yields $y_{\\rm eff}$, does not give the model enough freedom to fit the data.  An early period of infall is both necessary and plausible.  \\noindent {\\sl Model (a):} In this model even the metal-rich bulge clusters form starting from initially unenriched gas.  However, the {\\sl infalling} gas for stage 2 must compensate by being substantially pre-enriched (to about one-quarter Solar) to produce the correct secondary peak in the MDF. For comparison, the hierarchical merging models \\citep[see][]{col00,bea03} show that the ``quiescent'' metal-poor mode starts with unenriched material but can continue upward to near-Solar abundance as long as it is not interrupted.  The halo clusters might have belonged only to the first part of this phase.  \\noindent {\\sl Model (b):} In this model -- the poorest fit of the four -- the {\\sl infalling} gas is always primordial, so $Z_0$ for stage 2 must start quite high to directly produce the secondary MDF peak (leaving the rather unphysical spike at $Z = 0.34 Z_{\\odot}$). Here, almost none of the bulge clusters fall in the metal-poor regime; there is almost no overlap between the two parts of the MDF.  \\noindent {\\sl Model (c):} This model starts with perhaps the most plausible assumption for both stages:  the infalling gas has the same abundance as the initial $Z_0$.  The job of producing the secondary MDF peak at the right place is then shared more or less equally between the initial $Z-$level and the subsequent enrichment cycles from star formation.  The final maximum enrichment level from stage 1 is fairly close to the initial $Z_0$ for stage 2, thus making it easier for the halo clusters and field stars to form over similar time periods, unlike model (a) above.  \\noindent {\\sl Model (d):} This model, which allows an unconstrained solution for both parametric abundances $Z_0$ and $Z_{\\rm add}$, ends up almost identical with model (a) and similar to (c) as well."
    },
    "0311/gr-qc0311055_arXiv.txt": {
        "abstract": "Contrary to what is often stated, a fundamental spacetime discreteness need not contradict Lorentz invariance.  A causal set's discreteness is in fact locally Lorentz invariant, and we recall the reasons why.  For illustration, we introduce a phenomenological model of massive particles propagating in a Minkowski spacetime which arises from an underlying causal set.  The particles undergo a Lorentz invariant diffusion in phase space, and we speculate on whether this could have any bearing on the origin of high energy cosmic rays. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311012_arXiv.txt": {
        "abstract": "The signal produced in neutrino observatories by the pair-annihilation neutrinos emitted from a 20~$M_{\\odot}$ pre-supernova star at the silicon burning phase is estimated. The spectrum of the neutrinos with an average energy $\\sim$2~MeV is calculated with the use of the Monte Carlo method. A few relevant reactions for neutrinos and anti-neutrinos in modern detectors are considered. The most promising results are from $\\bar{\\nu}_e + p \\longrightarrow n + e^{+}$ reaction. During the Si-burning phase we expect 1.27 neutrons/day/kiloton of water to be produced by neutrinos from a star located at a distance of 1~kpc. Small admixture of effective neutron-absorbers as e.g. NaCl or GdCl$_{3}$ makes these neutrons easily visible because of Cherenkov light produced by electrons which were hit by $\\sim$8~MeV photon cascade emitted by Cl or Gd nuclei. The estimated rate of neutron production for SNO and Super-Kamiokande is, respectively, 2.2 and 41 events per day for a star at 1~kpc. For future detectors UNO and Hyper-Kamiokande  we expect 5.6 and 6.9 events per day even for a star 10~kpc away. This would make it possible to foresee a massive star death a few days before its core collapse. Importance of such a detection for theoretical astrophysics is discussed. ",
        "introduction": "The detection of solar neutrinos by the R.~Davis's chlorine detector and supernova 1987A neutrinos by the Kamiokande water Cherenkov observatory are milestones of modern science, {\\it Nobel Prize} honored in 2002. Currently built and operating neutrino detectors are larger, much more efficient and are able to detect all neutrino species in a wide range of energy. There are already proposed new experiments with next-generation megaton-scale detectors, like Hyper-Kamiokande \\cite{hk} and UNO \\cite{uno}.  \\begin{table}[b] \\begin{tabular}{ c c c } \\hline & The Sun & Neutrino-cooled phase of 20~M$_\\odot$ star \\\\ \\hline Lifetime & $10^{10} $ years & 300 years \\\\ Photon luminosity & L$_\\odot$ & 10$^5$L$_\\odot$ \\\\ Maximum $\\nu$ luminosity & 0.02 L$_\\odot$ & 10$^{12}$ L$_\\odot$ \\\\ Average neutrino energy & 0.3 MeV & 0.7-2 MeV \\\\ Total $\\nu$ energy released & 10$^{49}$ erg & 10$^{51}$ erg \\\\ \\hline \\end{tabular} \\caption{Differences between the Sun (main sequence star) and the neutrino-cooled 20~M$_\\odot$ star (\\cite{woosley}, Table~1). Approximate values (orders of magnitude) are given. The total energy carried by neutrinos for a 20~M$_\\odot$ neutrino-cooled star  corresponds only to late stages of nuclear burning (without supernova neutrinos).} \\label{table1} \\end{table}   Solar neutrinos were detected with sub-kiloton detectors, because of relatively small distance to the Sun and continuous emission. Supernova neutrinos carry enormous energy of 10$^{53}$ ergs released in gravitational collapse of the stellar core. This makes their detection possible from quite large distances, but such events are rare. We consider here the feasibility  of detection of another giant astrophysical sources of neutrinos, namely massive (M$>$8M$_\\odot$) stars at late stages of their evolution. Such stars, during C, Ne, O and Si burning phases, emit neutrinos copiously and are referred to as neutrino-cooled stars (\\cite{arnett}, Sect. 10)  or pre-supernova stars \\cite{langanke}. The structure of such evolutionary advanced stars is different from main sequence stars like the Sun. The most important differences relevant to the neutrino detection possibility are summarized in Table~\\ref{table1}.   The solar neutrino luminosity is $7.8 \\cdot 10^{31}$ erg/s \\cite{bahcall}. The maximum neutrino luminosity during Si burning exceeds $3 \\cdot 10^{45}$ erg/s  (\\cite{woosley},~\\cite{zimmerman}) -- a value larger by a factor of $3.85 \\cdot 10^{13}$. The neutrino energy flux at Earth from such a star and from the Sun are equal for a star $\\sqrt{3.85 \\cdot 10^{13}} = 6.2 \\cdot 10^6 $AU, i.e. 30~pc away. This indicates that neutrino-cooled stars could be detected at astronomical distances. Actually, because of a different spectrum (cf.~Fig.~\\ref{solar}) and the presence of anti-neutrinos (cf. Table~\\ref{f_factor}), we are able to detect them from kiloparsec distances.  Massive (neutrino-cooled) stars are believed to end as {\\it core-collapse} supernovae. SN~1987A confirmed this theory. This close connection may be used to estimate chance of neutrino-cooled stage observation. If some of massive stars die other way, rate of neutrino-cooled events may be higher than of SNe, but it is reasonable to use estimates made for supernovae. At least three different methods exist there: historical records \\& remnants counts \\cite{vandenbergh}, extragalactic observations \\cite{capellaro_turrato} and population synthesis (i.e. simulated stellar evolution of the Galaxy) methods  \\cite{piran}. Only first method gives {\\it local rate} directly: time between events less than 5kpc away is estimated to be 175 years \\cite{strom}. Two other methods give rate for entire Galaxy. Extragalactic counts give 40\\ldots 200 years \\cite{capellaro_turrato} of average time interval between supernova explosions. Population Synthesis give 10 years \\cite{piran}. To compare {\\it local rates} we may use Galaxy model of e.g Bahcall\\&Soneira \\cite{galaxy_model}. For the solar neighbourhood we found 0.5\\%, 10\\% and 50\\% of disc stars closer than 1kpc, 5kpc and 10kpc respectively. Assuming that the Sun is not in privileged position close to e.g. giant star forming region, we get time between supernovae closer than 5kpc from 100 (Population Synthesis) to 400\\ldots2000 (extragalactic counts) years. Apparent disagreement between these two methods is usually explained by a large number of unobserved events. Clouds of interstellar gas and dust obscuring optical detection are obviously transparent for neutrinos, so we may conclude that neutrino-cooled events closer than 5kpc may be observed even more frequent than one per century. Close candidates for such a detection apparently do exist, with Betelgeuse ($\\beta$~{\\it Ori}) being the most popular example \\cite{apod}. In Sect.~5 we combine rates discussed above with our results to estimate a chance of successful detection.   As an example we consider a 20~M$_\\odot$ star model with properties reported in Table~1 of Ref. \\cite{woosley}, with {\\it explicite} given total neutrino luminosities of various stages of nuclear burning. As our scope here is to study the feasibility of neutrino detection, not to find state-of-art details of emission\\footnote{ Detailed picture of late stages of nuclear burning is still a subject of active research as there are unanswered questions about e.g. hydrodynamical nature of nuclear burning \\cite{2dburning}. Neutrinos itself are still enigmatic objects and e.g. the role of spin-flip interactions and the influence of right-handed neutrinos~\\cite{rightnu} is investigated. }, we assume that the entire neutrino flux is resulting from $\\mathrm{e}^+\\mathrm{e}^-$ annihilation. The values from Table~\\ref{table2} of the temperature, density and the electron chemical potential at the stellar center are assumed to approximate conditions in the stellar core. For simplicity we treat neutrino propagation without oscillation, however the effect of neutrino oscillations may change significantly the number of observed events in a given experiment \\cite{rubbia}. To calculate the relevant neutrino cross sections the standard Weinberg-Salam theory \\cite{gws} is used.  In Sect.~2 we calculate the neutrino spectrum with the use of Monte Carlo method, in Sect.~3 we briefly present neutrino-induced reactions and outline of selected detectors.  Sect.~4 contains the expected neutrino signal from a star located at a distance of 1~kpc. In Sect.~5  astrophysical implications of the pre-supernova neutrino detection are discussed.    \\begin{table*} \\begin{tabular}{ c c c c c c c } \\hline Burning  & $T_c$ & $\\rho_c$ & $\\mu_e $ & $L_{\\nu}$ & Duration & Total energy \\\\ Phase &[MeV]& [g/cc] & [MeV] & [erg/s] &  $\\tau$ & emitted [erg] \\\\ \\hline C  & 0.07 &$2.7 \\cdot 10^5$& 0.0& $7.4\\cdot 10^{39}$ & 300 yrs & $ 7 \\cdot 10^{49} $ \\\\ Ne & 0.146 &$4.0 \\cdot 10^6$& 0.20& $1.2 \\cdot 10^{43}$ & 140 days & $1.4 \\cdot 10^{50}$\\\\ O   & 0.181 &$6.0 \\cdot 10^6$& 0.24& $7.4 \\cdot 10^{43}$ & 180 days& $1.2 \\cdot 10^{51} $\\\\ Si &   0.319  &$4.9 \\cdot 10^7$&  0.84 & $3.1 \\cdot 10^{45}$ & 2 days & $5.4 \\cdot 10^{50}$\\\\ \\hline \\end{tabular} \\caption{Properties of a 20~M$_\\odot$ star according to Ref.~\\cite{woosley}. We have calculated the total energy radiated in neutrinos as a product $\\tau L_{\\nu}$. Actually, the neutrino emission is expected to be a function of time.} \\label{table2} \\end{table*} ",
        "conclusions": ""
    },
    "0311/astro-ph0311538_arXiv.txt": {
        "abstract": "{  We have measured the spectrum of the Cosmic X-ray Background (CXB) in the 2-8 keV range with the high throughput EPIC/MOS instrument onboard XMM-Newton. A large sample of high galactic latitude observations was used, covering a total solid angle of 5.5 square degrees. Our study is based on a very careful characterization and subtraction of the instrumental background, which is crucial for a robust measurement of the faintest diffuse source of the X-ray sky. The CXB spectrum is consistent with a power law having a photon index $\\Gamma=1.41\\pm0.06$ and a normalization of 2.46$\\pm$0.09 photons cm$^{-2}$ s$^{-1}$ sr$^{-1}$ keV$^{-1}$ at 3 keV ($\\sim$11.6 photons cm$^{-2}$ s$^{-1}$ sr$^{-1}$ keV$^{-1}$ at 1 keV), corresponding to a 2-10 keV flux of (2.24$\\pm$0.16)$\\times10^{-11}$ erg cm$^{-2}$ s$^{-1}$ deg$^{-1}$ (90\\% confidence level, including the absolute flux calibration uncertainty). Our results are in excellent agreement with two of the most recent CXB measurements, performed with BeppoSAX  LECS/MECS data (Vecchi et al. 1999) and with an independent analysis of XMM-Newton EPIC/MOS data (Lumb et al. 2002), providing a very strong constrain to the absolute sky surface brightness in this energy range, so far affected by a $\\sim$40\\% uncertainty. Our measurement immediately implies that the fraction of CXB resolved by the recent deep X-ray observations in the 2-10 keV band is of $80\\pm7$\\% (1$\\sigma$), suggesting the existence of a new population of faint sources, largely undetected within the current sensitivity limits of the deepest X-ray surveys. ",
        "introduction": "\\label{intro} The discovery of a diffuse Cosmic X-ray Background (CXB) radiation dates back to the birth of X-ray astronomy: it was a serendipitous result of the same rocket experiment which detected the first extra-solar X-ray source, Scorpius X-1 (Giacconi et al. 1962). The problem of the nature of the CXB immediately became one of the most debated topics in astrophysics.  In the late 70's the first broad-band measurement of the CXB spectrum was obtained with the HEAO-1 satellite. In the 3-50 keV range the data were found to follow a thermal bremsstrahlung distribution (kT$\\sim$40 keV), well approximated below 15 keV by a simple power law with photon index $\\Gamma\\sim1.4$ (Marshall et al. 1980). In the same years several pieces of evidence (see e.g. the review by Fabian \\& Barcons 1992) led to the understanding that the bulk of the CXB above the energy of $\\sim$1 keV is extragalactic in origin. COBE FIRAS observations (Mather et al. 1990) ruled out models based on diffuse emission from an hot intergalactic medium, strongly supporting the hypothesis that the CXB is made up from the integrated emission of faint discrete sources, with a dominant contribution from Active Galactic Nuclei (AGNs) (Setti \\& Woltjer 1989).  This picture has been confirmed by the results from imaging X-ray observatories. Starting from the early observations with the Einstein satellite (Giacconi et al. 1979), coming to the recent deep surveys with Chandra (Moretti et al. 2003 and references therein) and XMM-Newton (Hasinger at al. 2001), a higher and higher fraction of the CXB, up to 80-90\\% in the overall 0.5-10 keV range, has been resolved into discrete sources, mainly obscured and unobscured AGNs. Indeed, the final solution of the origin of the CXB seems today to be quite close. The wealth of information coming from the deep pencil-beam X-ray surveys (Chandra Deep Field North, Brandt et al. 2001; Chandra Deep Field South, Giacconi et al. 2002), medium-deep wide angle X-ray surveys (HELLAS2XMM, Baldi et al. 2002) and their multiwavelength follow-up campaigns, combined with synthesis models (e.g. Gilli et al. 2001), are defining the big picture, explaining which are the sources of the CXB and constraining their cosmological properties.  One of the main uncertainties involved in the problem is the CXB intensity itself. Since the HEAO-1 experiment, several measurements of the CXB spectrum have been obtained at energies below 10 keV. While the results on the spectral shape confirmed a power law with $\\Gamma\\sim1.4$, the normalization of the CXB remained highly uncertain as a consequence of large discrepancies (well beyond the statistical errors) among the different determinations. A difference as large as $\\sim40$\\% is found from the highest measured value (Vecchi et al. 1999 using SAX data) to the lowest one (the original HEAO-1 experiment, Marshall et al. 1980).  Barcons et al. (2000) pointed out that two different causes are required to explain the large scatter among the different determinations of the CXB intensity. First, cosmic variance: spatial variations of the CXB intensity are expected as a consequence of its discrete nature. This problem may be reduced by measurements covering large solid angles. Second, systematic errors, including cross-calibration differences, must have some role. In any case, the measurements published after the analysis of Barcons et al. (2000), namely Lumb et al. (2002) with XMM-Newton EPIC and Kushino et al. (2002) with ASCA/GIS, while differing by $\\sim$15\\% only, either because of the small covered solid angle (Lumb et al. 2002), or because of large uncertainties in the stray light assessment (Kushino et al. 2002), do not allow to constrain the value of the CXB normalization to a much narrower range.  An uncertain value of the intensity represents a very severe limitation to any understanding of the ultimate nature of the CXB. Even a basic information such as the resolved fraction of the CXB cannot be firmly evaluated, leaving largely unsolved the problem of ``what is left'' beyond the detection limits of the deepest X-ray surveys: is there any fainter population still waiting to be resolved by even deeper observations? is there room for truly diffuse emission?  In this paper we present a new measurement of the CXB spectrum in the 2-8 keV range performed with the high throughput EPIC intrument onboard XMM-Newton. Our study is based on a large sample of high galactic latitude pointings for a total solid angle of $\\sim$5.5 square degrees, reducing the effects of cosmic variance. Our analysis includes a very robust characterization of the instrumental background properties (including the issue of low-energy particle contamination - so far neglected in XMM-Newton data analysis), and particular care was devoted to the study of the possible sources of systematics.  The paper is organized as follows: in Sect.~\\ref{nxbgeneral} we give an overview of the EPIC cameras instrumental background components and of the different approaches required for their correct subtraction. In Sect.~\\ref{recipe} we describe in detail the algorithm which we used to extract the CXB spectrum starting from the raw EPIC data. In Sect.~\\ref{results} we present our results; a detailed analysis of the possible sources of uncertainty is given in Sect.~\\ref{syst}. In Sect.~\\ref{discuss} our findings are discussed and compared to previous works. Our study of the EPIC instrumental background is quite technical and complex, but it may be useful for the study of extended sources such as  Clusters of Galaxies; for these reasons it is reported in detail in the Appendices~\\ref{qnxb} and ~\\ref{aqnxb}. In the figures throughout the paper error bars represent 1$\\sigma$ uncertainties, unless otherwise specified. ",
        "conclusions": "\\label{discuss} Our measurement of the intensity of the Cosmic X-ray Background has been plotted in Fig.~\\ref{cxbmeas} (adapted from Moretti et al. 2003), together with previous determinations. To be conservative, in order to identify the most probable range of sky surface brightness constrained by our data, we have used an estimate of the accuracy of the cross-calibration in absolute flux among different recent instruments (see Sect.~\\ref{absolute}) as a measure of the absolute flux calibration uncertainty of EPIC.  Our results are in full agreement with the BeppoSAX LECS/MECS analysis of Vecchi et al. (1999) and with the study of Lumb et al. (2002) using XMM-Newton EPIC/MOS data. The ASCA GIS measure of Kushino et al. (2002) is marginally consistent with ours, while the HEAO-1 value (Marshall et al. 1980) is significantly lower.  As discussed by Barcons et al. (2000), part of the dispersion of the different measurement can be ascribed to cosmic variance. The above authors estimated that differences up to $\\sim$10\\% in the CXB intensity are expected when the solid angle is of order 1 square degree. They concluded therefore that systematic errors and cross-calibration differences among different instruments must be present.  The HEAO-1 measurement, performed on$\\sim10^4$ square degrees, is the most robust as for solid angle coverage. However, all of the subsequent measurements yielded systematically higher values, casting at least some doubts on the absolute flux calibration of the HEAO-1 instruments. We note that a recent reanalysis of the HEAO-1 data (Gruber et al. 1999), as pointed out by Gilli (2003), showed differences of order 10\\% among different detectors in the overlapping energy ranges.  Two new measurements have been published since the analysis of Barcons et al. (2000). Lumb et al. (2002) used a combination of 8 XMM-Newton EPIC/MOS pointings covering $\\sim1.2$ square degrees. Their approach for the instrumental background subtraction (based on a model of the {\\em out FOV} spectrum) is markedly different from ours and therefore their determination is to be considered largely independent from the measurement presented in this work. The ASCA/GIS study of Kushino et al. (2002) was performed over a large solid angle ($\\sim$50 square degrees) to minimize the effects of cosmic variance. However, the large stray light component ($\\sim$40\\% of the collected flux) affecting ASCA represents a severe limit to the absolute flux measurement of the CXB, which they estimated to be uncertain at the $\\sim$10\\% level.  In our analysis, we used a compilation of sky pointings covering $\\sim5.5$ square degrees. The effects of cosmic variance (roughly scaling as $\\Omega^{- 1/2}$) on our measure of the CXB intensity should be rather small, being of the same order of the overall quoted uncertainty ($\\sim4$\\%). Our measurement, performed with the well-calibrated EPIC/MOS instrument, relies on a very robust characterization and subtraction of the instrumental background; possible sources of systematics were carefully studied and their impact on the determination of the CXB spectrum was evaluated. In conclusion, these considerations, coupled to the excellent agreement of our findings with the results of two of the most recent investigations, lead us to believe that the measurements of the CXB intensity have finally converged to a well constrained value, significantly higher than the former result from HEAO-1 data, assumed more than 20 years ago as a reference.  It is very interesting to compare our new measurement of the CXB intensity with the source number counts derived from the recent deep X-ray observations. Moretti et al. (2003) have built a Log$N$/Log$S$ function using a very large source compilation, including the results from six different surveys, both pencil-beam and wide-field, performed with ROSAT, Chandra and XMM-Newton. We assume here their results for the integrated flux of the 2-10 keV source counts down to the sensitivity limit of the 1 Ms Chandra Deep Fields ($\\sim1.7\\times10^{-16}$ erg cm$^{-2}$ s$^{-1}$). Our measurement of the CXB intensity, F$_{CXB}$=(2.24$\\pm$0.10)$\\times10^{-11}$ erg cm$^{-2}$ s$^{-1}$ deg$^{-1}$ ($1\\sigma$ error, including the absolute flux uncertainty), implies that 80$^{+7}_{-6}$\\% of the cosmic X-ray background has been resolved into discrete sources in the 2-10 keV band. Even extrapolating the Log$N$/Log$S$ down to fluxes of $\\sim10^{-17}$ erg cm$^{-2}$ s$^{-1}$ (a factor of 10 below the detection limit of the deep surveys) the integrated flux would rise only to $\\sim$84\\% of the CXB (due to the slow growth of the faint branch of the Log$N$/Log$S$), not consistent with its total value. A new class of faint sources (possibly heavily absorbed AGNs, or star-forming galaxies, see Moretti et al. 2003 and references therein) could emerge at fluxes of a few  $10^{-16}$ erg cm$^{-2}$ s$^{-1}$, steepening the Log$N$/Log$S$ and accounting for the remaining part of unresolved CXB. Some contribution to the CXB could also be due to truly diffuse emission.  The analysis of the deepest 2 Ms Chandra Deep Field North and of the XMM deep pointing to the Lockman Hole and of their associated multiwavelength follow-up campaign, will shed light on the issue. On the basis of a well constrained value for the CXB intensity, the new observations will possibly allow for an ultimate understanding of the CXB nature.    \\begin{figure} \\centering  \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{fig_lc.ps}} \\caption{ Light curve of a typical sky field observation. Events are extracted from the {\\em in FOV} region in the energy range 0.5-12 keV; time bin is 30 s. The second part of the observation is affected by intense soft proton flares, the peak count rate is more than 200 times higher than the quiescent one. The selected threshold, in units of sigma of the quiescent count rate distribution (see text and Fig.~\\ref{histo}), is marked with a dashed line.} \\label{lc} \\end{figure}   \\begin{figure} \\centering  \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{fig_histo_log.ps}} \\caption{  Histogram of the count rate distribution for a typical sky field observation. The corresponding light curve is shown in Fig.~\\ref{lc}. The peak corresponds to the quiescent count rate Poisson distribution; points falling to the right correspond to the soft proton flares. The GTI threshold, in units of sigma of the quiescent count rate distribution, is marked with a dashed line. } \\label{histo} \\end{figure}     \\begin{figure} \\centering  \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{bestfit.ps}} \\caption{  The cosmic X-ray background spectrum in the 2-8 keV range is displayed, folded with the instrumental response. MOS1 data are represented in red, MOS2 in black. The best fit model is overplotted. The lower panel shows the residuals in units of sigma.} \\label{bestfit} \\end{figure}     \\begin{figure} \\centering  \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{fig_slope_vs_gti.ps}} \\caption{ The best fit photon index of the CXB is plotted as a function of the selected GTI threshold (in units of sigma of the quiescent count rate distribution - see Sect.~\\ref{gti}). Errors (at 90\\% confidence level for a single interesting parameter) include the uncertainty on the renormalization of the quiescent NXB spectrum. The adopted 3.3$\\sigma$ threshold is marked by the vertical solid line. No significant variations in the CXB photon index are obtained as a result of different choices of the GTI threshold within the range $\\sim2.9-3.8\\sigma$, marked by the dotted vertical lines.} \\label{gtisystslope} \\end{figure}      \\begin{figure} \\centering  \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{fig_norm_vs_gti.ps}} \\caption{ Same as Fig.~\\ref{gtisystslope}, the best fit CXB normalization is shown. Units are photons cm$^{-2}$ s$^{-1}$ sr$^{-1}$ keV$^{-1}$ at 3 keV. The contributions from out-of-field scattered light have not been subtracted. Errors (at 90\\% confidence level for a single interesting parameter) include the uncertainty on the renormalization of the quiescent NXB spectrum. The adopted GTI threshold is marked by the vertical dashed line. It is evident that no significant variations are obtained as a result of small changes of the selected GTI threshold (within the range $\\sim2.9-3.8\\sigma$, marked by the dotted vertical lines). See text for further details.} \\label{gtisystnorm} \\end{figure}   \\begin{figure} \\centering  \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{fig_slope_vs_SPresid_m1.ps}} \\caption{ Best fit CXB photon index as a function of the maximum allowed ratio for the 8-12 keV surface brightness $\\Sigma_{IN}/\\Sigma_{OUT}$, a probe of the presence of residual soft proton NXB. The MOS1 case is shown (MOS2 is very similar). Errors (at 90\\% confidence level for a single interesting parameter) include the uncertainty on the renormalization of the quiescent NXB spectrum. No significant changes are obtained. The case of the CXB normalization is markedly different (see Fig.~\\ref{norm_vs_ratio}).} \\label{slope_vs_ratio} \\end{figure}    \\begin{figure} \\centering  \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{fig_norm_vs_SPresid_m1.ps}} \\caption{ Same as Fig.~\\ref{slope_vs_ratio}, the best fit CXB normalization (in units of photons cm$^{-2}$ s$^{-1}$ sr$^{-1}$ keV$^{-1}$ at 3 keV) is plotted here. The stray light correction have not been applied. Errors (at 90\\% confidence level for a single interesting parameter) include the uncertainty on the renormalization of the quiescent NXB spectrum. When dataset having large values of the ratio $\\Sigma_{IN}/\\Sigma_{OUT}$ are included, the normalization is systematically higher owing to the presence of a significant residual SP NXB component (see text). The adopted maximum allowed ratio ($R_{MAX}=1.3$) is marked by the vertical line. No significant changes are obtained within small changes of $R_{MAX}$. See Appendix ~\\ref{aqnxb} for further details.} \\label{norm_vs_ratio} \\end{figure}  \\begin{figure} \\centering  \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{fig_slope_vs_theta_m2.ps}} \\caption{ Best fit CXB photon index as a function of the off-axis angle. Each point represents an independent measure obtained from the analysis of a selected portion of the detector plane (see text). Errors (at 90\\% confidence level for a single interesting parameter) include the uncertainty on the renormalization of the quiescent NXB spectrum. The MOS1 case is shown (MOS2 is very similar). The constant trend is a result of the correct vignetting function calibration.} \\label{slope_theta} \\end{figure}      \\begin{figure} \\centering  \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{fig_norm_vs_theta_m2.ps}} \\caption{ Same as Fig.~\\ref{slope_theta}, the case of the CXB normalization is presented here. Units are photons cm$^{-2}$ s$^{-1}$ sr$^{-1}$ at 3 keV. The corrections for the light coming from out-of-field angles have not been applied. Errors (at 90\\% confidence level for a single interesting parameter) include the uncertainty on the renormalization of the quiescent NXB spectrum. The variations are always smaller than the uncertainty, confirming the overall correctness of the vignetting curve calibration.} \\label{norm_theta} \\end{figure}  \\begin{figure} \\centering  \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{cxb_measurements_bw.ps}} \\caption{CXB intensity measurements. The flux in the 2-10 keV band is represented as a function of the epoch of the experiment. The plot is an update of Fig.3 of Moretti et al. (2003) including the results of the present work.  From left to right the CXB values are from Marshall et al. (1980) with HEAO-1 data; McCammon et al. (1983) with a rocket measurement; Gendreau et al. (1995) with ASCA SIS data; Miyaji et al. (1998) with ASCA GIS; Ueda et al. (1999) with ASCA GIS and SIS; Vecchi et al. (1999) with BeppoSAX LECS/MECS; Lumb et al. (2002) with XMM-Newton EPIC/MOS; Kushino et al. (2002) with ASCA GIS. Finally, the hollow star mark our own CXB measurement with the EPIC MOS cameras onboard XMM-Newton. All the uncertainties are at 1$\\sigma$ level. The horizontal dotted lines mark the range of CXB intensity constrained by our measurement after including an extra 3.2\\% uncertainty (1$\\sigma$ confidence level, rescaled from the 5\\% uncertainty quoted in the text, which corresponds to 90\\% confidence level) on the absolute flux calibration of the EPIC cameras (it is indeed a measure of the cross-calibration accuracy among different instruments). Such range represents the most robust estimate of the true absolute sky surface brightness in the 2-8 keV band. } \\label{cxbmeas} \\end{figure}     \\appendix"
    },
    "0311/astro-ph0311362_arXiv.txt": {
        "abstract": "We present a spatially resolved spectral analysis of full \\asca\\/ observations of the remnant of the supernova of 1006 AD.  This remnant shows both nonthermal X-ray emission from bright limbs, generally interpreted as synchrotron emission from the loss-steepened tail of the nonthermal electron population also responsible for radio emission, and thermal emission from elsewhere in the remnant. In earlier work, we showed that the spatially integrated spectrum was well described by a theoretical synchrotron model in which shock acceleration of electrons was limited by escape, in combination with thermal models indicating high levels of iron from ejecta.  Here we use new spatially resolved subsets of the earlier theoretical nonthermal models for the analysis.  We find that emission from the bright limbs remains well described by those models, and refine the values for the characteristic break frequency.  We show that differences between the northeast and southwest nonthermal limbs are small, too small to account easily for the presence of the northeast limb, but not the southwest, in TeV gamma-rays.  Comparison of spectra of the nonthermal limbs and other regions confirms that simple cylindrically symmetric nonthermal models cannot describe the emission, and we put limits on nonthermal contributions to emission from the center and the northwest and southeast limbs. We can rule out solar-abundance models in all regions, finding evidence for elevated abundances. However, more sophisticated models will be required to accurately characterize these abundances. ",
        "introduction": "A rapidly growing number of galactic remnants exhibit nonthermal (and non-pulsar-related) emission in the X-ray band. Nonthermal emission can dominate over much or all of the remnant, as is the case with SN~1006 \\citep{Dyer2000}, G266.2--1.2 [RX J0852.0--4622] \\citep{Slane6}, G347.3--0.5 \\citep{Slane1999, Uchiyama2003}, and AX J1843.8-0352, \\citep{Ueno2003}. Alternatively an SNR can have signatures of nonthermal emission along with X-ray spectral lines, for example RCW 86 \\citep{Rho2003}, Cassiopeia A \\citep{Vink2003a}, and Tycho, \\citep{Hwang2002}. Finally some SNRs (Cassiopeia A, Kepler, Tycho, SN~1006 and RCW~86) reveal nonthermal emission in the hard-X-ray sensitivity of the \\rxte\\/ \\pca\\/ (up to 60 keV), suggesting nonthermal emission may be common among most young SNRs \\citep{Allen1999}.  In most of these instances, a strong case can be made that the nonthermal emission is synchrotron radiation.  The total number of SNRs with suspected nonthermal emission is now in the double digits and we expect, as new X-ray instruments improve in spectral and spatial resolution, many more SNRs will be found with varying amounts of X-ray synchrotron emission, previously undetected among thermal continuum and line emission. In the past this emission was missed because few analyses included models for nonthermal emission. Even of those that do, many do not take advantage of this unique opportunity to extract information about particle acceleration from the emission of ultra-relativistic electrons. We know that the spectrum extrapolated from radio frequencies must roll off before X-ray energies to avoid exceeding X-ray flux measurements (Reynolds \\& Keohane 1999; Hendrick \\& Reynolds 2001).  Therefore observations of X-ray synchrotron emission lie in the regime where the particle spectrum is changing rapidly. Shock-acceleration models that can reproduce this drop-off may be able to provide information about the remnant age, radiative losses of electrons or the spectrum of magnetohydrodynamic (MHD) waves near the shock.   In Dyer et al.~(2000; hereafter Paper I) we demonstrated that the total-flux spectrum from SN~1006 was well fit by a combination of new synchrotron models and plane-shock thermal models, and that the fit was an improvement over previous fits both due to the robustness of the synchrotron models (which use radio measurements to constrain the nonthermal emission) and due to the new accuracy of the thermal models, which, with the assistance of the synchrotron model, detected for the first time half a solar mass of iron. However, these models were fit to spectra summed over the entire remnant and it is well known that the remnant spectra can vary significantly across the face of a remnant.  In particular \\citet{Koyama6} settled a controversy about the X-ray spectrum of SN~1006 by exhibiting a {\\em difference} in spectra between the limb and center.  One of the most compelling reasons for spatially resolved studies of SN~1006 is the nonaxisymmetric TeV $\\gamma$-ray emission from SN~1006 \\citep{Tanimori1998}. In Paper I we interpreted this emission as inverse-Compton upscattering of cosmic-microwave-background photons and derived a preshock magnetic field of $\\sim$3 \\mg, implying a mean field in the remnant of about 9 \\mg. Since then, full analysis of the $\\gamma$-ray spectrum yielded a magnetic field of 4~\\mg\\/ \\citep{Tanimori2001}, in reasonable agreement with our previous results. Inverse-Compton is not the only possible emission mechanism to explain the $\\gamma$-rays -- recent work \\citep{Berezhko2002} attributes the TeV emission from SN~1006 to decay of $\\pi^0$ mesons produced by inelastic collisions between cosmic-ray protons and thermal gas. However, neither interpretation explains why $\\gamma$-rays are seen only in the northeast (NE) -- only spatially resolved studies could hope to answer this question.        The clear symmetry of SN~1006 about a northwest-southeast axis suggests that generalizations of the cylindrically symmetric model used in Paper I (\\sresc\\/) is a reasonable starting point for such studies. While the symmetric synchrotron model is an oversimplification, the results from the whole remnant fits (Paper I) were very encouraging, suggesting that we extend the model to find its limits.  We hope to answer a set of related questions:  Is the spatial distribution of emission from SN~1006 consistent with an axially symmetric model? How exact is the bilateral symmetry, and can any hint be found in the X-rays for the gross asymmetry shown by the TeV $\\gamma$-ray emission?  In the regions where nonthermal emission dominates can thermal emission be characterized by simple models?   There has been much recent work on SN~1006. \\citet{Vink2000} analyzed integrated spectra from \\beppo\\/ of SN~1006 with a new version of SPEX designed to simultaneously fit regions which overlap in the point spread function. They found an adequate two temperature fit to the thermal emission and supersolar abundances.  New high resolution \\chandra\\/ observations of the NW and NE provide the best spatial resolution images to date and several groups have been exploring the fine details. \\citet{Long2003} analyzed the first two \\chandra\\/ pointings and found no evidence for the halo predicted by \\citet{Reynolds1996}. The radio and X-ray features were found to be perfectly correlated in the NE although above 0.8 keV X-ray limb brightening is more pronounced. Clumps of thermal X-ray material were found interior to the shock in the NW and NE with super-solar abundances in the NW. \\citet{Bamba2003b} carried out spatial-spectral fits of cross sections of the \\chandra\\/ observation of the sharp fine filaments in the NE. Using an energy cut to separate thermal and nonthermal emission they found the nonthermal emission to be narrower than the thermal emission. From thermal fits they found super-solar abundances, including iron, and a profile consistent with Sedov dynamics. More recent \\chandra\\/ observations by Hughes use short exposures to image the entire SNR.    Here we attempt to extract the maximum useful information from all \\asca\\/ observations of SN~1006.  We present spatial subsets of the \\sresc\\/ synchrotron model of Paper I, and apply them to describe the nonthermal emission.  We believe that this study represents the fullest use of the substantial amount of observing time obtained with the \\asca\\/ satellite on SN~1006, and that further advances will require both better data and better thermal and nonthermal models. ",
        "conclusions": "\\label{conclusions}  \\subsection{Nonthermal Emission}    1.  From comparing the limb-to-center ratio in X-ray and radio wavelengths (Figure \\ref{slices}), we find that the morphology of SN~1006 is not solely due to limb brightening of a cylindrically symmetric source about a line parallel to the bright limbs.  This has the inescapable consequence that the observer is now required to have a particular vantage point with respect to the SNR.      2. \\sresc\\/ provides a useful framework to describe synchrotron emission in SNR blast waves. Spatially resolved subsets of the \\sresc\\/  model, \\cp, \\limb, and \\ctr, differ surprisingly little from the spatially integrated model, with the greatest differences occurring when the upstream magnetic field is in the plane of the sky.  The differences are unlikely to be detectable with the current generation of X-ray telescopes.  3.  The \\limb\\/  model provides a good description of the limbs of SN~1006, but the \\cp\\/  and \\ctr\\/ region models overpredict nonthermal emission in other regions of the remnant if \\rolloff\\/ is assumed to be constant everywhere.  4. We do find differences in the nonthermal spectrum between the NE and SW, but they are not sufficient to explain why only the NE is detected in TeV $\\gamma$-rays, assuming the $\\gamma$-rays are produced by inverse-Compton upscattering of the cosmic microwave background photons.  The usefulness of \\limb, \\cp, and \\ctr\\/ models for other SNRs is not obvious. Most SNRs have more complex morphology than SN~1006, and few present regions clearly appropriate for a \\ctr\\/ or \\cp\\/ model. While the \\limb\\/ model may be appropriate for filamentary structure at the edge of other SNRs, its curvature differs only very slightly from the \\sresc\\/ model. We recall that the model \\sresc, and its derivatives have several assumptions that make them unsuitable for SNRs of unknown type -- the models presume a Sedov-phase remnant encountering a uniform upstream medium with a constant magnetic field. \\srcut, as discussed in \\citet{Reynolds1999}, which describes synchrotron radiation with a minimum of assumptions, may be a better choice for SNRs about whose environment less is known.  \\subsection{Thermal Emission}   1. We can rule out solar abundances in the thermal models in all regions of the SNR, including regions dominated by nonthermal emission.  2. Simple thermal models are incapable of describing the \\asca\\/ data adequately.  We note that in much smaller regions of SN~1006, the analysis of \\chandra\\/ data encountered similar problems \\citep{Long2003}.  We believe that this is mainly caused by deficient Fe L-shell atomic data, where many important processes (such as inner-shell ionization of the Na-like Fe ion) are currently not included in NEI spectral codes.  It is also likely that the available spectral models are not adequate for modeling chemically-inhomogeneous SN ejecta. A full analysis of the combined thermal and nonthermal emission from SN~1006 will require \\xmm\\/ and \\chandra\\/ quality data with substantial improvements in atomic data and modeling.   3. To the extent that our thermal results can be trusted, we see elevated abundances which suggest the thermal emission contains substantial contributions from the reverse shock. Here the models agree with the silicon equivalent width, and with the results of other X-ray analyses of SN~1006."
    },
    "0311/astro-ph0311154_arXiv.txt": {
        "abstract": "{ We present mid-infrared spectro-imaging ($5 - 16\\,\\mu$m) observations of the infrared luminous interacting system Arp 299 (= Mrk\\,171 = IC\\,694+NGC\\,3690) obtained with the ISOCAM instrument aboard ISO. Our observations show that nearly 40\\,\\% of the total emission at 7 and $15\\,\\mu$m is diffuse, originating from the interacting disks of the galaxies.  Moreover, they indicate the presence of large amounts of hot dust in the main infrared sources of the system and large extinctions toward the nuclei.  While the observed spectra have an overall similar shape, mainly composed of Unidentified Infrared Bands (UIB) in the short wavelength domain, a strong continuum at $\\sim 13\\,\\mu$m and a deep silicate absorption band at $10\\,\\mu$m, their differences reveal the varying physical conditions of each component.  For each source, the spectral energy distribution (SED) can be reproduced by a linear combination of a UIB ``canonical'' spectral template and a hot dust continuum due to a $230-300$\\,K black body, after independently applying an extinction correction to both of them.  We find that the UIB extinction does not vary much throughout the system ($A_{\\rm V} \\ltapp 5$\\,mag) suggesting that most UIBs originate from less enshrouded regions.  IC\\,694 appears to dominate the infrared emission of the system and our observations support the interpretation of a deeply embedded nuclear starburst located behind an absorption of about 40 magnitudes. The central region of NGC\\,3690 displays a hard radiation field characterized by a [Ne{\\sc iii}]/[Ne{\\sc ii}] ratio $\\geq 1.8$. It also hosts a strong continuum from 5 to $16\\,\\mu$m which can be explained as thermal emission from a deeply embedded ($A_{\\rm V}\\sim60$\\,mag) compact source, consistent with the mid-infrared signature of an active galactic nucleus (AGN), and in agreement with recent X-ray findings. ",
        "introduction": "At a distance of 41\\,Mpc ($v_{\\rm hel} = 3080\\,{\\rm km\\,s}^{-1}$, $H_{0} = 75\\,{\\rm km\\,s}^{-1}\\,{\\rm Mpc}^{-1}$), Arp\\,299 (Mrk\\,171 = IC\\,694 + NGC\\,3690) is one of the nearest interacting galaxies. Because of its proximity, its spatial extent (8\\,kpc), and its high infrared (IR) luminosity\\footnote{We use the standard definition of $L_{\\rm IR}(8-1000\\,\\mu{\\rm m}) = 5.62\\times10^5\\,D^2_{\\rm Mpc}\\,(13.48f_{12} + 5.16f_{25} + 2.58f_{60} + f_{100})\\,L_{\\sun}$ \\citep[see][]{Sanders96}.} ($L_{\\rm IR} = 5.16\\times 10^{11}\\,L_\\odot$), it is one of the prime candidates for exploring the effects of triggered star formation activity as the tidal forces during an interaction lead to instabilities in the galactic disks and rapidly funnel large quantities of gas into the dynamical centers of the galaxies \\citep{Mihos96,Sanders96}.  Similar to the prototypical interacting pair of galaxies NGC\\,4038/39, this system is dynamically young, as shown by the prominent 180\\,kpc ($13^\\prime$) long tidal tail \\citep{Hibbard99}.  \\begin{figure*}[!ht] \\IfFileExists{pgallais_3926_f1.eps}{ \\resizebox{\\hsize}{!}{\\includegraphics{pgallais_3926_f1.eps}}}{\\framebox(13.5,6.7){}} \\caption[bidon]{{\\bf a)} The HST/NICMOS 2.2\\,$\\mu$m image of Arp\\,299 from \\citet{Alonso00} with the overlay of the ISOCAM 7\\,$\\mu$m emission. The different components of the system are indicated.  The 9 contour levels are set with logarithmic spacing between 1 and 33\\,mJy\\,arcsec$^{-2}$.  The image scale is displayed with a bar of 2\\,kpc (10{\\arcsec}) at the top.  {\\bf b)} Same as in {\\bf a)} but using the ISOCAM 15\\,$\\mu$m image as an overlay, having set the contour limits to 6 and 60\\,mJy\\,arcsec$^{-2}$.  Adapted from \\citet{Charmandaris02}.} \\label{fig:nicmos_CAM} \\end{figure*}  Arp\\,299 has been studied extensively at all wavelengths over the past thirty years \\citep[see the original work by][]{Weedman72}.  Most notably though, the mid-infrared and radio work of \\citet{Gehrz83}, the observations in the optical by \\citet{Augarde85}, in the near-infrared by \\citet{Gallais91} and \\citet{Wynn91}, along with the information on its molecular gas kinematics by \\citet{Aalto97} and \\citet{Casoli99} have provided ample evidence for the extreme characteristics of this system.  More recently, the high spatial resolution images obtained with the Hubble Space Telescope (HST) in the near-IR \\citep{Alonso00}, as well as in the mid-IR using Keck \\citep{Soifer01} have revealed numerous point-like sources that were undetected in the optical, and firmed up our understanding of the complex star-forming history of the galaxy \\citep[see][, and references therein]{Alonso00}.  Following the notation suggested by \\citet{Gehrz83}, the system is usually described by the galaxy IC\\,694 (source\\,A) to the east, the nuclear region of NGC\\,3690 (source\\,B) to the south-west and two regions of emission (sources\\,C and C$^\\prime$) to the north-west.  Source\\,B in NGC\\,3690 has been further resolved into several components, among which B1 and B2 clearly dominate the global emission in the near- and mid-IR (see Fig.\\,\\ref{fig:nicmos_CAM} for details).  Interestingly it was also recently shown by \\citet{Charmandaris02} that despite its inconspicuous appearance in the optical, IC\\,694 contributes nearly half of the infrared ($8-1000\\,\\mu$m) luminosity of the system.  One of the recurrent questions though in all studies of Arp\\,299 is whether this galaxy is a pure starburst as was proposed by \\citet{Weedman72} or whether some luminosity could be contributed by Active Galactic Nucleus (AGN). The fact that it is one of the most X-ray luminous galaxies \\citep[$L_{0.2-4.0\\,{\\rm keV}}\\sim 10^{42}\\,{\\rm erg\\,s}^{-1}$, see][]{Fabbiano92,Zezas98} along with the flatness of its radio continuum emission had given some indications that an enshrouded AGN could be lurking in one of the nuclei.  Only recently though the hard X-ray ($10-40$\\,keV) Beppo-SAX observations of \\citet{Dellaceca02} have provided unambiguous evidence that a deeply buried ($N_{\\rm H} = 2.5 \\times 10^{24}\\,{\\rm cm}^{-2}$) AGN is located in one of the interacting components of the system. Furthermore, the same authors also demonstrated that the intrinsic AGN luminosity in both the UV and X-rays is almost an order of magnitude less than the infrared emission of the system calculated by \\citet{Charmandaris02}, indicating that the bulk of the infrared emission of the galaxy is indeed due to massive star formation as proposed by \\citet{Laurent00}.  The poor spatial resolution of the Beppo-SAX can not pinpoint the exact location of the AGN. However, preliminary analysis of the Chandra data \\citep{Zezas03}, which provide sub-arcsecond resolution, suggests that the AGN might be located within source B1 of NGC\\,3690.  Motivated by this recent activity, we decided to analyze in detail our $5-16\\,\\mu$m ISOCAM \\citep{Cesarsky96} spectrophotometric data on this system.  Our observations provide good spatial resolution with unprecedent sensitivity and enable us to examine in detail the dust properties of the obscured nuclear regions of Arp\\,299.  The observations and data reduction methods are described in Sect.\\,2, a brief description of the global mid-IR morphology of the system is presented in Sect.\\,3 and the analysis of the mid-IR spectra is shown in Sect.\\,4.  In Sect.\\,5, we discuss the implications of our findings on the physical characteristics of the various infrared sources and our conclusions are summarized in Sect.\\,6. ",
        "conclusions": "Our mid-IR spectroscopy of Arp\\,299 over the range 5 to $16\\,\\mu$m has revealed new information on the properties of this intriguing interacting system.  Based on our analysis, we conclude that:  \\begin{enumerate}  \\item Large quantities of dust are present in the system as indicated by the hot continuum observed above $10\\,\\mu$m in all sources and by the high extinction derived by our model.  We find that, assuming a uniform foreground screen model, we can reproduce the mid-IR SED and the $9.7\\,\\mu$m silicate absorption band with an $A_{\\rm V}$ of $\\sim 40$\\,mag in the central region of IC\\,694 (source A), $\\sim 10$\\,mag for source C, and $\\sim 50$\\,mag for C$^\\prime$, the dominant sources in the galactic disks overlap region.  A best fit with an $A_{\\rm V}$ of $\\sim 60$\\,mag is derived in the central region of NGC\\,3690 (source B1) using a homogeneous mixture of dust with radiating sources.  \\item The nuclear region of IC\\,694 harbors enshrouded massive star forming activity. No evidence for an active galactic nucleus is found in our mid-IR spectra.  \\item We confirm that C and C$^\\prime$ are sites of recent starburst events based on the elevated values of the [Ne{\\sc iii}]15.6$\\,\\mu$m flux and the high [Ne{\\sc iii}]/[Ne{\\sc ii}] ratio, which is similar to those of classical young starbursts.  \\item Source B is characterized by a unique mid-IR spectrum which is dominated by a hot dust continuum extending shortward of $5\\,\\mu$m, and which is the strongest of the Arp\\,299 system.  Given the ample presence of dust and the strength of the radiation field, this continuum is consistent with the recent Chandra/Beppo-SAX results which favor the presence of a Compton thick AGN at this location.  \\end{enumerate}"
    },
    "0311/astro-ph0311318_arXiv.txt": {
        "abstract": "We report the results of VLT and Keck adaptive optics surveys of known members of the $\\eta$ Chamaeleontis, MBM 12, and TW Hydrae (TWA) associations to search for close companions. The multiplicity statistics of $\\eta$ Cha, MBM 12, and TWA are quite high compared with other clusters and associations, although our errors are large due to small number statistics. We have resolved S18 in MBM 12 and RECX 9 in $\\eta$ Cha into triples for the first time. The tight binary TWA 5Aab in the TWA offers the prospect of measuring the dynamical masses of both components as well as an independent distance to the system within a few years. The AO detection of the close companion to the nearby young star $\\chi$$^{1}$ Orionis, previously inferred from radial velocity and astrometric observations, has already made it possible to derive the dynamical masses of that system without any astrophysical assumption. ",
        "introduction": "Nearby, young stellar associations offer unique advantages for detailed studies of star and planet formation. With diffraction limited adaptive optics (AO) near-infrared observations on 8 - 10 m class telescopes, their proximity facilitates sensitive studies of individual star systems down to 40 - 50 milliarcsecond scales (corresponding to physical scales of $\\sim$ 2 - 3 AU at 50 pc, or 6 - 9 AU at 150 pc). Resolutions of a few AU orbital radius imply orbital periods of only a few years, closing the gap between spectroscopic and visual binaries and thereby offering the prospect of a complete census of multiplicity among young pre-main-sequence (PMS) stars down to substellar masses. Resolving companions spatially is important since spectroscopic observations by themselves only yield the relative masses of binaries. By achieving resolutions down to a physical scale of a few AU, it is not only possible to determine the dynamical masses of the components, but also to derive an independent distance to the system. Thus, spatially resolved PMS binaries are essential tools for testing models of PMS stellar structure and evolution.  Searching for companions of PMS stars is also important since this allows us to address the ambiguity when placing them on evolutionary luminosity-color diagrams. Unresolved PMS binaries show luminosity (and possibly color) differences compared to the individual components, thereby biasing any conclusions drawn from evolutionary diagrams (see \\cite{math94} for an extensive review).  The nearby young associations $\\eta$ Chamaeleontis ($d$ $\\sim$ 97pc; \\cite{mam99}), MBM 12 ($d$ $\\sim$ 275pc; \\cite{luh01}), and TWA ($d$ $\\sim$ 55 pc; \\cite{kast97}, \\cite{per97}) are prime targets for detailed studies of close companions. Here, we report the results of an AO survey of known members of these associations. ",
        "conclusions": "We report the results of an AO survey of the known members of the MBM 12, TW Hydrae, and $\\eta$ Chamaeleontis associations to search for close companions. The main scientific conclusions can be summarized as follows:  \\begin{itemize} \\item Near-infrared AO observations on 8 - 10 meter class telescopes are closing the sensitivity gap to the spectroscopic surveys for binaries and are allowing direct component dynamical mass estimates\\\\ \\item The multiplicity statistics of $\\eta$ Cha, MBM 12, and TWA are quite high compared with other clusters and associations, although the significance of this result is low due to our small number statistics\\\\ \\item Within a few years, follow-up astrometric and spectroscopic observations of TWA 5A have the potential to provide both the dynamical masses of the individual components of this tight 54 mas binary, as well as an independent distance to the system\\\\ \\item AO observations have allowed the individual component dynamical masses of $\\chi$$^{1}$ Ori to be calculated for the first time without any astrophysical assumptions \\end{itemize}"
    },
    "0311/astro-ph0311632_arXiv.txt": {
        "abstract": "We present evidence for the dramatically different morphology between the rest frame UV and 7$\\mu$m mid-IR emission of VV\\,114 and Arp\\,299, two nearby (z$\\sim$0) violently interacting infrared luminous galaxies (LIRGs). Nearly all LIRGs are interacting systems and it is currently accepted that they dominate the IR emission at z$>$1. Luminous IR galaxies located at z=1--2 could easily be detected as unresolved sources in deep optical/near-IR ground based surveys, as well as in upcoming 24$\\mu$m surveys with the Space Infrared Telescope Facility. We demonstrate that the spatial resolution of these surveys will result in blending of the emission from unresolved interacting components. An increased scatter will thus be introduced in the observed optical to mid-IR colors of these galaxies, leading to a systematic underestimation of their dust content. ",
        "introduction": "Over the past decade several studies on the physical characteristics of high redshift galaxies have revealed a number of new exciting results. In particular it has been established that the star formation rate (SFR) of both individual galaxies, as well as the SFR per co-moving volume, increases by more than an order of magnitude as we move from systems of the local Universe to those found to a redshift of z$\\ge$2 \\citep[e.g.,][]{Steidel96,Madau98}. The details on exactly how the SFR evolves at even higher redshifts are still under debate as they are strongly affected by small number statistics on the properties of the population of the sources detected in sub-mm surveys, but more importantly by the accurate modeling of the effects of dust extinction. A number of deep surveys --- most notably the Hubble Deep Field --- have also shown that a large number of faint galaxies at intermediate and high redshift display irregular and/or disturbed morphology as well as an excess of blue light emission \\citep[see][]{Williams96,Abraham96,LeFevre00}. A way to explain this faint blue excess in the galaxy luminosity function is by introducing a rapidly evolving population of star forming galaxies at those redshifts \\citep{Glazebrook95,Metcalfe01}. What are the properties though, of this population? Several researchers suggest that these distant galaxies are a new class of systems which are physically small in size \\citep{Pascarelle96}, while others speculate that they are high redshift analogues of local interacting galaxies. The latter scenario can be understood since those distant galaxies may look peculiar because of the k-correction and band shifting of their short wavelength (UV) emission to the optical HST filters \\citep[][~and references therein]{Hibbard97}.  Two additional independent findings seem to suggest that galaxy interactions may indeed have significant implications to the blue excess in the galaxy luminosity function. The first one is the ample evidence that the galaxy interaction/merge rate increases as a power law function of redshift $\\sim$(1+z)$^m$ \\citep{Carlberg90,Lavery96}. Even though a direct observational measure of the interacting/merging galaxy pairs, and hence $m$, has been difficult due to a number of potential biases in surveys, current estimates indicate that $m$ can be as high as 3--4.5 \\citep[see][~for a review]{Lavery96,LeFevre00}. The second is that phenomena associated with galaxy interactions such as the presence of starburst, infrared luminous, and ultraluminous (LIRGs/ULIRGs) galaxies detected with IRAS also increase with redshift\\citep[e.g.,][]{Lonsdale90,Sanders96,Xu01}. Tidally triggered massive star formation in galaxies, which may eventually transform gas rich galaxies into their ultraluminous phase, will be responsible for increasing their rest frame UV flux and hence for the blue color excess of the galaxies detected in deep surveys. It is also the reason why large quantities ($\\sim$10$^9$M$_{\\sun}$) of molecular gas may be compressed over an area of few hundred pc around their nuclei \\citep{Downes98,Bryant99} eventually obscuring them behind large quantities of dust. As a result, wavelength dependent extinction by dust grains distorts the morphology of interacting galaxies, limiting our ability to accurately probe their properties .  A prime example of this phenomenon was displayed in the study of NGC\\,4038/39 (``The Antennae'') by \\citet{Mirabel98}. This system of interacting galaxies, located at distance of 20\\,Mpc is only moderately infrared luminous\\footnote{We use the standard definition of L$_{\\rm IR}$(8--1000\\,$\\mu$m) = 5.62$\\times$10$^5$\\,D(Mpc)$^2$\\,(13.48{\\it f}$_{12}$+5.16{\\it f}$_{25}$+2.58{\\it f}$_{60}$+{\\it f}$_{100}$)\\,L$_\\sun$ \\citep[see][]{Sanders96}, H$_0$\\,=\\,75\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, and q$_o$\\,=\\,0.5 to facilitate comparison with earlier work.}  (L$_{\\rm IR}$=4.9$\\times10^{10}$\\,L$_{\\sun}$) and still contains considerable quantities of molecular gas dispersed over a large area, which may become available for a more intense star formation event in the future \\citep{Wilson00}. Despite that, half of the mid-IR flux of the system originates from the optically obscured overlap region between the two nuclei \\citep[see][]{Mirabel98}.  It is thus reasonable to pose the following questions. How would the morphology of systems more IR luminous than NGC\\,4038/39 be affected by the presence of dust and what consequences this may have in measurements of the {\\em spatially unresolved} IR luminous galaxies which are readily discovered at high (z$\\gtrsim1.5$) redshift? More specifically, how might this affect the conclusions of color correlations between a number of optical and near-IR surveys that are performed from the ground \\citep[e.g.,][]{Jannuzi99}, and those scheduled to be performed with the Space Infrared Telescope Facility (SIRTF) \\citep[e.g.,][]{Rieke00,Lonsdale03}?  We will base our analysis on high resolution UV and mid-IR observations of Arp\\,299 and VV\\,114 that became recently available. Both galaxies have IR luminosities almost an order of magnitude higher than NGC\\,4038/39 (L$_{\\rm IR}\\sim$5$\\times10^{11}$L$_\\sun$ for Arp\\,299 and $\\sim$4$\\times10^{11}$L$_\\sun$ for VV\\,114) but the projected spatial separation between their two interacting components is still sufficiently large (8 and 6\\,kpc respectively) for us to be able to resolve them given their proximity (41 and 80\\,Mpc respectively). ",
        "conclusions": "Using archival images of VV\\,114 and Arp\\,299, two nearby violently interacting infrared luminous galaxies, we show the strikingly different morphology between their 0.2\\,$\\mu$m UV and 7\\,$\\mu$m mid-IR emission. Even though the redshifted rest frame emission from both wavelengths will be easily detected at z$\\sim$2 by ground based R-/I-band surveys and by the 24\\,$\\mu$m surveys of SIRTF respectively, the fact that we will not be able to spatially resolve the different components of each source will introduce a scatter in their measured colors. This scatter will not only depend on the relative contribution of each interacting component to the total IR luminosity, but it will also cause a systematic underestimate of the integrated rest frame mid-IR to UV colors leading to lower values of the total dust content of those systems."
    },
    "0311/astro-ph0311404_arXiv.txt": {
        "abstract": "\\noindent QSOs are believed to be powered by accretion onto massive black holes (BHs). In this paper, assuming that each central BH in nearby galaxies has experienced the QSO phase and ignoring BH mergers, we establish a relation between the QSO luminosity function (LF) and the local BH mass function (MF). The QSOLF is jointly controlled by the luminosity evolution of individual QSOs and the triggering history of the accretion onto seed BHs. By comparing the time integral of the QSOLF with that inferred from local BHs, we separate the effect of the luminosity evolution of individual QSOs from the effect of the triggering history. Assuming that the nuclear luminosity evolution includes two phases (first increasing at the Eddington luminosity with growth of BHs and then declining), we find that observations are generally consistent with the expected relation between the QSOLF and the local BHMF and obtain the following constraints on QSO models and BH growth: (i) The QSO mass-to-energy efficiency $\\epsilon$ should be $\\ga 0.1$. (ii) The lifetime (defined directly through the luminosity evolution of individual QSOs here) should be $\\ga 4\\times10^7\\yr$. The characteristic declining timescale in the second phase should be significantly shorter than the Salpeter timescale $\\tau\\Sp$, and BH growth should not be dominated by the second phase. (iii) The ratio of obscured QSOs/AGNs to optically bright QSOs should be not larger than 7 at $M_B\\sim -23$ and 3 at $M_B\\sim -26$ if $\\epsilon=0.31$, and not larger than 1 at $M_B\\sim -23$ and negligible at $M_B\\sim -26$ if $\\epsilon=0.1$. (iv) It is unlikely that most QSOs are accreting at super-Eddington luminosities. We point out that it is hard to accurately estimate the value of the QSO lifetime estimated from the QSOLF and/or the local BHMF, if it is longer than a certain value (e.g.,\\ $\\sim4\\tau\\Sp$ in this study). We discuss the importance of accurate measurements of the intrinsic scatter in the BH mass and velocity dispersion relation of local galaxies and the scatter in the bolometric correction of QSOs. We also discuss some possible applications of the work in this paper, such as to the study of the demography of QSOs and the demography of normal galaxies at intermediate redshift. ",
        "introduction": "\\label{sec:intro}  \\noindent The exploration of the relation of QSOs with massive black holes (BHs) in nearby galaxies has been of considerable interests since the discovery of QSOs. On the one hand, QSOs are suggested to be powered by gas accretion onto massive BHs \\citep{S64,ZN64}. This (currently widely accepted) model suggests that a population of massive BHs as ``dead'' QSOs exist in nearby galactic centers \\citep{Lynden69}. Furthermore, the total mass density of these remnant BHs and the typical BH mass in nearby galaxies can be inferred from the total energy density radiated in photons by QSOs (\\citealt{S82}; see also \\citealt{R84}). These simple and elegant arguments motivate the search for BHs in nearby galaxies (e.g.,\\ \\citealt{KR95,Magorrian98,Gebhardt03,K03}). On the other hand, as a result of the endeavor in the past two decades, not only has the existence of the massive dark objects (which are presumably BHs here) in most nearby galactic centers been confirmed, but also dramatic progress on their demography has been recently achieved (e.g.,\\ \\citealt{KG01,Tremaine02} and references therein). Comparisons between the properties of local BHs and those inferred from the QSO model may shed new light on our understanding of the BH growth, the accretion physics, the mechanisms to trigger and quench nuclear activities, the formation and evolution of galaxies etc.\\ (e.g.\\ \\citealt{CP89,SB92,FI99,Salucci99,Elvis02,MS02,YT02,Fab03}).  The comparison between the local BH mass density and the local energy density in QSO photons has suggested that local BHs acquire most of their mass through accretion during the QSO/AGN phase \\citep{YT02,AR02,Fab03}. The mass distribution of local BHs as the remnant of nuclear activities is therefore controlled by the triggering history of the accretion onto seed BHs and the luminosity evolution of individual triggered nuclei with a mass-to-energy conversion efficiency $\\epsilon$ (i.e., the growth of individual BHs due to accretion). The triggering rate is usually believed to be related to the formation and evolution of galaxies and is a function of cosmic time, and the luminosity evolution of individual nuclei is believed to contain information on the accretion process in the vicinity of the BH and is a function of the physical time spent since the triggering of the accretion onto seed BHs. The triggering history and the luminosity evolution jointly control the QSO luminosity function (LF) as a function of luminosity and redshift; however, each cannot be uniquely determined from the QSOLF itself, partly as a result of the mixing of their effects on the QSOLF. In this paper, by using the local BH mass function (MF) as an additional constraint on QSO models and establishing a new relation between the QSOLF and the local BHMF, the luminosity evolution of individual QSOs is isolated from the triggering history of the QSO population in their effects on the QSOLF, and then we use observations to test QSO models and provide constraints on the QSO luminosity evolution and BH growth.  A relation between the QSOLF and the local BHMF has been achieved by \\citet{YT02} under the assumptions that the seed BH mass is negligible and the luminosities of QSOs are only an increasing function of their central BH mass. In the study of this paper, we relax these two assumptions. In addition, the relation in \\citet{YT02} includes the effect of BH mergers. However, BH mergers are ignored in the relation established in this paper for the following reasons: (i) BH mergers are not shown to play a significant role or not necessarily required at least for growth of high-mass ($\\ga 10^8\\msun$) BHs, if $\\epsilon\\simeq 0.1$ \\citep{YT02}; (ii) currently, the BH merger process and rate are very uncertain; and (iii) comparison of observations with the expectation obtained by ignoring BH mergers may also provide considerable insights in the role of BH mergers. The detailed difference between these two relations will be further discussed in this paper (see \\S~\\ref{sec:degeneracy} and \\ref{sec:parameters}).  This paper is organized as follows. In \\S~\\ref{sec:relation}, by studying the continuity equation for the BH mass and nuclear luminosity distribution, we establish the relation between the QSOLF and the local BHMF. The triggering history of the accretion onto seed BHs is (implicitly) considered in the continuity equation but circumvented in the relation between the QSOLF and local BHMF. Only the luminosity evolution of individual QSOs is incorporated in this relation. In \\S~\\ref{sec:obs} we obtain the local BHMF by using the Sloan Digital Sky Survey (SDSS) observation results on both early-type and late-type galaxies and the empirical relations on the demography of galaxies and BHs (the BH mass and velocity dispersion relation, the Tully-Fisher relation, etc.). We also review the QSOLF obtained from large optical surveys. In \\S~\\ref{sec:comparison} we combine the observations with the relation obtained in \\S~\\ref{sec:relation} to provide accurate observational constraints on the luminosity evolution of individual QSOs and BH growth, such as the QSO lifetime, the mass-to-energy conversion efficiency, the role of obscuration in BH growth, etc. The results are discussed in \\S~\\ref{sec:discussion}. In \\S~\\ref{sec:application} we further discuss some possible applications of the framework established in \\S~\\ref{sec:relation}, such as to the study of the demography of QSOs and the demography of normal galaxies at intermediate redshift. Finally, our conclusions are summarized in \\S~\\ref{sec:conclusion}.  In this paper we set the Hubble constant as $H_0=100h\\kms\\Mpc^{-1}$, and if not otherwise specified, the cosmological model used is $(\\Omega_{\\rm M},\\Omega_{\\Lambda},h)=(0.3,0.7,0.65)$ \\citep{Wang00}. ",
        "conclusions": "\\label{sec:conclusion} \\noindent Assuming that each massive BH in nearby galactic centers has experienced the QSO phase and becomes quiescent at present, we have established a relation between the QSOLF and the local BHMF by studying the continuity equation for the BH mass and nuclear luminosity distribution and ignoring BH mergers. This relation compares the time integral of the QSOLF and that inferred from the local BHMF and only incorporates the luminosity evolution of individual QSOs. The triggering history of the accretion onto seed BHs is (implicitly) considered in the continuity equation but is circumvented in the relation between the QSOLF and the local BHMF. For comparison, the old relations between QSOs and local BHs on the total/partial BH mass densities (see eqs.~\\ref{eq:soltan} and \\ref{eq:YT}; \\citealt{S82,YT02}) include the effect of BH mergers, but the seed BH mass is ignored. The relation on the total BH mass density (eq.~\\ref{eq:soltan}; \\citealt{S82}) is unrelated with the luminosity evolution of individual QSOs, and the relation on the partial BH mass density (eq.~\\ref{eq:YT}; \\citealt{YT02}) assumes that the luminosity of QSOs is only an increasing function of their central BH mass (e.g.,\\ the Eddington luminosity in the calculation in \\citealt{YT02}). The new relation on the time integral of the QSOLF in this paper can be used to explore the luminosity evolution of individual QSOs (see \\S~\\ref{sec:intro} and \\ref{sec:degeneracy}).  By applying observations into the relation established in this paper and assuming that the nuclear luminosity evolution includes two phases (first increasing at the Eddington luminosity with the BH growth and then declining), we find that the time integral of the QSOLF is generally consistent with that inferred from local BHs and obtain the following observational constraints on the QSO luminosity evolution and BH growth: (i) The QSO mass-to-energy efficiency $\\epsilon$ should be $\\ga 0.1$ (see Fig.~\\ref{fig:eff}). (ii) The lifetime (defined directly through the luminosity evolution of individual QSOs) should be longer than $\\tau\\Sp$($\\simeq 5\\times 10^7\\yr$) if $\\epsilon=0.1$ and $0.2\\tau\\Sp$($\\simeq 4\\times 10^7\\yr$) if $\\epsilon=0.31$ (see Figs.~\\ref{fig:model1sta} and \\ref{fig:eff}). The characteristic declining timescale in the second phase should be significantly shorter than $\\tau\\Sp$, and BH growth should not be dominated by the second phase (when QSOs are accreting at sub-Eddington luminosities; see Fig.~\\ref{fig:model3}). (iii) The upper limit of the ratio of obscured QSOs/AGNs to optically bright QSOs is provided, which should be not larger than 7 at $M_B\\sim -23$ and 3 at $M_B\\sim -26$ if $\\epsilon=0.31$ and not larger than 1 at $M_B\\sim -23$ and negligible at $M_B\\sim -26$ if $\\epsilon=0.1$ (see Figs.~\\ref{fig:model1sta} and \\ref{fig:eff}). (iv) It is unlikely that most QSOs are accreting at super-Eddington luminosities (see Fig.~\\ref{fig:supedd}). The constraints above are obtained by assuming that the two accretion phases appear only once in the luminosity evolution, although this assumption is not required in the relation established in this paper. The possibility of more complicated accretion patterns deserves further investigation.  We find that if the QSO lifetime is longer than a certain value (e.g.,\\ $\\sim 4\\tau\\Sp$; see Fig.~\\ref{fig:model1sta}), the time integral of the nuclear LF inferred from local BHs becomes insensitive to the value of the QSO lifetime, and thus it is difficult to provide an accurate estimate on the QSO lifetime unless observations extend to fainter luminosities or precise measurements of the QSOLF and local BHMF are available (e.g.,\\ with error much less than 10\\%). We also point out that this difficulty would also exist in many other methods to estimate the QSO lifetime by using the QSOLF, as a result of the sharp decrease of the QSOLF at the bright end and the limited luminosity range in observations.  We show the importance of accurately measuring the intrinsic scatter in the relation between the BH mass and velocity dispersion of local galaxies and the scatter in the distribution of bolometric corrections of QSOs to precise understanding of the physics behind the QSO phenomenon and BH growth. Both of the scatters affect the shape and values of the time-integral of the nuclear LF, especially at the bright end.  With the upcoming more precise measurement on QSOs (including both unobscured and obscured AGNs) and the demography of local BHs and galaxies (by SDSS, Chandra, XMM, etc.), the method presented in this study would help to further explore the nuclear activity triggering and quenching mechanisms, obscuration of QSOs/AGNs, the demography of QSOs/AGNs, and the demography of normal galaxies at intermediate redshift and finally understand the physics behind the QSO phenomenon and the formation and evolution of galaxies."
    },
    "0311/astro-ph0311297_arXiv.txt": {
        "abstract": "We present high resolution mid-infrared (mid-IR; 11.7 and 17.65 $\\micron$) maps of the central 400 pc region of the starburst galaxy M82.  Seven star forming clusters are identified which together provide $\\sim$ 15\\% of the total mid-IR luminosity of the galaxy. Combining the mid-IR data with thermal radio measurements and near- and mid-IR line emission, we find that these young stellar clusters have inferred masses and sizes comparable to globular clusters.  At least 20\\% of the star formation in M82 is found to occur in super-star clusters. ",
        "introduction": "M82 (NGC 3034) is often considered the archetypical starburst galaxy since it has a derived star formation rate ($\\sim$10 M$_{\\odot}$ yr$^{-1}$; \\citealt{om78}) that would deplete the observed molecular gas in $<$ 10$^8$ yrs (i.e. short on a Hubble timescale; \\citealt{lord96}).  The starburst phenomena traces recent star formation since it is the massive stars ($>$8 M$_{\\odot}$), which have short lives, that dominate the energetic output of the host galaxy. The nuclear starburst of M82 dominates the infrared (IR) luminosity of the galaxy - essentially all of the galaxy's L$_{IR} \\sim$ 3.6 $\\times$ 10$^{10}$ L$_{\\odot}$\\footnote{\\citet{telesco80} published L$_{IR}$ = L(1-300$\\micron$) = 3.0 $\\times$ 10$^{10}$ L$_{\\odot}$ assuming a distance to M82 of 3.3 Mpc.  We have updated L$_{IR}$ for the distance of 3.6 Mpc \\citep{freedman94, sakai99} used in this paper.} comes from the central kpc \\citep{telesco80}.  Optical and near-IR imaging of M82 has revealed numerous super-star clusters in its active star-forming nucleus.  Ground-based optical imaging of the central region detected eight young knots distributed throughout the region \\citep{om78}.  Further study of these knots broke them into smaller star clusters (half light diameters $\\sim$3-4 pc) and suggested cluster dynamical masses 10$^4$-10$^6$ M$_{\\odot}$ \\citep{smith01, degrijs01, mccrady03}.  M82's current starburst is thought to have been triggered by its interaction with M81 $\\sim$10$^8$ yrs ago \\citep{cottrell77,achtermannlacy95}.  It has been argued that the conditions resulting from interactions and mergers of galaxies are favorable for globular cluster formation (eg. \\citealt{taniguchi99} and ref. therein).  UV studies of global properties of starbursting galaxies have shown that as much as 20\\% of the light is produced in the luminous knots \\citep{meurer95, zepf99}, suggesting a high efficiency of cluster formation in starbursts.  Understanding the star formation occurring in M82's nuclear region can provide insight into both the general process of star formation in starburst environments and also the process and efficiency of forming super-star clusters.  The nearly edge-on geometry of M82 combined with heavy optical extinction (A$_V \\sim$ 5-25 mag; \\citealt{lester90, telesco91, larkin94, satyapal95}) has made the galaxy's central 400 pc difficult to study.  Even in the near-IR where A$_{2.2 \\micron} \\sim$ 1/10A$_V$ \\citep{rieke85}, it is difficult to directly measure the deepest star forming regions.  The mid-IR region of the spectrum is essential for probing deep into the central regions of M82 and revealing the details of the intense star formation occurring in the central regions.  Previously, the highest resolution mid-IR map of the central region of M82 was a 12.4 $\\micron$ map by \\citet{telescogezari92}, which had a resolution $\\sim$1.1$\\arcsec$.  Mid-IR maps of M82 with lower resolution have also been published by \\citet{rieke80}, \\citet{telesco89}, \\citet{deitz89}, \\citet{telesco91} and \\citet{forsterschreiber03}.  In this paper, we improve upon previous observations by presenting higher resolution mid-IR maps of M82.  We discuss evidence suggesting the sources in the maps are young counterparts to globular clusters and estimate the efficiency of super-star cluster formation in M82. ",
        "conclusions": "This paper presents mid-IR (11.7 and 17.65 $\\micron$) maps with $\\sim$0.5$\\arcsec$ resolution of the central 400 pc of the starbursting galaxy M82.  We find 7 resolved sources in this region of M82 with luminosities summing to 15\\% of the total IRAS flux of the entire galaxy.  The mid-IR maps exhibit features comparable to those found in maps of [Ne II] emission, Br$\\gamma$ emission, and thermal free-free emisson.  We present evidence implying that the mid-IR sources are giant H II regions in which globular cluster-sized star clusters are forming.  Our data imply that $\\gtrsim$20\\% of the star formation in M82 is occurring in super-star clusters."
    },
    "0311/astro-ph0311068_arXiv.txt": {
        "abstract": "Titanium carbide (TiC) nanocrystals were recently proposed as the carrier of the mysterious 21$\\mum$ emission feature observed in post-asymptotic giant branch stars, based on their close spectral match and the presolar nature of meteoritic TiC nanograins (which reveals their stellar ejecta origin). But we show in this {\\it Letter} that the Kramers-Kronig dispersion relations, which relate the wavelength-integrated extinction cross section to the total dust mass, would impose a lower bound on the TiC mass. This Kramers-Kronig lower limit exceeds the maximum available TiC mass by a factor of at least $\\simali$50, independent of the absolute value of the ultraviolet/visible absorptivity of nano TiC. The TiC model is therefore readily ruled out by the Kramers-Kronig physical principle. ",
        "introduction": "} A prominent broad (with a FWHM of $\\simali$2$\\mum$) emission feature at about 21$\\mum$\\footnote{% \\label{ftnt:iso21um} High resolution spectra obtained with the {\\it Short-Wavelength Spectrometer} (SWS) instrument (with a resolution of $\\lambda/\\Delta\\lambda \\approx 2000$) on board the {\\it Infrared Space Observatory} show that this feature actually peaks at $\\simali$20.1$\\mum$ (Volk, Kwok, \\& Hrivnak 1999). But for historical reasons, this feature continues to be referred to as the ``21$\\mum$ feature''. } was discovered by Kwok, Volk, \\& Hrivnak (1989) in an analysis of the {\\it Infrared Astronomical Satellite} (IRAS) {\\it Low Resolution Spectrometer} (LRS) spectra of four carbon-rich post-asymptotic giant branch (AGB) stars, and subsequently confirmed by both ground-based and airborne {\\it Kuiper Airborne Observatory} and ISO observations (see Kwok, Volk, \\& Hrivnak 1999 and references therein). So far, this feature has been detected in twelve post-AGB stars (commonly termed as the ``21$\\mum$ sources''; Kwok et al.\\ 1999)\\footnote{% A weak 21$\\mum$ feature was recently reported for three planetary nebulae (PNs; Hony, Waters, \\& Tielens 2001; K. Volk 2003, in preparation). }, with little shape variation found between different sources. These stars have quite uniform properties: they are mostly metal-poor carbon-rich F and G supergiants with strong infrared (IR) excesses and overabundant $s$-process elements (see Kwok et al.\\ 1999).  The carrier of this mysterious 21$\\mum$ feature remains unidentified, although many candidate materials have been proposed including iron oxides Fe$_2$O$_3$ or Fe$_3$O$_4$, hydrogenated amorphous carbon, hydrogenated fullerenes, % hydrogenated nanodiamond, amides (thiourea or urea ${\\rm OC\\left[NH_2\\right]_2}$), SiS$_2$, % oxygen-bearing side groups in coal (see Kwok et al.\\ 1999 and references therein), and more recently titanium carbide (TiC) nanoclusters (von Helden et al.\\ 2000), SiC (Speck \\& Hofmeister 2003), and stochastically-heated silicon core-SiO$_2$ mantle nanograins (Smith \\& Witt 2002; Li \\& Draine 2002).  The nano-TiC model seems attractive. While bulk TiC does not show any noticeable feature near 20.1$\\mum$ (Henning \\& Mutschke 2000), laboratory absorption spectra of TiC nanocrystals containing 27--125 atoms exhibit a distinct feature at $\\simali$20.1$\\mum$, closely resembling the astronomical 21$\\mum$ emission feature both in peak position, width, and in spectral details (von Helden et al.\\ 2000). This model further gains its strength from the identification of presolar TiC grains (with radii $\\simali$100$\\Angstrom$) in primitive meteorites as inclusions embedded in micrometer-sized presolar graphite grains (Bernatowicz et al.\\ 1996).  Since Ti is a rare element, it is important to know whether the amount of TiC dust required to account for the observed 21$\\mum$ feature is a reasonable quantity. Since the optical properties of nano TiC in the ultraviolet (UV)/visible wavelength range are unknown, one has to rely on those of bulk TiC as a starting point. Such an attempt has recently been made for the post-AGB star $\\HD$ by Hony et al.\\ (2003). They found that to explain the observed 21$\\mum$ feature, the TiC model requires a much higher UV/visible absorptivity than that of bulk TiC (by a factor of $\\simgt$20). However, we will demonstrate in this {\\it Letter} how the TiC absorption integrated over a finite wavelength range places a lower limit on the TiC dust mass through the Kramers-Kronig dispersion relations. For $\\HD$, this Kramers-Kronig lower limit exceeds the maximum available TiC mass by a factor of $\\simgt$50, independent of the absolute value of the UV/visible absorptivity. Therefore, increasing the UV/visible absorptivity is unable to avoid the TiC abundance problem. ",
        "conclusions": "\\label{sec:discussion} In this {\\it Letter} we have assumed that the TiC dust follows the spatial distribution of the bulk carbon dust (see \\S\\ref{sec:obs}). But this is not critical. Even if we assume that all the TiC dust accumulates at the inner edge of the envelope, to explain the observed 21$\\mum$ feature, one still needs $\\simali$42\\% of the total TiC mass derived in \\S\\S\\ref{sec:qabs},\\ref{sec:kk}, since for the latter the average starlight intensity to which the dust is exposed is just $\\simali$42\\% higher. Therefore, local density enhancement of TiC would not overcome the TiC abundance problem. High-resolution (FWHM 0.4$^{\\prime\\prime}$) imaging at 20.8$\\mum$ revealed two bright blobs at $\\simali$1$^{\\prime\\prime}$ from the star, but their surface brightness is at most two times higher than the average (Kwok, Volk, \\& Hrivnak 2002), suggesting the insignificance of local density enhancement of the 21$\\mum$ feature carrier.  Von Helden etal.\\ (2000) postulated that TiC nanograins could have formed during the extreme conditions (i.e., high density and high pressure) associated with the short, high mass-loss superwind phase where AGB stars lose the remaining stellar envelope, terminating their life on the AGB and starting the transition to the PN phase. However, Kwok et al.\\ (2002) argued that there is no evidence that the 21$\\mum$ emission is created by a sudden ejection at the end of the AGB, implying that the necessary physical conditions required to form TiC in the ejecta of carbon-rich evolved stars may not be met.  In summary, although we are lack of experimental knowledge of the UV/visible absorption properties of nano TiC, a physical argument based on the Kramers-Kronig dispersion relations implies that the nano-TiC model requires too much Ti to reconcile with the observed photospheric Ti abundance.\\footnote{% If these grains are coated by a mantle of graphite (see Chigai et al.\\ 2003), the situation would be even worse since the large IR emissivity of graphite together with the fact that the increase in grain size may prohibit them from single-photon heating would jointly suppress emission in the 21$\\mum$ feature. } Recently, the nano-TiC model was also challenged by Chigai et al.\\ (2003), who found that the abundance ratio of Ti to Si needed to reproduce the observed flux ratios of the 11.3$\\mum$ SiC feature to the 21$\\mum$ TiC feature of post-AGB stars must be at least 5 times larger than the solar abundance ratio. Finally, we should note that laboratory absorption spectra of TiC nanoclusters also show prominent features at other wavelengths which are not seen in post-AGB stars (e.g. Ti$_{14}$C$_{13}$ also has a strong band at $\\simali$16$\\mum$; van Heijnsbergen et al.\\ 1999)."
    },
    "0311/astro-ph0311542_arXiv.txt": {
        "abstract": "Using a novel approach we have reanalized the question of whether the extreme star forming galaxies known as HII galaxies are truly young or rejuvenated old systems.  We first present a method of inversion that applies to any monotonic function of time describing the evolution of independent events. We show that, apart from a normalization constant, the ``true'' time dependence can be recovered from the inversion of its probability density function.  We applied the inversion method to the observed  equivalent width of \\hb\\ (EW(\\hb )) distribution for objects in the Terlevich and collaborators Spectrophotometric Catalogue of HII galaxies and found that their global history of star formation behaves much closer to the expectations of a continuos star formation model than to an instantaneous one. On the other hand, when the inversion method is applied to samples within a restricted metallicity range we find that their history of star formation behaves much closer to what the instantaneous model predicts.  Our main conclusion is that, globally, the evolution of HII galaxies seems consistent with a succession of short starbursts separated by quiescent periods and that, while the emission lines trace the properties of the present burst, the underlying stellar continuum traces the whole star formation history of the galaxy. Thus, observables like the EW(\\hb ) that combine an emission line flux, i.e. a parameter pertaining to the present burst, with the continuum flux, i.e. a parameter that traces the whole history of star formation, should not be used alone to characterize the present burst. ",
        "introduction": "HII galaxies are dwarf emission line galaxies undergoing a burst of star formation.  They are characterized by strong and narrow emission lines originated in a giant star forming region which dominate their {\\em observable} properties at optical wavelengths. Most HII galaxies are in fact Blue Compact Galaxies (BCG's), but due to the different selection criteria only a small percentage of BCG's,  i.e. those with the largest emission line equivalent widths, are HII galaxies. We will stick to the name \\,``HII galaxies\\,'' to refer to the star forming systems selected from objective prism surveys and  having strong narrow emission lines.  Various studies of the spectroscopic properties of HII galaxies in optical wavelengths revealed systems of very low heavy element abundances and high rates of star formation. Earlier morphological studies have suggested that a large proportion of the sample of HII galaxies observed are compact and isolated  (Melnick 1987). This, together with the spectroscopic properties indicating low to very low metallicity plus a very young stellar content, led workers in the field ever since their discovery, to question whether these systems are truly young galaxies or made up by a few bursts separated by long quiescent periods in the lifetime of the galaxy. Reviews of the general properties of HII galaxies can be found in  Melnick (1987), Terlevich (1988), the {\\em \\,``Spectrophotometric Catalogue of HII Galaxies\\,''}  (hereafter SCHG, Terlevich \\etal 1991), Stasi\\'nska \\&  Leitherer (1996), Telles \\& Terlevich (1993; 1995; 1997), Telles, Melnick \\& Terlevich (1997), and more recently in the excellent review by Kunth \\& \\\"{O}stlin (2000).  More than 20 years ago Dottori (1981) suggested the use of the EW(\\hb ) to detect differences in age among HII regions. Dottori \\& Bica (1981) applied the EW(\\hb ) method to 31 LMC and SMC HII regions and found an uneven distribution of ages with most HII regions clustering at an EW(\\hb ) of about 70\\AA\\ and only 5 HII regions  with EW(\\hb ) $>$120 \\AA . This result led them to suggest that a galaxy-wide burst of star formation did occur 6 to 7 Myr ago in the Magellanic clouds.  \\begin{figure*} \\setlength{\\unitlength}{1cm} \\begin{picture}(7,8) \\put(-5.5,-3.0){\\special{psfile=fig1a.ps hoffset=0 voffset=0 hscale=45.0 vscale=45.0 angle=0}} \\put(3.5,-3.0){\\special{psfile=fig1b.ps hoffset=0 voffset=0 hscale=45 vscale=45 angle=0}} \\end{picture} \\caption{ The equivalent width distribution of HII galaxies from the SCHG is plotted on the left. Clearly, and contrary to model expectations, there are only a few systems with EW(\\hb ) $>$ 150\\AA. The predicted time evolution of a coeval population for a Salpeter IMF with 100\\Msol\\ and 0.1\\Msol\\  upper and lower mass limit respectively is plotted on the right. For a random population with a constant production rate and considering that after 6.5 Myr a system will no longer be considered an HII galaxy, about equal number of systems with ages below and above 3.2Myr are expected, i.e. about equal number of systems with EW(\\hb ) above and below 150-200 \\AA. } \\label{figure1} \\end{figure*}   The SCHG compiles, among other parameters, line ratios and equivalent widths for several hundred HII galaxies, and gives the opportunity of investigating the ages of actively star forming galaxies to find whether there is among them a truly young system. The SCHG is particularly appropiate for this task because, due to the selection criterion used to find them, the HII galaxies in the SCHG are probably the youngest systems that can be studied in any detail. This is due to the fact that the SCHG samples the narrow emission line galaxies with the highest equivalent width in their emission lines, and is particularly biased in the local universe (at z$<$0.04) towards strong and compact emission in [OIII] 5007\\AA. This bias is introduced by the technique used in all the 3 sources of data, the Cambridge, Tololo and University of Michigan surveys, that searched for strong emission using a  Schmidt camera (either the UK or the Tololo Schmidt at the AAT and CTIO respectively) equipped with low dispersion objective prisms and IIIaJ emulsion, a combination that produces a sensitivity range from 3500\\AA\\ to 5300\\AA .   One of the early surprising results of the SCHG was that even in this strong emission line biased sample, there are less than 10$\\%$ of the systems with EW(\\hb ) $>$150 \\AA\\ (Fig.~1, left). In contrast, as shown in Fig.~1, right, a quick analysis of the predictions for a single burst indicates that if star formation proceeeds uniformly with time, about half of the regions of star formation  should have an EW(\\hb ) larger than about 150 - 200 \\AA\\ depending on model assumptions (all our synthesis computations use the SB99 code of Leitherer and collaborators (1999)).  Many explanations have been put forward for the lack of systems with high EW(\\hb ). Chiefly among them, time dependent escape of ionizing photons, dust affecting either the ionizing radiation or the visibility occurring preferentially during the first 3 Myrs, the presence of an underlying old population, uncertainties with the models, etc. The wide variety of possibilities has had  the negative effect of almost halting the research in this important area.   We decided to take a fresh and different approach based as much as possible on unsupervised analysis of samples. We describe here a new method that, under the assumption that the catalogue samples a population of bursting galaxies at different ages, allows the reconstruction of the time evolution of the burst Balmer emission lines equivalent width, from their observed distribution.  \\begin{figure*} \\setlength{\\unitlength}{1cm} \\begin{picture}(7,10.5) \\put(-4.5,-0.3){\\special{psfile=fig2.eps hoffset=0 voffset=0.0 hscale=61.0 vscale=61.0 angle=00}} \\end{picture} \\caption{Panel a shows the sample of random numbers homogeneously distributed in time, which we associate with the starburst birthrate. Panel b shows the adopted model, same as figure 1. Panel c shows the computed distribution function of equivalent widths. Panel d shows the inversion of panel c  distribution function using equation 3. Note that the inversion closely reproduces the input model. } \\label{panels} \\end{figure*} ",
        "conclusions": "We have developed a simple inversion tool that allows to reconstruct from observed probability density distributions of some monotonical parameter, its time evolution.   We applied the inversion method to the EW(\\hb ) distribution of a sample of 217 extreme star forming galaxies from the SCHG. We have  shown  that, considering the sample as a whole, its EW(\\hb ) evolution is not well described by a coeval burst model, and that HII Galaxies seem to have a star formation history that is closer to that predicted by a continuous star formation model.  The simplest interpretation is that while the observed emission lines track the present burst, the underlying continuum contains the whole history of star formation of the HII galaxy.  Even though HII Galaxies are selected by their strong emission lines, they are not truly young galaxies. Most of them have undergone substantial star formation probably during the previous 100 - 1000 My to the present burst.   The situation changes when the analysis is restricted to a sub-sample covering a narrow range in metallicities ($7.50<12+log\\,O/H<8.25$). In this case the EW(\\hb ) evolution seems well described by a coeval burst model with an IMF having an upper mass limit around 80~\\Msol.    Clearly it would be very interesting to extend the method to the analysis of larger samples of galaxies like, e.g., that one provided by the Sloane Digital Sky survey."
    },
    "0311/astro-ph0311519_arXiv.txt": {
        "abstract": "{ A small fraction of all quasars are strongly lensed and multiply imaged, with usually a galaxy acting as the main lens. Some, maybe all of these quasars are also affected by microlensing, the effects of stellar mass objects in the lensing galaxy. Stellar microlensing not only has photometric effects (the apparent magnitudes of the quasar images vary independently due to the relative motion between source, lens and observer), it also affects the observed position of the images. This astrometric effect was first explored by  Lewis and Ibata (1998): the position of the quasar -- i.e. the center-of-light of the many microimages -- can shift by  tens of microarcseconds due to the relatively sudden (dis-)appearance of a pair of microimages when a caustic is being crossed.  We explore this effect quantitatively for different values of the lensing parameters  $\\kappa$ and  $\\gamma$ (surface mass density and external shear) covering most of the known multiple quasar systems. We show examples of microlens-induced quasar motion and the corresponding light curves for different quasar sizes. We evaluate statistically the occurrence of large shifts in angular position and their correlation with apparent brightness fluctuations. We also show statistical relations between positional offsets and time from random starting points. As  the amplitude of the astrometric offset depends on the source size, astrometric microlensing signatures of quasars -- combined with the photometric variations -- will provide very good constraints on the size of quasars as a function of wavelength. We predict that such signatures will be detectable for realistic microlensing scenarios with near future technology in the infrared/optical (Keck-Interferometry, VLTI, SIM, GAIA). Such detections will show that not even high redshift quasars define a ``fixed'' coordinate system. ",
        "introduction": "Gravitational lensing acts on quasars in a number of ways. The best known effect is     strong lensing which       produces multiply imaged quasars. Only a small fraction of all quasars (about 1 in 500) are multiply imaged with galaxies acting as the main lenses (for an updated list, see the CASTLES web page {\\tt http://cfa-www.harvard.edu/glensdata/}). Most of these quasars -- and possibly some ``single'' quasars as well -- are also affected by microlensing: the coherent effects of stellar mass objects in the lensing galaxy. Microlensing is well known for affecting the apparent brightness of the quasar, which changes due to the relative  motion between lens, source and observer (see, e.g., Paczy\\'nski 1986, Wambsganss, Paczy\\'nski \\& Schneider 1990, Lewis et al. 1998, Wyithe et al. 2000a, Wambsganss 2001, Wo\\'zniak et al. 2000a,b). However, stellar microlensing has yet another effect, which was first explored by  Lewis and Ibata (1998): when the quasar crosses a (micro-)caustic, a new pair of microimages is created or destroyed. The relatively sudden (dis-)appearance of a highly magnified image pair does  not just produce a strong fluctuation in apparent brightness,     it can also shift the center-of-light of the quasar (the weighted sum of all the microimages) by       tens of microarcseconds. Such a positional change is in principle detectable with current VLBI techniques in the radio regime, and will be measurable in the infrared and optical wavebands as well with (near) future technology (Keck-Interferometry, VLTI, SIM, GAIA).  A number of studies have explored astrometric microlensing in a different regime, namely stars in the Milky Way (or dark matter objects in its halo) acting on background stars either in the Magellanic Clouds or in the bulge of the Milky Way (e.g., Miyamoto \\& Yoshii 1995; Miralda-Escude 1996; Mao \\& Witt 1998; Boden, Shao \\& van Buren 1998; Goldberg \\& Wo\\'zniak 1998; Paczy\\'nski 1998; Han, Chun \\& Chang 1999; Han \\& Kim 1999; Han \\& Jeong 1999; Safizadeh, Dalal, Griest 1999; Dominik \\& Sahu 2000; Gould \\& Han 2000; Han \\& Kim 2000; Salim \\& Gould 2000; Delplancke et al. 2001; Belokurov \\& Evans 2002; Dalal \\& Lane 2003). The effects of Milky Way stars on background quasars were studied by Hosokawa, Ohnishi \\&  Fukushima (1997), Sazhin et al. (1998) and Honma \\& Kurayama (2002). Williams \\& Saha (1995) had discussed large image shifts produced by substructure in the lensing galaxy. However, not much work has yet been done on cosmological astrometric microlensing -- stars in lensing galaxies acting on even more distant quasars -- beyond Lewis \\& Ibata (1998), with the recent exception of Salata \\& Zhdanov (2003). The present paper aims to explore this further.  We study quantitatively eight different cases with various   values of the surface mass density,  with and without external shear. We present example microlensing lightcurves with the corresponding center-of-light shifts and  we investigate the correlation between high-magnification photometric events and large-offset astrometric events.  As first pointed out by Lewis \\& Ibata (1998), large offset events are   typically correlated with high-magnification events, whereas the inverse is much less true. The reason for this is that the location of the newly appearing bright image pair during a caustic crossing is unrelated to the  previous  center-of-light of the microimages. In rare cases, the new image pair may appear at or close to the center-of-light of the pre-existing  microimages. Such a situation would correspond to a large change in brightness with very little change in position. In most cases, however, the new bright image pair will appear at a location which is unrelated to the previous center-of-light, and hence produce a sudden jump. Since the positional offset is preferentially perpendicular to the direction of the external shear, the shear direction can be inferred this way.  After introducing the microlensing length and time scales (Section 2), we describe our simulations (Section 3). In Section 4, we  illustrate the effect of astrometric and photometric fluctuations and present statistical correlations between positional offsets, magnitude fluctuations and time intervals between the measurements for four microlensing situations with different surface mass densities ($\\kappa = 0.2, 0.4, 0.6, 0.8$), with and without external shear ($\\gamma = \\kappa$ or  0). Finally  we discuss the possibilities  of real detections of this phenomenon in the near future. ",
        "conclusions": "We analyzed the shifts in the center-of-light positions of gravitationaly lensed quasar images and the corresponding flux variability due to the microlensing effect of stars in the lensing galaxy. We found the following results:  \\begin{enumerate}  \\item The center-of-light of multiply imaged quasars is expected to vary on time scales of months to years, due to astrometric microlensing effects.  \\item This effect    depends on the size of the lensed quasar, which we modelled with a Gaussian profile: it is larger for smaller sources. This means that measuring centroid shifts will allow  us to constrain the size of the quasar.  \\item We studied eight cases, four without external shear ($\\kappa = 0.2, 0.4, 0.6, 0.8$; $\\gamma = 0.0$) and four with shear ($\\kappa = \\gamma = 0.2, 0.4, 0.6, 0.8$). The strongest effects are seen in the cases $\\kappa = \\gamma = 0.4$ and 0.6, which in fact are very typical parameter values for many lensed quasar  images.  \\item The effect of center-of-light shift of quasars is not limited to multiply imaged quasars: single quasars may occasionally be microlensed too (as was suggested by Hawkins 1996). The offsets should be relatively small, however, due to the fact that the values of $\\kappa$ and $\\gamma$ will be lower for single quasars than for multiply imaged ones. Our simulation with $\\kappa = 0.2$ comes closest to describing these cases with very low convergence.  \\item The centroid shift effect is statistically strongly correlated with changes in the apparent brightness of the quasar. Both occur during the crossing of a (micro-)caustic: photometric fluctuations $\\ge 0.5$ mag produce offsets $\\ge 2 \\theta_E$ in 50\\% of the cases, and $\\ge 5 \\theta_E$ in 5\\% of the cases. These values correspond to about 15 and 35 microarcseconds, respectively, for the lensed quasar Q2237+0305.  \\item The best strategy for measuring centroid shifts is therefore to monitor the apparent brightness of multiple quasars with ground based telescopes and to determine the relative positions of the various quasar images occasionally, ideally before, during and after a high magnification event. \\item In the cases with external shear, the center-of-light shift depends strongly on the direction of the shear, which can hence be determined in a statistical sense.    \\end{enumerate} With the next generation of astrometric instruments providing an accuracy of order 10 microarcseconds, the astrometric microlensing effect of stars acting on background quasars will become detectable."
    },
    "0311/astro-ph0311205_arXiv.txt": {
        "abstract": "We present a snapshot \\emph{Hubble Space Telescope (HST)} image of the galaxy cluster A\\,1201 ($z{=}0.169$), revealing a tangential arc 2\\,arcsec from the brightest cluster galaxy (BCG).  Keck--ESI spectroscopy confirms that the arc is gravitational in nature and that the source galaxy lies at $z{=}0.451$.  We construct a model of the gravitational potential of the cluster that faithfully reproduces the observed arc morphology.  Despite the relaxed appearance of the cluster in the \\emph{HST} frame, the best fit ellipticity of the total matter distribution is $\\epsilon_{\\rm total}{\\ge}0.5$, in contrast to the light distribution of the BCG ($\\epsilon_{\\rm BCG}{=}0.23{\\pm}0.03$) on $2''$ scales.  Further deep optical observations and pointed X--ray spectro--imaging observations with \\emph{Chandra} are required to determine whether this elongation is due to a single elongated dark matter halo, or a more complex distribution of matter in the cluster core.  We compare the arc with a sample drawn from the published literature, and confirm that it is unique among tangential systems in the small physical scales that it probes (${\\sim}6$\\,kpc).  In anticipation of a more thorough investigation of this cluster across a broad range of physical scales, we use our fiducial lens model to estimate the projected mass and mass--to--light ratio of the cluster within a radius of 6\\,kpc, obtaining: $M{=}(5.9^{{+}0.9}_{{-}0.7}){\\times}10^{11}{\\rm M_\\odot}$, $M/L_V{=}9.4^{{+}2.4}_{{-}2.1}(M/L)_\\odot$.  Overall our results confirm the importance of \\emph{HST} snapshot surveys for identifying rare lensing constraints on cluster mass distributions.  In combination with follow--up optical and X--ray observations, the arc in A\\,1201 should help to increase our understanding of the physics of cluster cores. ",
        "introduction": "Galaxy clusters are important laboratories in which to study physical processes that are generally inaccessible in other environments.  For example the radial density profile and the projected ellipticity of clusters on the sky may offer valuable clues into the nature of dark matter (e.g.,\\ Spergel \\& Steinhardt 2001; Sand et al.\\ 2002; 2003; Miralda-Escud\\'e 2002; Arabadjis et al.\\ 2002).  Complications often arise in cluster--based studies of dark matter due to the presence of baryons (e.g.,\\ Allen 1998; Smith et al.\\ 2001; Lewis et al.\\ 2003). However, from a broader perspective such complications provide us with important clues into the physics of gas cooling, and inter--play between baryons and dark matter, both of which are central to attempts to understand the physics of galaxy formation (e.g.,\\ Cole et al.\\ 2000).  Progress towards these goals requires detailed study of the distribution of mass in clusters.  Strong gravitational lensing offers a direct and precise probe of cluster mass distributions (e.g.,\\ Kneib et al.\\ 1996; Smith 2002; Smith et al.\\ 2003).  Complementary constraints can also be obtained from X--ray observations (e.g.,\\ Allen et al.\\ 2002), weak lensing (e.g.,\\ Kneib et al.\\ 2003) and the three--dimensional distribution of cluster galaxies (e.g.,\\ Czoske et al.\\ 2002).  A combination of these techniques is necessary for a comprehensive understanding of mass in clusters.  Armed with the results from such multi--wavelength studies, robust constraints on the dark matter particle and gas cooling may ultimately flow.  We have conducted a snapshot survey of 55 X--ray luminous galaxy clusters with the WFPC2 camera on--board \\emph{HST} (PID's 8301 \\& 8719; PI Edge).  A key goal of this survey is to uncover new cluster lenses with which to explore the questions outlined above.  The snapshot observing strategy is well--suited to identifying clusters containing rare and powerful constraints such as radial arcs (e.g.,\\ RXJ\\,1133 -- Sand et al.\\ 2003).  In this letter we present an \\emph{HST} observation of A\\,1201 ($z{=}0.169$; $(\\alpha,\\delta){=}11^h13^m01.1^s{+}13^\\circ25'40''$~[J2000]; $L_X{=}(3.7{\\pm}0.8){\\times}10^{44}$\\,erg/s [0.1--2.4\\,keV] -- Ebeling et al.\\ 1998).  These data reveal a tangential arc $2''$ from the optical centroid of the BCG.  Spectroscopic observations at the Keck Observatory confirm that the arc is a gravitationally--lensed galaxy at $z{=}0.451$ (\\S2).  The best--fit gravitational lens model faithfully reproduces the arc morphology, however the total matter distribution of this model appears to be much more elongated than the optical isophotes of the BCG (\\S3).  We also discuss the uniqueness of this small--scale gravitational arc by comparing it with other known cluster lenses, and outline how follow--up optical and X--ray observations will help to fully exploit this powerful new constraint on the matter distribution in cluster cores (\\S4).  We assume $H_0{=}65\\,{\\rm km~s^{{-}1}Mpc^{{-}1}}$, $\\Omega_{\\rm M}{=}0.3$ and $\\Omega_\\Lambda{=}0.7$.  In this cosmology, $1''{\\equiv}3.1{\\rm kpc}$ at $z{=}0.169$ and $1''{\\equiv}6.2{\\rm kpc}$ at $z{=}0.451$. ",
        "conclusions": "\\subsection{Is A\\,1201 Unique Among Cluster Lenses?}  A\\,1201 is drawn from a sample of $55$ X--ray luminous clusters observed with \\emph{HST} in the Edge et al.\\ (PIDs: 8301 \\& 8719) snapshot survey of BCGs.  This snapshot survey doubles the number of clusters that have been observed to date with either the WFPC2 or ACS cameras.  Among these ${\\sim}100$ clusters, A\\,1201 is the only system with a tangential arc on scales as small as $R{=}2''$.  This rarity underlines the importance and efficiency of snapshot surveys with \\emph{HST} to discover such small--scale probes of the mass distribution in clusters.  We also investigate the uniqueness of the tangential arc in A\\,1201 among multiple--image systems in spectroscopically confirmed cluster lenses.  The deflection angle of a  gravitational lens depends on the angular diameter distance ratio $D_{\\rm LS}/D_{\\rm OS}$ where $D_{\\rm LS}$ is the distance from the lens to the source and $D_{\\rm OS}$ is the distance from the observer to the source.  We plot the distribution of distance ratios for known multiple--image systems in Fig.~\\ref{hist}, based on an extensive review of the \\emph{HST} archive and the published literature (see Sand et al.\\ 2004, in prep.\\ for more details).  The cluster sample upon which this histogram is based is heterogeneous.  We therefore also plot (as the dashed histogram) the distribution of distance ratios for the multiple--image systems identified thus far in a well--defined sample of X--ray luminous clusters at $z{=}0.21{\\pm}0.04$ by Smith (2002; see also Smith et al.\\ 2001, 2002; Kneib et al.\\ 2003, in prep.; Smith et al.\\ 2003, in prep.).  A\\,1201 lies at the lower envelope of both multiple--image samples, with a value of $D_{\\rm LS}/D_{\\rm OS}{=}0.597$. From a lens geometry perspective, A\\,1201 is therefore unusual but not unique among the multiple--image systems discovered to date. Other low $D_{\\rm LS}/D_{\\rm OS}$ systems include famous lensing clusters, for example the giant arc ($z{=}0.724$; Soucail et al.\\ 1988) in A\\,370 ($z{=}0.370$).  However, the higher redshift of this, and other clusters with low distance ratios renders the observed deflection angle (${\\sim}10$--$20''$) and the physical scales probed (${\\sim}60$--120\\,kpc) much larger than that relevant to A\\,1201 ($R{=}6$kpc).  A\\,1201 is therefore unique in the small physical scales probed by its tangential arc.  \\begin{figure} \\medskip \\centerline{ \\psfig{file=f2.ps,width=70mm} } \\caption{The distribution of $D_{\\rm LS}/D_{\\rm OS}$ for observed multiple--image systems, drawn from an extensive search of the \\emph{HST} archive and the published literature (Sand et al.\\ 2004, in prep.).  The solid histogram shows the full sample of 41 multiple--image systems; the dashed histogram shows a sample drawn from a well--defined sample of X--ray luminous clusters (Smith 2002); the filled histogram marks the location of A\\,1201.} \\label{hist} \\end{figure}  \\subsection{Summary and Outlook}  \\emph{HST} snapshot imaging of A\\,1201 ($z{=}0.169$) with the WFPC2 camera reveals a tangential arc 2--arcsec from the center of this cluster.  Spectroscopy obtained with ESI on the Keck--II telescope confirms the gravitational nature of the arc, and places the source galaxy at $z{=}0.451$.  We construct a gravitational lens model that is able to reproduce the observed arc morphology.  The key feature of this model is that the total matter distribution is significantly more elongated ($\\epsilon_{\\rm total}{\\ge}0.5$) than the light distribution of the BCG on $2''$ scales ($\\epsilon_{\\rm BCG}{=}0.23{\\pm}0.03$). With the current data we are unable to determine whether the matter distribution really is more elongated than the stellar distribution, or if there is a significant amount of mass whose center of mass does not coincide with the BCG. This could indicate that the cluster is dynamically less mature than the optical data suggest.  The proximity of the arc to the center of this cluster is unique among cluster lenses, and provides an important constraint on the mass of the cluster on very small physical scales.  We measure the projected mass within the tangential arc to be $M(R{\\le}2''){=}(5.9^{{+}0.9}_{{-}0.7}){\\times}10^{11}{\\rm M_\\odot}$, and the $V$--band mass--to--light ratio to be $M/L(R{\\le}2''){=}9.4^{{+}2.4}_{{-}2.1}({\\rm M/L})_\\odot$, the angular scale of $2''$ corresponding to a physical scale of 6\\,kpc.  This constraint, in conjunction with complementary high--resolution space--based data from \\emph{Chandra} and multi--color follow--up with \\emph{HST}/ACS will lead to substantial progress in understanding the distribution of mass in this cluster.  Extending the unique small--scale (6\\,kpc) mass constraint out to larger scales (${\\sim}50$--500\\,kpc) also promises an important role for A\\,1201 in the quest to understand the physical processes at play in galaxy cluster cores."
    },
    "0311/astro-ph0311496_arXiv.txt": {
        "abstract": "To address the problem concerning the early formation of stars in massive galaxies, we present the results of a semi-analytic model of galaxy formation which includes a physical description of starbursts triggered by galaxy interactions. These originate from the destabilization of cold galactic gas occurring in galaxy encounters, which in part feeds the accretion onto black holes powering quasars, and in part drives circumnuclear starsbursts at redshifts $z\\approx 2-4$, preferentially in massive objects. This speeds up the formation of stars in massive galaxies at high redshifts without altering it in low mass galactic halos. Thus, at intermediate $z\\approx 1.5-2$ we find that a considerable fraction of the stellar content of massive galaxies is already in place, at variance with the predictions of previous hierarchical models. The resulting high-$z$ star formation rate and B-band luminosity functions, and the luminosity and redshift distribution of galaxies in K-band at $z\\lesssim 2$ are all in good agreement with the existing observations concerning the bright galaxy population. ",
        "introduction": "To describe the galaxy evolution in the hierarchical scenario, we adopt the SAM described in detail in Menci et al. (2002); here we recall the basic points. We consider both the host DM halos containing the galaxies (i.e., groups and clusters of galaxies with mass $M$, virial radius $R$ and circular velocity $V$), and the DM clumps (with mass $m$, radius $r$ and circular velocity $v$) associated with the individual member galaxies. The former grow hierarchically to larger sizes through repeated merging events (at the rate given in Lacey \\& Cole 1993), while the latter may coalesce either with the central galaxy in the common halo due to dynamical friction, or with other satellite galaxies through binary aggregations. The timescale for the dynamical friction and the binary merging processes, and so the probability for such processes to occur in each timestep, are given by eq. (2) and (4) in Menci et al. (2002).  We assume initially (at $z\\approx 10$) one galaxy in each host structure, with the latter following the Press \\& Schechter mass distribution. The probability for the merging processes (dynamical friction and binary aggregations) to occur during the hierarchical growth of the hosting structure yields the differential distribution function $N(v,V,t)$ of galaxies with given $v$ in hosts with circular velocity $V$ at the cosmic time $t$. From $N(v,V,t)$ we derive the number $N_T(V,t)$ of galaxies in a host halo (membership), and the overall distribution of galaxy circular velocity $N(v,t)$ irrespective of the host.  The properties of the baryons (gas and stars) contained in the galactic DM clumps are computed as follows. Starting from an initial amount $m\\,\\Omega_b/\\Omega$ of gas at the virial temperature of the galactic halos, we compute the mass $m_c$ of cold baryons within the cooling radius. The disk associated to the cold baryons will have a radius $r_d(v)$, rotation velocity $v_d(v)$, and dynamical time $t_d = r_d/v_d$, all computed after Mo, Mao \\& White (1998). From such a cold phase, stars are allowed to form at the rate \\begin{equation} \\dot m_* = {m_c\\over t_{dyn}}\\,\\Big({v\\over 200\\,{\\rm km\\,s^{-1}} }\\Big)^{-\\alpha_*} \\end{equation}  Finally, a mass $\\Delta m_h=\\beta\\,m_*$ is returned from the cool to the hot gas phase due to the energy fed back by canonical type II Supernovae associated to $m_*$; the feedback efficiency is taken to be $\\beta= (v/v_h)^{\\alpha_h}$. The values adopted for the parameters $\\alpha_*=-1.5$, $\\alpha_h=2$ and $v_h=150$ km/s fit both the local B-band galaxy LF and the Tully-Fisher relation, as illustrated by Menci et al. (2002).  At each merging event, the masses of the different baryonic phases are refueled by those in the merging partner. The further increments $\\Delta m_c$, $\\Delta m_*$, $\\Delta m_h$ from cooling, star formation and feedback are recomputed on iterating the procedure described above.  Thus, for each galactic $v$, the star formation defined by eq. (1) is driven by the cooling rate of the hot gas and by the rate of refueling of cold gas, which in turn is related to the progressive growth of the total galactic mass along the merging tree. The related brightening of galaxies is a gradual process, and the associated star formation history is often referred to as ''quiescent'' star formation.  The integrated stellar emission $S_{\\lambda}(v,t)$ at the wavelength $\\lambda$ is computed by convolving the SFR with the spectral energy distribution $\\phi_{\\lambda}$ obtained from population synthesis models: \\begin{equation}\\label{sed} S_{\\lambda}(v,t) = \\int_0^t\\,dt'\\,\\phi_{\\lambda}(t-t')\\,\\dot m_*(v,t')~. \\end{equation} In the following we adopt $\\phi_{\\lambda}$ taken from Bruzual \\& Charlot (1993), with a Salpeter IMF. The dust extinction is computed as described in Menci et al. (2002).  All computations are made in a $\\Lambda$-CDM cosmology with $\\Omega_0=0.3$, $\\Omega_{\\lambda}=0.7$, a baryon fraction $\\Omega_b=0.03$, and Hubble constant $h=0.7$ in units of 100 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "We have presented the results of a SAM which includes a physical model for starbursts. The latter originate from the destabilization of cold galactic gas occurring in galaxy encounters as described by Cavaliere \\& Vittorini (2000). In this picture, part of the  destabilized gas feeds the accretion onto BHs powering the QSOs (see Menci et al. 2003), while the complementary, larger fraction produces circumnuclear  starsbursts at redshifts $z\\approx 2-4$, preferentially in massive objects. This speeds  up the star formation in massive galaxies (favored by their larger cross sections for encounters) at high redshifts (when the higher densities and the slower relative velocities of galaxies in common halos favor strong  interactions). The LFs resulting in our model match those of {\\it luminous} ($M_B<-21$) Lyman-break galaxies observed at $z=3$ and 4, while leaving almost unchanged the model predictions for fainter objects. Thus, at $z\\approx 2$, a larger fraction of the stellar content of the massive galaxies is already {\\it in place}, at variance with the predictions of previous hierarchical models; the resulting K-band LFs and $z$- distributions match the existing data.  Compared to Somerville, Primack, \\& Faber (2001) our model gives qualitatively similar results, but a better fit (see fig. 3) to the observed $z$-distribution of bright $K>20$ galaxies at $z\\gtrsim 1.5$. Such a larger rate is mainly due to the novel inclusion of fly-by events which, although producing starbursts with lower efficiency ($f\\sim 0.1-0.2$ at $z\\approx 3$), are more frequent than the bound merging events. In our picture, the UV luminosity, and hence the star formation rate, of massive galaxies at high redshift is in part contributed by the few powerful starbursts occurring in major merging events, and in part by a more widespread, though less powerful, brightening of galaxies almost continuously stimulated by the frequent encounters. In this regime, a considerable fraction close to $0.25$ of the present stellar content of massive ($10^{10}\\,M_{\\odot}<m_*<10^{11}\\,M_{\\odot}$) galaxies is formed by $z=2$, see fig. 4. As a result, at lower redshifts $z\\approx 1.5$ a considerable fraction of massive galaxies contains a stellar population evolved enough to provide the bright $K$-band luminosities required to match the observed LFs and counts.  In sum, we have shown that making early massive galaxies takes not only the early collapse of large DM halos favored in the $\\Lambda$-CDM cosmology, but also baryons forced to shine starlight by galaxy interactions.  We stress that our results are achieved on the basis of a physical -- rather than phenomenological -- model to derive the burst rate and the amount of cold gas converted into stars during the bursts. Moreover, our model naturally connects with the accretion onto BHs and with the corresponding QSOs emission. Such a unified description is thus able to {\\it link} the independent observations concerning the luminosity distributions of QSOs and of the galaxies.  For $z<2$ the merging and the encounter rates decline, and no longer affect the evolution of the cold gas reservoirs of massive galaxies.  So, these galaxies specifically enter a phase of nearly passive evolution, with redder colors which in some cases -- whose probability we shall give in detail elsewhere -- correspond to those of many observed extremely red objects (EROs see, e.g., Daddi et al. 2002 and references therein). In our model, the beginning and the development of such a phase is naturally matched to the dramatic drop of the QSO luminosities; these are entirely triggered by the gas destabilized in the encounters, and so are particularly sensitive to the current merging and interaction rates of their host galaxies."
    },
    "0311/astro-ph0311425_arXiv.txt": {
        "abstract": "Here we report on combined X-ray and radio observations of SN 2002ap with XMM-Newton ToO observation and GMRT observations aided with VLA published results. In deriving the X-ray flux of SN 2002ap we account for the contribution of a nearby source, found to be present in the pre-SN explosion images obtained with Chandra observatory. We also derive upper limits on mass loss rate from X-ray and radio data. We suggest that the prompt X-ray emission is non-thermal in nature and its is due to the repeated compton boosting of optical photons. We also compare SN's early radiospheric properties with two other SNe at the same epoch. \\vspace{1pc} ",
        "introduction": "SN 2002ap was a type Ic SN discovered on Jan 29.4, 2002 in NGC 628. The explosion energy of the SN was larger than usual type Ic SNe (4-10$\\,\\times 10^{51}$ ergs). SN 2002ap was of great interest because it was one of the closest extragalactic SN (D=7.3 Mpc); secondly it was a Type Ic SN and GRB association of type Ic is found in SN 1998bw (with GRB 980425) and SN 2003dh (with GRB 030329). SN 2002ap  has shown SN 1998bw like features like high velocity in early spectroscopy of optical data. Study of type Ic SNe is interesting because they are devoid of H and He core and hence can probe closer to the  central engine. ",
        "conclusions": "We suggest that repeated scattering of optical photons can account for the prompt X-ray emission from SN 2002ap. X-ray sphere (7$\\times\\,10^{14}$ cm) lies between optical (3.5$\\times\\,10^{14}$)and the radio photosphere (4$\\times\\,10^{15}$ cm). Both Wolf-Rayet and interacting binary case BB are capable of providing the adequate \uf074 optical depth and are viable scenarios for the progenitor star. The properties of different kind of SNe differ significantly from each other. In fact the same type of SNe can show the significantly different properties.\\\\ \\newline {\\bf Acknowledgment} We thank staff of XMM-Newton and GMRT (NCRA-TIFR) that made these observations possible."
    },
    "0311/astro-ph0311339_arXiv.txt": {
        "abstract": "This is the first paper about the fragmentation and mass outflow in the molecular cloud by using  three-dimensional MHD nested-grid simulations. The binary star formation process is studied paying particular attention to the fragmentation of a rotating magnetized molecular cloud. We assume an isothermal rotating and magnetized cylindrical cloud in hydrostatic balance. Non-axisymmetric as well as axisymmetric perturbations are added to the initial state and the subsequent evolutions are studied. The evolution is characterized by three parameters: the amplitude of the non-axisymmetric perturbations, the rotation speed, and the magnetic field strength. As a result, it is found that non-axisymmetry hardly evolves in the early phase, but begins to grow after the gas contracts and forms a thin disk. Disk formation is strongly promoted by the rotation speed and the magnetic field strength. There are two types of fragmentation: fragmentation from a ring and that from a bar. Thin adiabatic cores fragment if a thickness is smaller than 1/4 of the radius. For the fragments to survive, they should be formed in a heavily elongated barred core or a flat round disk. In the models showing fragmentation, outflows from respective fragments are found as well as those driven by the rotating bar or the disk. ",
        "introduction": "It is important to understand how a binary star system forms from a molecular gas cloud because most main-sequence stars are in binary or multiple systems \\citep[e.g.][]{abt83,duquennoy91}. Furthermore, recent observations indicate that companions may be even more common for pre-main sequence stars \\citep{richichi94}. Numerical simulations show that it is difficult for single stars to capture another one and to form a binary system in the main sequence phase \\citep{kroupa01}. Accordingly, most stars seem to be born as binaries rather than single stars in a gas cloud, and molecular bipolar outflows are often observed in star forming regions. It has been shown using two-dimensional magnetohydrodynamical simulations that outflow plays an important role in the star formation process because the excess angular momentum, which is conserved from the host cloud, is removed by these outflows from the formed core \\citep{tomisaka00, tomisaka02}. These outflows are caused by the twisted magnetic field lines made by the rotation near the adiabatic core. In order to investigate the cloud fragmentation and later binary star formation, three-dimensional calculation is required. In this letter, we discuss fragmentation in the course of star formation which leads to binary star formation.  Fragmentation of isothermal cloud is studied by many authors using three-dimensional hydrodynamical simulations \\citep[e.g.][]{bodenheimer00}. The criterion for an initially uniform cloud to fragment is described by $\\alpha_0$ (thermal-to-gravitational energy ratio) of less than $\\simeq 0.2-0.5$ if we restrict ourselves to the isothermal gas. The cloud with $\\alpha_0 \\ga 0.2 -0.5$ experiences the run-away collapse \\citep{larson69,norman80}. When the gas density increases up to $n \\, \\ga \\, 10^{10}$~cm$^{-3}$, the isothermal gas approximation breaks down. Recently, \\citet{bate02,bate03}, \\citet{cha03} and \\citet{matsu03b} have studied the cloud evolution from isothermal to adiabatic state using a barotropic equation of state and they clarified the evolution of the cloud with no magnetic field. However, there are few papers on the fragmentation of the magnetized cloud. \\citet{boss02} studied the fragmentation taking account of radiative transfer with the Eddington approximation instead of barotropic equation of state, but simplified the effect of magnetic field. In his studies, the angular momentum transfer by the magnetic field is ignored. However, the angular momentum is largely removed by magnetic braking and molecular outflow in the cloud evolution process \\citep{tomisaka00}. Therefore, to explore the fragmentation of the magnetized cloud, full three-dimensional magnetohydrodynamical MHD simulations are necessary.  Here, we employ the nested grid MHD code, which always maintains sufficient spatial resolution in the central region. The nested grid method is powerful to study the star formation process, because this requires a large dynamic range in spatial dimensions. In the simulations, we use the barotropic equation of state instead of inclusion of radiative transfer. In this study, we calculate the non-axisymmetric dynamical contraction of the cloud from $4\\times10^2$  cm$^{-3}$ to $1 \\times10^{17}$  cm$^{-3}$ in number density, and  investigate how fragmentation proceeds and outflow occurs in the course of binary star formation. ",
        "conclusions": "As mentioned in \\S~\\ref{sec:results}, modes of fragmentation and shapes of the outflow depend on the growth of non-axisymmetry. This shows that the cloud evolution can be classified by this growth. In order to characterize the shape of the central region, we define oblateness ($\\varepsilon_{\\rm{ob}} $) and axis ratio ($\\varepsilon_{\\rm{ar}} $) as follows: \\begin{eqnarray} \\varepsilon_{\\rm{ob}}  &\\equiv& \\sqrt{h_l\\cdot h_s }/h_z , \\\\ \\varepsilon_{\\rm{ar}}  &\\equiv& h_l/h_s , \\end{eqnarray} where $h_l$, $h_s$ and $h_z$ mean the length of the long-, short-, and $z$-axis for the gas with $\\rho>(1/10) \\rho_{\\rm c}$, respectively \\citep[see][for definition of $h$]{matsu99}.  In Fig.~\\ref{fig:2}, the oblateness and the axis ratio of the central region are plotted for all models at the beginning of the accretion phase ($\\rho_c =\\rho_{\\rm{cri}}=10^{10} \\ \\rm{cm^{-3}}$). Curves in this figure illustrate the evolutionary tracks for four different models ($A_{\\phi}, \\alpha, \\omega$)= (0.01, 0.01, 0.5),  (0.2, 0.1, 0.5), (0.2, 1, 0.5), and (0.2, 5, 0.5). The clouds evolve from the lower-left corner, which means the early shape is almost spherical, to the upper-right region via the upper-left region. First, the clouds evolve vertically upward. This indicates that the oblateness increases. Then, the curves change their directions to the right for three models with $A_{\\phi} =0.2$. Evolutionary tracks of the high-density region in Fig.~\\ref{fig:2} show very clearly that the non-axisymmetric perturbation grows only after the disk is formed. That is, the axis ratio increases after the oblateness grows considerably large ($\\varepsilon_{\\rm{\\rm{ob}} } \\ge4$). Because the gas is partially supported in the radial direction by the magnetic pressure and the centrifugal force, the radial contraction is more delayed than that of the $z$-direction. As a result, the disk is formed ($\\varepsilon_{\\rm{ob}}  > 4$) earlier for models with larger $\\alpha$ and $\\omega$. In such models, the non-axisymmetry grows sufficiently and forms a large bar in the isothermal collapse phase. However, the non-axisymmetry hardly evolves in the model with small $\\alpha$ and $\\omega$, because it takes a long time for disk formation.  Adiabatic cores are classified into three types: $core$ (Fig.\\ref{fig:1}a: $\\varepsilon_{\\rm{ob}} <4$ and $\\varepsilon_{\\rm{ar}} <4$), $disk$ (Fig.\\ref{fig:1}b$:  \\varepsilon_{\\rm{ob}} >4$  and $\\varepsilon_{\\rm{ar}} <4$), and $bar$ (Fig.\\ref{fig:1}c: $ \\varepsilon_{\\rm{ob}} >4$ and $\\varepsilon_{\\rm{ar}} >4$). Some $disks$ evolve into the rings, in which the central density is lower than that outside, even if it has the same oblateness and the axis ratio. Models are divided into three kinds of evolution: thin-line symbols mean that the models form cores which continue to collapse, while thick-line and filled symbols represent the models in which fragmentation occurs. Although the fragments survive against mutual merger in the models of thick-line symbols, they merge each other in the models of filled symbols.  Since the models with thick-line or filled symbol are distributed only in the region $\\varepsilon_{\\rm{ob}} >4$, it is concluded that the necessary condition for fragmentation is $\\varepsilon_{\\rm{ob}} >4$, which means that a thin disk or bar is necessary for fragmentation. Even in the region of $\\varepsilon_{\\rm ob}>4$, there are exceptional three models in which the adiabatic core evolves to form an almost axisymmetric thin disk or ring because non-axisymmetry is counteracted in the isothermal collapse phase. The shapes of the $bar$ and $ring$ are also seen in \\citet{boss02} and \\citet{li02}. It is concluded that for the adiabatic core to fragment, the central gas needs to form a sufficiently thin disk or bar in the isothermal collapse phase. The thick-line symbols are distributed in two regions enclosed by circles. The fragmented cores in the circles seem to evolve into a binary star because the length of the semi-major axis expands or oscillates. Models in which the fragments merge with each other are distributed outside the domain. Outside of the domain, a few fragments survive at the end of the calculation but the separation between the two cores shrinks with time. These will not grow into binary stars.  In conclusion, if the following criteria are fulfilled at the end of the isothermal collapse phase, fragmentation occurs and fragments survive to form binary or multiple stars. \\begin{enumerate} \\item $\\varepsilon_{\\rm{ob}} >4$ : fragmentation condition. \\item $0<\\varepsilon_{\\rm{ar}} <2$, or  $\\varepsilon_{\\rm{ar}} >10$ : survival condition against merger. \\end{enumerate} The symbol in this figure indicates the shape of the adiabatic core when the fragmentation occurs or the calculation ends. Fig.~\\ref{fig:2} confirms the idea that the fragmentation occurs from a ring ($\\bigcirc$) or bar ($\\bigtriangleup$) and not from a core ($\\sq$).  We showed the domain which may evolve into binary stars in Fig.~\\ref{fig:2}. Outside the binary-forming domain, the core formation model ($\\sq$) results in a compact massive core at the center of cloud. The bar-fragmentation model ($\\bigtriangleup$) in which the mutual merger of the fragments occurred also leads to a disk-like core at the center of cloud. This configuration seems to lead to single star formation."
    },
    "0311/astro-ph0311563_arXiv.txt": {
        "abstract": " ",
        "introduction": "The inflationary paradigm in which the inflaton sector is responsible for the density perturbations is economical.  One potential serves many purposes: driving inflation, ending inflation, and generating the observed spectrum of density perturbations.  The resulting models of inflation are predictive, and falsifiable by experiment.  The flip side of the coin, though, is that they are restrictive too, and it has proven extremely hard to build realistic models of inflation. Often a considerable amount of fine-tuning is needed for the model to satisfy all constraints. For this reason it may be worthwhile to explore alternatives.  Inflationary models can be ``liberated'' if its task list is reduced~\\cite{liberation}. This is the idea behind both the curvaton scenario~\\cite{early,curv} and the inhomogeneous reheating scenario~\\cite{FDR}, in which not the inflaton field but some other field is responsible for the density perturbations.  The inflaton sector merely serves to drive and end inflation, and is considerably less constraint.  Of course, the price to be paid is that a new field has to be introduced into the theory.  But the costs are minimized if the new field and/or the new scales introduced are already present in our models of particle physics.  There is already an extensive literature exploring this possibility in the context of the curvature scenario~\\cite{curv_ext,curv_MSSM}, but little has been done in the context of the inhomogeneous reheating scenario (but see~\\cite{FM,anupam}).  In this paper we will concentrate on the inhomogeneous reheating scenario.  In particular, we will address the following two questions. Is low scale inflation --- with the Hubble scale during inflation of the order of the gravitino mass --- possible, such that the inflaton sector can be naturally identified with SUSY breaking sector?  Can any of the flat directions in the minimal supersymmetric standard model (MSSM) be responsible for the density fluctuations?  Low scale inflation is hard to realize in the conventional setting in which the inflaton is responsible for the density fluctuations.~\\footnote {The exception are ``new inflation'' models, in which the inflaton starts out close to the origin \\cite{sarkar}.  When Taylor expanded around the inflaton VEV, the linear term is anomolous small, making this different behaviour possible.} In that case the density fluctuations are proportional to $H_*/\\mpl$, with the subscript $*$ denoting the quantity evaluated at the time observable scales leave the horizon, some 60 e-folds before the end of inflation. Unless the slow roll parameter $\\epsilon= (\\mpl V'/V)^2 \\ll 1$, which often requires some amount of fine tuning, high scale inflation is needed to get the observed size of density perturbations. In the curvaton scenario the density perturbations are proportional to $H_* / \\sigma_*$ with $\\sigma$ the vacuum expectation value (VEV) of the curvaton field.  Since the VEV can be much lower than the Planck scale, low scale inflation is more natural in the curvaton scenario. However, in~\\cite{lyth} it was shown that for the scenario to work the scale of inflation cannot be too low: $H_* \\gtrsim 10^7 \\GeV$.  In the inhomogeneous reheating scenario the density perturbations are proportional to $H / S_*$ or $H / M$, depending on the form of the decay rate, where $S$ is the VEV of the flat direction field responsible for the density perturbations and $M$ is some cutoff scale.  Since both $M$ and $S_*$ can be much smaller than the Planck scale, also here, low scale inflation seems natural.  We will see, that indeed, low scale inflation with $H_*$ of the order of the gravitino mass is possible in this context, although only in a small part of parameter space.  The fluctuating decay rate scenario can then be married with any of the interesting inflaton models, in which the inflaton originates from the SUSY breaking sector~\\cite{low_inflation}.  In both the curvaton and inhomogeneous reheating scenario a scalar field other than the inflaton is responsible for the density perturbations. This scalar has to be light compared with the scale of inflation to be able to fluctuate freely, and produce the scale invariant perturbation spectrum observed.  Obvious candidates for this scalar field are the MSSM flat directions.  Within the curvaton scenario this possibility has been studied in~\\cite{curv_MSSM}.  It was found that only under very special conditions can the MSSM scalar produce the density perturbations.  The main problem is that the curvaton has to (nearly) dominate the energy density before decay. This means large initial VEVs and thus potential problems with non-renormalizable operators, and this means late decay and thus potential problems with disastrous finite temperature effects. In this respect, the inhomogeneous reheating scenario offers better prospects, as the MSSM flat direction does not need to dominate the energy density at the time of decay.  Consequently, its VEV can be small enough to avoid problems with the non-renormalizable potential, and the density perturbations may be generated before thermal effects become important.  This paper is organized as follows.  In the next section, we provide the back ground material.  We describe those features of the inflaton and MSSM sector important for the inhomogeneous reheating scenario. In section \\ref{S_constraints} we turn to the description of the various constraints on the inhomogeneous reheating scenario.  Section \\ref{S_scales} discusses the various time scales in the problem, giving further insights in the nature of especially the thermal constraints. We address the prospects for model building, concentrating on the possibility of low scale inflation and using MSSM flat directions, for a polynomial decay rate in section \\ref{S_pol}, while section \\ref{S_const} discusses the same issues in the context of an (approximate) constant decay rate.  In section \\ref{S_psi} we discuss the varying mass scenario, which is a variation of the inhomogeneous reheating scenario, and point out its virtues and drawbacks for model building.  Finally, we conclude in section \\ref{S_concl}. ",
        "conclusions": "\\label{S_concl}  In the inhomogeneous reheating scenario not the inflaton but some other field is responsible for the observed density perturbations.  It would be economical if the new fields and scales introduced in this scenario could be identified with the fields and scales appearing in our models of particle physics.  In particular, we have addressed the following two questions.  Is low scale inflation possible, with the Hubble scale of the order of the gravitino, such that the inflaton sector can be naturally identified with the SUSY breaking sector?  Can any of the MSSM flat directions be responsible for the density fluctuations?   We discussed various decay rates, obtained from non-renormalizable couplings, renormalizable couplings, and from phase space effects. For the last two examples, the observed density perturbations can only be obtained if the Hubble constant, the flaton VEV, and the cutoff scale all lie within 4 decades.  This requires some tuning, especially since Gaussianity of the perturbations is only assured for $S_* \\gtrsim 10 H_*$.  For the phase space decay rate the cutoff is not a fundamental scale in the theory. This has the advantage that there is no need to explain the origin of this scale.  The disadvantage is that this model is quite constraint, and for example, prompt decay is excluded.  For the non-renormalizable decay rate there is more freedom, since the density perturbations do not depend on the cutoff.   After inflation, already before inflaton decay has completed, there is a dilute plasma.  Fields coupling to the flat direction field, whether they are in equilibrium with the thermal bath or not, will induce a thermal mass for $S$. Such thermal corrections will lead to early induced oscillations of the flat direction field, which are generically fatal for the inhomogeneous reheating scenario.  The thermal effects can be avoided if all flat direction couplings are small, which excludes MSSM flat directions.  For a polynomial decay rate, which can originate from non-renormalizable couplings, in addition the flaton mass has to be much smaller than the Hubble scale during inflation.  The thermal constraints can be trivially avoided if the inflaton decays promptly at the end of inflation, so that there is no time for thermal effects to act.  In this case the flat direction mass and couplings are unconstrained, except for the requirement that $m_S \\lesssim 0.1 H_*$.  The reheat temperature is high, and there is a potential gravitino problem.  Prompt decay however, is only consistent with a Gaussian perturbation spectrum for a small window of inflaton masses and couplings.  The constraints are stronger for a constant decay rate. In particular, decay rates which inherit their $S$ dependence from phase space effects are incompatible with prompt decay.  Large inflaton mass and couplings are needed, inconsistent with one field models of inflation.  Finally, we discussed a variation of the inhomogeneous reheating scenario, in which not the inflaton field but another field has a fluctuating decay rate.  After inflation, this field first has to come to dominate the energy density, and then decay.  The prospects for model building are bleak.  There is no analog of prompt decay: thermal effects are always there, even before the field comes to dominate the energy density.  Only with small couplings can one avoid the disastrous consequences of thermally induced masses.  This excludes the possibility of identifying either the decaying field or the flat direction field with an MSSM field.  Moreover, there is an extra step in the model, making it less advantageous for low scale inflation."
    },
    "0311/astro-ph0311049_arXiv.txt": {
        "abstract": "We present a kinetic theory for the evolution of the phase-space distribution of dark matter particles in galaxy halos in the presence of a cosmological spectrum of fluctuations.  This theory introduces a new way to model the formation and evolution of halos, which traditionally have been investigated by analytic gravitational infall models or numerical N-body methods.  Unlike the collisionless Boltzmann equation, our kinetic equation contains nonzero terms on the right-hand side arising from stochastic fluctuations in the gravitational potential due to substructures in the dark matter mass distribution. Using statistics for constrained Gaussian random fields in standard cosmological models, we show that our kinetic equation to second-order in perturbation theory is of the Fokker-Planck form, with one scattering term representing drift and the other representing diffusion in velocity-space.  The drift is radial, and the drift and diffusion coefficients depend only on positions and not velocities; our relaxation process in the quasilinear regime is therefore different from the standard two-body relaxation.  We provide explicit expressions relating these coefficients to the linear power spectrum of mass fluctuations and present results for the currently favored cold dark matter model with a nonzero cosmological constant.  Solutions to this kinetic equation will provide a complete description of the cold dark matter spatial and velocity distributions for the average halo during the early phases of galaxy halo formation. ",
        "introduction": "\\label{sec:intro}  Most mass in the universe is in the form of dark matter, and most dark matter is gravitationally clustered in the form of halos.  Dark matter halos are therefore the building blocks of the universe.  Knowledge of the formation and evolution history of dark matter halos is essential for an understanding of the halos' structure and dynamics, as well as the environment that harbors galaxies and clusters.  Designs for dark matter search experiments also rely on the predicted spatial and velocity distributions of dark matter particles in our own galaxy.  According to the theory of structure formation in current cosmological models, dark matter halos arise from tiny primordial fluctuations seeded in the matter and radiation fields in the early universe at, for example, the end of an inflationary era.  These small inhomogeneities grow via gravitational instability, with the slightly overdense regions becoming denser and the slightly underdense regions becoming emptier.  At a given time, there exists a spectrum of density fluctuations over a wide range of length and mass scales, whose spatial distribution is characterized by the power density in Fourier space $P(k,t)$, the power spectrum.  Dark matter halos therefore rarely form and evolve in isolation, growing instead by frequent accretion of smaller-mass halos and occasional major mergers with other halos of comparable mass.  The earliest theoretical insight into cosmological halo formation came from the spherical radial infall model, which considers an isolated spherical overdense volume and describes the secondary infall of bound but initially expanding collisionless matter onto this region \\citep{gg72}.  If the initial density perturbation spectrum is scale-free and the velocity distribution follows the Hubble law, the infall solution is found to approach a self-similar form in an Einstein-de Sitter universe.  The resulting density profile asymptotically approaches $\\rho \\propto r^{-\\alpha}$ with $\\alpha\\ge 2$ \\citep{fg84, bert85}. The exact value of $\\alpha$ depends on the index of the initial perturbation spectrum $P(k)\\propto k^n$: it is isothermal ($\\alpha=2$) for $n\\le 1 $ and for $n>1$ steepens to $\\alpha\\approx 3(n+3)/(n+4)$ \\citep{hs85} or $3(n+3)/(n+5)$ \\citep{sw98,nusser99}.  The collapse of halos therefore retains memory of the initial conditions in the radial infall model.  By contrast, modern cosmological N-body simulations find dark matter halos over a wide range of mass scales to follow a density profile that is nearly independent of the initial power spectrum and is shallower near the halo centers (with $\\alpha\\sim 1$) than that predicted by the radial infall model \\citep{nfw96,nfw97,moore98}. Furthermore, the dark matter in the simulated halos is not entirely smoothly distributed spatially; instead, at least 10\\% of the halo mass is in the form of hundreds to thousands of smaller, dense satellite halos of varying mass \\citep{ghigna98, klypin99, moore99}. Thus even though the analytic similarity solutions derived from one-dimensional infall onto a smooth overdensity have provided valuable insight into the evolution of self-gravitating systems, realistic halo formation is clearly more complicated.  Numerical simulation has been the tool most frequently used in cosmology to study structure formation. It has provided powerful insight and occasionally surprising results (e.g. universal density profile) unexplainable by the simulations themselves.  Our understanding of structure formation can certainly be enhanced if a complementary tool is developed. Complementary approaches are also especially needed at this time to help us better understand and perhaps resolve the ``cuspy halo'' and ``dwarf satellite'' problems of the cold dark matter model.  In this paper we offer a new perspective on the subject of dark matter halo formation by proposing an approach different from both the radial infall model and N-body simulations.  We focus on the phase-space distribution of dark matter in galaxy halos instead of on the motion of a mass shell (as in infall models) or the orbit of a simulation particle (as in N-body).  We attempt to understand the evolution of the phase-space distribution with a statistical description based on kinetic theory.  Our method is to be contrasted with the standard method where the individual dark matter particles are described by the collisionless Boltzmann equation; N-body simulations are in essence a Monte Carlo method for solving this equation. Instead, we consider the phase space of dark matter particles in an average halo and develop a cosmological kinetic theory for halo evolution in the presence of a spectrum of density fluctuations.  The essential features of our approach are represented by the fluctuation-dissipation theorem for time-dependent stochastic processes.  To elaborate on this concept, we recall that stochastic modelling was first developed to explain the Brownian motion of pollen particles suspended in a fluid.  (See \\S~\\ref{sec:brownian} for a more detailed discussion.)  This phenomenon is now understood to arise from a fluctuating force on a Brownian particle due to many rapid collisions with the molecules in the surrounding fluid.  These fluctuations cause successive small changes in the particle velocity, hence to a random walk.  The consequence of many such random walks is diffusion, a form of dissipation.  Due to the large number of molecules, it is impossible to solve the coupled equations of motion for the trajectories of all molecules and the pollen particle in the fluid. Fortunately, a stochastic treatment suffices.  One mathematical prescription is provided by the Langevin equation \\citep{langevin}, which models the forces arising from the fluctuations as a stochastic term that is added to all other deterministic forces in the equation of motion for a particle.  An alternative prescription is given by the Fokker-Planck equation \\citep{fokker, planck}, which is an equation of motion for the probability density or distribution function of particles undergoing random walks.  In this case, the collective changes in the velocities due to the fluctuations in the system are represented by a drift coefficient and a diffusion coefficient arising from the collision terms in the collisional Boltzmann equation.  Stochastic phenomena occur beyond molecular scales.  In astrophysics, the best studied example is the evolution of dense stellar systems such as globular clusters \\citep{chandra, spitzer, bt88}.  A star in a globular cluster experiences a fluctuating force from many gravitational two-body encounters and suffers a succession of small velocity changes.  The fluctuations in these systems arise from granularity in the mass distribution at the locations of the stars. The Fokker-Planck approximation is valid for these weak encounters. We explore this example in greater detail and contrast it with our model of galaxy halo evolution in \\S~\\ref{sec:globular}.  The key feature common to all of these systems is the presence of stochastic fluctuations that result in dissipation. In the cosmological setting addressed by this paper, the fluctuations arise from cosmological perturbations produced in the early universe.  At early times, these perturbations are Gaussian and fully characterized by their power spectrum; at late times, the substructures in halos seen in recent simulations are the nonlinear counterparts of this spectrum of early ripples.  These fluctuations scatter particles.  If one considers an ensemble of halos, the masses and locations of the substructures (or their Gaussian predecessors) are random variables. The substructures thereby introduce ``noise'' into the evolution of dark matter halos.  As we show in this paper, the result is dissipation in the form of drift and velocity-space diffusion.  Recent work has explored the use of the Fokker-Planck equation in cosmological problems.  \\cite{evans97}, for example, searched for steady-state solutions of the Fokker-Planck equation assuming the fluctuations arise from two-body encounters.  It also assumed a simple isotropic phase space distribution dependent only on the total energy and a radial power law form for the potential and density. Under these assumptions, \\cite{evans97} found that $\\rho\\propto r^{-4/3}$ was a steady-state solution under the effects of encounters.  There was no dynamics in this work, however: the system was assumed to be in collisionless equilibrium and the nature of the clumps producing the collisions was left arbitrary.  It also remains to be determined whether galaxy halos, either dark matter or stellar, actually evolve to $\\alpha=\\frac{4}{3}$ cusp.  \\cite{w00} examined the dynamics of halos due to stochastic fluctuations from three different noise processes for the Fokker-Planck scattering term: the merging clumps were assumed to fly by with constant velocity, to follow decaying orbits due to dynamical friction, or to orbit inside halos.  The results of this study were promising and supported the notion that the density profiles of halos were affected by mergers and substructures. The fluctuations in this work, however, were constructed by hand instead of being self-consistently calculated, and the results about a central cusp were inconclusive.  Our objective here is to construct a kinetic equation in which the drift and diffusion terms are derived from realistic density fluctuations in current cosmological models.  We will also retain the full phase space information by, in effect, using $f(E,J,v_r,t)$ without averaging over the angular momentum or radial velocity. We also do not assume collisionless equilibrium; instead we base our approach on exact dynamics.  The organization of this paper is as follows:  In \\S~2, we derive the kinetic equation relevant for the evolution of an average halo, starting with the collisionless Boltzmann equation for individual dark matter particles.  This procedure gives rise to a collision term on the right-hand side of the kinetic equation.  This term represents a correlated force density and arises from fluctuations in the ensemble averaged gravitational potential due to clustering of the matter distribution.  In \\S~3, we introduce the concepts and physical variables needed to apply this equation to cosmological problems.  The key step is in relating the phase space density $f(\\vec r,\\vec v,t)$ to the probability densities of two standard variables in cosmology: the mass density $\\rho$ and velocity $\\vec v$.  We then focus on the quasi-linear regime in which the fluctuations about the mean density and Hubble flow of an isotropic and homogeneous universe are small. This focus allows us to relate $\\rho$ and $\\vec v$ to the density and displacement fields $\\delta$ and $\\vec\\psi$, which are Gaussian random fields in standard cosmological models.  We leave the most mathematical part of the paper to Appendix~\\ref{sec:pgauss}, where we describe how to use the statistics of constrained Gaussian random fields to express the probability densities of \\S~3 in terms of the (constrained) means and covariances of the density and displacement fields.  The idea of a constrained field is central to this study because our aim is to construct a kinetic theory for dark matter particles that will later reside in a collapsed halo rather than at any random location in the universe. The relevant perturbations are for a density field constrained to have certain properties, e.g., a specified height or gradient at some point in space.  The spectrum for such constrained fields is related but not identical to the familiar linear power spectrum $P(k)$ for unconstrained Gaussian density fluctuations.  In \\S~\\ref{sec:fpeq}, we put the final expression for the correlated force derived in Appendix~\\ref{sec:pgauss} back into the kinetic equation, and show that the resulting equation is of the Fokker-Planck form where one term represents a drift force and the other term represents diffusion in velocity-space. We work out these diffusion coefficients for the currently-favored cold dark matter plus a cosmological constant ($\\Lambda$CDM) model, and provide explicit expressions for these coefficients in terms of the linear matter fluctuation power spectrum.  We also present the results for the coefficients computed for this model.  \\S~\\ref{sec:discuss} discusses further the physical meaning of the results in \\S~\\ref{sec:fpeq}.  We interpret the diffusion terms by taking velocity moments of the kinetic equation.  To place our cosmological kinetic equation in a broader perspective, we discuss in more detail the classical examples of Brownian motion and globular clusters, and point out the similarities and differences among the three cases.  \\S~\\ref{sec:conclu} provides a summary and conclusions. In Appendix~\\ref{sec:stationary} we investigate the relaxation to stationary solutions of our cosmological Fokker-Planck equation. ",
        "conclusions": "\\label{sec:conclu}   In this paper we have developed a cosmological kinetic theory, valid to second-order in perturbation theory, to describe the evolution of the phase-space distribution of dark matter particles in galaxy halos. This theory introduces a new way to model the early phases of galaxy halo formation, which has traditionally been studied by analytic infall models or numerical N-body methods.  The key physical ingredients behind our kinetic description are stochastic fluctuations and dissipation caused by substructures arising from a spectrum of cosmological density perturbations. The kinetic equation that we have obtained at the end of our derivation, Eqs.~(\\ref{fplin})--(\\ref{difflin}), has the standard form of a Fokker-Planck equation.  The diffusion coefficients $\\vec{A}$ and {\\bf D} represent acceleration due to a drift force and velocity-space diffusion, respectively.  To second order in perturbation theory, these coefficients are related to various covariance matrices of the cosmological density and velocity fields given by Eqs.~(\\ref{driftlin})--(\\ref{difflin}).  These expressions can in turn be written in terms of the familiar linear power spectrum $P(k)$ of matter fluctuations, as shown in Eqs.~(\\ref{coeffs}) and (\\ref{sumfun}).  The results for $\\vec{A}$ and {\\bf D} from second-order perturbation theory are shown in Figures~1--4 for the currently favored $\\Lambda$CDM model, and in Figures~5 and 6 for various scale-free models with power-law $P(k)\\propto k^n$. Our results indicate that dissipative processes are important during the initial collapse of a dark matter halo, and it is not valid to model halos using the spherically symmetric Vlasov (collisionless Boltzmann) equation. Furthermore, we find that the diffusivity is initially negative close to the center of a halo, indicating a thermodynamic instability.  The source of this instability is gravity: perturbations enhance gravitational collapse of dark matter halos.  We emphasize that our derivation leading to the Fokker-Planck equation given by Eqs.~(\\ref{fplin})--(\\ref{difflin}) is exact to second order in cosmological perturbation theory.  We do not follow the frequently used approach (e.g. Binney \\& Tremaine 1988) that begins by expanding the collisional terms of the non-exact master equation in an infinite series of changes in phase-space coordinates (i.e. $\\Delta \\vec{w}$), and arrives at the Fokker-Planck equation by assuming weak encounters (i.e. small $|\\Delta \\vec{w}|$) and truncating the series after the second-order terms.  Instead, our starting point is Eq.~(\\ref{boltzeq}), the first BBGKY hierarchy equation for a smooth one-particle distribution function.  This is an exact kinetic equation; solving it requires specifying the one-particle distribution at an initial time and the two-particle correlation function at all times.  We have shown in \\S~\\ref{sec:dfgauss} and Appendix~\\ref{sec:pgauss} that if the cosmological density field $\\delta$ and displacement (or velocity) field $\\psi$ are Gaussian random fields, the resulting two-particle correlation function depends on the one-particle distribution in a simple way given by Eq.~(\\ref{f2cexact}). Eq.~(\\ref{f2cexact}) is exact to second order in perturbation theory and provides the closure relation for the BBGKY hierarchy in Eq.~(\\ref{boltzeq}).  Combining Eq.~(\\ref{f2cexact}) with Eq.~(\\ref{boltzeq}) results in Eq.~(\\ref{fplin}), which is a Fokker-Planck equation.  The cosmological dissipative processes identified in this paper are quite different from the standard two-body relaxation, which also leads to a Fokker-Planck equation.  We find that there is no dynamical friction (i.e. terms proportional to $-\\vec{v}$ in the diffusion coefficients) in second-order cosmological perturbation theory. Instead, there is a radial drift toward (or away from) the center of a halo due to the clustering of matter within the halo. The usual treatment of two-body relaxation has no such drift term because it assumes the medium to be homogeneous.  Thus we have identified a new relaxation process that must affect the phase-space structure of dark matter halos.  Although our derivation is exact to second order in the density fluctuations, it is valid only when the fluctuations are small. We are therefore describing only the early stages of dark matter halo formation.  The expressions we obtain for the diffusion coefficients are plainly invalid in the nonlinear regime, where we expect dynamical friction to be present as well as radial drift. In a later paper we will investigate the nonlinear generalizations of the exact second-order results derived in this paper.  The Fokker-Planck description should still apply in the nonlinear regime provided that the fluctuating force on a particle can be modelled as a Markov process consisting of a random walk of many steps.  We conjecture that this description will be approximately valid when the matter distribution is modelled as a set of clumps (i.e., the halo model) which scatter individual dark matter particles away from the orbits they would have in a smooth, spherical potential. The Fokker-Planck description should apply not only to the initial collapse and virialization of a dark matter halo containing substructure, but also to the subsequent evolution under minor mergers.  Our inability to solve exactly the nonlinear dynamics, however, means that the diffusion coefficients will have to be calibrated using the nonlinear halo model or numerical N-body simulations.  Calibrating the diffusion coefficients is not equivalent to reproducing the N-body simulations.  N-body simulations provide a Monte Carlo solution of the BBGKY hierarchy, which is much more general than the Fokker-Planck equation.  If the simulation results are describable by Fokker-Planck evolution, this already represents a significant new result.  Moreover, even if functional forms for the radial dependence of the diffusion coefficients have to be calibrated with simulations, the solution of the Fokker-Planck equation contains far more information because it gives the complete phase-space density distribution, not simply radial profiles.  Another important step will therefore be to actually solve the Fokker-Planck equation.  In deriving it we have identified and quantified the relaxation processes affecting halo formation and evolution, but this is only the first step.  In Appendix~\\ref{sec:stationary} we presented solutions in unrealistically simple cases just to highlight the roles played by radial drift, drag, and diffusion. An outstanding question is whether these processes actually are rapid enough to erase the memory of cosmological initial conditions sufficiently so that dark matter halos relax to a universal profile.  Once we have a nonlinear model for the diffusion coefficients, we will address this important question.  The work presented in this paper has a number of implications. First, we have highlighted the importance of the phase space density for understanding dark matter dynamics, and have developed a method for calculating its evolution during the early phases of structure formation.  Second, we have identified a dissipative process that may erase memories of initial conditions during the early phases of galaxy halo evolution.  This erasure is necessary for universal halo density profiles, although further investigation is needed to quantify how this dissipative process affects the density profiles.  Third, the statistical description of the potential fluctuations caused by substructure that we have developed can also be added to investigate related problems such as the heating of galactic disks \\citep{benson03} and the merging of central black holes when galaxy halos merge \\citep{kh00,hk02,hb03}."
    },
    "0311/astro-ph0311094_arXiv.txt": {
        "abstract": "{ The excitation of internal gravity waves by an entropy bubble oscillating in an isothermal atmosphere is investigated using direct two-dimensional numerical simulations. The oscillation field is measured by a projection of the simulated velocity field onto the anelastic solutions of the linear eigenvalue problem for the perturbations. This facilitates a quantitative study of both the spectrum and the amplitudes of excited $g$-modes. ",
        "introduction": "The problem of the excitation of internal gravity waves (hereafter IGWs) in stably stratified media is often studied in connection with the possible detection of solar $g$-modes (see e.g.\\ the latest attempt of $g$-mode detection in the GOLF data by Gabriel et al.\\ 2002). As an example, two- and three-dimensional numerical simulations of penetrative convection have shown that it is possible to excite such waves below the convection zone from the penetration of strong downward plumes into the stable radiative zone located below (Hurlburt et al.\\ 1986; Hurlburt et al.\\ 1994; Brandenburg et al.\\ 1996; Brummell et al.\\ 2002; Dintrans et al.\\ 2003).  Furthermore, IGWs play a major role in stellar evolution as they can transport angular momentum and/or chemical elements over long distances in the stellar interior. This transport mechanism has been proposed to explain the quasi-solid rotation profile of the solar core as revealed by helioseismology (Kumar et al.\\ 1999) as well as the lithium depletion observed in low mass stars (Talon \\& Charbonnel 1998). However, the correct amount of IGWs excited by, say, penetrative convection remains still unknown as numerical studies were essentially led from a qualitative point of view, so that the mode amplitudes were not deduced.  The main goal of this paper is to show that it is possible to determine quantitatively the amplitudes of gravity waves propagating in a stable zone of a numerical simulation using the anelastic subspace. This subspace is built from the solutions of the associated anelastic linear eigenvalue problem for the perturbations. The projection of the simulated velocity field onto this basis yields the IGW amplitudes. We will present here the application of this method to the simple problem of $g$-mode oscillations in an isothermal atmosphere. The dynamics of such an atmosphere is well known (Brandenburg 1988) so that this problem is ideal to illustrate and test the validity of our method, before applying it to the more difficult problem of IGWs excited by penetrative convection, where the atmosphere is in general non-isothermal.  After presenting the hydrodynamical model describing our isothermal atmosphere (Sect.\\ \\ref{themodel}), we apply two classical and widely used methods to detect the $g$-modes excited by an oscillating entropy bubble and show their limitations (Sect.\\ \\ref{classic}). Next, we introduce our new method based on the anelastic subspace and give the analytic forms we found for both the eigenfrequencies and eigenvectors of the anelastic set of linear equations for the perturbations (Sect.\\ \\ref{newmethod}). We then apply this formalism to the same simulation used to test the two classical methods discussed in Sect.\\ \\ref{classic} and show that we now have access to both the spectrum and amplitudes of the IGWs (Sect.\\ \\ref{results}). In particular, we illustrate the usefulness of our method by studying the influence of the box geometry on the mode amplitudes. Finally, we conclude in Sect.\\ \\ref{conclu} by introducing the next step of this work, which will consist of the application of the anelastic subspace method to the detection of IGWs that are stochastically excited by penetrating convective plumes. ",
        "conclusions": "\\label{conclu}  We have presented a new and accurate method to measure the internal gravity waves propagating in a numerical simulation. This is of prime importance when one deals with the problem of the generation of IGWs in a stably stratified medium such as the radiative zone of a star. The knowledge of the vigour of the velocity field in such a zone is the main input of any theory of wave transport of angular momentum and/or chemicals in stellar radiative zones (e.g.\\ Schatzman 1993).  However, we have shown that classical methods based only on Fourier transforms in space and/or time of the simulated velocity field do not allow a quantitative determination of the mode amplitudes. The reason is mainly that the contribution of the propagating $g$-modes to the total energy in the simulation is diluted among the other ones due to $p$-modes and turbulent motions.  Our method is based on the projection of the simulated velocity field onto the anelastic subspace built from the solutions of the linear problem for the perturbations. Once the resulting time-dependent coefficients of the velocity field are computed, one can deduce the amplitudes of every $g$-mode and measure their kinetic energy at every timestep of the simulation.  We have applied this formalism to the measurement of IGWs excited by an entropy bubble oscillating in an isothermal atmosphere. We have shown that the mode amplitudes follow the same law than those of a linear damped harmonic oscillator, that is, the periodic amplitudes decrease exponentially with time, with a time scale that is simply proportional to the inverse of the viscosity; see Eq.~(\\ref{harm}). By using a Parseval-Bessel relation valid in the anelastic subspace, we have deduced the contribution of every $g$-mode to the time-averaged kinetic energy and have shown that large-scale motions are mainly composed of $n=0$ and $n=1$ modes. Such a result is in general impossible to obtain in a more complicated settings using a classical $(k_x-\\omega)$ diagram.  The fact that the mode amplitudes in the anelastic basis follow the same law than those of a linear damped oscillator is important in the case of a stochastic excitation of IGWs, such as the one encountered when one deals with numerical simulations of penetrative convection. Indeed, in this case the mode amplitude $c_{\\ell n}$ evolves either chaotically around zero when the mode is not excited, or in a periodic fashion following Eq.~\\eq{harm} when the mode is excited. By computing time-frequency diagrams of that time-dependent amplitude, one can extract the time intervals during which the mode has been really excited. In a forthcoming paper, we will present the application of the anelastic subspace method we tested here to the problem of IGWs stochastically excited by penetrative convection (Dintrans et al.\\ 2004)."
    },
    "0311/astro-ph0311577_arXiv.txt": {
        "abstract": "{ Recent Cosmic Microwave Background (CMB) observations indicate that the temperature anisotropies arise from quantum fluctuations in the inflationary scenario. In the simplest inflationary models, the distribution of CMB temperature fluctuations should be Gaussian. However, non-Gaussian signatures can be present. They might have different origins and thus different statistical and morphological characteristics.  In this   context and motivated by   recent and future CMB experiments, we search for, and discriminate between, different non-Gaussian signatures. We analyse simulated maps of three cosmological sources of temperature anisotropies: Gaussian distributed CMB anisotropies from inflation, temperature fluctuations from cosmic strings and anisotropies due to the kinetic Sunyaev-Zel'dovich (SZ) effect both showing a non-Gaussian character. We use different multi-scale methods, namely, wavelet, ridgelet and curvelet transforms. The sensitivity and the discriminating power of the methods is evaluated using simulated data sets.  We find that the bi-orthogonal wavelet transform is the most powerful for the detection of non-Gaussian signatures and that the curvelet and ridgelet transforms characterise quite precisely and exclusively the cosmic strings. They allow us thus to detect them in a mixture of CMB + SZ + cosmic strings.  We   show that not one method only should be applied to understand non-Gaussianity but rather a set of different robust and complementary methods should be used.  } ",
        "introduction": "The search for non-Gaussian signatures in the Cosmic Microwave Background (CMB) temperature fluctuation maps  furnished by MAP\\footnote{http://map.gsfc.nasa.gov/} \\cite{komatsu2003}, and to be furnished by PLANCK\\footnote{http://astro.estec.esa.nl/SA-general/Projects/Planck/}, is of great interest for cosmologists.  Indeed, the non-Gaussian signatures in the CMB can be related to very fundamental questions such as the global topology of the universe \\cite{riazuelo2002}, super string theory, topological defects such as cosmic strings \\cite{bouchet88}, and multi-field inflation \\cite{bernardeau2002}.  The non-Gaussian signatures can, however, have a different but still cosmological origin. They can be associated with the Sunyaev-Zel'dovich (SZ) effect \\cite{sunyaev80} (inverse Compton effect) of the hot and ionised intra-cluster gas of galaxy clusters \\cite{gauss:aghanim99,cooray2001}, with the gravitational lensing by large scale structures \\cite{bernardeau2003}, or with the reionisation of the universe \\cite{gauss:aghanim99,castro2002}.  They may also be simply due to foreground emission \\cite{gauss:jewell01}, or to non-Gaussian instrumental noise and systematics \\cite{banday2000}.  It has now become clear that the detection of non-Gaussian signature is important and feasible (e.g. \\cite{detroia2003,komatsu2003}). Nevertheless, even if a non-Gaussian signal is detected the question of its origin might still be posed (e.g. \\cite{banday2000}). Hence, {\\it it remains to separate between several non-Gaussian contributions.  This is the most crucial issue for the use  the non-Gaussian signatures as cosmological tools. }  It is therefore not surprising that a large number of studies have recently been devoted to the subject of the detection of non-Gaussian signatures. Many approaches have been investigated such as the Minkowski functionals and the morphological statistics \\cite{gauss:novikov00,gauss:shandarin02}, the bispectrum \\cite{gauss:bromley99,gauss:verde00,gauss:phillips01}, the trispectrum \\cite{gauss:kunz01}, or wavelet transforms \\cite{gauss:aghanim99,gauss:forni99,gauss:hobson99,gauss:barreiro01,gauss:cayon01,gauss:jewell01}. Different wavelet methods have been studied, such as the \\`a trous algorithm \\cite{starck:book98} and the bi-orthogonal wavelet transform \\cite{ima:mallat98}.  A series of recent papers \\cite{cur:candes99_1,Curvelets-StMalo,starck:sta01_3}, however, argued that wavelets and related classical multi-resolution techniques are based on a limited dictionary made up of roughly isotropic elements present at all scales and locations. We view as a limitation the fact that those dictionaries do not exhibit highly anisotropic elements and that there is only a fixed number of directional elements, independent of scale. Despite the success of the classical wavelet viewpoint, there are objects, such as filamentary structures  or cosmic strings, that do not exhibit isotropic scaling and thus call for other kinds of multi-scale representations. Following this, new multi-scale systems like curvelets and ridgelets \\cite{Curvelets-StMalo} have been introduced which are very different from wavelet-like systems. Curvelets and ridgelets take the  form of basis elements which exhibit very high directional sensitivity and are highly anisotropic.  In two dimensions, for instance, curvelets are localised along curves, in three dimensions along sheets, etc. \\par\\bigskip The goal of the present study is to compare these new multi-scale representations to standard wavelet methods  on a set of simulated maps representing  two generic families of non-Gaussianities (spherical-like and string-like sources), and to answer the following questions: \\begin{enumerate} \\item Are there multi-scale methods that are better suited to detect the non-Gaussian signatures in CMB data? \\item Can we  go beyond the detection and extract some information about the nature of the non-Gaussian signal? \\end{enumerate}  The second point is particularly important and it will represent the heart of our present study. A non-Gaussian character can be due to many effects: Topological defects, the SZ effect, calibration problems, or even a combination of them. Here, we perform a study based on the idea that a given multi-scale transform is optimal to detect features which have the shape of its basis elements (this is quite similar to the matched filter approach).  Consequently, if different multi-scale methods are able to detect a non-Gaussian signal, one of them will outperform the others since the shape of the non-Gaussian features contained in the data will be close to its basis elements. In particular, using wavelet-like systems and curvelet and ridgelet-like systems, we expect to differentiate between the spherical-like and string-like sources of non-Gaussianity.  Sections~\\ref{sect_transforms} and 3 briefly review the different transforms and simulated maps considered in our analysis. We present the analysis together with our main results in section 4 and discuss them in section 5. We provide the main conclusions in section 6. ",
        "conclusions": "In the present study, we have tested the relative sensitivities of the commonly used multi-scale transforms to the detection of non-Gaussian signatures in CMB maps. We have used simulated maps of the primary anisotropies (assumed to be Gaussian) and we have taken into account the contributions from the kinetic SZ effect and the topological defects (cosmic strings); each of them represents one family of non-Gaussian contributors (spherical-like and string-like). The main results are:  \\begin{itemize} \\item the bi-orthogonal wavelet transform is the most sensitive to the non-Gaussian signatures associated with cosmic strings buried in the primary Gaussian CMB signal. \\item The ridgelet and curvelet transforms are well adapted to {\\it discriminate} between string-like and quasi-spherical like non-Gaussian features. \\end{itemize}  In order to study the non-Gaussian signatures in the CMB and use them as a cosmological tool to probe the early universe or the cosmic structures, it is not sufficient to detect them accurately. The most important step is to be able to separate the different contributions to the signal. In this context, we clearly show that not only one method should be applied but rather a set of different robust and well understood methods.  In the present study, we  use several multi-scale transforms: The isotropic wavelet transforms suited for spherical-like sources of non-Gaussianity, and a curvelet transform  representing well sharp and elongated structures. Each  provides an adapted non-Gaussian estimator, namely the normalised mean excess kurtosis. We show that the combination of these transforms through the product of the normalised mean excess kurtosis of isotropic wavelet transforms by normalised mean excess kurtosis of curvelet transforms highlights the presence of the cosmic strings in a mixture CMB+KSZ+CS. Such a combination   gives information about the nature of the non-Gaussian signals. Even though the detection power of the non-Gaussian character by the ridgelets and curvelets is not important compared to that of the wavelet, their sensitivity to a particular shape makes them a very strong discriminating tool.  This is a first step towards the separation of the statistical contributions to the CMB signal and it can help in the more general context of component separation."
    },
    "0311/astro-ph0311607_arXiv.txt": {
        "abstract": "{ In several nearby $L \\sim L_*$ early-type galaxies, recent observations at large radii have shown a indications of a lack of dark matter, substantially at odds with the  prediction from the Cold Dark Matter (CDM)  hierarchical merger models. Here we discuss a pilot observational project for the study of the internal kinematical and dynamical properties of this remarkable sample of galaxies. Using the VIMOS-IFU in its high spectral resolution mode, it would be possible to investigate the regions up to $\\sim ~1.2 ~R_{\\rm e}$, taking advantage of the much larger field of view and telescope diameter. This will allow to disclose the presence of any kinematical substructures which could affect the conclusion on the mass modeling and definitely clarify the inner structure of this particular class of early-type galaxies. ",
        "introduction": "Intermediate luminosity early-type galaxies ($L\\sim L_*$) are interesting laboratories where to test both the galaxy formation theories and the presence of the dark matter. Their structural and kinematical properties make them a different class with respect to the brightest systems. Indeed, from the photometric point of view, bright galaxies have usually boxy isophotes and flat inner cores in the light distribution while fainter galaxies are mostly disky with power-law inner light profiles (Nieto \\& Bender 1989, Capaccioli et al. 1992, Faber et al. 1998).\\\\ These isophote shapes appear to be correlated with a variety of properties of  these systems (Kissler-Patig 1997, Pellegrini 1999 and references therein), in particular with the kinematics: bright/boxy systems are slowly rotating in the inner parts, while faint/disky galaxies are rotationally supported at least within $1~R_{\\rm e}$ (the effective radius $R_{\\rm e}$ encloses half of the galaxy projected light), the typical observational limit for kinematical studies by integrated light long-slit spectroscopy (Bertola \\& Capaccioli 1975, Davies et al. 1983, Scorza \\& Bender 1995).   Recently, some lines of evidence has been found about the persistence of such a dichotomy also in the dark matter properties of these systems. Capaccioli et al. (2003) and Napolitano et al. (2003) have discussed the correlation between the structural parameters and the radial trend in the mass-to-light ratios (M/L) using archive kinematic data up to $\\sim 6 R_{\\rm e}$ for an heterogeneous sample of ellipticals: faint/disky galaxies show small or null M/L radial gradients, i.e., almost constant M/L, while bright/boxy objects have larger gradients in agreement with a substantial amount of dark matter in the outskirts of those systems. In particular the evidence of a dearth of dark matter in the intermediate luminosity galaxies has been confirmed by Romanowsky et al. (2003) using planetary nebulae (PNe) kinematics out to 4 $R_{\\rm e}$. Confirmation of these preliminary results by new observations with larger PNe samples or new techniques is very important, since a possibly tight connection between structural properties and density environment (i.e., the dark matter halo) where the galaxies have grown would deeply affect our understanding of galaxy formation and evolution.   The formation and evolution of these intermediate luminosity galaxies with almost null M/L gradients is still not clear. As a class, disky galaxies can be considered close to the bulge dominated S0 galaxies since disky isophotes could  reflect the presence of an embedded faint stellar disk (Capaccioli 1987; Bender et al. 1988). Scorza \\& Bender (1995) have shown that the information on both photometry and line profile of a sample of disky galaxies is compatible with simple disk+bulge models where the two subcomponents have parallel angular momenta. This latter evidence could be an hint of a contemporary evolution for these galaxies rather than a late accretion or merger events.\\\\ However, disky isophotes are also expected in hierarchical merging scenarios in a CDM cosmology (Meza et al. 2003): boxy/disky isophotes are found for inclinations which minimize/maximize $V_{\\mathrm{rot}}/\\sigma$. Furthermore, clues on angular momentum alignment could change when looking at the halo regions of these systems. Preliminary results on PNe radial velocity fields of NGC~3379 and NGC~4494 from the Planetary Nebula Spectrograph, (PN.S, Douglas et al. 2002), have revealed that there is a tilt of the major kinematical axis of the PNe sample with respect to the stellar long-slit data as it is clearly shown in Fig. 1. In the same figure a quite good alignment is found for NGC821 as well as it has been found also in earlier studies (NGC~4697: Mend\\`ez et al. 2001). Incidentally, the latter are viewed quite edge-on, while NGC~3379 and NGC~4494 are possibly nearly face-on. Furthermore, at least for the four systems mentioned above, the rotation velocity curves increase in the inner parts (up to $\\sim 1~R_{\\rm e}$) and rapidly drop in the outer parts down to zero at the outermost points. Is this an indication of some secondary evolution processes in some of these galaxies (i.e., tidal interactions or merging), or is this the result of different stellar population inhabiting inner (disk population) and outer (bulge/halo population) regions, which appear to dominate depending on the viewing angle?  The geometry can play a major role even in modeling the mass distribution, in particular for the systems seen almost face--on (Magorrian \\& Ballantyne 2001). If the  presence of a disk is kinematically relevant up to large radii, then we could expect radial orbits to dominate in projection, causing an apparent decreasing velocity dispersion and then a mass underestimate, i.e., we would erroneously conclude that there is a lack of dark matter.   So far traditional long-slit spectroscopy has been employed to address these problems, resulting in one dimensional information mostly performed along the main principal photometrical/kinematical axes of galaxies (Bender, Saglia \\& Gerhard 1994, Statler \\& Smecker-Hane 1999). Nowadays, integral field unit spectrographs are opening a new era in the understanding of the galaxy structure and evolution. We discuss here an observational project to study a preliminary sample of intermediate luminosity/disky galaxies using the Integral Field Unit (IFU) of VIMOS. The most important wide field integral field spectroscopy instruments used to study nearby elliptical galaxies are shortly described in Sect. 2, guideline of the project are briefly described in Sect. 3, together with the main expected results.  \\begin{figure*} \\resizebox{\\hsize}{!} {\\includegraphics[]{Fig1.ps}} \\caption{Gaussian smoothed PNe velocity field from PN.S data (R+03), darker to brighter regions are negative to positive radial velocities. Dashed line is the P.A. of the inner stellar major kinematical axis. Misalignment of the PNe axis with respect the inner stellar kinematics, kinematical substructures and twist of the isovelocity contours are evident. Signatures of such features in the inner stellar population can be investigated by VIMOS-IFU in great detail.} \\label{figlabel} \\end{figure*} ",
        "conclusions": "The ubiquity of dark matter in early-type galaxies is a debated topic: already Saglia et al. (1993), using stellar kinematics observations, hinted at the possible existence of two classes of early-type galaxies, according to the velocity dispersion trend at large radii. They studied several early-type galaxies with decreasing velocity dispersion which required no dark matter (see also Capaccioli et al. 2003). However, long-slit spectroscopical stellar kinematics studies are generally confined to the inner luminous regions ($R < R_{\\rm e}$). A remarkable improvement has been possible using PNe as kinematical tracers in the outskirts of early-type galaxies with dedicated instruments like the PN.S: preliminary results have already been published, and a larger sample of early-type galaxies is being observed with the same technique. However, in order to improve our knowledge of the kinematics of the inner luminous regions, deep IFS observations are mandatory, and they are already being performed on large sample of nearby galaxies by the SAURON team.   Here we discussed an observational project to survey the full 2D kinematics of intermediate luminosity/disky galaxies, up to about 1.2  $R_{\\rm e}$, in order to overlap the region where PNe data are available. The preliminary sample includes galaxies which have been modeled with nearly constant mass-to-light ratios up to about 6 $R_{\\rm e}$ using the information from PNe discrete radial velocity fields (Romanowsky et al. 2003). This finding is extremely interesting, as it seems to violate the expectations from CDM hierarchical merger models (e.g., Navarro, Frenk \\& White 1997), which predict a substantial amount of dark matter in the outskirt of ellipticals.   As discussed above, there are still some indeterminacies due to the intrinsic structure of disky systems (like the presence of faint disks or multicomponent with misaligned kinematics) and geometry (i.e., inclination): the use of the detailed 2D kinematics and high velocity resolution is therefore necessary to clarify any presence of biases in the dynamical modeling due to kinematical substructures.  IFU-VIMOS is a very suitable instrument for such a purpose as it allows to combine an adequate spatial and spectral resolutions with a large survey field and the advantage of an 8m telescope collecting power."
    },
    "0311/astro-ph0311431_arXiv.txt": {
        "abstract": "We work in the framework of an inflationary cold dark matter universe with cosmological constant, in which the cosmological inhomogeneities are considered as gravitational lenses for the CMB photons. This lensing deforms the angular distribution of the CMB maps in such a way that the induced deformations are not Gaussian. Our main goal is the estimation of the deviations with respect to Gaussianity appeared in the distribution of deformations. In the new approach used in this paper, matter is evolved with a particle-mesh N-body code and, then, an useful ray-tracing technique designed to calculate the correlations of the lens deformations induced by nonlinear structures is applied. Our approach is described in detail and tested. Various correlations are estimated at an appropriate angular scale. The resulting values point out both deviations with respect to Gaussian statistics and a low level of correlation in the lens deformations. ",
        "introduction": "Introduction}  In the absence of a reionization, the photons of the Cosmic Microwave Background (CMB) are not scattered from decoupling to present time, except inside galaxy clusters (Sunyaev-Zel'dovich effect); nevertheless, the propagation direction of these photons changes due to the gravitational action of the cosmological inhomogeneities (lensing).  The so-called primary CMB anisotropies were produced by linear structures at high redshifts. The corresponding temperature contrast, $\\Delta_{_{P}} \\equiv  (\\Delta T/T)_{_{P}}$, is a statistical field which can be expanded in spherical harmonics as follows: $\\Delta_{_{P}} = \\sum_{m = -\\ell}^{+\\ell} a_{\\ell m} Y_{\\ell m}$, where the $a_{\\ell m}$ quantities are statistically independent variables depending on the sky realization we are expanding. If many sky realizations are averaged, the resulting $a_{\\ell m}$ coefficients have zero means and variances $\\langle \\mid a_{\\ell m} \\mid^{2} \\rangle = C_{\\ell}$. Since $\\Delta_{_{P}}$ is a homogeneous and isotropic statistical field, quantities $\\langle \\mid a_{\\ell m} \\mid^{2} \\rangle$ do not depend on $m$. In the model under consideration, scalar energy density fluctuations are initially Gaussian, and they remain Gaussian during linear evolution, by this reason, the primary CMB anisotropies --which were produced by linear inhomogeneities-- also are Gaussian.  Gaussian primary anisotropy appears to be superimposed to secondary anisotropies generated well after decoupling, and also to the anisotropy of the microwave radiation emitted by both our galaxy and some extragalactic sources (contaminant foregrounds). These foregrounds are sub-dominant non Gaussian components contributing to the total observable signal at microwave wavelengths. In the model we are considering here, there are various secondary anisotropies. A few comments about them are worthwhile: the late integrated Sachs-Wolfe anisotropy was produced by quasi-linear inhomogeneities \\cite{ful00} and, consequently, deviations with respect to Gaussianity are not expected to be very important; another secondary gravitational anisotropy was produced by strongly nonlinear structures as galaxy clusters and substructures (Rees-Sciama effect), this component is sub-dominant and non Gaussian, and its deviations with respect to Gaussianity were studied by various authors \\cite{mol95,mun95,ali02}; finally, the Sunyaev-Zeldovich anisotropy was produced by hot plasma in galaxy clusters, and it is also non Gaussian \\cite{coo01}.  The main goal of this paper is the analysis of the non Gaussianity caused by lensing, and such an analysis can be performed using the operative ray-tracing procedure designed in next sections. In the absence of lensing, there are directional frequency shifts which cause an anisotropic temperature distribution in the CMB. The lens effect does no cause frequency shifts, but it produces angular deviations of the propagation directions. This means that the lens effect does not produce any alteration of the CMB temperature. It changes the propagation direction of the CMB photons and, then, the angular distribution of the CMB temperature changes accordingly. In other words, the lens effect deforms the maps of the CMB temperature which appear as a result of pure frequency shifts. Any numerical estimate of the alterations of the CMB maps produced by lensing (including non Gausianity generation) involves three steps: (i) maps of the CMB temperature distribution are built up for the model under consideration (without any lensing), (ii) these maps are deformed taking into account the deviations of the propagation directions produced by lensing, and (iii) the resulting deformations of unlensed maps (hereafter called either lens deformations or lens effect) are analyzed. Let us now consider each of these steps in more detail.   In order to built up unlensed maps of the CMB, we assume that, in the absence of lensing, the Gaussian primary anisotropy dominates and, consequently, the maps to be deformed can be built up using the angular power spectrum of this primary anisotropy, which has been calculated using CMBFAST \\cite{seza96} for the $\\Lambda$CDM model under consideration. As it is discussed below, only small Gaussian maps having sizes of a few degrees and resolutions of the order of one arc-minute are necessary and they are built up using the Fourier Transform (FT) and the mentioned spectrum ($C_{\\ell}$ quantities). See \\cite{sae96} and references cited therein for details. The resulting maps have spots with different sizes and amplitudes (angular structure). Spots with angular size $\\pi / \\ell $ have a mean amplitude proportional to $C_{\\ell}^{1/2}$; hence, spots with angular size close to one degree ($\\ell \\sim 200$) have large amplitudes, whereas spots corresponding to one arcminute ($\\ell \\sim 10000$) have negligible amplitudes as a result of the smallness of $C_{10000}$. For the spectrum under consideration, significant spots have sizes of various arcminutes.      Cosmological inhomogeneities produce lensing and this effect is associated to the deviation field $\\vec {\\delta}$. The anisotropy observed in the $\\vec {n}$ direction is the primary anisotropy corresponding to the direction $\\vec {n}_{0}$ in the absence of lensing, where $\\vec{n} = \\vec {n}_{0} - \\vec {\\delta} $; hence, \\begin{equation} \\Delta (\\vec {n}) = \\Delta_{_{P}} (\\vec {n}_{0}) = \\Delta_{_{P}} ( \\vec {n} + \\vec {\\delta} )  \\ . \\label{defor} \\end{equation} The unit vectors $\\vec {n}$ and $\\vec {n}_{0}$ point towards two points of the last scattering surface, and the deviation field $\\vec {\\delta}$ gives the angular excursion on this surface due to lensing. We must estimate the deviation field $\\vec {\\delta} $ (see below) to get the deformed observable map $\\Delta$. If the $\\vec {\\delta} $ deviations are smaller than the size of the smallest significant spots (various minutes), Eq. (\\ref{defor}) can be expanded in $\\vec {\\delta} $ powers to get \\begin{equation} \\Delta (\\vec{n}) \\simeq \\Delta_{_{P}} (\\vec{n}) + \\frac {\\partial \\Delta_{_{P}}}{\\partial \\vec{n}} \\cdot \\vec{\\delta} \\label{dmap} \\end{equation} We see that map deformations, namely, the differences in the temperature contrasts corresponding to associate directions (second term of the r.h.s.) are the product of two statistical fields: $\\Delta_{_{P}}$ and $\\vec {\\delta}$ and, even if these field are Gaussian and statistically independent, the resulting product would be non Gaussian.  Several authors have studied angular correlations in lensed CMB maps. In the most theoretical papers on this subject, some correlations have been estimated without using simulations; the approach used in these papers only applies for those correlations which can be written in terms of the time varying form of the matter power spectrum. In the most basic of these papers \\cite{sel96}, the mentioned spectrum was modelled beyond the linear regime and the second order correlations of lensed CMB maps were estimated. A similar method was used by Bernardeau \\cite{ber97} to estimate fourth order correlations corresponding to various sets of four directions (with distinct configurations); furthermore, this last author pointed out that third order correlations vanish. Other methods are necessary to estimate correlations of order $n$ for $n>4 $, and also to simulate CMB maps with the true statistics induced by lensing, namely, maps having the true correlations at any order (not only at second and fourth orders). The most powerful technique to study lens distortions in CMB maps is the use of {\\em ray-tracing through N-body simulations}. Two methods to apply this tecnique are described in references \\cite{jai00} and \\cite{whi01} (other relevant papers in this field are cited in these two references). Another new method is proposed here. We agree with White \\& Hu \\cite{whi01}, who stated that: {\\em ..... On subdegree scales, a full description of week lensing therefore requires numerical simulations, the most natural being N-body simulations ......}; in fact, using ray-tracing through N-body simulations, CMB photons are deviated by fully nonlinear structures (as galaxy clusters) which move in the simulation box. Nowadays, only N-body simulations lead to a proper description of a distribution of strongly nonlinear structures. Unfortunately, these simulations constraint us to work in a periodic universe, and relevant problems associated to periodicity must be solved. Each method --designed to use ray-tracing through N-body simulations-- corresponds to a different way for preventing periodicity effects. In this paper, one of these methods is described and tested.   Hereafter, quantity  h is the reduced Hubble constant $h=10^{-2}H_{0}$, where $H_{0}$ is the Hubble constant in units of $Km s^{-1} Mpc^{-1}$, the density parameters corresponding to baryons,  dark matter, and vacuum, are  $\\Omega_{b}$, $\\Omega_{d}$ and  $\\Omega_{\\Lambda}$, respectively, the total density parameter is $\\Omega = \\Omega_{b} + \\Omega_{d} + \\Omega_{\\Lambda}$, and the matter density parameter is $\\Omega_{m} = \\Omega_{b} + \\Omega_{d}$. In the flat inflationary universe under consideration, the above parameters take on the following values: $h=0.65$, $\\Omega_{b}=0.05$, $\\Omega_{d}=0.25$, and $\\Omega_{\\Lambda}=0.7$ and, then, according to Eke et al. \\cite{eke96}, the power spectrum of scalar energy density perturbations must be normalized with $\\sigma_8=0.93$ in order to have cluster abundances compatible with observations. Units are chosen in such a way that $c=8\\pi G =1$, where $c$ is the speed of light and $G$ the gravitation constant. Whatever quantity \"$A$\" may be, $A_{_{L}}$ and $A_{0}$ stand for the $A$ values on the last scattering surface and at present time, respectively. The scale factor is $a(t)$, where $t$ is the cosmological time, and its present value, $a_{0}$, is assumed to be unity, which is always possible in flat universes. ",
        "conclusions": "(i) a size of $128 \\ Mpc$ is preferable against larger sizes including extended linear structures (which would produce a significant lensing) and, (ii) resolutions of $0.5 \\ Mpc$ and $1. \\ Mpc $ lead to similar correlations; hence, results converge as resolution increases. Is this convergence real or apparent? As it has been pointed out in  Sec. V, the resolution of our simulations must be increased to answer this question. How much should we increase the resolution? We should consider --at least-- a resolution of $0.25 \\ Mpc$ and, perhaps another of $0.125 \\ Mpc $. These resolutions lead to severe computational problems because many simulations are necessary for statistical purposes. Higher-resolution codes --for example tree or  $P^{3}M$ codes-- could be more appropriate than our PM one. Fortunately, our simulations lead to boxes containing nonlinear structures which allow us to apply and test our ray-tracing procedure and, furthermore, our best resolution ($0.5 \\ Mpc $) should lead to a rather acceptable description of clusters, because the central region, where the density is more peaked (distances $\\sim 1 \\ Mpc$ from the center), is covered by about $2^{3} $ pairs of cells. We can state that our main goal (description of a new ray-tracing procedure and estimations of nongaussianity from nonlinear structures) has been reached, whereas results from higher resolutions are out of the scope of this work, and they will be developed in due time and presented elsewhere.  In our model, photons move in a periodic universe which is the repetition of the simulation box. We have verified that periodicity is important, but photons do not feel periodicity when they move along directions close to a preferred one (see Sec. IV); this fact constraints us to calculate deformations for this type of directions. Fortunately, these directions suffice to find the required correlations. It is worthwhile to emphasize some novel aspects of our calculations: the basic equations are (\\ref{funda1}) and (\\ref{funda2}); the r.h.s of Eq. (\\ref{funda1}) has the form of a Fourier transform, and it gives --at the same time-- the deviation field $\\vec{\\delta }(\\vec{x}_{_{P}},\\vec{n})$ for all the observers located in the nodes ( $\\vec{x}_{_{P}}$ ) of the simulation grid. This fact has been probed to be relevant when averages are performed. The r.h.s. of Eq. (\\ref{funda2}) is a time integral which can be analytically calculated after using the interpolation formula (\\ref{interpol}), this procedure simplifies the calculations. The use of directions close to the preferred one is the last novel aspect deserving attention. For these directions and fixed boxes, photons enter the n-th box across a certain region  and the (n+1)-th box through a different independent zone; hence, roto-traslations are not necessary to avoid periodicity effects; furthermore, it is worthwhile to emphasize that --without roto-traslations-- function $ \\vec {\\nabla}_{\\bot } \\phi$ is continuous at the points where photons cross boundary boxes (periodicity). However, if the boxes are moved (rototralations), discontinuities at these points are unavoidable and, consequently, the function to be integrated in Eq. (5) has artificial finite discontinuities which influence deviation calculations. Perhaps these discontinuities do not produce any relevant total effect when many boxes are crossed, but potential problems with discontinuities are clearly surmounted by our method in the absence of roto-traslations, which seem not to be either necessary or appropriate for us.   Since periodicity effects are prevented along any direction forming an angle of a few degrees with the preferred one, our method can be used: (1) to simulate squared maps with sizes of a few degrees, for example, $4^{\\circ} \\times 4^{\\circ}$ maps with the preferred direction pointing towards their centers, and (2) to compute angular correlations $C_{n}$ for $n>4 $. In previous sections, correlations $C_{3}$ and $C_{4}$ have been computed using sets of three and four directions close to the preferred one and, evidently, using the same methods and other sets of these directions, the correlation $C_{6} $ and higher ones can be computed. On account of these considerations, future applications of the proposed method are certainly promising.   Hereafter, we only make reference to the lens effect (or lens deformations) produced by nonlinear structures. As it was predicted by Bernardeau \\cite{ber97}, the lens effect does not create any significant $C_{3}$ correlation ($r_{3} \\simeq 1.17 \\times 10^{-4}$). Correlations $C_{2}$ and $C_{4}$ appear to be weaker than those of the primary maps (low correlation level, see Sec. V). The resulting ratio $3 C_{2}^{2}/C_{4} \\simeq 1.43$ points out a deviation from Gaussianity. In the observational maps, the lens deformations are superimposed to both the dominant primary anisotropies and other contributions described in the introduction (foregrounds and nonlinear effects). Let us analyze the superposition of primary maps and lens deformations taking into account that:  (1) the correlation $C_{2}$  ($C_{4}$) of primary maps has appeared to be greater than that of lens deformations by a factor close to $1700$ ($4 \\times 10^{6}$), (2) the correlation $C_{3}$ vanishes in primary maps and it has been found to be negligible for deformations, and (3) the two components under consideration are statistically independent. Simple calculations give the correlations of the superposition of components, which are almost identical to those of the primary maps. Deviations with respect to Gaussianity are almost negligible in this superposition. In the observational maps, there are other deviations due to the presence of other components, which would contribute to hide the small deviations associated to lens deformations.  In future, our calculations could be repeated using other N-body codes and resolutions, the lens effect due to linear structures contained in boxes with different sizes could be studied in more detail; the present study could be extended to other components of the observational CMB maps, in particular to the Rees Sciama component \\cite{ali02}; maps of lens deformations and other components could be simulated --with our method-- and these maps (which would have the true statistical properties at any order of correlation) could be analyzed using various estimators of nonGaussianity proposed in the literature \\cite{tak01,ber98}; our simulations could be also used to study the correlation between galaxy shear and CMB temperature distortions \\cite{van00} and, finally, using other N-body codes and greater resolutions, it would be worthwhile the study of very small (large) angular scales ($\\ell $ values), for which, the $C_{\\ell} $ coefficients of the primary anisotropy are very small. Some of these studies and other possible ones could be useful to improve on the method proposed in this paper; nevertheless, a good version has been already implemented and numerical codes are operative."
    },
    "0311/astro-ph0311357_arXiv.txt": {
        "abstract": "Spectroscopic observations of distant quasars have resulted in the detection of molecular hydrogen in intervening damped Lyman-$\\alpha$ absorption clouds (DLAs).  We use observations compiled from different experimental groups to show that the molecular hydrogen abundance exhibits a dramatic increase over a cosmological time period corresponding to 13\\% to 24\\% of the age of the universe.  We also tentatively show that the heavy element abundances in the same gas clouds exhibit a faster and more well-defined cosmological evolution compared to the general DLA population over the same time baseline.  We argue that this latter point is unsurprising, because the general DLA population arises in a wide variety of galaxy types and environments, and thus a spans broad range of ISM gas-phases and abundances at the same cosmic time.  DLAs exhibiting H$_2$ absorption may therefore circumvent this problem, efficiently identifying a narrower class of objects, and provide a more sensitive probe of cosmological chemical evolution. ",
        "introduction": "\\label{sec:intro}  The detection of cosmological chemical evolution is a cornerstone of modern cosmology, providing empirical details about galaxy formation and evolution, and independent supporting evidence of the big bang. High-precision measurements of chemical abundances at early cosmological epochs come from damped Lyman-$\\alpha$ absorption systems (DLAs), arising in distant galaxies which intersect the lines of sight to high-redshift quasars. DLAs are easily detected in quasar spectra due to their particularly high neutral hydrogen column densities [$\\log_{10} N($H{\\sc \\,i}$) \\ge 20.3{\\rm \\,atoms\\,cm}^{-2}$ by definition, though Voigt profile damping wings are seen at lower $N($H{\\sc \\,i}$)$].  Generally, they comprise a clumpy distribution of clouds, spread typically over a few hundred ${\\rm km\\,s}^{-1}$.  Detection levels for heavy element transitions in DLAs are $\\sim$\\,7 orders of magnitude below that for $N($H{\\sc \\,i}$)$. High resolution optical spectra reveal transitions from many redshifted UV species, allowing detailed heavy element abundance patterns to be studied. The metallicity of DLAs -- defined as the heavy element abundance with respect to hydrogen, relative to that of the solar neighbourhood: ${\\rm [M/H]} \\equiv \\log_{10}[N({\\rm M})/N({\\rm H})] - \\log_{10}[N({\\rm M})/N({\\rm H})]_\\odot$ -- ranges over 2 orders of magnitude and is usually $\\sim$\\,0.1\\,solar or lower.  The presence of dust in DLAs is revealed by the relative depletion of refractory elements (e.g.~Fe) with respect to those thought not to readily deplete onto dust grains \\citep[e.g.~Zn;][]{PettiniM_97a}, and by the suppression of blue light from the background quasar \\citep[e.g.][]{FallS_89a}.  These spectroscopic observations demonstrate that DLAs arise along lines of sight through distant galaxies but they do not disclose details such as the galaxy's morphology, luminosity, mass or age.  Evidence for an increase in DLA metallicity with cosmic time has emerged only gradually over the past few years from a number of different studies \\citep{PettiniM_95a,LuL_96b,VladiloG_00a,KulkarniV_02a,ProchaskaJ_03b}. The latter reference provides the most compelling statistical evidence so far since, using the full sample of 125 DLAs, one sees a strong anti-correlation between [M/H] and absorption redshift, $z_{\\rm abs}$: Kendall's $\\tau$ is $-0.32$ with an associated probability of $P(\\tau)\\la 10^{-7}$. The need for such a large sample is due to the $\\sim$\\,2\\,dex scatter in [M/H] at a given epoch.  The types of galaxies responsible for DLAs can be investigated using direct imaging methods or comparison of spectroscopic data with models. The kinematics of optical resonance absorption transitions in low redshift DLAs appear consistent with 2-component models: a rotating disk plus an extended halo \\citep{BriggsF_85a}.  A recent detailed radio study of one particular DLA at $z_{\\rm abs} = 0.437$ clearly shows that at least {\\it some} DLAs arise in the rotating disks of ordinary spiral galaxies \\citep{BriggsF_01a}.  High redshift ($z_{\\rm abs} \\sim 3$) DLAs have been interpreted similarly, using detailed kinematic modelling to argue that they are rotating ensembles of clouds with further evidence for this interpretation coming from comparison of abundance ratios in DLAs and the thick component of our Galaxy's disk \\citep{ProchaskaJ_97b,WolfeA_01a}.  On the other hand, direct imaging of DLAs at $z_{\\rm abs} < 1.5$ reveals the hosts to be a mix of irregulars, spirals and low surface-brightness galaxies \\citep[LSBs; e.g.][]{LeBrunV_97a,RoaS_03a}, a result further borne out by a blind 21-cm emission survey at $z=0$ \\citep{Ryan-WeberE_03a}. The observed number of DLAs per unit redshift interval and the $N($H{\\sc \\,i}$)$ distribution suggest that DLAs at $z_{\\rm abs} < 2$ are a roughly equal mix of LSBs and spirals; at higher redshift they are more likely to be dwarfs (that subsequently merge) since there are too few precursors of present-day disk galaxies to explain the data \\citep{BoissierS_03b,BoissierS_03a}. Fitting simple chemical evolution models to an observed $N($H{\\sc \\,i}$)$--[Zn/H] distribution suggests that at redshifts $z_{\\rm abs} \\sim 2.5$, about half of all DLAs are due to dwarf galaxies, while at $z_{\\rm abs} < 1$, giant galaxies may dominate the population \\citep{BakerA_00a}. H{\\sc \\,i} 21-cm absorption measurements also hint that dwarfs could dominate at high-$z$ \\citep{KanekarN_03a}.   Further work is clearly needed to generate an accurate inventory of DLA host galaxy types as a function of cosmic epoch.  Nevertheless, the evidence for a broad mix of galaxy types is already compelling, if not conclusive. ",
        "conclusions": ""
    },
    "0311/astro-ph0311161_arXiv.txt": {
        "abstract": "A previously inconspicuous knot in the M87 jet has undergone a dramatic outburst and now exceeds the nucleus in optical and X-ray luminosity. Monitoring of M87 with the {\\it Hubble Space Telescope} and {\\it Chandra} X-ray Observatory during 2002-2003, has found month-timescale optical variability in both the nucleus and HST-1, a knot in the jet $0.82''$ from the nucleus. We discuss the behavior of the variability timescales as well as spectral energy distribution of both components.  In the nucleus, we see nearly energy-independent variability behavior.  Knot HST-1, however, displays weak energy dependence in both X-ray and optical bands, but with nearly comparable rise/decay timescales at 220 nm and 0.5 keV. The flaring region of HST-1 appears stationary over eight months of monitoring.  We consider various emission models to explain  the variability of both components.  The flares we see are similar to those seen in blazars, albeit on longer timescales, and so could, if viewed at smaller angles, explain the extreme variability properties of those objects. ",
        "introduction": "M87 is among the nearest galaxies with a bright radio/optical/X-ray jet. Its proximity (distance 16 Mpc, Tonry 1991, giving a scale of 78 pc/arcsec) allows features to be studied with unparalleled spatial resolution.     The dynamic nature of the M87 jet was first  recognized in the X-rays by Harris et al. (1997, 1998), who found variability on  $\\sim 1$ year timescales. In the optical, Biretta et al. (1999) found superluminal motion throughout the inner $10''$ of the jet, with speeds up to  $6c$, as well as gradual changes in flux. These observations did not explore timescales $\\lsim 6$ months, however. More recently, the jet's optical spectrum was found to harden in knots, consistent with local particle acceleration (Perlman et al. 2001a).  And {\\it Chandra} observations found spectral indices   $\\alpha_x > 1$  $(S_{\\nu} \\propto \\nu^{-\\alpha})$, broadly consistent with an extrapolation of the radio-optical synchrotron emission (Marshall et al. 2002, Wilson \\& Yang 2002).  X-ray synchrotron emission implies radiative lifetimes $\\sim$ 1-10 years, and thus requires {\\it in situ} acceleration.  All this evidence suggested the possibility of shorter-timescale variability in the M87 jet.  Indeed, recently Harris et al. (2003, hereafter paper I) detected X-ray flaring on timescales $\\sim 1 $ month during 2002 in both the nucleus and HST-1, a knot in the jet. Here we discuss observations of the M87 jet with both {\\it HST} and the {\\it Chandra} X-ray Observatory, which find flares in both the nucleus and knot HST-1 on timescales of $\\sim 2$ months. A second paper (Biretta et al. 2003, hereafter paper III) discusses the longer timescale variability of the M87 jet.   \\section {Observations and Data Reduction}  We make use of {\\it Chandra} and {\\it HST} observations of M87 during November 2002 - June 2003.  During this interval we observed M87  roughly every 20-40 days.  Table 1 lists the  {\\it HST} observations; the {\\it Chandra} data will be discussed in an upcoming paper (Harris et al., in prep.).  Our {\\it Chandra} data reduction procedures were detailed in Paper I. All {\\it HST} observations were reduced in IRAF and PyRAF with the best available flat fields, biases, darks and illumination correction images.  Dithered images were combined using PYDRIZZLE, which combines drizzling (Fruchter \\& Hook 2002) with cosmic ray and hot pixel removal, geometric correction and mosaicing. Galaxy subtraction was done on the F475W and F814W data using ELLIPSE, BMODEL and IMCALC.  Identical galaxy models were used at all  epochs.    Flux-calibrated images were obtained using SYNPHOT.  We applied a standard Galactic extinction law and assumed $N(H)=2.4 \\times 10^{20} {\\rm cm^{-2}}$ to correct the fluxes for reddening.  Polarized light images were combined into Stokes' parameters in AIPS, using standard formulae.  Considerable effort was expended to estimate uncertainties. We accounted for several error sources, including zero-point and slope errors in the SYNPHOT calibration, flat-fielding errors, and Poisson errors in the data and in modelling and subtraction of the galaxy.  Errors were propagated by adding in quadrature.  The cumulative error in the {\\it HST} photometry is typically 3-4\\%, and is dominated by the flat-fielding and zero-point errors.  Lightcurves were extracted from our data at three UV/optical wavelengths - 220 nm, 475 nm and 814 nm, and three X-ray bands (paper I) - 0.2-0.75 keV (``soft''; nominal 0.5 keV), 0.75-2 keV (``medium''; nominal 1.4 keV), and 2-6 keV (``hard''; nominal 4 keV). The 2003 March 31 and 2003 May 10 (part of a snapshot program led by D. Maoz) images were taken at a slightly longer wavelength; to place them on a common 220 nm flux scale we assumed $\\alpha_o=0.6$ (Perlman et al. 2001a).  Finally, we have also made use of flux points from 2001-2002 {\\it HST} data which are discussed in more detail in Paper III.  \\section {Results}    Figure 1 shows the 2003 April 17 F220W image of the inner $4.5''$ of the M87 jet,  as well as optical and X-ray lightcurves for the nucleus and HST-1 respectively.  As can be seen, both components vary on timescales of a few months in the optical, with typical month-to-month variations being about 10-15\\% at 220 nm. Other regions do not vary significantly from month to month and so are not shown in Figure 1.  For the nucleus, flux was extracted from $0.2''$ and $0.6''$ square boxes, which showed identical variability.  In Figure 1 we use the  $0.2''$ box.  For knot HST-1, we extracted flux from a $0.6''$ box.  The extraction regions for our {\\it Chandra} data were described in paper I. Our data confirm Tsvetanov et al.'s  (1998) claim of month-timescale nuclear optical variability.  However, this is the first such report for any jet component.  \\begin{deluxetable}{ccc} \\tablewidth{0pt} \\tablecolumns{3} \\tablecaption{Log of HST Observations} \\tablehead{ \\colhead{Date} & \\colhead{Instrument \\& Bands} & \\colhead{Program}} \\startdata 2002-11-30 \t& ACS F220W, F475W, F814W & 9705\\\\ 2002-12-07 \t& ACS F606W + POLVIS & 9705\\\\ 2002-12-22\t&ACS  F220W, F475W, F814W & 9705\\\\ 2003-02-02\t& ACS  F220W, F475W, F814W & 9705\\\\ 2003-03-06\t&  ACS F220W, F475W, F814W & 9705\\\\ 2003-03-31\t& ACS F250W, F330W & 9454 \\\\ 2003-04-17\t& ACS F220W, F475W, F814W & 9705\\\\ 2003-05-10\t&ACS  F250W, F330W & 9454 \\\\ 2003-06-07\t& STIS F25QTZ & 9474 \\enddata \\end{deluxetable}      \\begin{figure} \\centerline{\\includegraphics*[scale=0.45]{fig1.ps}}  \\caption[]{At top, we show the 2003 April 17 HST/ACS F220W image of the inner $4.5''$ of the M87 jet. The image has been rotated so that the jet appears along the {\\it x}-axis.  The middle and bottom panels show lightcurves for the nucleus and HST-1.  Some of the optical/UV lightcurves are multiplied by arbitrary factors to separate them from the X-ray lightcurves. Flare events discussed in the text are enumerated.  All error bars are at the $1\\sigma$ level.}  \\end{figure}  At 220 nm, $>95\\%$ of the flux in both varying regions is contained within a distribution consistent with the PRF.  Thus both are  unresolved, and so $<0.02''$ (1.5 pc) in size.  This implies light-travel times within a factor 3-20 (depending on the Doppler factor $\\delta=[\\Gamma(1-\\beta \\cos \\theta)]^{-1}$) of the variability timescales (\\S4, Paper I).  The nucleus's variable region is $<0.01''$ from the galactic center,  with HST-1's varying region being $0.82''$ (64 pc) distant, locations identical to $\\pm~ 0.01''$ of those measured with {\\it Chandra} (Paper I).  We detect no motion in HST-1's flaring region;  therefore, the flare likely occurred in the nearly stationary component at its upstream end (Biretta et al. 1999), rather than one that propagates at the velocities seen in moving components within HST-1, namely $v_{app} = 6c$ ($0.013''$/year, easily detectable  in these data).   Figure 2 shows the F606W polarized light image. As can be seen, the nucleus is essentially unpolarized.  HST-1 shows fractional polarization $P=0.46$ at its flux peak, a maximum $P=0.68$ (nearly the theoretical maximum for synchrotron radiation from an ordered magnetic field) at its  upstream end $\\sim 0.72''$ from the nucleus, and a minimum $P=0.23$ $0.92''$ from the nucleus.  The magnetic field vectors in HST-1 are perpendicular to the jet direction, consistent with a shock.  The alternative interpretation of a tightly wound helix would require cycles separated by $<0.02''$ to be consistent with the observed morphology. A much higher bulk $\\Gamma$ would also be needed (because of the longer path) to be consistent with $v_{app}=6c$ (Biretta et al. 1999).  The polarization at HST-1's flux peak is much higher than in 1995 ($P=0.14$, Perlman et al. 1999). This is in line with the properties of BL Lac objects, which often have higher optical polarizations in high states (e.g., Hagen-Thorn et al. 1998).  \\begin{figure}  \\centerline{\\includegraphics*[angle=270,scale=0.35]{fig2.ps}}  \\caption[]{An image of the innermost $1.5''$ of the M87 jet, with polarization vectors (magnetic field) superposed.  The magnitude of the vector is proportional to  the percent polarization, with a $0.02''$ vector representing 100\\% polarization. Contours are shown at ( 0.2, 0.4, 0.6, 0.8, 1, 1.5, 2, 4, 6, 10, 20) ADU/sec}  \\end{figure}   \\section {Constraints from the Lightcurves}  To constrain the physics, we compute doubling and halving timescales $\\tau_2, \\tau_{1/2}$, given in Table 2 for timeperiods referenced in Figure 1.  We can compare these with predictions from models with $E^2$ losses, which are of the form $\\tau\\propto \\nu^\\eta$.   If the emission is synchrotron radiation and synchrotron cooling dominates, $\\eta=-0.5$, (Kirk, Mastichiadis \\& Rieger 1998; Paper I).  As shown in B\\\"ottcher et al. (2003),  if synchrotron self-Compton (SSC) losses dominate, the dominant factor in $d\\gamma/dt$ is the synchrotron radiation energy density $U_{sy}$, so that $\\eta= (q-4)/2$ = --1 to --0.75 for electron spectral index $q=2.5$ to 2 [$N(\\gamma) \\propto \\gamma^{-q}$, which is related to $\\alpha$ via $\\alpha=(q-1)/2$].  \\begin{deluxetable*}{llllcccccc} \\tablewidth{0pt} \\tablecolumns{10} \\tablecaption{Flare Timescales} \\tablehead{ & & & & \\multicolumn{6}{c}{Doubling or Halving Timescale in Band (days)} \\\\ \\colhead{Component} & \\colhead{Band} & \\colhead{Dates} & \\colhead{Type (Flare \\#)} & \\colhead{814 nm} & \\colhead{475 nm} & \\colhead{220 nm} & \\colhead{0.5 keV} & \\colhead{1.4 keV} & \\colhead{4 keV} } \\startdata Nucleus & X-rays & June-July 2002 & Rise (2) & ... & ... & ... & $135 \\pm 39$ &  $608 \\pm 474$ & $83 \\pm 20$  \\\\ Nucleus & X-rays & March-May 2002 & Fall (1) & ... & ... & ...& $78 \\pm 11$ & $71 \\pm 6$ & $67 \\pm 9$   \\\\ Nucleus & Optical & Dec. 2002-Feb. 2003 & Rise (2)& $317 \\pm 108$ & $206 \\pm 47 $ & $ 256 \\pm 82$ & ... & ... & ... \\\\ Nucleus & Optical & Mar.-Apr. 2003 & Fall (2) & $81 \\pm 12$ & $83 \\pm 12$ & $274 \\pm 106$&  ... & ... & ...  \\\\ HST-1 & X-rays & Feb.-Apr. 2003 & Rise (2) & ... & ... & ... & $223 \\pm 49$ & $118 \\pm 13$ & $ 61 \\pm 9$  \\\\ HST-1 & X-rays & Nov.-Dec. 2002 & Fall (1) & ... & ... & ... & $149 \\pm 55$ & $91 \\pm 17$ & $54 \\pm 10$ \\\\ HST-1 & Optical & Feb.-Apr. 2003 & Rise (2) & $359 \\pm 82$ & $ 414 \\pm 108$ & $243 \\pm 36$ & ... & ... & ... \\\\ HST-1 & Optical & Dec. 2002 - Feb. 2003 & Fall (1) & $260 \\pm 133$ & $217 \\pm 91$ & $139 \\pm 37$& ... & ... & ... \\\\ \\enddata \\end{deluxetable*}  Table 2 shows several differences between the variability behavior of the nucleus and HST-1.   The nucleus does not show strong energy-dependence in either the X-rays ($\\eta_{x,rise}=-0.24 \\pm 0.21$ and $\\eta_{x,fall}=-0.07 \\pm 0.10$) or optical ($\\eta_{o,rise}=-0.09 \\pm 0.43$ and $\\eta_{o,fall}=0.52 \\pm 0.32$, but the departure from zero in the latter measurement is due entirely to the lower significance of the fall in F220W during 2003 Feb.-Apr.; further, note that the drop accelerated after monitoring in F475W and F814W ceased).   By comparison, HST-1 does show energy-dependent behavior in the X-rays, with $\\eta_{x,rise}=-0.62 \\pm 0.13$ and $\\eta_{x,fall}=-0.49 \\pm 0.21$. But in the optical, our data are less constraining, giving  $\\eta_{o,rise}=-0.35 \\pm 0.22$ and $\\eta_{o,fall}=-0.51 \\pm 0.78$.  HST-1's behavior in the X-rays agrees with the predictions of a simple synchrotron model.  But such a model cannot easily account for the comparable optical and X-ray variability timescales.  Two other models can explain such behavior.  The first is that the flare was caused by adiabatic  compression, followed by expansion on a dynamical timescale of  $\\tau_{dyn}\\sim 200$ days. This would allow radiative losses dominate at high energies, where  $\\tau_{syn} <  \\tau_{dyn}$, while expansion losses would dominate at low energies. The X-ray variability behavior then requires $B \\approx 2 ~\\delta^{-1}$ mG. However, this model has difficulty explaining X-ray and optical increases and decreases that do not exactly coincide in time (Figure 1).  Alternately, the X-ray flare could be triggered by shock compression (as suggested by the polarimetry, \\S 3), with the optical emission representing the shocked plasma's downstream propagation.  This model requires a $\\sim 10\\times$ stronger magnetic field downstream of the shock to explain the similar 0.5 keV and 220 nm variability timescales.  The nearly energy-independent variability behavior of the nucleus  disagrees with simple synchrotron models, although some cannot be excluded  formally because of the large error bars. One could account for the energy independence either by adiabatic compression and expansion, or a helical trajectory, where flux would change with viewing angle  (suggested by Urry et al. 1993 to explain the nearly energy-independent variability of PKS 2155$-$304 in 1991).   But perhaps the most attractive model is that the varying region has $R/c \\approx 70 ~\\delta$ light-days, so that the light-crossing time controls the variability behavior (e.g., Chiaberge \\& Ghisellini 1999).  In this case we would require $B\\lsim 20 ~\\delta^{-1}$ mG.  \\section {Comparison to Blazar Flares}  The flares of the nucleus and HST-1 (Figure 2) resemble those of blazars, albeit on longer timescales. Can we explain blazar flares ($\\tau \\sim 0.2-5$ days; Ulrich, Maraschi \\& Urry 1997; Pian 2002), as highly beamed versions of these events?  Timescales Lorentz-transform as $\\tau^\\prime \\propto \\tau \\delta^{-1}$, so $\\tau=5$ days requires $\\delta_{M87}/\\delta_{blazar}\\approx 1/12$ (also consistent with the luminosities seen in blazar flares, e.g., Giebels et al. 2002), while $\\tau=1$ day requires $\\delta_{M87}/\\delta_{blazar}\\approx 1/60$. For reasonable $\\Gamma$, this requires $\\delta_{M87} \\lsim 2$, at the lower end of the range considered for HST-1 in Paper I,   and implies $\\delta_{blazar} = 20-100$.   Consistency with HST-1's $v_{app}=6c$ (Biretta et al. 1999), then requires $\\Gamma > 10$ and $\\theta = 15-19^\\circ$,  much tighter constraints than in previous works. A timescale $\\tau=5$ days then requires $\\Gamma=10$ and $\\theta=1^\\circ$ if $\\delta_{M87}=1.6 ~ (\\Rightarrow \\Gamma_{M87}=10, \\theta_{M87}=19^\\circ$), while $\\tau=1$ day requires $\\Gamma>25$ and $\\theta<1^\\circ$, plus smaller $\\delta_{M87}$. Unified schemes usually constrain $\\Gamma$ and $\\theta$ less tightly (e.g., $\\Gamma= 5-30$ and $\\theta < 10-30^\\circ$, Urry \\& Padovani 1995).  However, one can imagine a range of $\\delta$ or stronger magnetic fields (e.g., Li \\& Kusunose 2000, Kataoka et al. 2000, B\\\"ottcher et al. 2003).  Blazars can show either energy-dependent or energy-independent behavior, even within the same campaign.  Often, energy-dependent behavior includes one band leading the other. With our data, any discussion of delays is premature, although there are possible indications in the data.  Monitoring of M87 continues, and we will discuss this subject in a later paper.    \\begin{figure}  \\centerline{\\includegraphics*[scale=0.45]{fig3.ps}}  \\caption[]{Spectral Energy Distributions for the nucleus (top) and knot HST-1 (bottom).  Three epochs are shown:  2002 December, 2003 February, and 2003 April.}  \\end{figure} ",
        "conclusions": ""
    },
    "0311/astro-ph0311211_arXiv.txt": {
        "abstract": "High redshift supernovae and Cosmic Microwave Background data are used to constrain the Cardassian expansion model (Freese \\& Lewis, 2002), a cosmology in which a modification to the Friedmann equation gives rise to a flat, matter-dominated Universe which is undergoing a phase of accelerated expansion. In particular, the precision of the positions of the Doppler peaks in the CMB angular power spectrum provided by the Wilkinson Microwave Anisotropy Probe (WMAP) tightly constrains the cosmology. The available parameter space is further constrained by various high redshift supernova datasets taken from Tonry et al. (2003), a sample of 230 supernovae collated from the literature, in which fits to the distance and extinction have been recomputed where possible and a consistent zero-point has been applied. In addition, the Cardassian model can also be loosely constrained by inferred upper limits on the epoch at which the Cardassian term in the modified Friedmann equation begins to dominate the expansion ($z_{eq}$). Using these methods, a Cardassian cosmology is constrained at the 2$\\sigma$ level to 0.19$<\\Omega_m$\\lsim0.26, 0.01$<n<$0.24 for the Cardassian expansion parameter, $n$, and 0.42$<z_{eq}<$0.89, in contradiction to the previous constraints of Sen \\& Sen (2003a). There is also a large discrepancy between the 1$\\sigma$ confidence regions defined by the CMB and tightest Supernova constraints, with the CMB data favouring a low-$\\Omega_m$, high-$n$ Cardassian cosmology and $z_{eq}>$1, as opposed to the Supernova data which supports a high-$\\Omega_m$, low-$n$ cosmology. ",
        "introduction": "Over the last ten years, the $m$-$z$ relation below $z\\approx$1 has been constrained through observations of distant type Ia supernovae, and indicates that the Universe is currently in a phase of accelerated expansion (e.g. Perlmutter et al. 1997; Riess et al. 1998). This is supported by recent measurements of the Cosmic Microwave Background angular power spectrum and the galaxy power spectrum. The standard model provides an interpretation of these results by proposing that the Universe consists of only $\\sim$4\\% baryonic matter with the remainder in the form of Cold Dark Matter and a dark energy component which drives the expansion. This cosmology provides a good fit to the Wilkinson Microwave Anisotropy Probe (WMAP) CMB power spectrum and high redshift supernova data, and is consistent with local large-scale structure and cluster baryon fraction measurements of the matter density. This cosmology is slightly problematic however as it invokes the use of two unobserved forms of energy. While the evidence for Cold Dark Matter is considerable, evidence for the existence of a dark energy component remains indirect.  An interesting alternative to the standard model of cosmology recently proposed by Freese \\& Lewis (2002) invokes an additional term in the Friedmann equation, proportional to $\\rho^n$, which results in a Universe which is consistent with the recent evidence for an increasing expansion rate, but is both flat and matter dominated. The need for a dark energy component is therefore removed, and the expansion is driven solely by the new $\\rho^n$ term in the modified Friedmann equation. The theoretical motivation for this Cardassian model of cosmology (see Freese \\& Lewis 2002) is fairly speculative however, and the implications for Einstein's equations of General Relativity are currently undetermined.  In order to constrain the Cardassian model, it is critical that the observational data used do not assume a prior cosmology. The simplest constraints can therefore be imposed by the CMB angular power spectrum, high redshift supernova data and local large-scale structure, since in a Cardassian cosmology these observables depend on $\\Omega_m$, the Cardassian parameter $n$, and the redshift at which the Cardassian term in the modified Friedmann equation begins to dominate, $z_{eq}$.  Previously, Sen \\& Sen (2003a) used predictions for the locations of the first and third Doppler peaks in the CMB power spectrum for a Cardassian cosmology, and Archeops and Boomerang data to constrain possible values for $\\Omega_m$ and $n$. Supernova data from the Calan-Tololo project and the Supernova Cosmology Project (Perlmutter et al. 1997) were also used to constrain these parameters, since the modification of the luminosity distance in a Cardassian cosmology alters the predicted apparent magnitude for supernovae. These two constraints restricted the parameters to 0.31\\lsim$n$\\lsim0.44 and 0.13\\lsim$\\Omega_m$\\lsim0.23. Zhu \\& Fujimoto (2003) also used these supernova data to constrain the Cardassian model, and found best fit parameters of $n$=-1.33 and $z_{eq}$=0.43 assuming $\\Omega_m$=0.3. However, it is not clear how the fit is affected for different values of $\\Omega_m$. More recently Sen \\& Sen (2003b) used WMAP and Boomerang measurements of the locations for the first, second and third Doppler peaks to further constrain the parameter space.  In this paper, the parameter space is constrained using the WMAP and Boomerang measurements for the locations of the first, second and third Doppler peaks in the CMB angular power spectrum used in Sen \\& Sen (2003b), and various high redshift supernova datasets taken from Tonry et al. (2003). An orthogonal constraint is also applied using an assumed upper limit of $z_{eq}$=1; this value is used in order to be consistent with local large-scale structure (Freese \\& Lewis, 2002). In Section 2, the Cardassian model is outlined. The supernova analysis method and the resulting constraints for various samples are presented in Section 3, along with the CMB constraints and the additional constraint imposed by the inferred upper limit on $z_{eq}$. The discussion and conclusions follow in Section 4.  \\begin{figure} \\begin{center} \\centerline{\\epsfxsize = 3in \\epsfbox{fig1.ps}} \\caption{Contours of $\\Omega_m$ with the Cardassian parameters $n$ and $z_{eq}$. $\\Omega_m$=0.1 (dot-dashed), $\\Omega_m$=0.2 (dashed), $\\Omega_m$=0.3 (solid) and $\\Omega_m$=0.4 (dotted) tracks are shown.} \\label{fig:f_fact} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\centerline{\\epsfxsize = 3in \\epsfbox{fig2.ps}} \\caption{Contours of $z_{eq}$ with the Cardassian parameter, $n$, and $\\Omega_m$. Contours for $z_{eq}$=1.2 (solid), $z_{eq}$=1.0 (large-dashed), $z_{eq}$=0.8 (dotted), $z_{eq}$=0.6 (small-dashed), $z_{eq}$=0.4 (dot-dashed) and $z_{eq}$=0.2 (dot-dot-dashed) are shown.} \\label{fig:z_eq} \\end{center} \\end{figure} ",
        "conclusions": "The Cardassian expansion model utilises an additional term in the Friedmann equation, proportional to $\\rho^n$, such that the Universe is flat, matter-dominated and in a phase of accelerated expansion. The need for a dark energy component is therefore removed, as the expansion observed in distant supernovae is driven solely by this additional component. In order to constrain the parameters of the Cardassian model $n$, $\\Omega_m$ and $z_{eq}$ (the epoch at which the Cardassian term begins to dominate the expansion), it is necessary to select observational data free from any cosmological priors.  In this paper, various samples of type Ia supernova data from Tonry et al. (2003) were used to constrain $\\Omega_m$ and $n$. Fig.~\\ref{fig:sne_scp} shows the 1$\\sigma$, 2$\\sigma$ and 3$\\sigma$ confidence levels derived from the 54 supernovae included in Fit C of Perlmutter et al. (1997). This is in general agreement with the constraints imposed by Sen \\& Sen (2003a) for the same sample of supernovae, although the constraints are slightly broader, perhaps due to their use of contouring for absolute values of $\\chi^2$. Since the errors are uncertain, more robust confidence levels can be derived when contouring around a minimum value of $\\chi^2$ as used for the supernova data in this paper. Sen \\& Sen (2003a) also use a different method of analysis in which a different value of ${\\mathcal{M}}_B$ is used with an uncertain  associated error.  The tightest constraints from the supernova data are provided by the 172 supernovae sample (Fig.~\\ref{fig:sne_fit}) for which an extinction limit of $A_V$=0.5 and a redshift cut of $z>$0.01 have been applied in order to remove objects with potentially large effects from perculiar velocities at low redshift and large host galaxy extinctions. The resulting 1$\\sigma$ constraints favour a value of 0.19$<\\Omega_m<$0.32 and $n<$0.13. However, the available parameter space within the 3$\\sigma$ confidence region is large, although this represents a significant improvement over previous constraints.  To constrain the available parameter space further, WMAP (Hinshaw et al. 2003) and Boomerang (Hu et al. 2001) measurements for the first three Doppler peaks in the CMB angular power spectrum are compared to the associated Cardassian prediction as in Sen \\& Sen (2003b) (Fig.~\\ref{fig:map}). These three constraints are in remarkably good agreement at $\\Omega_m\\sim$0.15, but begin to diverge at the values of $\\Omega_m$ and $n$ suggested by the 172 supernovae sample.  In Fig.~\\ref{fig:all}, the combined constraints from the first three Doppler peaks and the 172 supernovae sample are combined. Also included is the $z_{eq}$=1.0 track which represents an approximate upper limit arising from the effect of the Cardassian term in the modified Friedmann equation on local large-scale structure (Freese \\& Lewis, 2002). The 1$\\sigma$ constraint imposed by the combined CMB data indicates values of 0.15\\lsim$\\Omega_m$\\lsim0.18, 0.28$<n<$0.36 and $z_{eq}>$1.0. However, there is good agreement at the 2$\\sigma$ level with constraints of 0.19$<\\Omega_m$\\lsim0.26, 0.01$<n<$0.24 and 0.42$<z_{eq}<$0.89. This is consistent with the loose constraint of $z_{eq}<$1.0 imposed by local large-scale structure. However, it is in contradiction to the constraints derived by Sen \\& Sen (2003) of 0.13\\lsim$\\Omega_m$\\lsim0.23 and 0.31$<n<$0.44.  In conclusion, high redshift supernova data and recent mesaurements of the CMB power spectrum provide relatively tight constraints on the available parameter space of the Cardassian expansion model. There is a large discrepancy between the CMB and supernova 1$\\sigma$ confidence regions, with the CMB 1$\\sigma$ constraint implying a parameter space which is potentially at odds with local large-scale structure and is rejected by the supernova data at the $>$2$\\sigma$ level. However, there is good agreement at the 2$\\sigma$ level, with the associated confidence regions limiting 0.19$<\\Omega_m$\\lsim0.26, 0.01$<n<$0.24 and 0.42$<z_{eq}<$0.89."
    },
    "0311/astro-ph0311027_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1} The formation and evolution of active galactic nuclei (AGN) and the associated star-formation activity are conditioned by the rate at which gas is supplied from the surrounding host galaxy. Interactions and mergers are identified as one possible mechanism to provide the transport of gas towards the central regions of a galaxy (\\cite{1992ARA&A..30..705B} and references therein). The hypothesis of a merger-driven evolution of active galaxies from ULIRGs to unobscured QSOs is motivated by the spectral energy distributions of template objects showing a sequence of decreasing far-infrared excess and increasing UV brightness \\cite{1988ApJ...325...74S}. In accordance with the assumption that the evolution is triggered by interactions, the vast majority of ULIRGs show evidence for ongoing merger processes \\cite{1999Ap&SS.266..331S}. Merger signs are less clear in the case of objects close to the QSO phase. These transition objects are important sources to confirm the evolutionary sequence. As a nearby representative ($z=0.0611$, \\cite{1997ApJ...478..144S}) of a template transition object in the original evolutionary sequence (Fig.~17 in \\cite{1988ApJ...325...74S}), I~Zw~1 is suitable to address the merger properties, the structural composition, and the starburst activity using near-infrared (NIR) and molecular line observations at sub-kpc angular resolution. ",
        "conclusions": ""
    },
    "0311/gr-qc0311060_arXiv.txt": {
        "abstract": "The recent observations of type Ia supernova strongly support that the universe is accelerating now and decelerated in the recent past. This may be the evidence of the breakdown of the standard Friedmann equation. We consider a general modified Friedmann equation. Three different models are analyzed in detail. The current supernova data and the Wilkinson microwave anisotropy probe data are used to constrain these models. A detailed analysis of the transition from  the deceleration phase to the acceleration phase is also performed. ",
        "introduction": "The recent observations of Type Ia supernova indicate that the expansion of the universe is speeding up \\cite{sp99}. Observational results also provide the evidence of a decelerated universe in the recent past \\cite{agr}. On the other hand, the cosmic background microwave (CMB) observations indicate that the universe is spatially flat as predicted by the inflationary models \\cite{pdb00}. A dark energy component with negative pressure behavior which dominates the universe, was proposed to explain a flat and accelerating universe. One simple candidate of dark energy is the cosmological constant. However, there are some problems with the cosmological constant although cold dark matter cosmological constant models are consistent with the current observations. Why is the cosmological constant so small and not zero? Why does the cosmological constant become significant now? The quintessence models avoid some of the problems \\cite{cds98}. There are also other models, like tachyon filed as dark energy \\cite{padmt03}. But the property of dark energy is still mysterious. One logical possibility is that the standard Friedmann equation may need to be modified. In this scenario, the universe is dominated by ordinary pressureless matter, but the law of gravity and the standard Friedmann equation are modified. The idea of modifying the law of gravity is not new. The modified Newtonian Dynamics (MOND) was first used to explain the rotation curve in place of dark matter \\cite{mond,mond1}. In MOND, the Newtonian gravity $M/r^2$ is replaced with $M/r^2+\\sqrt{M}/r$. Since $M/r^2$ gives the standard Freidmann equation $H^2\\sim \\rho$, $M/r^2+\\sqrt{M}/r$ may provide a modified Friedmann equation $H^2\\sim \\rho^{2/3}\\ln \\rho+ \\rho^{2/3}$ \\cite{mond1}. Recall that the brane cosmology gives a non-standard Friedmann equation $H^2\\sim \\rho +\\rho^2$ \\cite{bdl,rs99,gong00}. Along this line of reasoning, Freese and Lewis recently proposed the Cardassian expansion in which the universe is dominated by the ordinary matter and the Friedmann equation becomes $H^2\\sim \\rho+\\rho^n$ \\cite{freese02}. The Cardassian model was later generalized to a more general form $H^2\\sim g(\\rho)$ \\cite{chung}. In addition, several authors modified the Friedmann equation as $H^2+H^\\alpha \\sim \\rho$ motivated by theories with extra dimensions \\cite{dvali00,lue03a}. In this paper, we first consider a model which is equivalent to the generalized Chaplygin gas model \\cite{kamenshchik} in terms of dynamical evolution. Then we consider the generalized Cardassian model. At last we consider a model proposed by Dvali, Gabadadze and Porrati (DGP) \\cite{dvali00}. We use the ten new supernovae at $z=0.36-0.86$ \\cite{raknop03} and the the first year Wilkinson Microwave Anisotropy Probe (WMAP) temperature (TT) and temperature polarization cross correlation (TE) data \\cite{ghinshaw} to constrain these models. We also investigate the transition from the decelerated phase to the accelerated phase.  For a spatially flat, isotropic and homogeneous universe with both an ordinary pressureless dust matter and a minimally coupled scalar field $Q$ sources, the Friedmann equations are \\begin{eqnarray} \\label{cos1} H^2=\\left({\\dot{a}\\over a}\\right)^2={8\\pi G\\over 3}(\\rho_{\\rm m}+\\rho_{\\rm Q}),\\\\ \\label{cos2} {\\ddot{a}\\over a}=-{4\\pi G\\over 3}(\\rho_{\\rm m}+\\rho_{\\rm Q}+3p_{\\rm Q}),\\\\ \\label{cos3} \\dot{\\rho_{\\rm Q}}+3H(\\rho_{\\rm Q}+p_{\\rm Q})=0, \\end{eqnarray} where dot means derivative with respect to time, $\\rho_{\\rm m}=\\rho_{\\rm m0}(a_0/a)^3$ is the matter energy density, a subscript 0 means the value of the variable at present time, $\\rho_{\\rm Q}=\\dot{Q}^2/2+V(Q)$, $p_{\\rm Q}=\\dot{Q}^2/2-V(Q)$ and $V(Q)$ is the potential of the quintessence field. The modified Friedmann equations (MFE) for a spatially flat universe are \\begin{eqnarray} \\label{cosa} H^2=H_0^2g(x),\\\\ \\label{cosb} {\\ddot{a}\\over a}=H^2_0g(x) -{3H^2_0x\\over 2}g'(x)\\left({\\rho+p\\over \\rho}\\right),\\\\ \\label{cosc} \\dot{\\rho}+3H(\\rho+p)=0, \\end{eqnarray} where $x=8\\pi G\\rho/3H^2_0=x_0(1+z)^3$ during the matter dominated epoch, $1+z=a_0/a$ is the redshift parameter, $g(x)=x+\\cdots$ is a general function of $x$ and $g'(x)=dg(x)/dx$. From phenomenological point of view, four dimensional gravity is modified so that we get a general function $g(x)$. Chung and Freese argued that a general $g(x)$ is possible if our universe as a three brane is embedded in five dimensional spacetime \\cite{chung}. For example, $g(x)\\sim x+x^2$ in Brane cosmology, $g(x)\\sim x+x^n$ in Cardassian model. On the other hand, we can think the additional terms $g(x)-x$ as dark energy component. For the Cardassian model, the additional term $x^n$ can be mapped to a dark energy component with constant equation of state parameter $\\omega_{\\rm Q}=n-1$. In general, we get the following relationship between the dark energy equation of state parameter and $g(x)$ \\begin{equation} \\label{darkom} \\omega_{\\rm Q}={xg'(x)-g(x)\\over g(x)-x}. \\end{equation} The equivalent dark energy potential can be found from the following equations \\begin{eqnarray} \\dot{Q}^2=\\rho_{\\rm m}[g'(x)-1],\\\\ V(Q)={3H^2_0\\over 8\\pi G}[ g(x)-0.5x-0.5xg'(x)]. \\end{eqnarray} For instance, if $g(x)\\sim x +x^n$, we find that $V(Q)\\sim [\\sinh(AQ+B)]^{2n/(n-1)}$. Note that the universe did not start to accelerate when the other terms in $g(x)$ started to dominate. The linear density perturbation of this model is given by \\begin{equation} \\label{perturb} \\ddot{\\delta}+2\\bar{H}\\dot{\\delta}=4\\pi G\\bar{\\rho}\\delta[g'(\\bar{x})+3\\bar{x}g''(\\bar{x})]. \\end{equation} For the matter dominated flat universe, $\\rho=\\rho_{\\rm m}$ and $p=p_{\\rm m}=0$. Let $\\Omega_{\\rm m0}=8\\pi G\\rho_0/3H^2_0$, then $x_0=\\Omega_{\\rm m0}$, $g(x_0)=1$. In general, $x=\\Omega_{\\rm m0}(1+z)^3+\\Omega_{\\rm r0}(1+z)^4$, where $\\Omega_{\\rm r0}=8.35\\times 10^{-5}$ is the current radiation component \\cite{DNSpergel}. ",
        "conclusions": "One attractive feature of the generalized Chalygin gas model is that it can be considered as a unified dark matter and dark energy model. In the unified scenario, both the dark energy and dark matter components are modelled by the generalized Chaplygin gas. In this paper, we are interested in modelling the generalized Chaplygin gas as dark energy only. We found the best fit parameters in the generalized Chaplygin gas model are $\\Omega_{\\rm m0}=0.26$, $A_s=0.999$ and $\\beta=1.43$. Cunha, Alcaniz and Lima found that $A_s>0.73$ by using the combined x ray data of galaxy clusters and supernova data and taking a prior $\\Omega_{\\rm m0}=0.3$ \\cite{kamenshchik}. Amendola \\etal found that $1\\le\\beta< 1.2$ and $0.8< A_s< 1$ by using WMAP TT data \\cite{kamenshchik}. In the context of MFE, the generalized Chaplygin gas model tends to be the $\\Lambda$-model in order to recover the standard cosmology at early times. Our results show that the generalized Chaplygin gas model is almost the same as the $\\Lambda$-model. By using the best fit parameter, we get $z_{q=0}=0.78$.  For the Cardassian model, the WMAP TT and TE data give that $\\Omega_{\\rm m0}=0.25$ and $n=0.02$. The 10 supernova data gives that $\\Omega_{\\rm m0}\\sim 0$ and $n=0.51$. The error from supernova data fit is also large. Zhu and Fujimoto also found low $\\Omega_{\\rm m0}$ from supernova data \\cite{zhu03}. Sen and Sen found that $0.31< n < 0.44$ and $0.13< \\Omega_{\\rm m0}< 0.23$ from supernova data and peaks in CMB power spectrum, $n\\le 0.3$ from WMAP data \\cite{sen03}. Godlowski and Szydlowski found that $\\Omega_{\\rm m0}=0.48^{+0.08}_{-0.13}$ and $n=-0.4^{+0.77}_{-1.24}$, and $\\Omega_{\\rm m0}=0.51^{+0.05}_{-0.06}$ and $n=-1.2^{+0.77}_{-1.06}$ from different sets of supernova data \\cite{smcw}. Frith found that $0.19<\\Omega_{\\rm m0}<0.26$ and $0.01<n<0.24$ from supernova data \\cite{frith}. In general, different analysis gives different results, but they are still consistent with each other at 99\\% confidence level. Dev, Alcaniz and Jain found that the best fit to the gravitational lensing effect is $n=0.76$ and $\\alpha=2.4$ by assuming $\\Omega_{\\rm m0}=0.3$ for the generalized Cardassian model \\cite{dev03}. Taking the best fit result to WMAP data, we get $z_{q=0}=0.82$.  For the DGP model, the best fit parameter to the WMAP TT and TE data is $\\Omega_{\\rm m0}=0.174$ and the best fit parameter to the ten supernova data is $\\Omega_{\\rm m0}=0.21$. Deffayet \\etal found that $\\Omega_{\\rm m0}=0.18^{+0.07}_{-0.06}$ from supernova observations and $\\Omega_{\\rm m0}=0.3$ from CMB constraints \\cite{deffayet}. However, the constraint from supernova data and CMB data are still consistent with each other at $1\\sigma$ level. If we take the bigger value $\\Omega_{\\rm m0}=0.21$, then we get $z_{q=0}=0.81$. Multamaki, Gaztanaga and Manera considered the effects of Cardassian model and DGP model on large scale structure and found that these effects are different from that of $\\Lambda$-model \\cite{maltamaki}.  Due to the uncertainties in the observational data, it is still difficult to discriminate different dark energy models and alternative models. Wang \\etal showed that future supernova data such as SNAP \\footnote{http://snap.lbl.gov/} could differentiate various models if $\\Omega_{\\rm m0}$ is known with 10\\% accuracy \\cite{wang}. On the other hand, if we can determine the dynamical evolution of equation of state parameter of dark energy with high precision using model independent method, then we can discriminate different models. Alam \\etal showed that future SNAP data will provide important insights on the nature of dark energy at high redshifts \\cite{alam}. Turner and Riess showed that $z_{q=0}\\sim 0.5$ \\cite{agr}, Daly and Djorgovski found that $z_{q=0} >0.3$ \\cite{daly}. In conclusion, the above models are consistent with current observations.  In terms of model building, we can assume a particular scale factor which manifests early deceleration and later acceleration, then find out the form of $g(x)$. For example, we take $a(t)=a_0[\\sinh(t/t_0)]^\\beta/\\alpha$, then we find $g(x)=Ax^{2/3\\beta}+B$. This model recovers the standard cosmology when $\\beta=2/3$.  \\ack The author Gong is thankful to Padmanabhan T, Finelli F, Godlowski W, Doran M and Sahni V for comments. The work of Gong is fully financed by Chongqing University of Post and Telecommunication under grants A2003-54 and A2004-05."
    },
    "0311/astro-ph0311333.txt": {
        "abstract": " ",
        "introduction": "Possibly the most complete representative sample of the earliest local historical record is the system of globular clusters. For the majority of the clusters, we know within varying degrees of uncertainty their chemical abundances, motions, positions, and their ages. Unfortunately we still do not know how they formed. Because the evolution of individual stars is the only known way of enriching the gas from which the clusters formed, there must be a population of stars which formed with a distribution of heavy element abundance ending at that of a particular cluster. Further, because we know the heavy element distribution among clusters from observations, given a theory of chemical evolution we can calculate the abundance distribution of these accompanying stars. If as we shall assume below, that the clusters are formed in the shocked gas caused by collisions between isolated chemically evolving clumps for example, when the cluster is formed any remaining gas will lose energy but not angular momentum while the previously formed stars and accompanying dark matter will continue their dynamical evolution relatively unscathed.  As discussed below there appear to be two groups of clusters in the Galaxy. For convenience we shall refer to them as the blue and red groups. The blue group is relatively metal poor, has low angular momentum, is spatially extended far out into the halo, and is generally nearly as old as the universe itself. The red group is relatively metal rich, shows rapid rotation at all radii, and possesses a disk like spatial distribution with a concentration towards the Galactic center. The diversity between the two groups suggests separate stages of formation and in our chemical evolution model each group is evolved separately. In the following model, the main thrust is consideration of the effects of globular cluster formation on chemical evolution and not the formation mechanism itself. For the latter subject  we refer to the more sophisticated modelling of Cote et al.\\ (2000,2002), Beasley et al.\\ (2002), and Kravtsov \\& Gnedin (2003). Like the above cited works, ours is a bottom-up picture with the formation mechanism most resembling that for the metal rich clusters in the Beasley et al. work (i.e. collisions between gas rich clumps). For simplicity we assume that all clusters are formed in this way. In order to account for the bimodality, we argue that the collapse was anisotropic.  The objective of this work is to reconstruct the cosmic star formation history from local observations and in doing so to gain further insight into how galaxies might have formed and evolved. As will be shown, this reconstruction is successful, comparing favorably with that deduced from high redshift observations (the Madau plot). This implies that our model provides a scenario for the chemical evolution of a representative sample of the universe. Further, a proposed physical framework incorporates the results from this chemical evolution model and predicts a major new stellar halo component, which we identify in the Milky Way with the thick disk.   \\subsection{The Apparent Universality of Globular Cluster Bimodality}  Classical globular clusters like M92 belong to a group of relatively blue, metal poor clusters which are old and are found in the outer halo of the Galaxy. It was early work by Becker (1950), Baade (1958), Morgan (1959) and Kinman (1959) which  directed attention to the disk-like spatial distribution of a second group of globular clusters which are redder, more metal rich, some what younger, and concentrated towards the center of the Galaxy. An important further difference is that whereas the metal poor group shows very little rotation about the Galactic center, the metal rich group shows significant rotation at all Galactocentric radii (Cote, 1999, Zinn 1985). A clear separation of the two groups shows in the frequency histogram of [Fe/H] values for 133 Milky Way globular clusters in Fig. 8 of Cote (1999) along with best fit Gaussians for each group.  Recent spectroscopic studies of the globular clusters in M31 by Perrett et al.\\ (2002), and earlier work by Barmby et al.\\ (2000) and Huchra et al.\\ (1991) also show a bimodal distribution in [Fe/H] with Gaussian parameters similar to those of the Milky Way. In fact, the bimodality of globular cluster populations appears to be common in more distant luminous galaxies where it is usually observed as a bimodality in color distribution (e.g. Ashman \\& Zepf 1998, Gebhardt \\& Kissler-Patig 1999, Forbes et al. 2001, Kundu \\& Whitmore 2001, Goudfrooij et al. 2003, and references therein).  \\subsection{A New Chemical Evolution Model to Accommodate Globular Cluster Formation}  The main ingredients of the model presented in this paper are the globular clusters. Unfortunately, it is still not understood in detail how these clusters are formed. One interesting possibility, outlined by McCrea (1982), suggests that clusters form in the compressed region between colliding streams or clumps of gas. Gunn (1980) uses a similar argument as a means to attain globular cluster sized Jeans masses at temperatures of a few thousand degrees. Some observational evidence for such a formation mechanism also exists (Schweitzer 1987). Given our present lack of understanding of how the clusters formed, we will assume the following phenomenological scenario. The starting point consists of a large number of protogalactic clumps (baryons inside dark matter halos) which, although spatially isolated, chemically evolve as in the simple mass loss model (a `quiescent' mode of star formation). The gas in these clumps is assumed to have been initially enriched in a previous Population III phase. When these clumps collide, globular clusters are assumed to form with the metal abundance of the gas in the colliding clumps (a `burst' mode of star formation). By design, the mass in cluster stars is small compared to the total mass of all stars. After most collisions, quiescent star formation in each colliding clump ceases and the remaining gas, having lost energy but not angular momentum in the collision, falls towards the center of the proto-galaxy to be recycled later. The stars previously formed quiescently (now considered field stars) and the dark matter continue their previous dynamical evolution, while the newly formed clusters are not expected to have the same kinematics or spatial distribution as these field stars. Relatively few collisions could also lead to cluster formation and coalescence of the interacting clumps to form a satellite galaxy.  The context of the new chemical evolution model is as follows. In the well known mass loss model of chemical evolution described below one starts with a mass of gas which is slowly turned into stars while enriched gas is lost at a rate proportional to the star formation rate. The constant of proportionality determines the `effective' yield. The evolution ends when the gas is exhausted due to mass loss. In the new model we imagine a large number of `chemical evolution boxes' which individually evolve as above but may now collide with a neighbor. This collision is assumed to create a globular cluster and halt the star formation in both colliding clumps. As discussed above, the colliding gas dissipates energy and falls to the center or as close to the center as the conserved angular momentum will allow. Meanwhile the stars and dark matter pass through and continue their previous dynamical evolution. In addition to the original mass loss process which can be thought of as `supernova driven' and hence is assumed to remove gas from further star formation, we now have a second mass loss process-the collisions-which cause the gas to lose energy and thus be available for later recycling. Because the clusters are assumed to form in collisions and because the abundance distribution of each group can be represented by a normal distribution, in the model we make the collision rate normally distributed in [Fe/H] (i.e. the role of globular cluster formation is to effect a sudden cessation of star formation by removal of the remaining gas from individual clumps in a Gaussian distributed manner in [Fe/H]).   Quantitatively, in the simple mass loss model of chemical evolution (Hartwick 1976, see also Pagel 1997, p.234 and  Binney \\& Merrifield 1998, p.308) the three components are the mass in stars, M$_{s}$, the mass in remaining gas, M$_{g}$, and the mass lost due to star formation, here denoted M$_{WHIM}$ in anticipation of later discussion. As the chemical evolution proceeds, the gas is gradually depleted and similtaneously enriched in heavy elements due to ongoing star formation, while also being lost at a rate proportional to the star formation rate. The equations governing each component for an initial mass of gas M$_{t}$ of metallicity Z$_{0}$, can be written as \\begin{equation} \\frac{dM_{s}}{dlog Z}={ln10}\\times {Z}\\times {M_{g}}/p \\end{equation} \\begin{equation} M_{g}=M_{t}\\times{exp(-(Z-Z_{0})(1+c)/p)} \\end{equation} \\begin{equation} M_{WHIM}={c}\\times{M_{s}} \\end{equation} \\noindent where M$_{s}$(Z$_{0}$)=0 and M$_{g}$(Z$_{0}$)= M$_{t}$, p is the yield of heavy elements and c is the ratio of mass loss rate to star formation rate and also determines the `effective' yield (i.e. $p_{eff}=p/ (1+c)$).  In our new model, it is assumed that the collisions between clumps which halt the chemical evolution are distributed normally with mean log Z$_{c}$ and standard deviation $\\sigma$. As the collisions occur, the mass in surviving clumps diminishes. This surviving clump mass with metallicity greater than log Z is M$_{t}\\times$~f~($>$~log Z) where the initial mass of gas in all clumps with metallicity Z$_{0}$ is M$_{t}$ and \\begin{equation} f(>~log Z)\\simeq\\frac{1}{\\sqrt{2\\pi}\\sigma}\\int_{log Z}^{\\infty} exp\\left(- \\frac {(log Z^{\\prime}-log Z_{c})^{2}}{2\\sigma^{2}}\\right)~d log Z^{\\prime} \\end{equation} \\noindent which is valid when Z$_{0}$ $\\ll$ Z$_{c}$. Note that for computing purposes the right hand side of (4) is one half of the complementary error function i.e. $f~(> log Z)=0.5\\times erfc~\\left(\\frac{(log Z-log Z_{c})}{\\sqrt{2}\\sigma} \\right)$. By replacing M$_{t}$ with M$_{t}\\times$~f, equations (1) and (2) can be written as \\begin{equation} \\frac{dM_{s}}{dlog Z}={ln10}\\times {Z}\\times {M_{g}}/p \\end{equation} \\begin{equation} M_{g}={M_{t}}\\times{0.5}\\times{erfc~\\left(\\frac{(log Z-log Z_{c})}{ \\sqrt{2}\\sigma}\\right)}\\times{exp(-(Z-Z_{0})(1+c)/p)} \\end{equation} \\noindent Note that equations (1) \\& (3) remain the same. (Equation (1) remains the same because the metallicity of the gas lost by collisions is the same as that in the clumps at the time of the collision). Because we now have an additional mass loss component, we define a new variable M$_{ml}$, which represents the gas lost when the clumps collide. Whereas M$_{WHIM}$ is considered here to be lost to the `warm hot intergalactic medium', the new M$_{ml}$ component, having lost energy in the collision, is assumed to fall to the center of the proto-galaxy to be recycled. Figuratively, an individual classical chemical evolution `box' is totally purged by having gas drop out of the bottom on collision as well as having been blown continuously out of the top as a result of previous ongoing star formation! The equations governing the evolution of these two mass loss components are \\begin{equation} M_{WHIM}={c}\\times{M_{s}} \\end{equation} \\noindent and \\begin{equation} \\frac{dM_{ml}}{dlog Z}={M_{t}}\\times{exp(-(Z-Z_{0})(1+c)/p)}\\times{ \\frac{1}{\\sqrt{2\\pi}\\sigma}}\\times{exp\\left(-\\frac{(log Z-log Z_{c})^ {2}}{2\\sigma^{2}}\\right)} \\end{equation} We limit the number of parameters by letting Z$_{c}$ define the effective yield, i.e. \\begin{equation} c=(p/Z_{c}-1) \\end{equation}  Equation (5) represents the MDF of all stars formed (including those in globular clusters). We have not attempted to model globular cluster formation. It is assumed that the clusters are formed in bursts during collisions between clumps. We do attempt to make the chemical evolution consistent by ensuring that stars are formed which, if assembled into a cluster system, would have the requisite mass and (by design) the appropriate Gaussian abundance distribution. To this end we express the amplitude of the cluster MDF as the product of the $total$ baryonic mass (M$_{t}$) (e.g. McLaughlin, 1999) and an efficiency factor $\\eta$ so that \\begin{equation} \\frac{dM_{GC}}{dlogZ}=\\eta M_{t}\\frac{1}{\\sqrt{2\\pi}\\sigma} {exp\\left(-\\frac{(log Z-log Z_{c})^{2}}{2\\sigma^{2}}\\right)}, \\end{equation} \\noindent where $\\eta=0.00068$. This value of $\\eta$ was chosen in order to ensure the assembly of $\\sim100$ metal poor (blue) clusters in the model to be discussed later. For the parameters adopted in the model to follow, all but the most metal rich 3\\% of the the blue clusters and 99.8\\% of the same number of metal rich (red) clusters can be accommodated before the gas is exhausted. Since the M$_{ml}$ component eventually gets turned into stars, the ratio of the integrated mass in clusters to that in stars (i.e. (M$_{GC}$/(M$_{s}$+M$_{ml}-$M$_{GC})$) is $\\sim0.001$ which is within a factor of two of observational determinations of the same quantity (e.g. McLaughlin, 1999).  We reiterate that in the context of this model the formation of globular clusters only determines the extent of the chemical evolution within individual clumps. It is the stars produced that enriched the gas from which the clusters formed and the enriched gas lost both due to star formation and collisions which are the dominant end products in this model and not the individual clusters.   Equations (5)-(10) along with boundary conditions M$_{s}(Z_{0})=0$, M$_{g} (Z_{0})$=M$_{t},$ M$_{WHIM}(Z_{0})=0,$ M$_{ml}(Z_{0})=0$ and M$_{GC} (Z_{0})=0$ and parameters $Z_{c}$, $Z_{0}$, $\\sigma$ and $\\eta$ define the new chemical evolution model. Figure 1 shows the evolution of the independent quantities for an illustrative set of parameters. The parameters were chosen to provide observable separation of the four components, but it should be emphasized that the model was motivated such that log~Z$_{c}$ and $\\sigma$ are determined by a particular $globular$ $cluster$ abundance distribution. As a result of the loss of gas due to collisions, the peak in the accompanying $stellar$ metal abundance distribution is found to be $\\sim0.25$ dex $lower$ than the corresponding globular cluster peak (log~Z$_{c}$) for values of $\\sigma$ of order 0.2-0.3.  \\subsection{Colliding Boxes and the Stellar Halo of M31}  Durrell et al.\\ (2001) have probed the halo of M31 at a distance of $\\sim$20 kpc from the center and have determined the metallicity distribution function (MDF) photometrically. This study provides an excellent `test' sample for the new model, since all stars are essentially at the same distance from the Milky Way and the photometry is sufficiently deep that there should be no serious observational selection effects. Their results are reproduced as the histogram in Figure 2, where for consistency their metallicity scale [m/H]  has been converted to [Fe/H] by subtracting 0.3 dex (as described in their paper). In a recent spectroscopic study of a large sample of globular clusters in M31, Perrett et al.\\ (2002) confirm earlier work that the abundance distribution is bimodal. These authors fit Gaussians with [Fe/H]$=-1.44~(\\sigma^{2}=0.22)$ and [Fe/H]$=-0.5~(\\sigma^{2}=0.13)$ to their observations. The solid line in Figure 2 shows the MDF predicted by the new model assuming the sum of contributions from a log$Z_{c}/Z_{\\odot}=-0.55~(\\sigma=0.2)$ component and from a log$Z_{c}/Z_{\\odot}=-1.5~(\\sigma=0.3)$ component in the initial ratio of 2.4:1 and normalized to the same area as the observations. The fit is reasonable while  possessing the feature of unifying the clusters and stars of each population. Note the offset in the peak of the stellar distribution from the assumed value of log$Z_{c}$ of the metal-rich component.  The simple phenomenological picture outlined above now allows us to model early galaxy evolution based on the near universal appearance of two groups of globular clusters. However, in order to use the model to determine galaxy evolution as a function of time and/or redshift and to calculate star formation rates, we need a relationship between metal abundance and age. ",
        "conclusions": "Starting with a new chemical evolution model and with the goal of accounting for the bimodal distribution of globular clusters, the cosmic star formation rate density history was successfully modelled with a minimum number of free (i.e. non-observationally determined) parameters. The resulting model compares favorably with the Madau plot and thus provides a scenario for the chemical evolution of a representative sample of the universe. The model was obtained by constructing a composite galaxy based on observational input from the Milky Way and M31. An important aspect of this result as presented is that it is independent of the effects of dust extinction. While it might be considered presumptuous to expect to describe the entire universe with a single galaxy, the summary of ancillary quantities listed in Table 3 shows that many local cosmological constraints are also satisfied in varying degrees.  \\begin{deluxetable}{cccc} \\tablenum{3} \\tablecolumns{4} \\tablecaption{Summary of Results} \\tablehead{ \\colhead{Parameter} & \\colhead{Model} & \\colhead{Observed} & \\colhead{Reference} } \\startdata $\\Omega_{\\star}$ & $0.0056\\pm0.0007$ & $0.0037^{+0.0019}_{-0.0013}~h^{-1}_ {0.66}$ & FHP (1998) \\nl $\\Omega_{WHIM}$ & $0.0048\\pm0.0007$ & $\\gtrsim0.0046~h^{-1}_{0.66}$ & Tripp et al.\\ (2000) \\nl $\\Omega_{diffuse}$ & $0.044\\pm0.002$ & $\\geq0.0164~h^{-2}_{0.66}$ & Stocke et al.\\ (2003) \\nl $\\left<[Fe/H]\\right>_{WHIM}$ & $\\sim-1$ & $\\geq-1$ & Mulchaey (2000) \\nl M$_{WHIM}$/M$_{\\star}$ & $0.9\\pm0.2$ & $>1$? & Mulchaey (2000) \\nl M$_{\\star,disk}$/M$_{\\star,spheroid}$ & $0.9\\pm0.2$ & $\\sim1$ & Schechter \\& Dressler (1987) \\nl $\\rho(M_{t})/(\\Omega_{m}\\rho_{crit})$ & $0.04\\pm 0.01$ & $\\geq0.04$ & Mulchaey (2000) \\nl $\\rho_{\\bullet}$(M$_{\\odot}$/Mpc$^{3}$) & $2.6\\pm0.5\\times10^{5}$  & $2.8\\pm 0.4\\times10^{5}~h^{2}_{0.66}$ & Yu \\& Tremaine (2002) \\nl  \\enddata  \\end{deluxetable}  A physical framework which incorporates the above results is presented. The picture involves a large scale anisotropic collapse of quiescently evolving, isolated protogalactic clumps. During the first collapse (transverse to the rotation axis), only the low angular momentum clumps collide resulting in the formation of the blue clusters and a reservoir of low angular momentum gas which later is turned into bulge stars. As a result of the second collapse (along the rotation axis), red clusters are formed which become the backbone of a thick disk component. The stars already formed in the clumps before they collide pass through and back into the halo. An important prediction of the model is that the Galactic halo should contain a relatively metal rich ([Fe/H]$_{peak}\\sim-0.8$) population of stars with disk-like angular momentum. Galactic counterparts of this red halo component and model thick disk are identified and are predicted to partially overlap in metal abundance, age and spatial distribution. This overlap might explain why the properties of the `thick disk' have been so difficult to define up to now (e.g. Norris, 1999)."
    },
    "0311/astro-ph0311332_arXiv.txt": {
        "abstract": "Fast particles diffusing along magnetic field lines in a turbulent plasma can diffuse through and then return to the same eddy many times before the eddy is randomized in the turbulent flow.  This leads to an enhancement of particle acceleration by large-scale compressible turbulence relative to previous estimates in which isotropic particle diffusion is assumed. ",
        "introduction": "Particle acceleration by magnetohydrodynamic (MHD) turbulence has been investigated by a number of authors (e.g., Fermi 1949, Hall \\& Sturrock 1967, Wentzel 1968, Kulsrud \\& Ferrari 1971, Eichler 1979, Achterberg 1981, Schlickeiser \\& Miller 1998), and has been studied in a number of astrophysical settings, including the interstellar medium (Ptuskin 1988, Schlickeiser \\& Miller 1998), solar flares (LaRosa et~al 1994, Miller et~al 1996, Miller et~al 1997, Miller 1998, Schlickeiser \\& Miller 1998, Miller 1998), active galactic nuclei (Blackman 1999, Gruzinov \\& Quataert 1999), diffuse intracluster plasma in galaxy clusters (Eilek \\& Weatherall 1999, Weatherall \\& Eilek 1999, Dogiel 1999, Brunetti et~al 1999, Ohno et~al 2002, Fujita et~al 2003), and the lobes of extragalactic radio sources (Eilek \\& Henriksen 1984, Manolakou et~al 1999).   The way in which magnetohydrodynamic (MHD) turbulence accelerates energetic particles depends on the relative magnitudes of $\\lambda_\\parallel$ and $\\lambda_{\\rm mfp}$, where $\\lambda_\\parallel$ is the length of a turbulent eddy (or the wave length of a weakly damped wave) measured along the magnetic field, and $\\lambda_{\\rm mfp}$ is the energetic-particle scattering mean free path.  In the weak-scattering limit, $\\lambda_\\parallel \\ll \\lambda_{\\rm mfp}$, a particle's velocity~$v_\\parallel$ along the background magnetic field is approximately constant as the particle travels many wave lengths along the magnetic field (assuming that the amplitudes of the fluctuations are sufficiently small), and the particle interacts resonantly with a weakly-damped wave when the wave frequency in the frame of reference moving along the magnetic field at speed~$v_\\parallel$ is an integer multiple of the particle's gyrofrequency. In the strong-scattering limit, $\\lambda_\\parallel \\gg \\lambda_{\\rm mfp}$, a particle travels diffusively across an eddy of length~$\\lambda_\\parallel$, and acceleration is non-resonant.  Particle acceleration in the strong-scattering limit is the focus of this paper.  Ptuskin (1988) studied particle acceleration in the strong-scattering limit by a power-law spectrum of acoustic or magnetosonic waves extending from a large scale~$l$ to a smaller scale~$l_1$, with rms fluctuation velocity~$u_{\\rm rms}$ dominated by the waves at scale~$l$. He assumed that particle diffusion in space is isotropic with diffusion coefficient~$D$, as opposed to primarily along field lines, and that some type of small-scale wave distinct from the acoustic/magnetosonic waves scatters the particles, with $\\lambda_{\\rm mfp} \\sim D/v \\ll l_1$, where $v$ is the particle speed. He found that for $D \\gg v_{\\rm w} l$, where $v_{\\rm w}$ is the wave phase velocity, particles diffuse in momentum space according to the equation \\begin{equation} \\frac{\\partial f}{\\partial t} = \\frac{1}{p^2} \\frac{\\partial }{\\partial p} \\left( p^2 D_p \\frac{\\partial f}{\\partial p}\\right), \\label{eq:fp0} \\end{equation} where $f$ is the particle distribution function averaged over the waves, and the momentum diffusion coefficient~$D_p$ is given by \\begin{equation} D_p = \\frac{2 p^2 u_{\\rm rms}^2}{9 D}, \\label{eq:dp1} \\end{equation} where $p$ is the particle momentum. The condition $D\\gg v_{\\rm w} l$ implies that particles diffuse a distance~$l$ in a time much shorter than the wave period at scale~$l$.  In this limit, the largest-scale waves make the dominant contribution to~$D_p$.  This paper extends Ptuskin's analysis to strong compressible MHD turbulence taking into account the reversible motion of particles along magnetic field lines.  It is assumed that the rms turbulent velocity $u_{\\rm rms}$ equals or exceeds the Alfv\\'en speed~$v_{\\rm A}$, and that there is a single dominant length scale (outer scale) denoted~$l$ for both velocity and magnetic fluctuations.  Only the fluctuations at scale~$\\sim l$ are considered. As in Ptuskin's (1988) study, it is assumed that some type of small-scale wave scatters particles, and that $\\lambda_{\\rm mfp} \\ll l$. It is assumed that eddies are randomized in a time~$\\tau_{\\rm rand} \\sim l/u_{\\rm rms}$, and that $D_\\parallel \\gg u_{\\rm rms} l$, where~$D_\\parallel$ is the diffusion coefficient for particle motion along field lines. Thus, a particle diffuses a distance greater than~$l$ during a time~$\\tau_{\\rm rand}$, and can diffuse through and then return to an eddy multiple times before the eddy is randomized in the turbulent flow. It is shown in section~\\ref{sec:dp} that this increases the coherence of stochastic particle acceleration, effectively increasing~$D_p$ above the value in equation~(\\ref{eq:dp1}) by a factor of $\\sim N$, where $N \\sim \\sqrt{D_\\parallel/u l}$ is the typical number of times a particle diffuses through and then returns to an eddy before the eddy is randomized, assuming a particle is tied to a single field line. The discussion of section~\\ref{sec:dp} consists of estimates based on physical arguments; these estimates are tested and validated by Monte Carlo particle simulations in section~\\ref{sec:MC}.  In section~\\ref{sec:num} the results of the paper are applied to particle acceleration in the interstellar medium and the lobes of radio galaxies.    It is known that small cross-field motions and the divergence of neighboring magnetic field lines in turbulent plasmas allow a particle to escape from its initial field line by traveling sufficiently far along the magnetic field.  This places an upper limit on~$N$, which is explained in the Appendix. It is shown that this upper limit is not important in the astrophysical examples considered in section~\\ref{sec:num} given the assumed values of~$D_\\parallel$. However, for sufficiently large values of~$D_\\parallel$, the upper limit on~$N$ becomes important. ",
        "conclusions": "\\label{sec:conc}  It is shown that the reversible motion of energetic particles along the magnetic field enhances the rate of particle acceleration by large-scale compressible turbulence relative to the isotropic-spatial-diffusion case. When this effect is taken into account, the time scale for acceleration of relativistic electrons by large-scale compressible turbulence in the lobes of radio galaxies is $\\sim 5\\times 10^7$~yr, and the time scale for acceleration of 1-GeV protons by large-scale compressible turbulence in the coronal plasma of the ISM is ~$\\sim 3-10\\times 10^7$~yr. The main source of uncertainty in these estimates is the uncertainty in~$D_\\parallel$; other sources include the uncertainty in $l$, ~$u_{\\rm rms}$, and the dimensionless constants $a$ and $b$ that relate $\\tau_{\\rm rand}$ and $(\\nabla \\cdot {\\bf u})_{\\rm rms}$ to $u_{\\rm rms}$ and~$l$.      I thank Eliot Quataert, Steve Spangler, Torsten Ensslin, and Eric Blackman for helpful comments. This work was supported by NSF grant AST-0098086 and DOE grants DE-FG02-01ER54658 and DE-FC02-01ER54651.   \\appendix"
    },
    "0311/astro-ph0311618_arXiv.txt": {
        "abstract": "We present results from a comprehensive number of relativistic, time-dependent, axisymmetric simulations of the runaway instability of {\\it non-constant} angular momentum thick discs around black holes. This second paper in the series extends earlier results where only constant angular momentum discs were considered. All relevant aspects of the theory of stationary thick discs around rotating black holes, necessary to build the equilibrium initial data used in our simulations, are presented in great detail. The angular momentum of the evolved discs is assumed to increase outwards with the radial distance according to a power law, $l=K r^{\\alpha}$, where $K>0$ corresponds to prograde discs (with respect to the black hole rotation) and $K<0$ to retrograde discs. The main simplifying assumptions of our approach are not to include magnetic fields and self-gravity in the discs (test-fluid approximation). Furthermore, the dynamics of the spacetime is accounted for by computing the transfer of mass and angular momentum from the disc to the black hole through the event horizon. In this approximation the evolution of the central black hole, which initially is non-rotating, is assumed to follow a sequence of Kerr black holes of increasing mass and spin. All discs we build slightly overflow the potential barrier at the cusp, departing from equilibrium, so that accretion is possible. In agreement with previous results based on stationary models we find that by allowing the mass and the spin of the black hole to grow, constant angular momentum discs rapidly become unstable on a dynamical timescale (a few orbital periods). The comparison with the results of our first paper shows that the effect of the angular momentum transfer from the tori to the black hole is to make constant angular momentum discs less unstable, increasing the timescale for the runaway instability to grow. However, we find that non-constant angular momentum discs are dramatically stabilized for very small values of the angular momentum slope $\\alpha$, much smaller than the Keplerian value $\\alpha=1/2$. Our fully relativistic and time-dependent simulations confirm, thus, the predictions of stationary studies concerning the stabilizing effect of non-constant angular momentum distributions. For the various disc-to-hole mass ratios considered we systematically find that the critical values of $\\alpha$ below which the runaway instabilty can exist are slightly smaller than those reported previously in the literature based on stationary studies. ",
        "introduction": "In a recent paper \\citep{font:02a} (Paper I hereafter) we presented the first time-dependent, hydrodynamical simulations in general relativity of the so-called runaway instability of geometrically thick discs (or tori) around a Schwarzschild black hole. The origin of this instability, discovered by \\citet{abramowicz:83}, is a dynamical process by which the cusp of a disc which initially is filling its Roche lobe penetrates inside the disc due to mass transfer from the disc to the accreting black hole. This process leads to the complete destruction of the disc on a dynamical timescale. \\citet{abramowicz:83} found that the runaway instability occurs for a large range of parameters such as the disc-to-hole mass ratio and the location of the inner radius of the disc. More detailed studies followed (see, e.g. Paper I for an up-to-date list), most of which were based on stationary models in which a fraction of the mass and angular momentum of the initial disc is instantaneously transferred to the black hole. The new gravitational field allows to compute the new position of the cusp, which controls the occurrence of the runaway instability. The self-gravity of the tori favours the instability, as shown from both, studies based on a pseudo-potential for the black hole \\citep{khanna:92,masuda:98} and from fully relativistic (stationary) calculations \\citep{nishida:96a}. However, there are some parameters in the models which have stabilizing effects. One of this is the rotation of the black hole \\citep{wilson:84,abramowicz:98}, but the most important one seems to be the radial distribution of the specific angular momentum, increasing with the radial distance as a power law \\citep{daigne:97,abramowicz:98}. Despite the abundant literature on the subject, it is commonly accepted that the very existence of the instability remains still uncertain.  In Paper I we initiated a program aimed at investigating the runaway instability through time-dependent hydrodynamical simulations of increasing complexity in general relativity. The starting point was simple. In the simulations reported in Paper I the distribution of the angular momentum in the disc was assumed to be constant. We also neglected the self-gravity of the disc as well as the presence of magnetic fields. Our model, however, accounted for the dynamics of the gravitational field of the black hole plus torus system in an approximate yet satisfactory manner. The evolution of the central black hole was assumed to follow a sequence of Schwarzschild black holes of increasing mass, whose increase rate was controlled by computing the mass flux from the disc to the black hole through the innermost point of the numerical grid. We notice that the ability to take into account such dynamics is of paramount importance to investigate the runaway instability at all.  Our study showed that by allowing the mass of the black hole to grow the runaway instability appears on a dynamical timescale, in agreement with previous estimates from stationary models. Since our study was mainly motivated by understanding the hydrodynamical evolution of the central engine of gamma-ray bursts (GRBs hereafter), we only considered the case of stellar mass black holes, leaving aside thick discs around supermassive black holes which can also be present in active galactic nuclei. For a black hole of $2.5\\mathrm{M_{\\odot}}$ and disc-to-hole mass ratios in the range 0.05 to 1, we found in Paper I that the timescale of the instability never exceeds $1\\ \\mathrm{s}$ for a large range of mass fluxes and it is typically a few 10 ms. We note in passing that the runaway instability seems to be a robust feature of non-self-gravitating, constant angular momentum discs irrespective of the way accretion is induced, as \\citet{zanotti:02} have recently reported.  Whatever the progenitor system which collapses to form a stellar mass black hole surrounded by a thick disc is, it is very unlikely that the initial distribution of the angular momentum in the disc is constant. As numerical simulations show this is not the case in discs formed after the coalescence of a close compact binary (consisting either of two neutron stars or of a black hole and a neutron star) \\citep{davies:94,ruffert:96,kluzniak:98,ruffert:99,lee:00,shibata2:03}, in the collapsar model of {\\it failed} supernova \\citep{macfadyen:99}, or, at different scales, in the collapse of a rotating supermassive star leading to a rotating supermassive black hole \\citep{shibata:02,shapiro:02}. For instance, in the Newtonian SPH simulations of \\citet{davies:94} the final configuration consists of a core of mass $2.44$M$_{\\odot}$ in uniform rotation surrounded by a differentially rotating thick disc of mass $0.36$M$_{\\odot}$. The distribution of the angular momentum in the disc is close to a power law of index 0.2. Similarly, in the most recent three-dimensional SPH coalescence simulations of \\citet{lee:00} the corresponding angular momentum distribution follows a power law of index 0.45, quite close to Keplerian. We note in particular that recent relativistic simulations of unequal mass neutron star coalescence \\citep{shibata2:03} yield a black hole surrounded by a thick disc, following tidal disruption of the less massive star. \\citet{shibata2:03} found that for a neutron star rest-mass ratio $\\sim 0.85$ the final mass of the disc may be several percents of the total mass of the system if the stars are not very compact. Furthermore, the angular momentum distribution in the disc that forms is not constant (M.~Shibata, private communication). In addition to the evidence provided by the results of numerical simulations, as the dynamical timescale on which the runaway instability occurs is much shorter than the viscous timescale necessary to achieve angular momentum transport, it seems much more realistic to assume a \\textit{non-constant} angular momentum distribution in the disc, increasing outwards from the cusp to larger radii.  Therefore, in the present paper we extend the study initiated in Paper I to such most interesting case of non-constant angular momentum discs. The main motivation of our work is to check through time-dependent simulations in general relativity whether such distributions have indeed the stabilizing effect previously found in non self-gravitating stationary models \\citep{daigne:97,abramowicz:98,lu:00}. We note that preliminary results of this investigation were already presented in a previous letter \\citep{font:02b} (Paper II hereafter). In the current paper we extend those results in two main directions: firstly, we present in great detail all relevant aspects of the theory of stationary, non-constant angular momentum thick discs around rotating black holes, which is necessary to build the equilibrium initial data we use in our simulations. Secondly, we report on the results of a comprehensive number of simulations which allows us to explore the dependence of the absence/presence of the runaway instability on various parameters of the model such as the disc-to-hole mass ratio, the potential barrier at the cusp (or in an equivalent manner, the initial mass accretion rate), or the efficiency of the angular momentum transfer from the disc to the black hole. The extended number of models considered here permits to draw some conclusive answers on the likelihood of the runaway instability of non-self-gravitating discs, as well as to pin-down the critical value of the exponent of the angular momentum distribution separating stable and unstable models.  As we showed in Paper II, and we further clarify in the present investigation, the runaway instability is most likely avoided in non rigidly rotating discs as those formed in the coalescence of a binary neutron star system or in the gravitational collapse of a massive star. This result is in agreement with recent perturbative analysis of vertically integrated discs performed by \\citet{rezzolla:03} who found that the innermost regions of discs tend to behave as regions of {\\it evanescent-wave propagation}, reducing the mass flux and, as a result, suppressing the runaway instability. If the growth of non-axisymmetric modes in the disc by the Papaloizou-Pringle instability is suppressed by the accretion process itself, as suggested by \\citet{blaes:87}, systems consisting of a Kerr black hole surrounded by a high density torus may then be long lived. The lifetime is probably controlled by the viscous timescale (a few seconds) rather than by the dynamical one. This may provide enough time for any plausible magneto-hydrodynamical process to efficiently transfer part of the energy reservoir of the system to a relativistic outflow. Therefore, the most favoured current models of GRBs (see \\citet{woosley:01}) can indeed be based on such central engines. This is especially important in the context of the so-called internal shock model for the prompt $\\gamma$-ray emission \\citep{rees:94}, where the observed lightcurve reflects the activity of the central engine, with no modification of the timescales other than the time dilation due to the redshift. In particular, the central engine has to survive for a duration at least comparable with the observed duration of the GRB.  The paper is organized as follows: Section~\\ref{sec:initialmodel} describes in detail the relevant aspects of the theory of relativistic, stationary, non-constant angular momentum thick discs necessary to build the initial state for our simulations. Part of this information is presented in three appendices at the end of the paper. Section~\\ref{sec:hydro} deals with a brief description of the numerical framework we use. Section~\\ref{sec:tests} shows specific tests for the Kerr spacetime that the code has successfully passed. The description of the simulations and the discussion of the results are presented in Section~\\ref{sec:simulations}. Finally Section~\\ref{sec:conclusions} summarizes our conclusions and presents possible directions for future work. As we did in Papers I and II we use geometrized units ($G=c=1$) unless explicitely stated. Usual cgs units are obtained by using the gravitational radius of the black hole $r_\\mathrm{g}=G M_\\mathrm{BH}/c^{2} = 1.5\\times 10^{5}\\ (M_{\\mathrm{BH}}/M_{\\odot})\\ \\mathrm{cm}$ as unit of length. Greek (Latin) indices run from 0 to 3 (1 to 3), and we use a spacelike signature for the metric (- + + +). ",
        "conclusions": "\\label{sec:conclusions}  We have presented a comprehensive numerical study aimed at clarifying the likelihood of the so-called runaway instability in axisymmetric thick discs orbiting around (initially non-rotating) black holes. Our main conclusion, already pointed out in a preliminary investigation \\citep{font:02b} is that the runaway instability of (non self-gravitating) thick discs around (rotating) black holes \\citep{abramowicz:83} is strongly suppressed when the distribution of angular momentum in the disc is assumed to be non constant. We have reached this conclusion using a numerical framework based on time-dependent hydrodynamical simulations in general relativity. This approach departs from the most common procedures followed in previous studies of the runaway instability in the literature, which are generally based on the analysis of stationary models (see. e.g. \\citet{font:02a} and references there in). Despite the novel approach it is remarkable that our main result is in complete agreement with such previous studies based on stationary sequences of equilibrium configurations, either using a pseudo-Newtonian potential \\citep{daigne:97} or in general relativity \\citep{abramowicz:98}. We further notice that our finding is also in good agreement with recent perturbative analysis of vertically integrated discs performed by \\citet{rezzolla:03}.  The present paper extends in a number of ways the preliminary results reported in our previous investigation \\citep{font:02b}. In the first half of the article (cf. Section~\\ref{sec:initialmodel}), we have shown how to build a series of thick accretion discs in hydrostatic equilibrium around a Schwarzschild or Kerr black hole. All relevant aspects of the theory of stationary thick discs around rotating black holes have been presented in great detail. In particular our description has been done for discs in which the specific angular momentum (per unit mass) is assumed to increase outwards with the radial distance according to a power law $l=K r^{\\alpha}$. All possible equipotential configurations for a Kerr black hole (including the limiting case of a Schwarzschild black hole) have been discussed in detail for both, prograde ($K>0$) and retrograde ($K<0$) rotation, and sub-Keplerian ($0\\le\\alpha < 1/2$), Keplerian ($\\alpha=1/2$), and super-Keplerian ($1/2<\\alpha<1$) discs. The various possible geometries of these configurations have been summarized in Table~\\ref{tab:geometry} and in Figs.~\\ref{fig:equi1pro} to \\ref{fig:equi3pro}. These figures and table extend earlier results presented in \\citet{font:02a} where only {\\it constant} angular momentum discs around Schwarzschild black holes were considered.  The second half of the article has been devoted to presenting the results of a comprehensive set of models that we have evolved with our numerical code. The main result of these simulations has been to show, quite generically, that a small increase outwards of the specific angular momentum distribution in the discs, well below the Keplerian limit, results in a strong stabilizing effect. In particular we have demonstrated that the case $\\alpha=0$ (constant angular momentum disc) is {\\it always} unstable when the black hole dynamics is taken into account, for both, $M_{\\mathrm D}/M_\\mathrm{BH}=1$ and 0.1, irrespective of the value of the potential barrier at the cusp (slightly overflowing the corresponding Roche lobe) and of the initial mass flux in the stationary regime. We note in passing that a similar conclusion has been recently obtained by \\citet{zanotti:02} for a Schwarzschild black hole and for different initial data, in which models that initially are in strict stable equilibrium are conveniently perturbed to trigger the accretion process. However, we have shown that discs with small non-zero values of $\\alpha$ are stable, and we have provided an estimate of the critical value of $\\alpha$ separating stable and unstable models, for various disc-to-hole mass ratios and mass fluxes in the stationary regime. More precisely we have found that for $M_{\\mathrm D}/M_\\mathrm{BH}=1$, the critical slope is $\\alpha_{\\mathrm{cr}}\\sim 0.085-0.09$ for $\\dot{m}_\\mathrm{stat}\\sim 3.2-3.6\\ \\mathrm{M_{\\odot} s^{-1}}$ and $\\alpha_{\\mathrm{cr}}\\sim 0.015-0.02$ for $\\dot{m}_\\mathrm{stat}\\sim 0.1\\ \\mathrm{M_{\\odot} s^{-1}}$. Correspondingly, for $M_{\\mathrm D}/M_\\mathrm{BH}=0.1$, the critical slope is $\\alpha_{\\mathrm{cr}}\\sim 0.055-0.06$ for $\\dot{m}_\\mathrm{stat}\\sim 1.7-1.9\\ \\mathrm{M_{\\odot} s^{-1}}$ and $\\alpha_\\mathrm{cr}\\sim 0.05$ for $\\dot{m}_\\mathrm{stat}\\sim 1.3\\ \\mathrm{M_\\mathrm{\\odot} s^{1}}$. These values are slightly smaller than the values found in stationary studies (0.14 in the relativistic case for a disc-to-hole mass ratio of $\\sim 0.15$ \\citep{abramowicz:98}). However these studies assume an instantaneous mass and angular momentum transfer from the disc to the black hole, and our study clearly indicates that the critical slope $\\alpha_\\mathrm{cr}$ increases with the mass flux $\\dot{m}_\\mathrm{stat}$. For the first time we have shown the suppression of the runaway instability in nonconstant angular momentum discs with a time-dependent relativistic calculation, and we have been able to estimate the lifetime of thick discs orbiting a black hole in both, the unstable and stable cases. The very strong stabilizing effect of the non-constant angular momentum distribution in the disc is demonstrated in our simulations by the exponential increase of the timescale for the instability to set in with the slope $\\alpha$.  The black hole plus thick discs systems we have simulated with $M_\\mathrm{D}/M_\\mathrm{BH}=0.1$ are very close to what numerical simulations of compact binary mergers and collapsars predict \\citep{ruffert:99,kluzniak:98,macfadyen:99,shibata:00,aloy:00}. Therefore, the results regarding this subset of models should be considered as, perhaps, much closer to what may be happening in realistic discs. However, even in this case, three important simplifications have been made: (i) the black hole is initially non-rotating whereas both in collapsars and binary mergers it is expected to form with an initial spin parameter $a/M_\\mathrm{BH}\\ga 0.5$; (ii) the self-gravity of the disc is not included; and (iii) magnetic fields, which are believed to play an important role on the dynamics of those objects, are also not included.  Certainly, along with analyzing the effects of varying the EoS of the matter in the disc, these are all issues that we plan to consider in the near future in order to extend the present study. Among those, the inclusion of an appropriate treatment to account for the self-gravity of the disc is the most difficult task despite its obvious relevance. Previous studies \\citep{khanna:92,nishida:96a,masuda:98}, based on either time-dependent simulations with pseudo-Newtonian potentials or stationary models in general relativity, indicate that self-gravity seems to have a strong destabilizing effect. It has been argued \\citep{nishida:96a} that in self-gravitating accreting tori the cusp is pushed towards the black hole by the gravitational force of the discs which, in principle, should act against the instability. However, during the accretion process, the interplay between the gravitational forces of the black hole (increasing) and of the disc (decreasing) make the cusp move closer to the centre of the torus than in non-self-gravitating models, thus favouring the instability. This important result is yet to be proved by means of time-dependent simulations in full general relativity. Unfortunately, successful attempts to long-term (numerically) stable simulations of thick discs orbiting black holes by solving the coupled system of Einstein and general relativistic hydrodynamic equations are still out of reach of present day numerical relativity codes (see, e.g. \\cite{fontlr} and references there in for an up-to-date list of such codes). Such simulations should provide a definite answer to the issue of whether the runaway instability exists or not.  We notice that neglecting the self-gravity of the disc is a more valid assumption the smaller the disc-to-hole mass ratio becomes. Therefore, with respect to the models analyzed in this paper we expect that for discs with $M_\\mathrm{D}/M_\\mathrm{BH} \\la 0.1$ the strong stabilizing effect of non-constant angular momentum distribution demonstrated in our study could still overcome the destabilizing effect of the much less relevant disc's self-gravity.  The incorporation of magnetic fields seems to be more easily conceivable. Recently, \\citet{devilliers:03} and \\citet{gammie:03} have presented magneto-hydrodynamical simulations of constant angular momentum tori in general relativity. These simulations assume much smaller accretion mass fluxes than in the present study and do not include any time evolution of the spacetime. However, they indicate a promising direction to include the effect of magnetic fields in future simulations of the runaway instability and hence suppress the third limitation listed above.  Nevertheless, the first aspect that we plan to consider in a forthcoming study is the time evolution of discs orbiting black holes which are {\\it initially} rotating. This will allow us to pinpoint the dependence of $\\alpha_{\\mathrm{cr}}$ with $a$. We notice that our hydrodynamics code has already proved in this study its ability to deal with the Kerr spacetime, and the method of construction of the initial state in the case of a rotating black hole has been fully described in this paper. Therefore, our upcoming study does not involve any new numerical development.  \\appendix"
    },
    "0311/astro-ph0311104_arXiv.txt": {
        "abstract": "By obtaining photometric redshift information, tomography allows us to cross-correlate galaxy ellipticities in different source redshift bins. The cross-correlation is non-vanishing because the different bins share much of the foreground mass distribution from which, over Gpc scales, the lensing signal is built. If the redshift bins are thick enough however, the cross-correlations are insensitive to contamination from the intrinsic alignments of galaxies since these fall off rapidly on scales larger than a few tens of Mpc. We forecast how lensing tomography using only the cross-power spectra can constrain cosmological parameters compared to tomography including the auto-spectra.  It is shown that the parameter errors are degraded by only ${\\cal O}(10\\%)$ for 5 or more source redshift bins. Thus, the cross-power spectrum tomography can be a simple, model-independent means of reducing the intrinsic alignment contamination while retaining most of the constraints on cosmology. ",
        "introduction": "As our cosmological knowledge and techniques mature, it becomes increasingly important to strive for observations and methods which are, as much as possible, immune to known systematic effects as well as being statistically powerful. Nowhere is this more true than in the field of weak gravitational lensing. Since its detection by several groups three years ago (Bacon, Refregier \\& Ellis 2000; Kaiser, Wilson \\& Luppino 2000; Van Waerbeke et al. 2000; Wittman et al. 2000), weak gravitational lensing by large-scale structure has become a well established technique, used already to set constraints on the mass density ($\\Omega_{\\rm m0}$) and the fluctuation amplitude $(\\sigma_8)$ (e.g., see Van Waerbeke \\& Mellier 2003 for the current status) and touted for its potential to constrain cluster scaling relations (Huterer \\& White 2002) and dark energy (Benabed \\& Bernardeau 2001; Huterer 2002; Hu 2002a,b; Heavens 2003; Abazajian \\& Dodelson 2003; Refregier et al. 2003; Jain \\& Taylor 2003; Bernstein \\& Jain 2003; Takada \\& Jain 2003, hereafter TJ03).  There are many sources of systematic errors which can affect lensing measurements. One source in particular, intrinsic ellipticity alignments of source galaxies (Heavens, Refregier \\& Heymans 2000; Catelan, Kamionkowski \\& Blandford 2001; Crittenden et al.~2001; Pen, Lee \\& Seljak 2001; Mackey, White \\& Kamionkowski 2002; Jing 2002) cannot be mitigated by improved instrumentation or reliably predicted by theory. While likely not important for the current generation of experiments (Heymans et al.~2003), it could prove to be a limiting uncertainty in more ambitious ongoing and future surveys such as the CFHT Legacy Survey, the Deep Lens Survey, DML/LSST, Pan-STARRS and SNAP.  In this {\\em Letter\\/}, we discuss a simple, largely (astrophysical) model independent, technique to reduce the susceptibility of weak lensing measurements to intrinsic alignments of galaxies. This technique assumes only that the intrinsic alignments fall off on scales above a few tens of Mpc while the lensing signal builds up over Gpc scales. We emphasize the advantages of splitting the galaxy distribution into multiple redshift bins and considering only the cross-power spectra. These have no shot noise bias, contain most of the cosmological information and should be totally insensitive to intrinsic alignments. While not as powerful as model dependent methods (Heymans \\& Heavens 2003; King \\& Schneider 2002, 2003; Heymans et al.~2003) these latter require modeling of intrinsic alignments, which involves many uncertain aspects of galaxy formation. ",
        "conclusions": "We have proposed lensing tomography that only uses the cross-spectra constructed from source galaxies in different redshift bins. Cross-spectrum tomography provides a simple, robust and model independent means of reducing systematics from intrinsic alignments which retains most of the cosmological information. For more than 4 redshift bins, cross-spectrum tomography yields errors on cosmological parameters only $\\simlt 15\\%$ larger than errors including the auto-spectra (see Table \\ref{tab:const}). The situation is further improved by combining with cross-bispectrum tomography.  As future experiments concentrate on subdividing the source galaxies into multiple redshift bins to better reconstruct the lensing effect as a function of distance, it becomes more important to guard against contamination by intrinsic alignments. Intrinsic alignments are difficult to predict, with calculations differing by an order of magnitude at present. However, simply neglecting intrinsic alignments could bias estimates of cosmological parameters from future lensing surveys (Heymans et al.~2003). Our suggestion is intermediate between detailed modeling of the effect and simple neglect. We believe it is worth exploring a 3D mass reconstruction technique using only the lensing effects on galaxies in different redshift bins even if this is only used as a `sanity check' during the analysis.  Finally we remark on some issues regarding photometric redshift determinations. In order to implement our procedure we require more redshift bins than is usual for tomographic techniques (which frequently saturate at 2 or 3 bins). This requires more accurate photo-$z$ determinations, but still well within the range of possible accuracies.  For example, {\\sl SNAP\\/} is designed to achieve accurate photo-$z$ measurement with a random error $\\sigma_{z}=0.03$ (Massey et al.~2003). This would allow us to perform tomography with up to 10 redshift bins in the range $0<z<3$. Of course it is still crucial to calibrate the photo-$z$ estimates to the percent level, since constraints on the cosmological parameters come from measuring the lensing efficiency to the percent level. This will require a spectroscopic survey to calibrate the photo-$z$ distribution, as discussed in Bernstein \\& Jain (2003). Such a measurement should also mitigate against another important systematic arising from the tails of the photo-$z$ distribution: the misidentification as source galaxies of a fraction of galaxies which actually lie near the lensing plane leading to a misestimation of the lensing efficiency. Both effects can be controlled introducing gaps, larger than the photo-$z$ error, in the galaxy distribution when computing the cross power spectra between neighboring bins. We have found that even gaps with $\\Delta z=0.2$, a typical photo-$z$ error from a five color survey, enlarge the errors from cross spectrum tomography with 5 redshift bins by less than 5\\%, as expected from our earlier discussion. While further study (including realistic photo-$z$ errors) is needed, this issue is not significantly different for cross-spectra and auto-spectra."
    },
    "0311/astro-ph0311274_arXiv.txt": {
        "abstract": "We present observational constraints on a scalar-tensor gravity theory by $\\chi^2$ test for CMB anisotropy spectrum. We compare the WMAP temperature power spectrum with the harmonic attractor model, in which the scalar field has its harmonic effective potential with curvature $\\beta$ in the Einstein conformal frame and the theory relaxes toward Einstein gravity with time. We found that the present value of the scalar coupling, i.e. the present level of deviation from Einstein gravity $(\\alpha_0^2)$, is bounded to be smaller than $5\\times 10^{-4-7\\beta}$ ($2\\sigma$), and $10^{-2-7\\beta}$ ($4\\sigma$) for $0< \\beta<0.45$.  This constraint is much stronger than the bound from the solar system experiments for large $\\beta$ models, i.e., $\\beta> 0.2$ and $0.3$ in $2\\sigma$ and $4\\sigma$ limits, respectively. Furthermore, within the framework of this model, the variation of the gravitational constant at the recombination epoch is constrained as $|G(z=z_{rec})-G_0|/G_0 < 0.05(2\\sigma)$, and $0.23(4\\sigma)$. ",
        "introduction": "The discovery of acoustic peaks on the anisotropy spectrum of the cosmic microwave background (CMB) opened the way to measure the early universe. The high quality data recently provided by the Wilkinson Microwave Anisotropy Probe (WMAP) mission \\cite{benn} makes such a challenge feasible. The measurement of the characteristic length of various processes at the recombination epoch reveals the physical process in the early universe \\cite{hss}. Length scales, which exhibit in the angular spectrum of the CMB as a characteristic feature, depend on the horizon length which is controlled by gravity. Usually, the horizon length at the recombination epoch is a function of the amount of non-relativistic matter. The strength of gravity caused by matter and radiation determines how much time consumed to make the universe cool enough to recombine hydrogens. However, the strength of gravity depends not only on the amount of matter but also on the coupling strength of gravitational interaction. Thus the CMB anisotropy spectrum contains the information about the magnitude of the gravitational constant at the recombination epoch.  Constancy of physical ``constants'' is the fundamental issue which has long history and has attracted much interest \\cite{CU}. The recent attempts toward unifying all elementary forces \\cite{GSW} predict the existence of scalar fields whose vacuum expectation values determine the physical ``constants''. In such context, scalar-tensor gravity theories, whose original version was proposed by Jordan \\cite{J} and Brans and Dicke \\cite{BD} and were extended in a more general framework later \\cite{BW}, are the most promising candidates among the alternatives of Einstein gravity. In scalar-tensor theories, the coupling of a massless scalar field to the Ricci scalar provides a natural framework of realizing the time-variation of the gravitational constant via the dynamics of the scalar field.  In the Jordan-Brans-Dicke theory \\cite{J,BD} (hereafter, we refer to it as the Brans-Dicke theory for simplicity) which is the simplest example of scalar-tensor theories, a constant coupling parameter $\\omega$ is introduced. In the limit $\\omega \\rightarrow \\infty$, the gravitational constant can not change and Einstein gravity is recovered. Although scalar-tensor theories including the Brans-Dicke theory are compatible with the Einstein gravity in several aspects, they have many deviations from it. Weak-field experimental tests in solar-system have constrained the post-Newtonian deviation from the Einstein gravity, $\\omega > 500$ \\cite{REA,CMW,robert,lebach}. According to some recent reports, this bound would be updated to several thousands \\cite{eubanks,CMWGQ,ALG}. Such a small deviation implies that the variability of the background scalar field is also small.  In more general scalar-tensor theories with non-trivial coupling functions $\\omega(\\phi)$, the small deviation is not the outcome of fine-tuning because the cosmological evolution drives $\\omega$ toward infinity in the late cosmological epochs and naturally reduces the present observable effects of $\\phi$ field \\cite{DN}. If such attractor mechanism takes place, the nature of gravity (the variability of the background scalar field, weak-field deviations, etc.) can be significantly different in the early universe. Hence, information on the different cosmological epochs may constrain such theories. A simple and natural extension of the Brans-Dicke theory to the attractor model is the harmonic attractor model in which the scalar field has a quadratic effective potential of a positive curvature in the Einstein conformal frame. The analysis of big-bang nucleosynthesis (BBN) in the harmonic attractor model \\cite{DP} restricts two parameters characterizing the potential (its curvature $\\beta$ and today's gradient $\\alpha_0$). It is concluded that the BBN limit on the possible deviation from Einstein gravity ($2\\omega_0+3={\\alpha_0}^{-2}$) is much stronger than the present observational limits in large $\\beta(>0.3)$ models. We can extract the information about the early universe also by the analysis of structure formation \\cite{LMB}. The advantage of the use of the CMB fluctuations is that the physics of CMB is well understood and that we now have very accurate observational data. The trace of primordial fluctuation can be seen clearly in the CMB anisotropy spectrum where the information on the early universe up to the last scattering time is projected on the acoustic peaks \\cite{CK,PBM,NCS}.  We compare the CMB temperature anisotropy spectrum measured by WMAP with that in the scalar-tensor cosmological model mentioned above. The formulations of background and perturbation equations and the detailed explanation of physical processes can be found in our previous paper \\cite{NCS}. The purpose of this work is to constrain the deviation of the scalar-tensor gravity from the Einstein gravity and to find, in the framework of this model, how large the variation of the gravitational constant is allowed.  This paper is organized as follows: In Section II, the scalar-tensor cosmological model and its observational consequences are explained. In Section III, we describe the overview of our analysis for constraining the scalar-tensor coupling parameters. In Section IV, the models are compared with the WMAP data and the result of $\\chi^2$ test is presented. Finally, some conclusions are in Section V. ",
        "conclusions": "In this paper, we have quantitatively compared the CMB temperature spectrum in the scalar-tensor flat $\\Lambda$CDM model to the WMAP data by $\\chi^2$ test of goodness of fit.  The present deviation from the Einstein gravity ($\\alpha_0^2$) must be smaller than $5\\times 10^{-4-7\\beta}$ ($2\\sigma$), and $10^{-2-7\\beta}$ ($4\\sigma$) for $0< \\beta<0.45$. This constraint is much stronger than the bound from the solar system experiments for large $\\beta$ models, i.e., $\\beta> 0.2$ and $0.3$ in $2\\sigma$ and $4\\sigma$ limits, respectively. Within the framework of the harmonic attractor model, the difference between the gravitational constant at the recombination epoch and at the present is constrained as $|G(z=z_{rec})-G_0|/G_0 < 0.05(2\\sigma)$ and $0.23(4\\sigma)$. This is the first-time bound on the variation of the gravitational constant from CMB anisotropy spectrum. While the present deviation from the Einstein gravity is severely constrained ($\\alpha_0^2< 10^{-4}$), larger deviation during radiation dominated epochs is compatible with CMB. Indeed ${\\alpha_{ini}}^2$ up to $2 \\times 10^{-2}$ can be within $2 \\sigma$ level. Although our analysis is limited to flat models, the further small scale CMB anisotropy data which will be provided in future \\cite{PLANCK} would break the degeneracy and would significantly improve the bound."
    },
    "0311/astro-ph0311042_arXiv.txt": {
        "abstract": "We measured the bar pattern speed in a sample of 7 SB0 galaxies using the Tremaine-Weinberg method. This represents the largest sample of galaxies for which bar pattern speed has been measured this way. All the observed bars are as rapidly rotating as they can be. We compared this result with recent high-resolution $N$-body simulations of bars in cosmologically-motivated dark matter halos, and conclude that these bars are not located inside centrally concentrated halos. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311568_arXiv.txt": {
        "abstract": "We examine the effects of low-temperature, or $\\Delta$$n = 0$, dielectronic recombination (DR) on the ionization balance of the Fe M-Shell (Fe~IX through Fe~XVI). Since $\\Delta$$n = 0$ rates are not available for these ions, we have derived estimates based on the existing rates for the first four ionization states of the CNO sequence and newly calculated rates for L-shell ions of 3rd row elements and Fe. For a range of ionization parameter and column density applicable to the intrinsic absorbers detected in {\\it ASCA}, {\\it Chandra}, and {\\it XMM-Newton} observations of Seyfert galaxies, we generated two grids of photoionization models, with and without DR. The results show that the ionization parameter at which the population of an Fe M-shell ion peaks can increase in some cases by factor $>$ 2 when these rates are included. More importantly, there are dramatic changes in the range in ionization parameter over which individual M-shell ions contain significant fractions of the total Fe (e.g. $>$ 10\\%) in the plasma. These results may explain the mismatch between the range of Fe ionization states detected in the X-ray spectra of Seyferts, identified by the energies of the M-Shell Unresolved Transition Array, and those predicted by photoionization models of the X-ray absorbers that reproduce lines of second and third row elements. The results suggest that care should be taken in using 3rd and 4th row ions to constrain the physical conditions in photoionized X-ray plasmas until accurate DR rates are available. This underscores the importance of atomic physics in interpreting astronomical spectroscopy. ",
        "introduction": "New X-ray observatories such as the {\\it Chandra X-ray Observatory} and {\\it XMM- Newton} have made high resolution spectroscopy of active galaxies possible for the first time, allowing the detection of phases of gas that were previously unknown. The physical state of this gas is interpreted by reference to large-scale numerical simulations, hence, it is critical that physical processes that occur in an ionized plasma be fully understood (the main topic of the conference ``Spectroscopic Challenges of Photoionized Plasmas'', held at the University of Kentucky, Lexington, Kentucky, USA on the 15 - 18 November, 2000; [Ferland \\& Savin 2001]).  The ionization structure and temperatures in cosmic plasmas depend strongly on ionization and recombination processes. Recombination can occur via radiative recombination, charge transfer, or dielectronic recombination (DR). The DR process involves the capture of a free electron by an ion, resulting in doubly-excited recombined species with the electron in an auto-ionizing state. This is followed either by auto-ionization (i.e., the reverse process) or the emission of a photon, the latter leaving the recombined ion in a stable state.  There are two varieties of DR. The first involves the capture of the electron to auto-ionizing states with high principal quantum number, with an accompanying core excitation with a change in principal quantum number $\\Delta$$n > 0$ (e.g. Burgess 1964). This form of DR dominates over radiative recombination at temperatures comparable with the ionization potential, i.e collisionally ionized plasmas. Nevertheless, terms of the same parity as the ground term exist between the ground state and the first excited state of the opposite parity, and thus a series of autoionizing states converge on these terms (see Nussbaumer \\& Storey 1983). This results in a second type of DR in which there is no change in the principal quantum number of the excited core, or $\\Delta$$n = 0$, which can proceed at low temperatures for which the thermal energy of the free electron is less than the excitation energy of first resonance of the recombining ion. Low-temperature DR is usually the dominant process in photoionized plasmas, with temperatures typically 5 - 10\\% of the ionization potential, such as those detected in X-ray spectra of Seyfert galaxies (e.g Reynolds 1997; George et al. 1998; Sako et al. 2000; Kaspi et al. 2001, 2002; Kinkhabwala et al. 2002; Turner et al. 2003).  While low-temperature DR rates have been calculated for a number of abundant 2nd and a few 3rd row elements (Nussbaumer \\& Storey 1983; Nussbaumer \\& Storey 1986; Nussbaumer \\& Storey 1987), the full set of ionization states have only been completed for 2nd row elements. One of the difficulties in these calculations is that the rates depend on the accurate determination of the low-level autoionizing states (see Gu 2003). Ali et al. (1991) estimated low-temperature DR rate coefficients for other ions by taking the mean values from the first four ions of C, and O calculated by Nussbaumer \\& Storey (1983), and demonstrated that the model-predicted emission-line strengths were strongly affected when these rates were introduced. Savin (2000) has shown that uncertainties in DR rates translate into similar uncertainties in predicted ionic column densities, which may hamper the determination of elemental abundances ratios in the intergalactic medium. There has been recent work in determining both high- and low- temperature DR rates for L-shell (the Li- through Ne-like iso-electronic sequences), both experimentally (Savin et al. 1999, 2002) and computationally (Gu 2003; Badnell et al. 2003), the results of which emphasize the importance of this process in photoionized plasma. Note that, due to the difficulties in calculating the energies of the autoionizing states, experimental results are critical.  On the other hand, there has not been any recent work published on $\\Delta$$n = 0$ rates for M-shell ions (the Na- to Ar-like sequence). This is unfortuitous, since M-shell Fe ions, Fe~IX to Fe~XVI, contribute strongly to the unresolved transition array (UTA) of inner shell $n = 2-3$ lines seen in absorption in high resolution X-ray spectra of several Seyfert galaxies observed with {\\it XMM-Newton} and {\\it Chandra} such as IRAS 13349+2438 (Sako et al. 2001), NGC 3783 (Blustin et al. 2002; Kaspi et al. 2002; Netzer et al. 2003), Mrk 509 (Pounds et al. 2001; Yaqoob et al. 2003), and NGC 5548 (Steenbrugge et al. 2003; Steenbrugge et al., in preparation). For N- and M-shell Fe, the UTA features lie between $\\sim$ 15.5 \\AA~ and 17 \\AA~, depending on which ions are present (Behar, Sako, \\& Kahn 2001). In principal, the UTA can be used to constrain the ionization state of the absorbers, which is especially valuable if the resonance lines from H-, He-, and Li-like C, N, and O are saturated. However, in the case of NGC 3783  (Netzer et al. 2003), the ionization state suggested by the UTA wavelengths seems lower than that predicted by photoionization models generated to fit CNO absorption features and the ionization state/column density of 3rd row elements such as Si and S. A similar discrepancy appears to be the case for NGC 5548 (Steenbrugge et al., in preparation). One of the possible causes Netzer et al. suggested to explain this mismatch is the lack of $\\Delta$$n = 0$ rates for the M-shell Fe ions, which leads to the models predicting the Fe in too high an ionization state (compared to lower Z elements).  In order to explore the effect of $\\Delta$$n = 0$ DR on the M-shell Fe ionization balance, we adopted the approach described in Ali et al. (1991), adding estimated rate coefficients to the development version of the photoionization code Cloudy (Ferland et al. 1998). In order  to gauge the magnitude of the effect, model grids were generated with and without these rates. We also examined the effects of uncertainties in the estimates rates. The model details and results are described below. ",
        "conclusions": "The results of this simple test indicate that the model predictions of the Fe ionization balance depend critically on accurate $\\Delta$$n = 0$ DR rate coefficients. Photoionization models of X-ray absorbers that match spectral features from abundant 2nd row elements such as C, N, O, and Ne will overpredict the average Fe ionization state if this process is not included. This can be easily demonstrated using Figures 1 -- 3, by selecting an ionization parameter to give the proper ratio of the O~VI, O~VII, and O~VIII column densities and reading off the relative Fe ionic columns for the models with and without the estimated DR.  The apparent discrepancy in the ionic states detected in the Fe M-shell UTA in NGC 3783 (Netzer et al. 2003), compared to those predicted by the absorber models, is likely due to the absence of $\\Delta$$n = 0$ rates. To illustrate this point, we regenerated the so-called low ionization model used by Netzer et al. to fit the soft X-ray absorption and inner shell Si and S lines detected in a 900 ks merged {\\it Chandra}/High Energy Transmission Grating spectrum of NGC 3783 (see, also, Kaspi et al. 2002). Following Netzer et al., we assumed the SED described in Kaspi et al. (2001) and set log U $=$ $-$0.26 and N$_{H}$ $=$ 10$^{21.9}$ cm$^{-2}$ (note that our conversion from the X-ray ionization parameter used by Netzer et al. to the form of U used in CLOUDY resulted in a model prediction of slightly higher relative ionization of Si and S than they predicted, which may be due to slight differences in the model SED; in order to reproduce the same columns of Si and S we adjusted the value of U downward by 10\\%). The predicted ionic columns for Fe~VIII through Fe~XII, without $\\Delta$n$=$0 DR, are given in Table 2, and the combined Fe~XI and Fe~XII column is $\\sim$ 3.5 greater than the combination of Fe~VIII and Fe~IX, while the peak ionization state is Fe~XI. Since the strongest UTA features for Fe~XI and Fe~XII lie at 16.16\\AA~ and 15.98\\AA~, respectively (Behar et al. 2001), the UTA would peak near 16.1\\AA~, as described by Netzer et al.. Also shown are the predictions for the same model with $\\Delta$n$=$0 DR included. In this case, the peak ionization state is Fe~X, and the total Fe~VIII and Fe~IX column are approximately equal to those of Fe~XI and Fe~XII combined. Since Fe~VIII and FeIX produce their strongest features at 16.89\\AA~ and 16.51\\AA, respectively (Behar et al.), and the profile would be dominated by Fe~X (16.34 \\AA), the predicted UTA will peak near 16.4\\AA~, as observed. It should also be noted that including the new L-shell rates for 3rd row elements such as S, Si, and Mg will alter the predicted ionization balance for those elements as well. As noted by Behar \\& Netzer (2002), inner shell absorption lines from these ions can be used as effective probes of high-column absorbers. However care must be taken in using photoionization models to determine ionization states and total column densities of the absorbers using these lines as constraints, at least until a full set of accurate L-shell DR rates have become available.  Seyfert galaxies are also strong soft X-ray emission-line sources (e.g., Sako et al. 2000; Kinkhabwala et al. 2002; Turner et al. 2003). It will be important to understand the connection between this component and forbidden ``coronal'' emission lines such as [Fe~X] $\\lambda$6374, [Fe~XI] $\\lambda$7892, [Fe XIV] $\\lambda$5303 and [S~XII] $\\lambda$7611 (e.g. Kraemer \\& Crenshaw 2000), especially since the optical/UV spectral images obtained with the {\\it Hubble Space Telescope}/Space Telescope Imaging Spectrograph provide tighter constraints on the spatial extent and kinematic properties of the emission-line gas. Accurate DR rates are crucial for modeling this component and exploring the link between the X-ray and optical emission lines.  While the models generated for this Letter clearly illustrate the magnitude of the problem, we are left with the question on how best to proceed in the absence of a full set of $\\Delta$$n = 0$ DR rate coefficients. Our Monte Carlo calculations (see Figure 5) provide an estimate of the error in the predicted column densities due to the uncertainties in the rate coefficients. In fact, our results suggest that comparisons between calculated and observed column densities be given error bars $\\sim$ 0.3 dex, assuming our rates are included in the calculations, due to the uncertainties in the rates and the unmodeled temperature dependence (as suggested by the variations in the published L-shell rates over the range in temperature predicted by our models, e.g.,  1.8 x 10$^{4}$K $\\leq$ T $\\leq$ 4.5 x 10$^{5}$ K). Clearly, using rates estimated from other ions is far better than assuming the process does not occur, which is what is done when it is neglected. Hence, we suggest that it is best to adopt these estimated rate coefficients with this stated error. Ultimately, new atomic data, in general, and DR rates, in particular, are needed to reliably compare the ionization of third and fourth row elements."
    },
    "0311/astro-ph0311038_arXiv.txt": {
        "abstract": "{We report the first detection of triply-deuterated methanol, with 12 observed transitions, towards the low-mass protostar IRAS 16293$-$2422, as well as multifrequency observations of $^{13}$CH$_3$OH, used to derive the column density of the main isotopomer CH$_3$OH. The derived fractionation ratio [CD$_3$OH]/[CH$_3$OH] averaged on a 10$''$ beam is 1.4\\%. Together with previous CH$_2$DOH and CHD$_2$OH observations, the present CD$_3$OH observations are consistent with a formation of methanol on grain surfaces, if the atomic D/H ratio is 0.1 to 0.3 in the accreting gas. Such a high atomic ratio can be reached in the frame of gas-phase chemical models including all deuterated isotopomers of H$_{3}^{+}$. ",
        "introduction": "Despite the relatively low elemental abundance of deuterium in space (a factor of $\\sim$~1.5$\\times$10$^{-5}$ less abundant than H; \\cite{Linsky98} 1998), extremely large amounts of doubly-deuterated formaldehyde (D$_2$CO/H$_2$CO~$\\sim$~10\\%) have been observed in the solar-type protostar IRAS~16293$-$2422 (hereafter IRAS~16293, \\cite{Ceccarelli98} 1998, \\cite{Loinard00} 2000, \\cite{Ceccarelli01} 2001), initiating the search for other multiply deuterated molecules. Subsequently, doubly deuterated formaldehyde, doubly deuterated hydrogen sulfide and multiply deuterated ammonia have been observed in other protostars and dark clouds from where protostars form (\\cite{Roueff00} 2000, \\cite{Loinard01} 2001, \\cite{Ceccarelli02} 2002, \\cite{vanderTak02} 2002, \\cite{Lis02} 2002, \\cite{Vastel03} 2003). These studies have been interpreted in terms of two different routes for formaldehyde, hydrogen sulfide and ammonia deuteration: active grain chemistry followed by at least partial desorption into the gas for formaldehyde and hydrogen sulfide on the one hand and gas-phase chemistry for ammonia on the other hand. However, ammonia may also be a grain surface product, provided a large D/H atomic ratio in the accreting gas. Recently, doubly-deuterated methanol was detected towards IRAS~16293 (\\cite{Parise02} 2002). This observation provided new constraints for chemical models. The observations of the deuterated methanols CH$_{2}$DOH and CHD$_{2}$OH were both consistent with the formation of methanol from successive hydrogenations of CO by reaction with atomic H on grain surfaces, but required an atomic D/H ratio of 0.2 to 0.3 in the accreting gas.  At the time of the observation of doubly-deuterated methanol, no gas-phase model was able to predict such a high atomic D/H ratio. Meantime, observations of doubly deuterated formaldehyde in a sample of pre-stellar cores showed that the degree of deuteration increases with increasing CO depletion (\\cite{Bacmann02} 2002, 2003). This deuteration of formaldehyde in pre-stellar cores may occur partially in the CO-depleted gas-phase and partially on the surface of dust grains, followed by some inefficient desorption mechanism. A further spectacular confirmation of enhanced deuteration in CO-depleted gas came from the detection of abundant H$_2$D$^+$, likely the most abundant ion, in the prestellar core L1544 (\\cite{Caselli03} 2003). \\cite{Phillips03} (2003) suggested that in CO-depleted gas, even the multiple deuterated forms of H$_3^+$ may be abundant and play a role in the molecular deuteration enhancement. The suggestion has been fully confirmed by the modelling of \\cite{Roberts03} (2003), which shows that including HD$_2^+$ and D$_3^+$ in the chemical network increases dramatically the molecular deuteration, and allows the production of the large atomic D/H ratio predicted by the methanol observations (\\cite{Parise02} 2002).  In this paper, we report the first detection of triply-deuterated methanol CD$_3$OH in space, performed towards the solar-type protostar IRAS~16293. We also present a multifrequency observation of $^{13}$CH$_3$OH, used to derive the column density of the main isotopomer CH$_3$OH. These observations provide yet another stringent test to confirm the validity of grain surface models. ",
        "conclusions": "The main result of this Letter is the first detection of triply-deuterated methanol in space, with 12 detected transitions. This discovery follows the detection of doubly-deuterated as well as singly-deuterated isotopomers towards the same object (\\cite{Parise02} 2002). Observations of multiple isotopomers of methanol represent a powerful constraining tool for chemical processes that lead to such a high deuteration.  It is interesting to compare these observations to the predictions of the simple grain chemistry scheme of \\cite{Rodgers02} (2002). If the D atoms are randomly distributed in the methanol isotopomers (i.e. this scheme does not consider any activation barrier for the reactions but rather assumes that all reactions are equiprobable), the fractionation ratios R of each isotopomer relative to CH$_3$OH should scale as follows: R(CH$_3$OD)\\,=\\,$\\alpha$, R(CH$_2$DOH)\\,=\\,3$\\alpha$, R(CHD$_2$OH)\\,=\\,3$\\alpha^2$ and R(CD$_3$OH)\\,=\\,$\\alpha^3$, where $\\alpha$ is the accreting atomic D over H ratio. The three independent observations of CH$_2$DOH, CHD$_2$OH and CD$_3$OH are consistent within the error bars with a value of 0.1$-$0.2 for the D over H accretion rate. Accounting for the different mass of the atoms, this ratio corresponds to an abundance ratio in the gas-phase of D/H\\,=\\,$\\sqrt{2}\\times$(0.1$-$0.2)\\,=\\,0.15$-$0.3. However, this simple scheme fails to explain the observed low abundance of CH$_3$OD.  More accurate grain chemical models accounting for different activation barriers for the reactions have been developed in the last few years. We compare in the following our observations with the model developed by \\cite{Stantcheva03} (2003). This model is based on the direct solution of the master equation and therefore gives essentially the same predictions as the Monte Carlo models described by \\cite{Caselli02} (2002) or \\cite{Charnley97} (1997). Fig.~\\ref{model} shows predictions for fractionation ratios of deuterated isotopomers of methanol relative to CH$_3$OH versus the atomic D/H ratio in the accreting gas (\\cite{Stantcheva03} 2003) when the various isotopomers are formed by active grain chemistry. In the limit of low temperature (10~K), this model essentially gives the ratios corresponding to a random distribution of deuterium atoms. Observed fractionation ratios with their error bars have been overlaid on each curve, allowing the derivation of the required atomic D/H ratio in the gas-phase. The CD$_3$OH, CHD$_2$OH and CH$_2$DOH observations are consistent with a  formation on grain surfaces with an atomic D/H abundance ratio of 0.1$-$0.2. Such a high atomic fractionation ratio in the gas phase is predicted by the recent gas-phase model of \\cite{Roberts03} (2003), which involves not only H$_2$D$^+$ but also D$_2$H$^+$ and D$_{3}^{+}$ as precursors for deuterium fractionation, when the density of gas is very high and heavy species such as CO are strongly depleted.  As can be seen in Fig.~\\ref{model}, CH$_3$OD appears to be under-deuterated when compared with the grain chemical predictions. It is possible that the CH$_3$OD fractionation may be affected in the warm gas; e.g., this isotopomer may be preferentially converted into CH$_3$OH when released in the gas-phase by protonation reactions followed by dissociative recombination with an electron (\\cite{Charnley97} 1997, \\cite{Parise02} 2002) : \\vspace{-0.2cm} \\begin{center} $ {\\rm CH_3OD + H^+_3 \\rightarrow  CH_3ODH^+ + H_2}$  ${\\rm  CH_3OHD^+ + e^- \\rightarrow  CH_3OH + D .} $ \\end{center} \\vspace{-0.2cm} The corresponding reactions with H$_2$D$^+$, HD$_2^+$ and D$_3^+$ come of little importance in view of their low abundance in the warm gas of the hot core. This hypothesis, which assumes that protonation reactions attack the oxygen end of the methanol only (Osamura, Roberts \\& Herbst, in prep), could be tested by observing the CH$_2$DOD isotopomer. This observation may be difficult due to the expected low intensity of the lines.   \\begin{figure} \\includegraphics[width=9cm]{fracvshd_20juil.ps} \\caption{The deuterium fractionation ratio for the various deuterated isotopomers of methanol is plotted against the abundance ratio of deuterium to hydrogen atoms in the gas phase. Dashed lines: model results of  Stantcheva et al. 2003. Thick lines: observations of IRAS~16293. } \\label{model} \\end{figure}"
    },
    "0311/astro-ph0311512_arXiv.txt": {
        "abstract": "We show that the recently-discovered orbits of massive stars with closest approach of $\\sim 10^3$ Schwarzschild radii from \\sgra, allow winds from these stars to provide the required mass deposition rate near the black hole horizon. The observed luminosity of \\sgra does not require viscous transport of angular momentum, as long as the total wind mass loss rate reaches $\\sim 10^{-6}M_\\odot~{\\rm yr^{-1}}$ from the close-in stars. The specific orbits of the nearest stars should cause modulation of the radio and infrared flux from \\sgra on a timescale of years in a predictable fashion. ",
        "introduction": "Early in the development of accretion flow theory, it had been realized that angular momentum can severely limit the accretion rate of gas onto a black hole relative to the spherical infall value (the so-called {\\it Bondi} (1952) {\\it rate}).  Indeed, most of the hot gas revealed by recent X-ray observations within $\\sim 0.1$ pc of the Galactic center (Baganoff et al. 2003), is apparently not flowing into \\sgra, the massive black hole there (Quataert 2003).  The latest models of the {\\it radiatively inefficient accretion flow} around \\sgra estimate the accretion rate of gas into the black hole horizon at $\\sim 10^{-8}M_\\odot~{\\rm yr^{-1}}$, several orders of magnitude below the Bondi rate (Yuan, Narayan, Quataert 2003). This low accretion rate is required by the frequency dependence of the linear polarization data (Bower et al. 2003), in combination with the synchrotron radio flux and the recently detected infrared emission from \\sgra (Genzel et al 2003b; Ghez et al. 2003c).  The absolute minimum accretion rate required by the bolometric luminosity of \\sgra ($\\sim 10^{36}~{\\rm erg~s^{-1}}$) is $\\sim 2\\times 10^{-10}M_\\odot~{\\rm yr^{-1}}(\\epsilon/0.1)^{-1}$, where $\\epsilon$ is the radiative efficiency of the accreting mass.  The disparity between the large gas reservoir available at the outer radius of influence of \\sgra and the small mass that actually flows into its horizon, has inspired a variety of accretion models in which the mass inflow rate diminishes with decreasing radius, due to convection (Quataert \\& Gruzinov 2000a,b; Narayan et al. 2000), turbulence (Coker \\& Melia 1997; Pen et al. 2003), or outflows (Begelman \\& Blandford 1999; Yuan et al. 2002). In this {\\it Letter}, we highlight the significance of an important ingredient that was omitted by these models, namely the powerful winds from known stars that are embedded within the dilute accretion flow towards Sgr A*.  \\sgra was recently found to be surrounded by a cluster of young, massive stars.  Over the past year, three of these stars, named SO-2, SO-16, and SO-19 by Ghez et al. (2003c; see parallel work by Sch{\\\" o}del et al. 2003), were inferred to have orbits with peri-centric distances of $\\sim 100$AU or $\\sim 10^3$ Schwarzschild radii from the central black hole. The detection of a Br$\\gamma$ and He I features in of the spectrum of SO-2 indicates that it is probably an O8-B0 star, while the other bright stars are also likely to be O-B stars, consistent with the lack of a CO absorption feature in their spectrum (Ghez et al. 2003a; Eisenhauer et al. 2003).  The winds of O-stars on the main sequence have typical mass loss rates of ${\\dot M}_w=10^{-7}$--$10^{-6}M_\\odot~{\\rm yr^{-1}}$ and speeds in the range $v_w=1$--$3\\times 10^3~{\\rm km~s^{-1}}$ (Puls et al. 1996; Repolust, Puls, \\& Herrero 2003). In \\S 2.1, we show that SO-2, SO-16, and SO-19 can supply the required low-angular momentum gas that fuels the emission from \\sgra {\\it without requiring viscous transport of angular momentum}. To demonstrate this point, we adopt the simplest model in which the radial gas streams move on ballistic Keplerian orbits. Gas with higher angular momentum is assumed, for simplicity, to be driven out of the system, as inferred to be case at the Bondi radius (Quataert 2003).  In \\S 2.3 we show that the characteristic ram-pressure of the radial streams is well above the thermal pressure of the diffuse hot gas detected by the {\\it Chandra X-ray Observatory} (Baganoff et al. 2003), making the ballistic approach viable. In \\S 2.2 we illustrate how the specific orbital parameters of the close-in stars can be used to predict the long-term modulation of the radio and infrared flux of Sgr A*. Finally, we discuss the implications of these results in \\S 3. ",
        "conclusions": "We have found that radial gas streams from the winds of massive stars with recently discovered orbits that approach within $\\sim 10^3$ Schwarzschild radii of the black hole at the Galactic center, can supply the required mass flow near the central black hole horizon and power the radio and infrared emission of Sgr A*. For O star winds with a total ${\\dot M}_w\\sim 10^{-6}M_\\odot~{\\rm yr^{-1}}$, we get that SO-2, SO-16, and SO-19 can supply as much as $\\sim 10^{-8}(\\eta/3)^2 M_\\odot~{\\rm yr^{-1}}$ of gas in orbits directed straight into the emission region around the black hole horizon. The total mass inflow rate would increase if angular momentum of other gas elements is transported away by magnetic or turbulent stresses. The surprising existence of massive stars in orbits with closest approach distances of $\\sim 100$ AU poses a major puzzle to theories of star formation (Genzel et al. 2003a), and was not incorporated in previous modeling of the gas flow around \\sgra (e.g. Coker \\& Melia 1997; Yuan et al. 2003).  The parameters of the winds could in principle be determined through detailed spectroscopic observations of the stars in the \\sgra cluster. Once the wind properties of the individual stars are constrained, our accretion model will be easily testable since it correlates the long-term ($>$month) variability of the radio or infrared emission from \\sgra with the orbital phases of the stars around it (see Fig. 1).  The simplified calculation presented in this {\\it Letter} can be improved using a full hydrodynamic simulation and compared against observational data on the long-term variability of \\sgra (Zhao et al. 2003) for different choices of $v_w$ and ${\\dot M}_w$ per individual star (for an alternative source of this variability, see Liu \\& Melia 2002b). A proper numerical treatment is challenging since only a small fraction of the gas mass follows the radial orbits of interest here, and since the inflow spans several orders of magnitude in radius - requiring a large dynamic range in spatial and temporal resolution. Previous accretion simulations (such as those by Coker, Melia, \\& Falcke 1997, Proga \\& Begelman 2003 or Rockefeller et al. 2003) did not meet the demanding requirements of this accretion problem.  \\sgra may be surrounded by fainter stars that have not been detected so far, even though they dominate the steady mass flow into the black hole. However, winds from massive stars are driven by radiation pressure, and so the stars dominating the fluctuations in the mass flow into the black hole horizon should have the highest luminosities and be included in the current census of bright infrared sources near Sgr A*.  Massive stars that are much further away also possess powerful winds (Krabbe et al. 1991; Najarro et al. 1997; Coker \\& Melia 1997; Quataert 2003), but their individual contributions to the radial mass flux near the horizon of \\sgra is reduced compared to the close-in stars\\footnote{Statistically, the net mass flux is dominated by the close-in stars as long as the number density run of wind sources is steeper than $\\propto r^{-1}$, as is the case around \\sgra (Genzel et al. 2003a).}, since ${\\dot M}_{\\rm BH}\\propto r^{-2}$ in equation~(\\ref{eq:mdot}). For example, IRS 16C with its ${\\dot M}_w\\approx 8\\times 10^{-5} M_\\odot~{\\rm yr^{-1}}$ and $v_w \\sim 650~{\\rm km~s^{-1}}$ at $r\\ga 2\\times 10^4$AU, and IRS 13E1 with its ${\\dot M}_w\\approx 8\\times 10^{-4} M_\\odot~{\\rm yr^{-1}}$ and $v_w \\sim 1000~{\\rm km~s^{-1}}$ at $r\\ga 4\\times 10^4$AU (Quataert, Narayan, \\& Reid 1999), produce at most a radial mass flux comparable to that of SO-2 if the latter has ${\\dot M}_w\\sim 10^{-6}M_\\odot~{\\rm yr^{-1}}$.  Nevertheless, the steady mass flux into \\sgra may still be dominated by undetected stars or by gas that originates at large distances and loses its angular momentum through viscous transport (Yuan et al. 2003). In this case, the winds from the close-in stars would only modulate a fraction of ${\\dot M}_{\\rm BH}$, but their fractional contribution could still be identified through the predictable short-term variability that their orbits introduce to the luminosity of Sgr A*.   \\bigskip \\noindent {\\bf Acknowledgements}  I thank Jean-Pierre Lasota, Eric Pfahl, Eliot Quataert, Mark Reid, George Rybicki, and Eli Waxman for comments on the manuscript.  This work was supported in part by NASA grant NAG 5-13292, and by NSF grants AST-0071019, AST-0204514."
    },
    "0311/astro-ph0311348_arXiv.txt": {
        "abstract": "In the inner parts of spiral galaxies, of high or low surface brightness, there is a close correlation between rotation curve shape and light distribution. For any feature in the luminosity profile there is a corresponding feature in the rotation curve and vice versa. This implies that the gravitational potential is strongly correlated with the distribution of luminosity: either the luminous mass dominates or there is a close coupling between luminous and dark matter. In a similar way, the declining rotation curves observed in the outer parts of high luminosity systems are a clear signature of the stellar disk which either dominates or traces the distribution of mass.  The notion that the baryons are dynamically important in the centres of galaxies, including LSBs, undermines the whole controversy over the cusps in CDM halos and the comparison with the observations. If the baryons dominate in the central regions of all spirals, including LSBs, how can the CDM profiles be compared with the observations? Alternatively, if the baryons do not dominate but simply trace the DM distribution, why, in systems of comparable luminosity, are some DM halos cuspy (like the light) and others (also like the light) are not? ",
        "introduction": "The relation between the distribution of light and the shape of rotation curves has been at the centre of the discussion on the interpretation of rotation curves and on the presence and the amount of dark matter in spiral galaxies. Burstein and Rubin (1985) discussed the systematic properties of rotation curves and concluded that for a given morphological type the shapes of rotation curves vary systematically with luminosity; these shapes are, however, similar for different morphologies and therefore the form of the gravitational potential is not correlated with the form of the light distribution or the morphological type. This was used as an argument in favour of the presence of dark matter also in the inner regions of galaxies.  In contrast to this conclusion, the success of the maximum disk hypothesis in the analysis of rotation curves (Kalnajs 1983, Van Albada and Sancisi 1986, Kent 1986, 1987, Broeils 1992, Palunas and Williams 2000) implies that within the optical disk either the mass of dark matter is small or that its distribution is very closely coupled to the distribution of luminous matter. The maximum disk analysis was carried out on observations of predominantly high surface brightness (HSB) systems. Recent work on low luminosity and low surface brightness (LSB) galaxies has shown that for the inner parts it is possible to obtain maximum disk fits as good as for the more luminous systems (Verheijen 1997, Swaters 1999). The main problem with the maximum disk in these systems is caused by the required large M/L ratios which would seem to argue in favour of DM being dominant everywhere.  Persic and Salucci (1991) and Persic, Salucci and Stel (1996) have discussed the relation between luminosity and rotation curves and have emphasized the strong dependence on luminosity for both the shape and the amplitude of the rotation curve. They have come to the conclusion, based on a large galaxy sample, that the observed curves have a universal shape, dependent only on total luminosity. However, Verheijen (1997) finds that roughly one third of the observed HI rotation curves from his Ursa Major sample of galaxies deviate noticeably from the \"universal rotation curve\" shape.  A new interest in the question of the shape of rotation curves and of the mass distribution in the central parts of spirals has recently been originated by the debate on the presence of cusps as predicted by the Cold Dark Matter simulations on structure formation (Navarro, Frenk and White 1997). Naturally the low luminosity and LSB galaxies, believed to be DM dominated, are taken as the best and most suitable for the test and their central rotation curves are thoroughly investigated. It is clear that if baryons dominate the mass distribution in the central regions of galaxies this would affect the whole discussion on CDM cusps and the comparison with the rotation curves; however, this could also reveal something about the distribution of DM. The problem is then to properly understand how cores are formed. If, on the other hand, DM dominates everywhere and closely follows the distribution of the light, then the question is why some dark halos are centrally peaked like the stars while others are not (again mimicking the light). ",
        "conclusions": ""
    },
    "0311/astro-ph0311381_arXiv.txt": {
        "abstract": "{ The evolution with redshift of the temper\\-ature-luminosity relation of X-ray galaxy clusters is a key ingredient to break degeneracies in the interpretation of X-ray clusters redshift number counts. We therefore take advantage of the recent measurements of the temperature-luminosity relation of distant clusters observed with XMM-Newton and Chandra satellites to examine theoretical number counts expected for different available X-rays cluster samples, namely the RDCS, EMSS, SHARC, 160deg$^2$ and the MACS at redshift greater than 0.3. We derive these counts without any adjustment, using models previously normalized  to the local  ($z\\sim 0.05$) temperature distribution function (TDF) and to the high-$z$ ($z \\sim 0.33$) TDF. We find that these models having $\\Omega_M$ in the range [0.85-1.] predict counts in remarkable agreement with the observed counts in the different samples. We illustrate  that this conclusion is weakly sensitive to the various ingredients of the modeling. Therefore number counts provide a robust evidence of an evolving population. A   realistic flat low density model  ($\\Omega_M = 0.3$), normalized to the local abundance of clusters is found to overproduce cluster abundance at high redshift (above $z \\sim 0.5$) by nearly an order of magnitude. This result is in conflict with the popular concordance model.  The conflict could indicate a deviation from the expected scaling of the $M-T$ relation with redshift. ",
        "introduction": "The evolution of the temperature distribution function (TDF) of X-ray clusters is known to be a powerful cosmological test of the density parameter of the Universe (Oukbir \\& Blanchard 1992). Indeed, the evolution with redshift of the mass function, on cluster scales, once normalized to present day, appears to be a robust  cosmological test, with an exponential sensitivity to $\\Omega_M$ through the gravitational growth rate of perturbations (Blanchard \\& Bartlett 1998). The high sensitivity of this test has allowed its first direct application from a sample  of ten (revised to 9) clusters at redshift 0.3 with measured temperature (Henry 1997). However, several authors have inferred somewhat different values from this single sample (Eke et al. 1998; Viana \\& Liddle 1999; Blanchard et al. 2000, hereafter B00). This might not be so surprising given the small size of the sample as well as the limited range of redshift. An alternative approach to track the evolution of the abundance of clusters is from their redshift distribution in  X-ray samples for which the selection function is known. This procedure, reducing the time telescope investment, has been applied to the EMSS sample (Oukbir \\& Blanchard 1997; Reichart et al. 1999) and to the RDCS sample (Borgani et al. 1999, 2001) but the conclusion on $\\Omega_M$ depends on the possible evolution of the $L-T$ relation (Sadat et al. 1998, hereafter SBO98; Borgani et al. 1999, Novicki et al. 2002). The various ingredients used  in this modeling were not necessary accurately known, most noticeably the temperature mass ($T-M$) relation, and the luminosity-temperature ($L-T$) relation and its possible  evolution with redshift. Therefore doubts have been raised up on  the applicability of this test  given these uncertainties (Colafrancesco et al. 1997; Rosati et al. 2002).  The XMM-$\\Omega$ project was conducted in order to provide an accurate estimation of the possible evolution of the luminosity temperature relation at high redshift  for clusters of medium luminosity which constitute the bulk of X-ray selected samples, allowing to remove a major source of degeneracy in the determination of $\\Omega_M$ (Bartlett et al. 2001). To maintain a better control on the resulting $L-T$ it is clearly preferable to use clusters homogeneously obtained from  X-ray selected samples. Therefore, the choice of the SHARC surveys in the XMM-$\\Omega$ project provides several advantages:  large angular coverage with  the Bright SHARC (Nichol et al. 1999; Romer et al. 2000),  deepness of the sample with the South SHARC (Burke et al. 1997), still keeping the number of clusters  to a  realistic size for comprehensive X-ray investigations.   The purpose of this letter is to examine the expected number counts in comparison with the observed counts in two different cosmological models, namely a concordance model and a high matter density flat universe,   and to examine the amplitude of the major sources of uncertainties: the statistical dispersion on the value of $\\sigma_8$ from the finite local sample, the systematic uncertainties in the $M - T$ and the $L - T$ relations at high redshift. In this modeling we take advantage of the accurate knowledge of the $L-T$ relation provided by the first results from the XMM-$\\Omega$ project (Lumb et al. 2003, hereafter L03) and from recent Chandra measurements of distant clusters (Vikhlinin et al. 2002).  \\begin{table} \\begin{tabular}{lllllc} \\hline  T$_{15}$   & $\\Omega_M$ & $\\sigma_8$ & $\\Gamma$ & Cosmological model  \\\\ (keV)      &            &            &          &    and ingredients \\\\ \\hline 4         &  1.      &    0.55   &   0.12     &  A: best model+BN98+SMT\\\\ 6.5      &  0.85       &    0.455   &   0.1     &  A: best model+M98+SMT\\\\ 4         &  0.3       &    1.  &   0.2      & B: Low $\\Omega_M$+BN98+SMT \\\\ 6.5      &  0.3       &    0.725   &   0.2     &  B: Low $\\Omega_M$+M98+SMT\\\\ \\hline \\end{tabular} \\caption{Models and parameters used in the number counts calculations. } \\label{tab2} \\end{table} ",
        "conclusions": "As we have seen a model which is normalized to the local as well as to the high redshift TDF, reproduces impressively well the redshift distribution of all the surveys we have investigated, without any adjustment and with little  uncertainties arising from the modeling. In contrast, we have checked that relaxing the request of matching the local abundance  a concordance model fitting the redshift distribution leads to a local  abundance of  X-ray clusters  $N(>4 {\\rm keV},z = 0.05)$  nearly 5 times lower than actually observed. This is a  strong indication that existing samples of clusters (namely the Henry sample, the RDCS, the EMSS, the Bright SHARC, the 160deg$^2$ and the MACS sample) draw the same picture, consistently pointing out toward the fact that the cluster abundance is significantly evolving with redshift, perhaps closing a long term controversy on this question (Henry et al. 1992). Furthermore, in standard hierarchical picture of structure formation such evolution points toward a high matter density universe with $\\Omega_M$ in the  range [0.85-1.], the precise value  depending on the $M-T$ normalization. This conclusion is clearly conflicting with the currently popular concordance model. However, it should be emphasized that this is entirely consistent with {\\em all previous existing analyzes} performed on the redshift distribution of  X-ray selected samples of  clusters performed with the same methodology: our conclusion corroborates previous analyzes of the EMSS clusters redshift distribution: SBO98 found $\\Omega_M=0.85\\pm{.2}$, Reichart et al. (1999) found $\\Omega_M=0.96\\pm{.3}$  as well as the RDCS redshift distribution for models normalized to the present-day X-ray clusters abundance. Indeed, from Eq. (14) in Borgani et al. (1999), $\\Omega_M=(\\beta+2)/3\\pm{1/3}$ we derive $\\Omega_M \\sim 0.88\\pm{.34} $  from our $\\beta \\sim 0.65\\pm{.21}$. Note that high $\\Omega_M$ models are also consistent with WMAP (Blanchard et al. 2003). A possible loophole could be a large systematic bias flux, but we have argued that this seems very unlikely. An other possibility would be that the scaling in the redshift of the $M-T$ relation (Eq. 1) is completely wrong, violating the basic scaling scheme. Voit (2000) has investigated such a possibility, but concluded to a  moderate effect. In  Fig. 2, we have plotted the predicted counts in a concordance model, assuming \\begin{equation} T = T_{15}(\\Omega_M \\Delta(z,\\Omega_M)/178)^{1/3}M_{15}^{2/3} \\end{equation} instead of Eq. 2. As one can see, such a modification reestablishes agreement of the concordance model with observations. It is well known that the $L-T$ relation cannot  be explained from simple scaling arguments. One may therefore argue that the redshift evolution of the $M-T$ relation may suffer from more  dramatic effect than usually assumed, although -- to our knowledge -- such a possibility has never been advocated and it is probably not obvious to find physical motivation leading to gas thermal energy in distant clusters ($z\\sim 1$) to be  reduced by a factor of two compared to  clusters in the local universe. We therefore conclude that the redshift distributions of present--day available X-ray clusters surveys, as well as the recent results on the $L-T$ relation of high redshift clusters, favor a high matter density universe unless the standard paradigm on clusters gas physics has to be deeply revised."
    },
    "0311/astro-ph0311454_arXiv.txt": {
        "abstract": "Winds from AGN and quasars will form large amounts of dust, as the cool gas in these winds passes through the (pressure, temperature) region where dust is formed in AGB stars.  Conditions in the gas are benign to dust at these radii.  As a result quasar winds may be a major source of dust at high redshifts, obviating a difficulty with current observations, and requiring far less dust to exist at early epochs. ",
        "introduction": "Dust is common in high redshift ($z$=4--6) quasars (Omont et al. 2001, Priddey \\& McMahon 2001), apparently implying that dust is widespread by this epoch. This presents a puzzle: how can so much dust be formed in the short time available? WMAP cosmology (Spergel et al. 2003) puts the age of the universe at z=6 to be only 0.95~Gyr, and the age of first reionization (and so the first stars) at an age of $\\sim$0.2~Gyr (z$\\sim$20).  In our galaxy the primary source of dust lies in the winds of red giant stars, specifically Asymptotic Giant Branch (AGB) stars. Only stars with masses greater than 2~$M_{\\odot}$ can reach the AGB in the less than $\\sim$1~Gyr available. Another source of dust is thus required.  Normally this extra source is taken to be type~1 (massive star) supernovae. Recent SCUBA results (Dunne et al. 2003) show that the 300~year old type~1 supernova remnant (SNR) Cas~A is rich in dust, lending credence to this picture. However, the possible destruction of SNR formed dust in their shocks may not be fully resovled. Here we point out another, seemingly inevitable, path for dust formation up to at least z=6: {\\em quasar winds}.  The fate of outflowing broad emission line (BEL) quasar gas constantly being ejected from quasars had previously not been considered. We examined that fate (Elvis, Marengo \\& Karovska 2002) in a manner that applies to any model in which the BEL clouds move outward. We find that dust creation is a natural consequence. ",
        "conclusions": "In summary, dust will inevitably be created from the free expansion of quasar broad emission line clouds in an outflowing wind. The association of dust with quasars is not then {\\em necessarily} linked with intense star formation around quasars, but may be a consequence of the quasar activity itself. The creation of dust in quasar winds may solve the puzzle of where the very first dust comes from, and in doing so suggests an unexpected role for quasars in cosmology."
    },
    "0311/hep-th0311077_arXiv.txt": {
        "abstract": "{\\small The distance between BPS branes in string theory corresponds to a flat direction in the effective potential. Small deviations from supersymmetry may lead to a small uplifting of this flat direction and to brane inflation. However, this scenario can work only if the BPS properties of the branes and the corresponding flatness of the inflaton potential are preserved in  the stable volume compactification. We present an ``inflaton trench'' mechanism that keeps the inflaton potential flat due to shift symmetry, which is  related to near BPS symmetry in our model. }  \\begin{document}  \\parskip 5pt ",
        "introduction": "It is difficult to find a mechanism for inflation in string theory. Brane inflation models \\cite{dbraneinflation,Herdeiro:2001zb,Dasgupta:2002ew,angles} provide a significant step towards this goal. However, the validity of  brane inflation models in string theory depends on the ability to stabilize the compactification volume. This means that an effective 4d theory has to be found in which there is a potential for the scalar field representing the volume of the internal space. This potential has to have a minimum in which the volume field is trapped. This potential should simultaneously describe the inflaton field corresponding to the distance between branes. The effective theory has to fix the volume modulus while keeping the potential  for the distance modulus flat.  Brane inflation models preserving supersymmetry should not have any potential for the modulus representing the distance between BPS branes. When supersymmetry is slightly broken, so that the configuration of branes is a near BPS state, a small attractive potential develops for the distance modulus, which in turn may lead to a slow-roll stage of inflation in the effective 4d theory. Brane inflation models of this type, where small deviations from the BPS configuration trigger inflation were studied in \\cite{Herdeiro:2001zb,Dasgupta:2002ew,angles}.  In the effective field theory, this situation may be characterized by a property of the inflaton field: its mass should be small and definitely much smaller than the Hubble parameter, $m^2_{\\rm infl}\\ll H^2$. In supergravity there are known mechanisms to provide this type of relation. Brane inflation models developed without an explicit moduli stabilization mechanism assumed that the stabilization mechanism would not lead to a large inflaton mass. However, the recent study of D3/D$\\bar{3}$ inflation in a warped background with fluxes and stable volume modulus \\cite{Kachru:2003sx} has shown that the danger of $m^2_{\\rm infl}\\sim H^2$ may persist, unless a fine-tuning can be performed.   The purpose of this paper is to suggest a mechanism which protects the inflaton field from acquiring a large mass in the context of stable volume brane inflation models. The mechanism can be intuitively understood as exhibiting the fact that BPS branes should be in neutral equilibrium. Compactification with a stable volume should preserve this fact under certain conditions.  We originally found this mechanism by trying  various potentials in Mathematica. We were looking for the possibility of having some kind of a flat direction: the inflaton should form a trench which gives an additional dimension to the narrow well trapping the volume modulus in the simplest KKLT model \\cite{Kachru:2003aw} or in more general models \\cite{egq}. We have dubbed the phenomena as the ``inflaton trench'' mechanism. Only later did we realize that the corresponding potentials allow the realization of an {\\it exact unbroken supersymmetry} in AdS space. One may try to keep this property for near BPS branes in near dS space, where it will be only approximately correct, but will preserve the $m^2_{\\rm infl}\\ll H^2$ relation.  As a specific example, we will work with the D3/D7 brane inflation model which  is related to the hybrid model \\cite{Linde:1991km} of D-term inflation \\cite{Binetruy:1996xj}. We will find that it is compatible with the KKLT mechanism of the volume stabilization. D-term inflation models in supergravity do not require a fine-tuning with regard to the inflaton mass since the source of inflationary energy is coming from the D-term while the F-term vanishes during inflation. Therefore the dangerous \\Ka\\ corrections, which could lead to $m^2_{\\rm infl}\\sim H^2$, vanish. However, the superpotential volume stabilization mechanism suggested in the KKLT model \\cite{Kachru:2003aw} requires some non-vanishing F-terms. In this paper we will suggest a consistent combination of these two  ideas which may provide the basis for stringy brane inflation models with a stable volume. ",
        "conclusions": "In this note we proposed an ``inflaton trench'' mechanism which may lead to a viable model of inflation with volume stabilization in string theory. The shift symmetry of the classical theory protects the small mass of the inflaton in the framework of a string compactification with a stabilized volume modulus. This symmetry is shown to follow from the description of BPS branes in a supersymmetric equilibrium. A small deviation from supersymmetry leads to a small violation of the shift symmetry which allows a slow-roll regime of inflation. A similar mechanism may be responsible for the present stage of acceleration of the universe.  A picture of a ``stringy landscape'' that has been developed in \\cite{Susskind:2003kw} takes into account, in particular, a profile of the KKLT model. Now it may be enriched  with the trenches of the kind shown in Fig. 2  where inflation or late-time acceleration of the universe may take place.     We are planning to develop these considerations in the context of the D3/D7 brane construction, with a hope that it may provide us with a realistic model of brane inflation with stable compactification. One can also expect that the ``inflaton trench'' mechanism may be useful for other models of brane inflation. We hope to return to the discussion of these possibilities in a separate publication.        \\subsection*"
    },
    "0311/astro-ph0311440_arXiv.txt": {
        "abstract": "We explore the dynamics of small planetesimals coexisting with massive protoplanetary cores in a gaseous nebula. Gas drag strongly affects the motion of small bodies leading to the decay of their eccentricities and inclinations, which are excited by the gravity of protoplanetary cores. Drag acting on larger ($\\gtrsim 1$ km), high velocity planetesimals causes a mere reduction of their average random velocity. By contrast, drag qualitatively changes the dynamics of smaller ($\\lesssim 0.1-1$ km), low velocity objects: (1) small planetesimals sediment towards the midplane of the nebula forming vertically thin subdisk; (2) their random velocities rapidly decay between successive passages of the cores and, as a result, encounters with cores typically occur at the minimum relative velocity allowed by the shear in the disk. This leads to a drastic increase in the accretion rate of small planetesimals by the protoplanetary cores, allowing cores to grow faster than expected in the simple oligarchic picture, provided that the population of small planetesimals contains more than roughly 1\\% of the solid mass in the nebula. Fragmentation of larger planetesimals ($\\gtrsim 1$ km) in energetic collisions triggered by the gravitational scattering by cores can easily channel this amount of material into small bodies on reasonable timescales ($< 1$ Myr in the outer Solar System), providing a means for the rapid growth (within several Myr at 30 AU) of rather massive protoplanetary cores. Effects of inelastic collisions between planetesimals and presence of multiple protoplanetary cores are discussed. ",
        "introduction": "\\label{sect:intro}   Formation of terrestrial planets and solid cores of giant planets is thought to proceed via the gravity-assisted merging of a large number of planetesimals --- solid bodies with initial sizes of roughly several kilometers. Despite a considerable progress made in this field since the pioneering works of Safronov (1969), a number of important problems still remain unsolved. One of the most serious questions has to do with the time needed for planets to complete their growth to present sizes. In the framework of conventional theory this time is rather long, especially in the outer parts of the protoplanetary nebula ($\\gtrsim 10^8-10^9$ yr), and it is likely that the gaseous component of the nebula dissipates much earlier (in $\\lesssim 10^6-10^7$ yr). This would make it very hard for the giant planets in our Solar System to accrete their huge gaseous envelopes via core instability (Mizuno 1980) which is otherwise considered to be an attractive scenario.  Wetherill \\& Stewart (1989) have identified a very rapid ``runaway'' regime of accretion of protoplanetary cores which at the time seemed like a solution of this problem. However, later on Ida \\& Makino (1993) and Kokubo \\& Ida (1996, 1998) have demonstrated that the runaway accretion would persist only through a rather limited interval of time and that the final growth of protoplanetary cores to isolation (which corresponds to roughly $10^{26}-10^{27}$ g at 1 AU) would proceed in a slow manner, making the formation of cores of giant planets rather problematic.  These studies usually implied that the gaseous component of the nebula plays only a secondary role in the planet formation process. Planetesimals were typically assumed to be rather massive ($10^{23}-10^{24}$ g) bodies weakly affected by the gas (Kokubo \\& Ida 1998). This allows gravity to excite energetic random motions of planetesimals leading to diminishing the role of gravitational focusing and reduction of accretion efficiency. The purpose of this paper is to relax this assumption and to see what impact is incurred on the planet formation picture by allowing most of the planetesimals to be small ($\\lesssim 10$ km) bodies immersed in a gaseous environment. Such planetesimals would be appreciably affected by the gas drag and we will demonstrate that this can bring qualitative changes to their dynamics in the vicinity of the protoplanetary cores and, consequently, to the behavior of the mass accretion rate of cores.  Throughout this study we will use the following approximation to the structure of the Minimum Mass Solar Nebula (MMSN): \\begin{eqnarray} && \\Sigma_g(a)\\approx100\\Sigma_p(a)\\approx 3000~\\mbox{g cm}^{-2}~a_{AU}^{-3/2},\\\\ && c_s(a)\\approx 1.2~\\mbox{km s}^{-1}~a_{AU}^{-1/4}, \\label{eq:MMSN} \\end{eqnarray} where $\\Sigma_p, \\Sigma_g$ are the particulate and gas surface densities correspondingly, $c_s$ is the gas sound speed, and $a_{AU}\\equiv a/(1~\\mbox{AU})$ is a distance from the Sun $a$ scaled by 1 AU. We will use the terms ``protoplanetary core'' and ``protoplanetary embryo'' interchangingly. Physical density of planetesimals $\\rho_p$ is always assumed to be $1$ g cm$^{-3}$.   The paper is organized as follows: after a discussion of different gas drag regimes in \\S \\ref{sect:dissip} we proceed to the description of planetesimal dynamics in the vicinity of protoplanetary cores in \\S \\ref{sect:scat}. The inclination of small planetesimals, a question very important for this study, is explored in \\S \\ref{subsect:inclination}. The separation of different gas drag and dynamical regimes in different parts of the nebula is described in \\S \\ref{subsect:separation} and lower limits on the random velocities of planetesimals are obtained in \\S \\ref{sect:lower}. We dwell upon the role of inelastic collisions between planetesimals in \\S \\ref{subsect:inelastic}. The role of small planetesimals in the growth of protoplanetary cores is studied in \\S \\ref{sect:accretion} and some important consequences for the planet formation picture are discussed in \\S \\ref{sect:discus}. ",
        "conclusions": "\\label{sect:concl}    We have explored the details of planetesimal dynamics near protoplanetary embryos in the presence of gas drag. We showed that large ($\\gtrsim 1$ km in size) planetesimals  are kept in the dispersion-dominated regime as a result of scattering by protoplanetary cores scattering, although their average velocities are reduced by the gas drag. Bodies smaller than roughly $0.1-1$ km (depending on the location in the nebula) interact with protoplanetary cores in the shear-dominated regime. Owing to the action of the gas drag these planetesimals settle into a geometrically thin layer near the nebular midplane;  between consecutive encounters with the embryos they experience strong velocity damping which allows them to approach embryos every time with the relative velocity comparable to the Hill velocity --- minimum velocity which can be achieved in a differentially rotating disk.  For different locations in the proto-Solar nebula we have determined  which planetesimals are only weakly affected by the gas drag and are in the dispersion-dominated regime, and which planetesimals are so strongly coupled to the gas that their velocities are below the Hill velocity of the protoplanetary cores. Dynamical peculiarities of the shear-dominated regime in the presence of gas lead to a very high efficiency of accretion of small bodies by the embryos. If the surface mass density of small bodies is high enough ($\\gtrsim 1\\%$ of the total surface density in solids) their accretion would dominate the embryo's growth rate (relative to the accretion of more massive, dispersion-dominated planetesimals).  Large planetesimals ($\\gtrsim 1$ km in size) likely containing most of the mass in solids have large random velocities and are not very efficiently accreted by embryos. However, they inelastically collide with each other at high velocities and fragment into smaller pieces contributing to the population of small bodies. Thus, the embryo's growth by accretion of small shear-dominated planetesimals can be regulated by the fragmentation of bigger, dispersion-dominated bodies. Planetesimal fragmentation would probably be easiest in the outer Solar System where colliding bodies are mostly composed of ices and are therefore internally weak and susceptible to easy destruction. The natural timescale for the growth of $10^{25}$ g protoplanetary embryos by accretion of small planetesimals turns out to be around several Myrs at $30$ AU from the Sun. Cores of ice giants can be formed in $\\sim 10^7$ yr after large planetesimals lose all their mass in catastrophic collisions to small debris which can be effectively accreted by massive embryos. This scenario would work only if the population of protoplanetary cores formed as the outcome of oligarchic growth can be kept on almost non-inclined and circular orbits. As we demonstrated in \\S \\ref{subsect:migration}, this is possible only if embryo masses do not exceed a specific limit dictated by the efficiency of planetesimal dynamical friction, which sensitively depends on the planetesimal mass fraction $\\chi$ locked up in small planetesimals and velocity of large dispersion-dominated planetesimals.  Future work should address the issues of the self-consistent evolution of the mass fraction $\\chi$ contained in small bodies; the role of inelastic collisions using improved fragmentation physics (can be an issue in the inner Solar System); the final fate of massive embryos which cannot be kept dynamically cold by the planetesimal dynamical friction, and so on."
    },
    "0311/astro-ph0311395_arXiv.txt": {
        "abstract": "A new ephemeris has been determined for the supersoft X-ray binary CAL~83 using MACHO photometry.  With an improved orbital period of 1.047568 days, it is now possible to phase together photometric and spectroscopic data obtained over the past two decades with new far ultraviolet spectra taken with FUSE.  We discuss the properties of the orbital and longterm optical light curves as well as the colors of CAL~83.  In the far ultraviolet the only well-detected stellar feature is emission from the O~VI resonance doublet.  The radial velocity of this emission appears to differ from that of He~II in the optical region, although we only have partial phase coverage for the O~VI line.  The FUSE continuum variations are similar to the optical light curve in phase and amplitude. ",
        "introduction": "The X-ray binary CAL~83 is one of a small number of supersoft sources which are characterized by an extremely soft X-ray spectrum with almost no radiation above $\\sim0.5$ keV.  This group includes a handful of sources in the Magellanic Clouds and the Galaxy (e.g. Greiner 1996, Hasinger 1996, Cowley et al. 1998), plus some less well studied objects in other nearby galaxies (e.g. Kahabka 1999; Swartz et al. 2002; Kong \\& Di Stefano 2003). CAL~83 itself is a member of the LMC, based on its position and radial velocity.  From radial velocity variations of the optical emission lines (especially He~II and H), Crampton et al.\\ (1987) discovered that CAL~83 is a binary system with a period near one day.  The emission lines are formed in an accretion disk that is viewed at a low inclination angle, so the resulting velocity amplitude is small (K$\\sim30$ km s$^{-1}$).  There is no spectroscopic evidence of either star, but the period and velocities suggest that they have masses well under 2M${\\odot}$ (see Crampton et al. 1987, Smale et al. 1988).  CAL~83 is known to exhibit irregular brightness variations of more than a magnitude over timescales of months (e.g. Smale et al. 1988, Alcock et al. 1997, Greiner \\& Di Stefano 2002).  Although CAL~83 is considered a nearly steady X-ray source, occasionally it drops into an X-ray-off state (e.g. Kahabka 1996, 1998).  Using MACHO data, Alcock et al.\\ and Greiner \\& Di Stefano have shown that these low X-ray states are followed within $\\sim30-50$ days by a low optical state.  The most likely explanation is that there is a change in the accretion flow and hence in the accretion disk which provides the optical luminosity.  In spite of these large changes in the mean brightness of CAL~83, Smale et al.\\ were able to detect small superimposed orbital variations with a full amplitude of $\\sim$0.2 mag and a period of P = 1.0436 days.  During the orbital period the light curve shows a sinusoidal variation with one maximum and one minimum per cycle.  Cowley et al.\\ (1991) later refined the orbital period to be P = 1.0475 days by combining all published and new spectroscopic data as well as the photometry of Smale et al.  Given the low inclination of the system, it is unlikely that the minimum is caused by an occulatation by the secondary star.  Furthermore, the phasing of the emission line velocities implies that the light minimum does not coincide with conjunction of the stars.  The variation is likely to arise from a partial obscuration of the bright inner accretion disk. Possible occulting regions might be a splash region raised by impact of the accretion stream or a gas stream which is significantly above the orbital plane.  Although we are uncertain what causes the orbital light curve, we have been able to use it to refine the orbital period to allow intercomparison of two decades of spectroscopic data. ",
        "conclusions": "To summarize, we have used MACHO photometry to refine the orbital period of the LMC supersoft X-ray source CAL~83.  With our new ephemeris, we have been able to compare previously published spectroscopic and photometric data using consistent phases.  In a high optical state, its color becomes bluer and the amplitude of the orbital variations increases.  New FUSE observations show that orbital variations in brightness in the far ultraviolet are similar to those in the optical region.  However, the velocity of the O~VI emission in the FUSE spectra indicates its behavior may differ from the optical He~II line."
    },
    "0311/astro-ph0311506_arXiv.txt": {
        "abstract": "This paper presents a method to reconstruct the position of the line of sight and orientation of the focal plane of the {\\it Planck} satellite, which may be expressed by the parameters of the effective boresight and roll angles, respectively. The accuracy to which these parameters must be determined to avoid compromising the reconstruction of the $C_{\\ell}$ power spectrum is assessed. It is expected that the pointing reconstruction will be performed using point source transits in the highest frequency channels, given their smaller beam sizes. Estimates are made for the number of galactic and extragalactic point sources visible by {\\it Planck} in these channels. The ability of the method presented here to reconstruct the pointing parameters to the required accuracies using these point sources is investigated. It is found that the scanning strategy employed and the actual orientation of the focal plane both independently influence the available number of suitable point sources, from which this method may successfully evaluate the pointing parameters.  While this paper focuses on the pointing reconstruction of the {\\it Planck} satellite, elements of the analysis presented here may be of use elsewhere. ",
        "introduction": "{\\it Planck} is a European Space Agency satellite designed to produce high-resolution temperature and polarisation maps of the CMB. It possesses detectors sensitive to a wide range of frequencies from 30 to 857 GHz, split between two instruments the HFI and LFI, the high and low frequency instruments, respectively. {\\it Planck} is scheduled for launch in 2007, when it will be inserted into a Lissajous orbit around the second Lagrange point of the Earth-Sun system. {\\it Planck} will spin about its axis once per minute, sweeping its detectors along an almost great circle, the line of sight being almost perpendicular to the spin axis. The spin axis will nominally be repositioned every hour, and the roughly 60 or so circles corresponding to a single spin axis positioning may be binned together to form a ring. The spin axis passes through the centre of the solar panels and is directed towards the sun, thus maintaining the rest of the satellite in a cone of shadow produced by the solar panels.  The scanning strategy is determined by the frequency and the positions of the spin axis repointings over the course of the mission. The baseline scanning strategy is such that the spin axis remains in the ecliptic plane throughout the course of the mission and it maintains its sun-wards direction by hourly respositionings of 2.5\\arcmin. In reality this scanning strategy will not be employed due to the lack of coverage at the ecliptic poles; the actual scanning strategy, however, has yet to be finalised.  This paper will investigate some of the issues concerning the evaluation of parameters required for the pointing reconstruction of the {\\it Planck} satellite, specifically those resulting from non-zero off-diagonal elements in the inertia tensor as defined in the nominal satellite reference frame. The slight misalignment between the nominal spin axis of the satellite and the nearest principal axis of the inertia tensor defines the effective boresight angle and the rotation of the focal plane. These parameters may only be established using the science data, and they will vary over the course of the mission due to changes in the inertia tensor as consumables are depleted. It is therefore important to assess what the accuracy requirements on these parameters are, and whether it will be possible to extract the parameters to those required accuracies on a sufficiently regular basis.  The variation of the misalignment over the course of the mission is not yet well defined and is currently under investigation. It is, however, expected that the frequency at which the roll and effective boresight angles must be evaluated will be of the order of once a week \\cite{puget01}.  In Section~\\ref{section_acc_req} the accuracy requirements of the effective boresight angle and the roll angle, the rotation of the focal plane, are discussed. Methods of evaluating the roll and effective boresight angles through the use of galactic and extragalactic point sources are discussed in Section~\\ref{section_methods}. The number of both galactic and extragalactic point sources likely to be detectable by {\\it Planck} are discussed in Section~\\ref{pt_sources_section}. The simulated data on which our methods are tested is discussed in Section~\\ref{section_simData}, the results of which are discussed in Section~\\ref{section_results}.   \\section[]{Accuracy requirements} \\label{section_acc_req}  \\begin{figure} \\begin{center} \\setlength{\\unitlength}{1cm} \\begin{picture}(7,7)(0,0) \\put(8.5,0){\\special{psfile=bw_cl_acc.eps vscale=37 hscale=37 angle=90 voffset=0 hoffset=0}} \\end{picture} \\end{center} \\caption[]{The figure shows the unsubtractable errors on the $C_{\\ell}$s which are dominated by cosmic variance to the left and the finite beam size to the right. The solid grey curves enclose the region corresponding to the error on each individual multipole, whereas the dashed grey curves correspond to the error on multipole bins of width 50. The black curves correspond to the additional errors on the reconstruction of the $C_{\\ell}$s due to unknown random pointing errors. A beam of FWHM 5\\arcmin, with a noise of 10.3\\uk per beam and a fractional sky coverage of 0.7 have been assumed.} \\label{acc_fig} \\end{figure}  The effect of unknown random pointing reconstruction errors on the recovered ${\\rm C_{\\ell}}$ values may be seen in Figure~\\ref{acc_fig}. This figure has been replicated from one first shown by~\\cite{lamarre01}, using equation~\\ref{c_ell_eqn} modified from~\\cite{knox95}: \\begin{equation} \\label{c_ell_eqn} \\frac{\\Delta C_{\\ell}}{C_{\\ell}}=\\sqrt{ \\frac{2}{(2\\ell+1) \\times f_{sky}}} \\left( 1 + \\frac{\\omega}{C_{\\ell} W_{\\ell}} \\right) \\end{equation} where $\\omega$ is the noise per beam, $f_{sky}$ is the fractional sky coverage and $W_{\\ell}$ is the window function. The sensitivity of an experiment to the $\\ell$ multipoles is given by its window function. For a map made with a Gaussian beam, the window function is given by: $W_{\\ell}={\\rm e}^{-\\ell^2 \\sigma_b^2}$. If the beam sigma, $\\sigma_b$, is known, the decreased sensitivity to the higher $\\ell$ multipoles may be accounted for by multiplying by the window function. It is this which results in the exponential increase in the noise for the high multipoles in Figure~\\ref{acc_fig}, as the noise is undiminished by the beam size. For low multipoles the noise is dominated by the first term in equation~\\ref{c_ell_eqn} which expresses the sample and cosmic variance.  Figure~\\ref{acc_fig} uses the expected beam size and noise levels for the 217 GHz frequency channel and assumes 70 per cent sky coverage. The grey curves enclose regions corresponding to the unsubtractable noise, given by equation~\\ref{c_ell_eqn}. The solid grey curves correspond to the error on each individual multipole, while the dashed grey curves the error on multipole bins of width 50.  The black curves, in Figure~\\ref{acc_fig}, correspond to the additional errors in the reconstruction of the  $C_{\\ell}$s due to the effects of unknown random pointing errors. This results in a decreased sensitivity to the higher $\\ell$ multipoles as is seen in Figure~\\ref{acc_fig}. The pointing uncertainty results in an effective smearing of the beam and hence an increase in the effective beam width. If this is well known, a new window function may be evaluated and the effects of the pointing errors will not be as dramatic as those shown.  However, in reality the pointing errors will be correlated, smearing the beam more in one direction, which will lead to asymmetric effective beams. The effect of smearing the beam predominately in one direction is to increase the sensitivity to higher $\\ell$ multipoles as compared to the uniform smearing shown inFigure~\\ref{acc_fig}. In other words, for a given magnitude, correlated pointing errors will lead to a smaller reduction in sensitivity at high $\\ell$, than random pointing errors. However, given that the extent of the correlation of the pointing errors is not known and that the reduction in the maximum allowable pointing error is only significant for highly correlated pointing errors, the pointing errors are assummed to be random in the subsequent analysis.  Ideally, the pointing accuracy should be such that the additional errors introduced are less than the unsubtractable noise.  For each frequency channel, with the exception of 857 GHz in which the $C_{\\ell}$s are not observable, the unsubtractable errors on the $C_{\\ell}$s, given the expected beam size and noise levels, were evaluated. Hence, the ideal pointing accuracy required for each frequency channel may be found. These are shown in Table~\\ref{pt_acc_table}, for the cases of 70 and 100 percent sky coverage.  \\begin{table}  \\caption{The requirement on the random pointing error versus frequency channel, given fractional sky coverage,  $f_{sky}$.} \\label{pt_acc_table}  \\begin{tabular}{|ccccc|} \\hline Frequency & FWHM & noise per & \\multicolumn{2}{c}{maximum random}  \\\\ Channel & &  beam & \\multicolumn{2}{c}{pointing error} \\\\ & & & $f_{sky}=0.7$ & $f_{sky}=1.0$ \\\\ (GHz) & (\\arcmin) & (\\uk) & (\\arcmin) & (\\arcmin) \\\\ \\hline 100 & 9.2 & 6.0 & 0.23 & 0.21\\\\ 143 & 7.1 & 6.5 & 0.18 & 0.16 \\\\ 217 & 5.0 & 10.3 & 0.16 & 0.15 \\\\ 353 & 5.0 & 40.9 & 0.29 & 0.26 \\\\ 545 & 5.0 & 46.3 & 0.31 & 0.28 \\\\ \\hline \\end{tabular} \\end{table}  The uncertainty in the roll angle should be such that the corresponding uncertainty in the position of the detector is less than the requirement set by the analysis shown in Table~\\ref{pt_acc_table}. The required accuracy of the roll angle is determined by a combination of the maximum allowable pointing error and the location of the detectors in the focal plane. The further from the centre of the focal plane the more the detectors will be affected by the roll angle. The required accuracy of the roll angle for each frequency channel is displayed in Table~\\ref{roll_acc_table}, again for both the cases of 70 and 100 percent sky coverage.  \\begin{table}  \\caption{The requirement on the error in the roll angle, given fractional sky coverage, $f_{sky}$.} \\label{roll_acc_table}  \\begin{tabular}{|ccc|} \\hline Frequency Channel& \\multicolumn{2}{c}{Roll Angle Error} \\\\ & $f_{sky}=0.7$& $f_{sky}=1.0$ \\\\ (GHz) & (\\arcmin) & (\\arcmin)\\\\ \\hline 100 & 6.9 & 6.3\\\\ 143 & 4.9 & 4.4\\\\ 217 & 6.7 & 6.2\\\\ 353 & 8.1 & 7.2\\\\ 545 & 8.3 & 7.5\\\\ \\hline \\end{tabular} \\end{table}  As shown in Table~\\ref{roll_acc_table}, the 143 GHz channel sets the requirement on the uncertainty in the roll angle to be 4.4\\arcmin; assuming no other pointing errors from other sources.  \\section[]{Methods} \\label{section_methods}  In this section methods for evaluating the roll angle and the effective boresight, using measurements of point sources, will be considered. The effect of a roll angle will be to change the positions of the detectors with respect to the scan and cross-scan directions, so any method to evaluate the roll angle will depend on the focal plane layout. Figure~\\ref{focal_fig} illustrates the layout of the focal plane. The four highest frequency channels have offset detectors to ensure full sampling in the cross-scan direction. The centres of adjacent detectors are offset by 1.5\\arcmin\\, in the nominal cross-scan, with every detector being in a pair with the same nominal cross-scan location. The highest frequency channels will be the most suited to extracting the roll angle and effective boresight angle, due to the smaller beam size and the larger numbers of point sources which will be detected at the higher frequencies. For this paper, unless otherwise stated, we will consider only the 857 GHz detectors. We will also define a positive roll angle as a clockwise rotation, and a negative rotation as any anti-clockwise rotation about a fiducial reference point, FRP, in the centre of the focal plane.  In the nominal case of zero roll angle there are two pairs of detectors with the same cross-scan position. Any separation found between the cross-scan positions of these detectors may be used to infer the roll angle, as shown in Figure~\\ref{def_axes_fig}.  \\begin{figure} \\begin{center} \\setlength{\\unitlength}{1cm} \\begin{picture}(8,8)(0,0) \\put(9,0){\\special{psfile=focal.eps vscale=37 hscale=37 angle=90 voffset=0 hoffset=0}} \\end{picture} \\end{center} \\caption[]{This figure displays the focal plane layout of the HFI, in terms of the nominal scan and cross-scan directions. The positive scan direction indicates the direction in which the focal plane traverses the sky. A source will therefore cross the focal plane from right to left in this figure. The exact location of the detectors are given by the small filled circles. The top four frequency channels have offset detectors to ensure full sampling in the cross-scan direction. The centres of adjacent detectors are offset by 1.5\\arcmin\\, in the cross-scan direction, with every detector being in a pair with the same cross-scan location. The roll angle is defined as a rotation about the FRP, which is shown by the cross, and the effective boresight is defined as the angle between the spin axis and the centre of the focal plane. In this paper we will define a positive roll angle as a clockwise rotation, and a negative rotation as any anti-clockwise rotation about the FRP.} \\label{focal_fig} \\end{figure}  In order to evaluate the roll angle a method is needed to find the cross-scan positions of the detectors, or the cross-scan separation of the detectors aligned in the nominal scan direction, directly from the amplitudes of a point source as observed by each detector. Once a method has been found it may be investigated as to what signal-to-noise ratio, SNR, is required to obtain the roll angle to the required accuracy and hence an estimate of the number of available point sources may be made. Methods for reconstructing the roll angle are discussed in Section~\\ref{subsec_roll_angle}.  \\begin{figure} \\begin{center} \\setlength{\\unitlength}{1cm} \\begin{picture}(11,8)(0,0) \\put(-1.8,11){\\special{psfile=fplane_rot.ps vscale=57 hscale=57 angle=-90 voffset=0 hoffset=0}} \\end{picture} \\end{center} \\caption[]{This figure shows the effect of a change in the orientation of the focal plane with respect to the scan direction. The dotted grey circles correspond to detector positions in the nominal orientation of the focal plane, whereas the solid black circles correspond to their positions after a rotation, $\\rho$, about the FRP. For the highest frequency channels the centres of alternate detectors are aligned in the nominal scan direction as shown in the figure by the dotted grey line. Any difference in the cross-scan position of these detectors is therefore indicative of a non-zero focal plane rotation. The nominal scan separation, ${\\rm NS_{sep}}$, of alternate detectors may be used with their cross-scan separation, ${\\rm C_{sep}}$, to calculate the roll angle.} \\label{def_axes_fig} \\end{figure}  The effective boresight is the angle between the spin axis and the scan circle which passes through the FRP, shown as $\\alpha$ in Figure~\\ref{sphere_fig}. In order to evaluate the effective boresight the cross scan position of a known point source relative to the FRP must be found. Once the roll angle has been evaluated, and hence the cross scan positions of the detectors are known, the evaluation of the cross scan positions of point sources becomes possible. Methods for evaluating the cross scan positions of point sources will be discussed in Section~\\ref{section_eff_boresight}.   These methods are applied to data from one or two pointings of the spin axis, as required. The data from each spin axis pointing may be used in terms of the individual revolutions, circles, or the accumulated revolutions corresponding to the spin axis pointing, a ring. In the text below a scan circle may be thought of as the path of the line of sight of the FRP on the sky in the case of a circle or the average path in the case of a ring. This differentiation is necessary as in reality the spin axis describes a small rotating ellipse on the sky. This nutation of the spin axis results in deviations from the expected path of the line of sight given a spin axis rotating around a fixed point. These deviations are small and will average out in the ring data as the period of the nutation is greater than the rotation rate of the satellite. Providing that the amplitude of the nutation remains small it may be safely ignored in all cases, should this no longer be the case the effects of the nutation may be corrected for using information on the nutation parameters as found from the Star Tracker data. Throughout this paper it is assumed that these methods will be applied to rings, given their lower noise levels and their greater robustness to the effects of nutation.  \\subsection{Reconstructing the Roll Angle} \\label{subsec_roll_angle}  \\subsubsection{Using a single ring} \\label{subsec_one_ring}  In the first instance one may consider evaluating the roll angle using a single ring, using the observed amplitudes from the detectors to solve for the amplitude of the point source and its cross-scan position. Here a ring consists of the data corresponding to a single repointing of the spin axis, and the observed amplitude and location in the scan direction of the point source is found for each detector from either binning or by phase-ordering~\\citep{leeuwen01} this data. Once the amplitude and cross-scan position of the point source has been evaluated the amplitude observed by each detector may be used to evaluate its cross-scan separation from the point source position. Given these relative cross-scan locations the roll angle may be evaluated. However, in practice the errors in the detected amplitudes result in errors in the cross-scan locations which may lead to a large uncertainty in the recovered roll angle. An example of this is shown in Figure~\\ref{degen_fig}, in which the errors in the detected amplitude have been converted into errors in the cross-scan position of the detector. In this figure it has been assumed that the amplitude and cross-scan position of the point source have been successfully recovered with negligible error; in this example the amplitude of the point source is 100 times the 1$\\sigma$ noise and its cross-scan position from the scan circle which contains the FRP is 0\\fdg552. For comparison two scenarios are plotted; the circles represent the detected amplitudes in the four detectors resulting from a 0\\fdg4 roll angle, while the triangles correspond to the amplitudes resulting from a 0\\fdg7 roll angle. The 1$\\sigma$ errors in the amplitudes have been converted into errors in the cross-scan position of the detectors. Working across the figure pairing up the adjacent triangles and circles which correspond to the same detector, it is possible to see that uncertainties in the cross-scan positions, for each of the roll angles, overlap for all of the detectors. Hence in this example it is not possible to successfully distinguish between roll angles of 0\\fdg4 to 0\\fdg7. Other cross-scan positions of the point source with respect to the FRP result in similar difficulties over different ranges of roll angle.  \\begin{figure} \\begin{center} \\setlength{\\unitlength}{1cm} \\begin{picture}(8,8)(0,0) \\put(9,0){\\special{psfile=degen.eps vscale=37 hscale=37 angle=90 voffset=0 hoffset=0}} \\end{picture} \\end{center} \\caption[]{This figure displays the amplitudes, as observed by each detector, for two different roll angles, given a point source with an amplitude of 100 times the 1$\\sigma$ noise and  cross-scan position from the FRP of 0\\fdg552. The circles represent the detected amplitudes given a roll angle of 0\\fdg4, while the triangles 0\\fdg7. The 1$\\sigma$ errors in the amplitudes have been converted into errors in the cross-scan position of the detectors. The corresponding 1$\\sigma$ uncertainties in the positions of the detectors overlap, it is therefore not possible to establish the true roll angle without a higher signal-to-noise point source.} \\label{degen_fig} \\end{figure}  In reality there will be errors in the recovered amplitude and cross-scan position of the point source, together with errors in the relative calibration between detectors which will increase the effective errors in the cross-scan positions of the detectors and hence the uncertainty in the recovered value of the roll angle. Indeed, it is not possible to evaluate the cross-scan position of the point source, without first knowing the value of the roll angle due to the lack of knowledge on the positions of the detectors.  \\subsubsection[]{Using Two Rings} \\label{two_rings}  \\begin{figure} \\begin{center} \\setlength{\\unitlength}{1cm} \\begin{picture}(10,10)(0,0) \\put(0,0){\\special{psfile=new_roll_method.eps vscale=63 hscale=63 angle=0 voffset=0 hoffset=0}} \\end{picture} \\end{center} \\caption[]{This figure shows two neighbouring, although not necessarily adjacent scans, which pass on either side of a bright point source. The detector positions relative to the scan are denoted by $Y_{j}$, while the $X_{ij}$ represents the offset from the point source of detector j in scan i. $C_{i}$ is the cross-scan position of the point source from scan i and $S_{12}$ is the distance between the two scans at the location (longitude) of the point source.} \\label{roll_fig} \\end{figure}  If there are no abrupt changes in the inertia tensor during the mission, such as fuel shifts, the roll and effective boresight angles will only change slowly over the duration of the mission. Assuming that this is indeed the case and hence the roll angle is a slowly varying parameter, its evaluation may be improved by including the additional information available from the detection of the same point source on a neighbouring ring. Figure~\\ref{roll_fig} shows schematically the relative locations of the detectors to the scan directions, corresponding to the two rings, and the point source position. The cross-scan position of the point source from scan 1 and scan 2 is given by  $C_1$ and $C_2$, respectively; $C_1$ has a negative sign so that the sign convention for the cross-scan position may be the same for both scans. As the roll angle is unchanged, the cross-scan positions of the detectors, $Y_{j}$, from the FRP, will therefore not have changed between the two neighbouring scans. The sum of the offset from the point source position of a detector in scan1, $x_{1j}$, and the same detector in scan2, $x_{2j}$ will therefore be equal to the sum of the angles $C_1$ and $C_2$. The angular separation between the two scans at the location of the point source, $S_{12}$, may also be related to the sum of the angles $C_1$ and $C_2$, however in this case this relationship is only an approximate. Though as will be shown below for neighbouring scans, it may be thought of as being exact: \\begin{equation} \\label{setup1} S_{12}\\approx-C_1+C_2. \\end{equation} The angle between the two scans, $S_{12}$, at the location of the point source depends on the separation of the mean spin axis positions corresponding to each of the scans, and the elevation, $\\psi$, of the point source from the plane that connects the mean spin axis positions.  Figure~\\ref{sphere_fig} shows a schematic of the two scans and the point source in question. The separation between the mean spin axis positions is given by $\\phi$, this is also the maximum separation between the two scans. The angles, $\\psi$ and $\\phi$, may be evaluated using additional information provided by the Star Trackers. In the case of $\\phi$, it may be evaluated from the two recovered mean spin axis positions. The evaluation of $\\psi$ also requires the time of first transit of the FRP though the reference circle which connects the mean spin axis position and the North Ecliptic Pole, NEP, as defined in \\cite{leeuwen01}.  The angle $S_{12}$, also has a dependence on the effective boresight, $\\alpha$, though given the expected value, $\\alpha \\approx $ 85\\dg, this dependence is weak: \\begin{eqnarray} \\label{s12_eqn} \\cos(S_{12}) & = & \\cos(\\phi) \\left(\\cos^{2}(\\alpha)+\\sin^{2}(\\alpha)\\cos^{2}(\\psi)\\right) \\nonumber \\\\ & & \\mbox{} + \\sin^{2}(\\alpha)\\sin^{2}(\\psi). \\end{eqnarray} As can be seen from Figure~\\ref{sphere_fig} using the value of $S_{12}$ for the sum of the two cross-scan angles to the point source, $-C_1+C_2$, will result in an underestimation of this value. This underestimation will be the greatest when the point source lies midway between the two scans. To quantify the magnitude of this underestimation, the angle between the two cross-scan directions, $\\beta$, given a point source which lies midway between the two scans may be found using: \\begin{eqnarray} \\label{beta_eqn} \\cos(\\beta) & = & \\cos(\\pi-\\phi) \\left(\\cos^{2}(\\psi)\\cos^{2}(\\alpha)+\\sin^{2}(\\alpha)\\right) \\nonumber \\\\ & & \\mbox{} - \\cos^{2}(\\alpha)\\sin^{2}(\\psi). \\end{eqnarray} Given that we are only interested in neighbouring scans the angle, $\\phi$ will be small and the angle, $\\delta$, between $ C_{1,2}$ and $ S_{12}$ may then be found using: \\begin{equation} \\label{delta_eqn} \\delta=\\frac{\\pi-\\beta}{2}. \\end{equation}\tFor $\\phi \\le$ 1\\dg, which corresponds to  $\\le$ 24 nominal repointings of 2.5\\arcmin,  the underestimation in the sum of the cross-scan angles is of the order of a thousandth of a percent. This corresponds to an overestimation for the magnitudes of the roll angle considered here of the order of a hundredth of an arc-second. Hence, for our present purposes the relation in equation~\\ref{setup1} may be considered to be exact.  \\begin{figure} \\begin{center} \\setlength{\\unitlength}{1cm} \\begin{picture}(10,16)(0,0) \\put(-1.5,18){\\special{psfile=sphere_fig1.ps vscale=55 hscale=55 angle=270 voffset=0 hoffset=0}} \\put(-1.5,10){\\special{psfile=sphere_fig2.ps vscale=55 hscale=55 angle=270 voffset=0 hoffset=0}} \\end{picture} \\end{center} \\caption[]{This figure shows the angles discussed in the text relating to two neighbouring scan circles with a point source located between them. The top diagram shows the plane in which the motion of the mean spin axis has occurred, and the angle between this plane and the point source, $\\psi$. The angle between the spin axis and the scan circle, $\\alpha$, is the effective boresight. The angle $\\psi$, may be evaluated once the reference phase, or time of first transit of the FRP is known, as described in \\cite{leeuwen01}. In the first instance this may be evaluated by the Star Trackers. The lower diagram shows the angle, $\\phi$, through which the spin axis has been repositioned. The angles $C_1$ and $C_2$ represent the cross-scan angular distance between the point source and scan circles 1 and 2, respectively. The angle $S_{12}$ is the angular separation between the intersection points of $C_1$ and $C_2$ on scan circles 1 and 2. The angle $\\phi$, between the two spin axis positions, has been exaggerated for clarity in this figure; in the analyis presented here $\\phi$ never exceeds 1\\dg. } \\label{sphere_fig} \\end{figure}  From Figure~\\ref{roll_fig} it is seen that the sum of the cross scan angles $C_1$ and $C_2$ is equal to the sum of the offsets, from the point source position, of the ${\\rm j^{th}}$ detector in scans 1 and 2, $x_{1j}$ and $x_{2j}$. The angle $S_{12}$ is therefore also equal to the sum of these offsets, via equation~\\ref{setup1}: \\begin{equation} \\label{setup2} S_{12}=x_{1j}+x_{2j}. \\end{equation}\tThe cross-scan position of the point source from the scan and the offsets of the detectors from the point source, may be related to the position of the detectors relative to the scan: \\begin{equation} \\label{setup3} -C_1=x_{1j}-Y_j\t, \\end{equation} \\begin{equation} \\label{setup4} C_2=x_{2j}+Y_j\t. \\end{equation}\tThe observations of the point source may be described by the following, the amplitude of the point source as seen by the ${\\rm j^{th}}$ detector in the ${\\rm i^{th}}$ scan is denoted by  $A_{ij}$: \\begin{equation} A_{ij}=A \\exp \\left( \\frac{-x_{ij}^2}{2 \\sigma_{j}^2} \\right). \\end{equation} Where $\\sigma_{j}$ is the $1\\sigma$ width of the assumed Gaussian beam for the ${\\rm j^{th}}$ detector.\\\\  The expression for the ratio of the amplitudes of the point source as observed in the two scans by the ${\\rm j^{th}}$ detector may be rearranged to find: \\begin{equation} x_{2j}^2-x_{1j}^2=2\\sigma_{j}^2 \\ln \\left( \\frac{A_{1j}}{A_{2j}} \\right), \\end{equation} from which, using equation~\\ref{setup2}, we find: \\begin{equation} \\label{shows_adv_eqn} x_{2j}-x_{1j}=\\frac{2\\sigma_{j}^2}{S_{12}} \\ln \\left( \\frac{A_{1j}}{A_{2j}} \\right), \\end{equation} where everything on the RHS of this equation is known observationally. Equation~\\ref{shows_adv_eqn} shows a major benefit of this method; as it is based on the ratio of the amplitudes seen by the same detector in different scans, uncertainties in the relative calibrations between detectors will be irrelevant.\\\\  An alternative expression for the difference between the offsets of the ${\\rm j^{th}}$ detector from the point source on each scan, may be found using equations~\\ref{setup3} and~\\ref{setup4}: \\begin{equation} \\label{observables} Y_j= \\frac{1}{2} \\left((C_2+C_1)- (x_{2j}-x_{1j}) \\right). \\end{equation}  The cross-scan separation, ${\\rm C_{sep}}$, of the aligned detectors may be found: \\begin{equation} Y_3 -Y_1 = \\frac{1}{2} \\left( (x_{21}-x_{11})-(x_{23}-x_{13}) \\right) , \\end{equation}  \\begin{equation} Y_4 -Y_2 = \\frac{1}{2} \\left( (x_{22}-x_{12})-(x_{24}-x_{14}) \\right) , \\end{equation}  \\begin{equation} {\\rm C_{sep}}=Y_3 -Y_1 = Y_4 -Y_2 , \\end{equation} and hence the roll angle, {$\\rm \\rho$}, may be evaluated, where the denominator, ${\\rm NS_{sep}}$, is determined by the separation in the scan direction of the aligned detectors in the case of zero roll angle, and is shown in Figure~\\ref{def_axes_fig}:  \\begin{equation} {\\rm \\sin(\\rho)} =\\frac{\\rm C_{sep}}{\\rm NS_{sep}}. \\end{equation}   \\subsection{Evaluating the Effective Boresight} \\label{section_eff_boresight}  The effective boresight, which is the angle between the mean spin axis position and the scan circle, may be found if a point source of known position is observed in the data from the scan circle. The cross scan position, C, of the point source is found with respect to the FRP, and hence the scan circle. Here the cross scan position of the point source is defined as positive if the point source lies on the opposite side of the scan circle to the spin axis position. The effective boresight, $\\alpha$, may then be found using: \\begin{equation} \\label{eff_B_eqn2} \\alpha = \\omega - C, \\end{equation}\twhere $\\omega$ is the angle between the point source position,  $(\\beta_{pt}, \\lambda_{pt})$, and the mean spin axis position, $(\\beta_{sa}, \\lambda_{sa})$, and is given by: \\begin{equation} \\label{eff_B_eqn} \\cos \\omega = \\cos \\beta_{pt}  \\cos \\beta_{sa} \\cos(\\lambda_{pt} - \\lambda_{sa}) + \\sin \\beta_{pt} \\sin \\beta_{sa}. \\end{equation}  The cross scan position of a point source may be found using data from a single ring or by an extension of the two ring method for the evaluation of the roll angle.  In the case of data from a single ring, the amplitudes found from the transits of the point source are fit by a Gaussian beam profile to establish best fit values for the position and amplitude of the peak. This method will only provide unbiased cross scan position for the point source if it lies between the range of cross scan positions of the detectors. Hence the value of the roll angle is required prior to the evaluation of the cross-scan position of the point source. Additionally, uncertainties in the relative calibration and/or beams of the detectors will increase the uncertainty in the recovered value of the cross-scan position and hence effective boresight angle.   Alternatively, by an extension to the method to extract the roll angle from two rings, the cross-scan position of the point source may be found using the ratio of the amplitudes observed by each detector in two neighbouring scans.  The sum over all detectors using equation~\\ref{observables}, may be used to evaluate the difference between the cross-scan positions of the point source to each scan: \\begin{equation} \\label{cross_dif_eqn} (C_2+C_1) = \\frac{1}{N} \\sum^N_j \\left( x_{2j}-x_{1j} \\right) + \\frac{2}{N} \\sum^N_j Y_j , \\end{equation}\twhere N=4 using a single frequency channel, or N=8 combining both the 857 and 545 GHz channels.  Once the roll angle is found the positions of the detectors, $Y_j$, are known, indeed the sum over the positions of the detectors is approximately a constant for small values of the roll angle. The cross scan position of the point source may be found for each scan by combining equations~\\ref{setup1} and~\\ref{cross_dif_eqn}: \\begin{equation} \\label{cross1_eqn} C_1 = - \\frac{S_{12}}{2}+ \\frac{1}{N} \\sum^N_j \\left( \\sigma_{j}^2 \\ln \\left( \\frac{A_{1j}}{A_{2j}} \\right) \\right) +  \\frac{1}{N} \\sum^N_j Y_j  , \\end{equation} \\begin{equation} \\label{cross2_eqn} C_2 = \\frac{S_{12}}{2}+ \\frac{1}{N} \\sum^N_j \\left( \\sigma_{j}^2 \\ln \\left( \\frac{A_{1j}}{A_{2j}} \\right) \\right) +  \\frac{1}{N} \\sum^N_j Y_j . \\end{equation}  Once the cross scan position of the point source is known with respect to the FRP, the effective boresight may be evaluated via equations~\\ref{eff_B_eqn2} and~\\ref{eff_B_eqn}, if the position of the point source is known. ",
        "conclusions": "\\label{section_conclusions}  We have investigated methods for the evaluation of parameters required for the pointing reconstruction of the {\\it Planck} satellite. The effect of pointing errors on the recovered $C_{\\ell}$s has been assessed and requirements on the accuracy of the pointing reconstruction found, such that the effect of random pointing errors is less than the unsubtractable noise. The accuracy at which the roll angle is required is determined so that the uncertainty in the positions of the detectors due to the error in the roll angle is less than the limit on the pointing error.   It is found that it is impractical to determine the roll angle from data from a single ring, as this would require an extremely high signal-to-noise ratio point source to successfully determine the roll angle. An alternative method using two neighbouring rings was developed, which allows the successful extraction of the roll angle to the required accuracy from a single point source, with a signal-to-noise ratio  $\\sim 50 $. This method also has the additional advantage of using the ratio of the amplitudes of a point source as seen by the same detector on different rings, hence this method does not require the relative calibration of the detectors. This method, however, does rely upon the assumption that the roll angle is a slowly varying parameter, and it should be remembered that any abrupt changes to the inertia tensor will invalidate this assumption. Once the roll angle is determined the cross scan position of any point source may be assessed. The effective boresight angle may then be evaluated using the cross scan position of a point source of known position.  Due to the depletion of consumables over the course of the mission and the resultant variation in the values of the pointing parameters, these parameters should be evaluated on a regular basis. Whether a method may be applied  successfully on a regular basis, will depend on the availability of suitable point sources over the course of the mission. In order to investigate this availability, the number of galactic and extragalactic point sources detectable by the {\\it Planck} mission was determined. It was found that there is an insufficient number of suitable extragalactic point sources to evaluate these parameters, the roll and effective boresight angles, on a weekly basis. Hence the number of galactic point sources available is crucial. The total number of galactic point sources observable by {\\it Planck} over the course of the mission, and more importantly the variation in the number observable on a weekly basis, depends upon the scanning strategy used. However, it was found that it was not possible to recommend a scanning strategy which will  maximise the number of galactic point sources available. This is due to other constraints on the scanning strategy such as maintaining the Earth aspect angle $\\le$15\\dg\\, within the L2 orbit.  The number of point sources available is also found to depend on the value of the roll angle. The method, presented here for the reconstruction of the roll angle from two neighbouring rings, is best suited for the evaluation of small magnitude roll angles. The larger the roll angle, the smaller the number of point sources available and the greater the chance of weeks with no suitable point sources.  For roll angles $\\sim +2$\\dg\\, and $\\sim -1$\\dg\\, it is likely that there will be an occasional week in which there is no suitable point source to evaluate the roll angle. As the magnitude of the roll angle increases above these values the number of suitable point sources diminishes until  $\\sim+4$\\dg\\, and $\\sim-3$\\dg\\, where the evaluation of the roll angle using the method presented here is no longer possible. This places a manufacturing requirement on the alignment of the focal plane with respect to the nominal scan direction.  While the evaluation of the roll angle does not require a known point source, the reconstruction of the effective boresight requires the position of the point source to be known. A known point source catalogue will therefore be required for the evaluation of the effective boresight angles, with positional errors less than 7\\arcsec\\, and ideally less than 2\\arcsec, in order to meet the accuracy requirements on the pointing parameters.  Given the success of this method depends on the numbers of bright point sources available, and for the reconstruction of the effective boresight on their accurate positions, it is important that the IRAS sources which are found here to be potentially suitable for the reconstruction of these parameters are observed prior to launch, to confirm their point source nature, to improve upon their flux estimates and acquire accurate positions. Any point source $\\geq 50$ times the noise expected in a binned ring in the 857 GHz channel is potentially useful in evaluating these parameters. Given the expected noise level this corresponds to point sources $\\geq 14.2$~Jy at 857 GHz.  The method presented here will allow the weekly monitoring of the roll angle, given the number of point sources determined here, under the expectation of small roll angles and a slowly varying inertia tensor. When combined with a catalogue of accurate point source positions, this method will also allow the weekly determination of the effective boresight angle. It is expected that the determinations of these two angles from the science data will be used to determine their systematic offsets from the equivalent angles determined more accurately by the Star Tracker."
    },
    "0311/astro-ph0311525.txt": {
        "abstract": "{ We report the angular diameter measurements of seven classical Cepheids, X\\,Sgr, $\\eta$\\,Aql, W\\,Sgr, $\\zeta$\\,Gem, $\\beta$\\,Dor, Y\\,Oph and $\\ell$\\,Car that we have obtained with the VINCI instrument, installed at ESO's VLT Interferometer (VLTI). We also present reprocessed archive data obtained with the FLUOR/IOTA instrument on $\\zeta$\\,Gem, in order to improve the phase coverage of our observations. We obtain average limb darkened angular diameter values of $\\overline{\\theta_{\\rm LD}}[{\\rm X\\,Sgr}] = 1.471 \\pm 0.033$\\,mas, $\\overline{\\theta_{\\rm LD}}[\\eta\\,{\\rm Aql}] = 1.839 \\pm 0.028$\\,mas, $\\overline{\\theta_{\\rm LD}}[{\\rm W\\,Sgr}] = 1.312 \\pm 0.029$\\,mas, $\\overline{\\theta_{\\rm LD}}[\\beta\\,{\\rm Dor}] = 1.891 \\pm 0.024$\\,mas, $\\overline{\\theta_{\\rm LD}}[\\zeta\\,{\\rm Gem}] =1.747 \\pm 0.061$\\,mas, $\\overline{\\theta_{\\rm LD}}[{\\rm Y\\,Oph}] = 1.437 \\pm 0.040$\\,mas, and $\\overline{\\theta_{\\rm LD}}[\\ell\\,{\\rm Car}] = 2.988 \\pm 0.012$\\,mas. For four of these stars, $\\eta$\\,Aql, W\\,Sgr, $\\beta$\\,Dor, and $\\ell$\\,Car, we detect the pulsational variation of their angular diameter. This enables us to compute directly their distances, using a modified version of the Baade-Wesselink method: $d[\\eta\\,{\\rm Aql}] = 276^{+55}_{-38}$\\,pc, $d[{\\rm W\\,Sgr}] = 379^{+216}_{-130}$\\,pc, $d[\\beta {\\rm Dor}] = 345^{+175}_{-80}$\\,pc, $d[\\ell\\,{\\rm Car}] = 603^{+24}_{-19}$\\,pc. The stated error bars are statistical in nature. Applying a hybrid method, that makes use of the Gieren et al.~(\\cite{gieren98}) Period-Radius relation to estimate the linear diameters, we obtain the following distances (statistical and systematic error bars are mentioned): $d[{\\rm X\\,Sgr}] = 324 \\pm 7 \\pm 17$\\,pc, $d[\\eta\\,{\\rm Aql}] = 264 \\pm 4 \\pm 14$\\,pc, $d[{\\rm W\\,Sgr}] = 386 \\pm 9 \\pm 21$\\,pc, $d[\\beta {\\rm Dor}] = 326 \\pm 4 \\pm 19$\\,pc, $d[\\zeta\\,{\\rm Gem}] = 360 \\pm 13 \\pm 22$\\,pc, $d[{\\rm Y\\,Oph}] = 648 \\pm 17 \\pm 47$\\,pc, $d[\\ell\\,{\\rm Car}] = 542 \\pm 2 \\pm 49$\\,pc. ",
        "introduction": "% For almost a century, Cepheids have occupied a central role in distance determinations. This is thanks to the existence of the Period--Luminosity (P--L) relation, $M = a\\,\\log P + b$, which relates the logarithm of the variability period of a Cepheid to its absolute mean magnitude. These stars became even more important since the {\\it Hubble Space Telescope} Key Project on the extragalactic distance scale (Freedman et al.\\,\\cite{freedman01}) has totally relied on Cepheids for the calibration of distance indicators to reach cosmologically significant distances. In other words, if the calibration of the Cepheid P--L relation is wrong, the whole extragalactic distance scale is wrong.  There are various ways to calibrate the P--L relation. The avenue chosen by the $HST$ Key-Project was to \\emph{assume} a distance to the Large Magellanic Cloud (LMC), thereby adopting a zero point of the distance scale. Freedman et al.\\,(\\cite{freedman01}) present an extensive discussion of all available LMC distances, and note, with other authors (see e.g. Benedict et al.\\,\\cite{benedict02}), that the LMC distance is currently the weak link in the extragalactic distance scale ladder. Another avenue is to determine the zero point of the P--L relation with Galactic Cepheids, using for instance parallax measurements, Cepheids in clusters, or through the Baade-Wesselink (BW) method. We propose in this paper and its sequels (Papers II and III) to improve the calibration of the Cepheid P--R, P--L and surface brightness--color relations through a combination of spectroscopic and interferometric observations of bright Galactic Cepheids.  In the particular case of the P--L relation, the slope $a$ is well known from Magellanic Cloud Cepheids (e.g. Udalski et al.\\,\\cite{udalski99}), though Lanoix et al.\\,(\\cite{lanoix99}) have suggested that a Malmquist effect (population incompleteness) could bias this value. On the other hand, the calibration of the zero-point $b$ (the hypothetic absolute magnitude of a 1-day period Cepheid) requires measurement of the distance to a number of nearby Cepheids with high precision. For this purpose, interferometry enables a new version of the Baade-Wesselink method, (BW, Baade\\,\\cite{baade26}, Wesselink\\,\\cite{wesselink46}) for which we do not need to measure the star's temperature, as we have directly access to its angular diameter (Davis\\,\\cite{davis79}; Sasselov \\& Karovska\\,\\cite{sasselov94}). Using this method, we derive directly the distances to the four nearby Cepheids $\\eta$\\,Aql, W\\,Sgr, $\\beta$\\,Dor and $\\ell$\\,Car. For the remaining three objects of our sample, X\\,Sgr, $\\zeta$\\,Gem and Y\\,Oph, we apply a hybrid method to derive their distances, based on published values of their linear diameters.  After a short description of the VINCI/VLTI instrument (Sect.\\,\\ref{inst_setup}), we describe the sample Cepheids that we selected (Sect.\\,\\ref{sample_cephs}). In Sect.~\\ref{data_proc} and \\ref{data_quality}, we report our new observations as well as reprocessed measurements of $\\zeta$\\,Gem retrieved from the FLUOR/IOTA instrument archive. Sect.~\\ref{angdiam} is dedicated to the computation of the corresponding angular diameter values, taking into account the limb darkening and the bandwidth smearing effects. In Sect.~\\ref{radius_curves_sect} and \\ref{puls_param}, we investigate the application of the BW method to our data, and we derive the Cepheid distances.  We will discuss the consequences of these results for the calibration of the Period-Radius (P--R), Period-Luminosity (P--L) and Barnes-Evans relations of the Cepheids in forthcoming papers (Paper II and III). ",
        "conclusions": "We have reported in this paper our long-baseline interferometric observations of seven classical Cepheids using the VINCI/VLTI instrument. For four stars ($\\eta$\\,Aql, W\\,Sgr, $\\beta$\\,Dor and $\\ell$\\,Car), we were able to apply a modified version of the BW method, resulting in an independent estimate of their distance. For all stars, we also derived their distances from lower order fitting methods, that use an {\\it a priori} estimate of their linear diameter from the P--R relation of Gieren et al.~(\\cite{gieren98}). We would like to emphasize that the order 0/1 and order 2 error bars are different in nature, and they should be treated differently in any further use of these results. While the order 2 error bars can be treated as statistical (i.e. reduced by averaging), the order 0/1 methods errors are dominated by the systematic uncertainty introduced by the {\\it a priori} estimation of the linear radius. The respective contributions of the statistical and systematic uncertainties are given separately in Tables~\\ref{results_cepheids0} and \\ref{results_cepheids1}. These values assume a constant value of the $p$-factor of 1.36, and can be scaled linearly for other values.  We will use these distances in Paper II, together with previously published measurements, to calibrate the zero points of the Period-Radius and Period-Luminosity relations %to a relative precision better than 2 and 5\\,\\%, respectively, and we find a %satisfactory agreement with previously published work. In Paper III, we will calibrate the surface brightness--color relation, with a particular emphasis on the evolution of $\\ell$\\,Car in this diagram over its pulsation. These three empirical relations are of critical importance for the extragalactic distance scale.  The direct measurement of the limb darkening of nearby Cepheids by interferometry is the next step of the interferometric study of these stars. It will allow a refined modeling of the atmosphere of these stars. This observation will be achieved soon using in particular the long baselines of the VLTI equipped with the AMBER instrument, and the CHARA array for the northern Cepheids. Another improvement of the interferometric BW methow will come from radial velocity measurements in the near infrared (see e.g. Butler \\& Bell~\\cite{butler97}). They will avoid any differential limb darkening between the interferometric and radial velocity measurements, and therefore make the resulting distances more immune to limb darkening uncertainties. % %The measured distances will allow us in a forthcoming paper to establish a new calibration %of the Cepheids P--L relation, as well as the P--R and surface brightness--color relations.  %___________________Table of fit results %\\begin{table} %\\caption{Summary of the best fit LD angular diameters and distances %of the seven Cepheids observed with VINCI. The detailed results for each %star and fitting order are presented in Tables \\ref{results_cepheids0} %to \\ref{results_cepheids2}. } %\\label{results_cepheids_summary} %\\begin{tabular}{lccll} %\\hline %\\noalign{\\smallskip} %Star & Order & $\\overline{\\theta_{\\rm LD}}$\\,(mas) & $d$\\,(pc) \\\\ %\\noalign{\\smallskip} %\\hline %\\noalign{\\smallskip} %X~Sgr & 0 & $1.471 \\pm 0.033$ & $324 \\pm 18$ \\\\ %\\noalign{\\smallskip} %$\\eta$~Aql & 2 & $1.839 \\pm 0.028$ & $276^{+55}_{-38}$ \\\\ %\\noalign{\\smallskip} %W~Sgr & 2 & $1.312 \\pm 0.029$ & $379^{+216}_{-130}$\\\\ %\\noalign{\\smallskip} %$\\beta$~Dor & 2 & $1.891 \\pm 0.024$ & $345^{+175}_{-80}$ \\\\ %\\noalign{\\smallskip} %$\\zeta$~Gem & 0 & $1.747 \\pm 0.061$ & $360 \\pm 25$\\\\ %\\noalign{\\smallskip} %Y~Oph & 1 & $1.438 \\pm 0.051$ & $648 \\pm 52$ \\\\ %\\noalign{\\smallskip} %$\\ell$~Car & 2 & $2.988 \\pm 0.012$ & $603^{+24}_{-19}$ \\\\ %\\noalign{\\smallskip} %\\hline %\\end{tabular} %\\end{table}"
    },
    "0311/astro-ph0311260_arXiv.txt": {
        "abstract": "We present a model for the distribution of void sizes and its evolution in the context of hierarchical scenarios of gravitational structure formation. We find that at any cosmic epoch the voids have a size distribution which is well-peaked about a characteristic void size which evolves self-similarly in time.  This is in distinct contrast to the distribution of virialized halo masses which does not have a small-scale cut-off.  In our model, the fate of voids is ruled by two processes. The first process affects those voids which are embedded in larger underdense regions:  the evolution is effectively one in which a larger void is made up by the mergers of smaller voids, and is analogous to how massive clusters form from the mergers of less massive progenitors. The second process is unique to voids, and occurs to voids which happen to be embedded within a larger scale overdensity:  these voids get squeezed out of existence as the overdensity collapses around them.  It is this second process which produces the cut-off at small scales.  In the excursion set formulation of cluster abundance and evolution, solution of the {\\it cloud-in-cloud} problem, i.e., counting as clusters only those objects which are not embedded in larger clusters, requires study of random walks crossing {\\it one-barrier}. We show that a similar formulation of void evolution requires study of a {\\it two-barrier} problem: one barrier is required to account for {\\it voids-in-voids}, and the other for {\\it voids-in-clouds}. Thus, in our model, the void size distribution is a function of two parameters, one of which reflects the dynamics of void formation, and the other the formation of collapsed objects. ",
        "introduction": "An overwhelming body of observational and theoretical evidence favours the view that structure in the Universe has risen out of a nearly homogeneous and featureless primordial cosmos through the process of gravitational instability. Almost all viable existing theories for structure formation within the context of this framework are hierarchical: the matter distribution evolves through a sequence of ever larger structures.  Hierarchical scenarios of structure formation have been successful in explaining the formation histories of gravitationally bound virialized haloes. They provide a basic framework within which more intricate aspects of the formation of a wide range of cosmic objects, ranging from galaxies to rich clusters, may be investigated. In particular, a fully analytical description of the collapse and virialization of overdense dark matter halos has been developed. The approach, originally proposed by Press \\& Schechter~\\cite{ps74}, and later modified by Epstein~\\cite{re83} and Bond et al.~\\cite{bcek91}, has led to simple and accurate models for the abundance of massive haloes which results from hierarchical gravitational clustering. This framework has come to be called the {\\it excursion set approach}.  The excursion set approach provides a useful framework for thinking about the formation histories of gravitationally bound virialized haloes in scenarios of hierarchical structure formation. It provides analytic approximations for the distribution of halo masses, merger rates, and formation times which are quite accurate (Lacey \\& Cole~1993), and can be extended to provide estimates of the distribution of the mass in randomly placed cells (Sheth~1998). A key ingredient in the original approach, inherited from the pioneering work of Press \\& Schechter~\\cite{ps74}, is the assumption that virialized objects form from a smooth spherical collapse. In reality the collapse can be quite different from spherical; recent work has shown that ellipsoidal collapse can be incorporated into the approach, with reasonable improvements in accuracy (e.g. Sheth, Mo \\& Tormen~2001).  Models based on spherical evolution are difficult to reconcile with the spatial patterns which characterize the cosmic matter distribution.  The observed world of galaxy redshift surveys, and the artificial world of numerical simulations of cosmic structure formation, are both characterized by filamentary and sheetlike structures. Such weblike patterns represent distinctly non-virialized structures for which gravitational contraction of initially aspherical density peaks has only been accomplished along one or two dimensions. At first sight, such weblike configurations would seem to be beyond the realm of the idealized excursion-set description.  \\begin{figure*} \\centering \\vskip -2.0cm \\epsfxsize=16.2cm \\mbox{\\hskip 0.8truecm\\epsfbox{voidsf1.ps}} \\vskip -1.0cm \\caption{Spatial structure in a void-like region selected from an N-body simulation of structure formation in the SCDM scenario. Three consecutive zoom-ins centered on the core of the void are shown: a $45\\hbox{Mpc}$ diameter particle sphere (top), $36\\hbox{Mpc}$ (bottom left) and $30\\hbox{Mpc}$ (bottom right). The existence of substructure within the void region is readily apparent, although it becomes more faint and tenuous towards the increasingly depleted interior of the void.} \\label{voidcdmcube} \\end{figure*}  Nevertheless, in this study we show that the formation and evolution of foamlike patterns can indeed be described by the excursion set analysis. This is accomplished by focusing on the evolution of {\\em underdense regions}, the {\\em voids}, rather than overdensities in the matter distribution. Whereas much of the mass in the universe is bound-up in virialized structures, most of the volume is occupied by large underdense voids: voids are {\\it the} dominant component of the Megaparsec-scale galaxy and matter distributions. In a void-based description of structure formation, matter is squeezed in between expanding voids, and sheets and filaments form at the intersections of the void walls (Icke~1984; Van de Weygaert~1991, 2002). Such a view is supported by Reg{\\H o}s \\& Geller~(1991), Dubinski et al.~(1993) and Van de Weygaert \\& Van Kampen~(1993), who give clear and lucid descriptions of how voids evolve in numerical simulations of gravitational clustering. We will stick to this basic framework in the present study.  We will argue that low density regions are the objects-of-choice for working out a succesful analytical description of cosmic spatial structure, if it is to be based upon the idealization of spherical symmetry.  This is because, in many respects, voids are ideally suited for an excursion set analysis based on a spherical evolution model. This is despite the fact that voids form from negative density perturbations in the initial fluctuation field, and neither maxima nor minima in the primordial Gaussian field, are spherical (see Bardeen et al. 1986). However, in marked contrast to the evolution of density peaks, primordial asphericity of negative density perturbations is quickly lost as they expand:  the generic evolution is towards an approximately spherical tophat geometry (Icke~1984). Moreover, the velocity structure of uniform density voids is simple to understand; an observer in the interior will observe a Hubble-type velocity field. All of this is discussed in some detail in Appendix~\\ref{stophat}, which describes the evolution of a single isolated void.  Although the image of a large scale matter distribution organized by expanding voids is appealing, in its basic form, the description essentially involves an extrapolation of single void characteristics to an entire random population of strictly distinct and non-interacting peers, each of them undisturbed smoothly expanding bubbles. This discards one of the most crucial and characteristic aspects of cosmic structure formation---that there are no isolated voids, nor smoothly unstructured ones. Any complete analysis will have to take into account the complications which arise from \\begin{itemize} \\item{} the substructure present within the primordial volume occupied by the void, and \\item{} the inhomogeneous matter distribution in its vicinity. \\end{itemize} The existence of internal void structure is not unexpected. The void shown in Figure~\\ref{voidcdmcube}, selected from a large N-body simulation of cosmic structure formation, shows the existence of structure on all scales. The figure shows three successive zoom-ins on the inner parts of the void; all exhibit some measure of internal structure, although substructure is less pronounced in the emptiest inner regions.  As was mentioned above, all viable cosmological structure formation scenarios imply a hierarchical mode of structural growth. The formation of any object involves the fusion of all substructure present within its realm, including the small-scale objects which had condensed out at an earlier stage. Underdensities are organized similarly---in the evolution of a void we may identify two, intimately related, processes: \\begin{itemize} \\item{} a bottom-up assembly, in which a void emerges as a mature and well-defined entity through the fusion and gradual erasure of its internal substructure, and \\item{} the interaction of the void with its surroundings, marking its participation in the continuing process of hierarchical structure formation. \\end{itemize} Considerable insight into the evolution of voids came from the rigorous and insightful study by Dubinski et al.~(1993).  Following an analytical study of (isolated) spherically symmetric voids by Blumenthal et al.~(1992), they used N-body simulations to study the evolution of the void hierarchy from a set of artificial and simplified initial conditions, consisting of various levels of hierarchically embedded spherical tophat voids. They showed that adjacent voids collide, producing thin walls and filaments as the matter between them is squeezed. Mainly confined to tangential motions, the peculiar velocities perpendicular to the void walls are mostly suppressed. The subsequent merging of voids is marked by the gradual fading of these structures while matter evacuates along the walls and filaments towards the enclosing boundary of the ``void merger''. The timescale on which the internal substructure of a void is erased is approximately the same as that when the void itself approached ``nonlinearity'' (Appendix~\\ref{stophat} gives a precise definition of what is meant by nonlinearity). At nonlinearity, smaller-scale voids collide and merge with one another, effectively dissolving their separate entities into one larger encompassing void. Only a faint and gradually fading imprint of their original outline remains as a reminder of the initial internal substructure. As this (re)arrangement of structure progresses to ever larger scales, the same basic processes repeat.  N-body simulations of voids evolving in more generic cosmological circumstances by Van de Weygaert \\& Van Kampen~(1993) (also see Van de Weygaert~1991) yielded similar results. This prompted them to suggest the existence of a natural {\\it void hierarchy}, in which small-scale voids embedded within a pronounced large-scale void gradually fade away. An illustration of such a void hierarchy process, within the context of the CDM scenario, is shown in Figure~\\ref{voidcdm6slc}. The major characteristics of an evolving void hierarchy, the gradual blending of small-scale voids and structures into a larger surrounding underdensity, is clearly visible in the sequence of six timesteps.  However, the artificial arrangement of voids embedded within voids represents only one aspect of reality---it misses a crucial component of the development of a void hierarchy. An evolving void hierarchy not only involves the merging of small voids into larger voids, but also {\\em the disappearance of small voids as they become embedded in larger-scale overdensities.} Thus, in contrast to the process of dark halo formation, the emerging void hierarchy is ruled by two processes instead of one. The main goal of this paper is to incorporate both processes into a model of the void hierarchy.  We do this by combining the spherical evolution model with the excursion set approach.  When used to describe the evolution of overdense clouds, the excursion approach requires consideration of a {\\it one-barrier} problem, the single barrier representing what is required for collapse in the spherical evolution model.  We show that the excursion set formulation of the void hierarchy requires consideration of {\\it two-barriers}:  one barrier is associated with the collapse of clouds, and the other with the formation of voids. The resulting framework is able to describe realistic settings of random density fields in which voids interact with their surroundings.  \\begin{figure*} \\centering \\vskip -5.0cm \\epsfxsize=17.cm \\mbox{\\hskip -0.truecm\\epsfbox{voidsf2.ps}} \\vskip -1.0cm \\caption{Void evolution. Six timesteps in the evolution of a void region in a 128$^3$ particle N-body simulation of structure formation in an SCDM model: top-left to bottom-right shows expansion factors $a_{\\rm exp}=0.1, 0.2, 0.3, 0.35, 0.4$ and $0.5$ (the present time has $a_{\\rm exp}=1.0$). Initial conditions were defined such that they would focus in on a $3 \\sigma (4h^{-1}\\hbox{Mpc})$ void, using a constrained random field code (Van de Weygaert \\& Bertschinger 1996). The sequence shows the gradual development of a large void of diameter $\\approx 25h^{-1}\\hbox{Mpc}$ as the complex pattern of smaller voids and structures which had emerged within it at an earlier time, merge with one another. This illustrates one aspect of the evolving void hierarchy:  the {\\em void-in-void} process.} \\label{voidcdm6slc} \\end{figure*}  This paper is organized as follows. Section~\\ref{voidevol} discusses important generic properties of isolated voids, which grow from depressions in the primordial density field, propelled by the perturbed gravitational field. The spherical model forms the core of further analytical considerations, and is discussed in some detail in Appendix~\\ref{stophat}. Section~\\ref{lss} discusses the generic effects of larger scale stucture on the evolution of voids.  {\\em Two} crucial processes which shape the void hierarchy are described: the {\\em void-in-void mode} and the {\\em void-in-cloud mode}. How these processes can be incorporated into the excursion set approach using two-barriers is the subject of Section~\\ref{excurform}. Section~\\ref{voiddistr} describes the associated distribution of void sizes, which is predicted to have a universal form, and to be peaked around a characteristic value. One of the results of Section~\\ref{voiddistr} is to show that peak-based models should be reasonably accurate for the largest voids, but, because they account neither for the {\\em void-in-void mode} nor for the {\\em void-in-cloud mode}, they predict many more small voids than does the excursion set approach. Appendix~\\ref{grfs} discusses the {\\em ``basic troughs model''}, which assumes there is a one-to-one identification between minima in the primordial Gaussian density field, with centres of voids in the evolved (and nonlinear) matter distribution.  This also serves to define notation for  the {\\em ``adaptive troughs model''} which is described in Section~\\ref{ajpeaks}.  Section~\\ref{voidhier} presents various other aspects of the hierarchically evolving void population. Global parameters, such as the fraction of mass in the cosmos contained within void regions, along with the fraction of space occupied by voids, are readily derived from the void size distribution. In addition, the formalism is applied towards a reconstruction of the {\\em ancestral} history of a given void, followed by an evaluation of the environmental influence on basic void properties. We also put forward suggestions towards an analytical treatment of the influence of the void environment on the galaxies that may form within. Finally, we indicate how an assessment of the evolution of dark matter clustering may be predicated on our formalism. In Section~\\ref{discuss}, we provide an overview of our results and seek to embed these in the wider context of the study of hierarchical structure formation.  We also comment on how our results for the distribution of voids in the dark matter distribution may be related to observations of voids in the galaxy distribution.  Although our model provides a useful framework, developing a more detailed model is beyond the scope of this work.  The results of numerical studies of void galaxies in semi-analytic galaxy formation models are described by Mathis \\& White (2002) and Benson et al. (2003). ",
        "conclusions": "\\label{discuss} Initially underdense regions expand faster than the Hubble flow. If they are not embedded within overdense regions, such regions eventually form voids which are surrounded by dense void walls. These voids expand with respect to the background Universe, and during their expansion tend to become more and more spherical (Figure~\\ref{voidcdm6slc}).  The outward expansion is differential, so most initial void configurations tend to evolve to distinct ``tophat'' density profiles (Figure~\\ref{sphmodel}). A description of the evolution of initially spherical tophat over- and underdense regions has been available for some time (Appendix~\\ref{stophat}). Although the spherical evolution model allows one to study the evolution of single isolated objects, a more complete theory must also describe void evolution within the context of a generic random density fluctuation field.  The evolving void hierarchy is determined by two processes: \\begin{itemize} \\item The {\\em void-in-void} process describes the evolution of a system of voids which are embedded in a larger scale underdensity; in this case small voids from an early epoch merge with one another to form a larger void at a later epoch (Fig.~\\ref{voidcdm6slc}). \\item The {\\em void-in-cloud} process is associated with underdense regions embedded within a larger overdense region; in this case the smaller voids from an earlier epoch may be squeezed out of existence as the overdense region around them collapses (Fig.~\\ref{voidcollnbody}). \\end{itemize} In contrast, the evolution of overdensities is governed only by the {\\em cloud-in-cloud} process; the {\\em cloud-in-void} process is much less important, because clouds which condense in a large scale void are not torn apart as their parent void expands around them.  This asymmetry between how the surrounding environment affects halo and void formation can be incorporated into the {\\em excursion set approach} by using one barrier to model halo formation and a second barrier to model void formation (Fig.~\\ref{excurcensus}). Only the first barrier matters for halo formation, but both barriers play a role in determining the expected abundance of voids. The resulting void size distribution is a function of two parameters (equation~\\ref{vfvoid}), which the model associates with the dynamics of expansion and collapse. The predicted distribution of voids is well-peaked about a characteristic size (Figs.~\\ref{fvlin} and~\\ref{frvoid})---in contrast, the distribution of halo masses is not. Comparison of the two-parameter family of void distribution curves (Figure~\\ref{fvoid}) with the void size distribution in numerical simulations of hierarchical clustering is the subject of work in progress (Colberg et al. 2004).  Five major observations about the properties of the void population result from the {\\em two-barrier} excursion set model: \\begin{itemize} \\item{} The {\\em void-in-cloud} mechanism (Fig.~\\ref{voidcollnbody}) is responsible for the demise of a sizeable population of small voids.  As a result, the void size distribution has a small-scale cut-off:  the void population is ``void'' of small voids (Section~\\ref{voiddistr}), in a way which our excursion set analysis quantifies. \\item{} The population of large voids is insensitive to this effect (Fig.~\\ref{fvlin}).  Therefore, the abundance of voids which are larger than the typical initial comoving sizes of clusters should be well described by peaks theory or its extensions described in Section~\\ref{ajpeaks}. \\item{} At any cosmic epoch there is a {\\em characteristic void size} which increases with time: the larger voids present at late times formed from mergers of smaller voids which formed at earlier times (e.g Fig.~\\ref{voidcdm6slc} and Section~\\ref{voidhier}). \\item{} At any given time the mass fraction in voids is approximately thirty percent of the mass in the Universe, and the voids approximately fill space (Section~\\ref{voidhier}). \\item{} As the size of most voids will be similar to the characteristic void size, the cosmic matter distribution resembles a {\\em foamlike packing of spherical voids of approximately similar size and excess expansion rate}. This may explain why simple models based on the Voronoi tesselation exhibit many of the features so readily visible in N-body simulations of hierarchical clustering. \\end{itemize}    \\subsection{Galaxies in Voids} It is with some justification that most observational attention is directed to regions where most of the matter in the Universe has accumulated. Almost by definition they are the sites of most observational studies, and the ones that are most outstanding in appearance. Yet, for an understanding of the formation of the large coherent foamlike patterns pervading the Universe, it may be well worth directing attention to the complementary evolution of underdense regions. These are the progenitors of the observed voids, the vast regions in the large-scale cosmic galaxy distribution that are practically devoid of luminous matter.  When extensive systematic redshift surveys began mapping the spatial galaxy distribution, voids were amongst the most visually striking features.  Since then, the role of voids as key ingredients of the cosmic galaxy distribution has been demonstrated repeatedly in extensive galaxy redshift surveys (see Kauffmann \\& Fairall~1991; El-Ad, Piran \\& da Costa~1996; El-Ad \\& Piran~1997; Hoyle \\& Vogeley~2002; Plionis \\& Basilakos~2002; Rojas et al. 2003). A number of studies also indicated that observed voids exhibit distinct hierarchical features. Van de Weygaert~(1991) suggested the existence of a {\\it ``void hierarchy''} when pointing out that the galaxy distribution in the CfA/SRSS2 redshift survey (Geller \\& Huchra~1989; da Costa et al.~1993) gave the impression of small-scale voids embedded in the less pronounced large-scale underdense region delimited by the ``Great Wall''. Even in the most canonical specimen amongst its peers, the Bo\\\"otes void, traces of a faint structured internal galaxy distribution were found (Szomoru et al.~1996).  The dynamical impact of voids has proven to be crucial for understanding the cosmic flow patterns in the Local Universe. Measured peculiar galaxy velocities imply reconstructions of the local cosmic density field in which the repulsive actions of voids are important (e.g. Bertschinger et al.~1990; Strauss \\& Willick~1995; Dekel \\& Rees~1995). And more locally, the void's influence on cosmic flows was established when Bothun et al.~\\cite{bothun92} studied galactic peculiar motions along a wall around the largest void in the CfA redshift sample (de Lapparent, Geller \\& Huchra~1986).  In all these respects, voids in the galaxy distribution are similar to those in the dark matter distribution.  However, although voids in the galaxy distribution are mostly roundish in shape, they have typical sizes in the range of $20h^{-1}-50h^{-1}$Mpc (e.g. Hoyle \\& Vogeley~2002; Plionis \\& Basilakos~2002; Arbabi-Bidgoli \\& M\\\"uller~2002). These sizes are considerably in excess of the typical void diameters in our model of voids in the dark matter distribution, but note that the typical void size in the galaxy distribution depends on the galaxies which were used to define the void.  The voids associated with rare luminous galaxies are larger in part because the number density of such galaxies is lower.  As we describe below, our excursion set analysis provides a framework for modeling this dependence.  In recent years, the possibility that void galaxies are a systematically different population has received considerable attention (see e.g. Szomoru et al.~1996; El-Ad \\& Piran~2000; Peebles~2001; Mathis \\& White~2002, Rojas et al.~2003; Benson et al.~2003). In the simplest models of biased galaxy formation (e.g. Little \\& Weinberg~1994) one would expect to find voids filled with galaxies of low luminosity, or galaxies of some other uncommon nature (e.g. Hoffman, Silk \\& Wyse~1992). Indeed, even though various studies were oriented towards establishing the properties of voids in galaxies (e.g., Kauffmann \\& Fairall~1991; El-Ad \\& Piran~1997,~2000; Hoyle \\& Vogeley~2002; Arbabi-Bidgoli \\& M\\\"uller~2002; Plionis \\& Basilakos~2002), and some focussed explicitly on the identity of galaxies inside voids (e.g. Szomoru et al.~1996; El-Ad \\& Piran~2000; Rojas et al.~2003), a clear picture of the relation between void galaxies and their surroundings is only just becoming available. This is in large part due to the fact the large scale surveys such as the SDSS (Abazajian et al.~2002) and 2dFGRS (Colless et al.~2003) now probe a sufficiently large cosmological volume that they contain a statistically significant number of large voids.  Recently, Mathis \\& White~\\cite{mathwhit2002} and Benson et al.~\\cite{bensetal2003} have identified and studied voids and void galaxies in semi-analytic galaxy formation models. In these models, the properties of galaxies are determined by the halos they inhabit.  Therefore, if one can model the halo population associated with voids, a model of the void galaxy population is within reach. The excursion set model developed here is phrased in the same language used in the simulations, so it represents the ideal framework within which to attempt such a model.  In particular, consider the {\\em cloud-in-void} process shown in the second series of panels in Fig.~\\ref{excurcensus}. Notice that the condition that the cloud exist in a void means that, on average, clouds in voids will be less massive than clouds in regions of average density (to represent a cloud, the walk must reach $\\delta_{\\rm c}$, and on average, it will take more steps to travel to $\\delta_{\\rm c}$ from $\\delta_{\\rm v}$ than from zero---more steps imply smaller masses).  For similar reasons, the clouds associated with the more massive halos should be more massive on average (this is also discussed more fully by Mo \\& White~1996; Sheth \\& Tormen~2002 and Gottl\\\"ober et al.~2003). Although we speak of the clouds as being within the voids, our discussion of how voids empty their mass into the ridge which surrounds them (c.f. Fig.~\\ref{sphmodel}) suggests it may be more appropriate to think of these clouds as being associated with the void walls. It seems natural to associate void galaxies with such clouds-in-voids. If low mass halos host lower mass galaxies, and less massive galaxies tend to be less luminous and bluer, then void galaxies should be fainter and bluer than field or cluster galaxies; our model allows one to quantify this trend.  Thus, the results presented here allow a more elaborate model for voids in the galaxy distribution and the galaxy population in voids than that discussed recently by Friedmann \\& Piran~\\cite{frpir2001}. Developing such a model is the subject of work in progress."
    },
    "0311/astro-ph0311110_arXiv.txt": {
        "abstract": "{ We present infrared and radio continuum observations of S 201 star forming region. A massive star cluster is observed in this region, which contains different classes of young stellar objects. The near-infrared colour-colour and colour-magnitude diagrams are studied to discuss the nature of these sources. We have discovered the knots of molecular hydrogen emission at 2.122 $\\mu m$ in the central region of S 201. These knots are clearly seen along the diffuse emission in north--west direction which are probably the obscured Herbig-Haro objects. High sensitivity and high resolution radio continuum images from GMRT observations at 610 and 1280 MHz show an interesting arc-shaped structure due to the interaction between the HII region and the adjacent molecular cloud. The ionization front at the interface between the HII region and the molecular cloud is clearly seen by comparing the radio, molecular hydrogen and Br$\\gamma$ images. The emission from the carriers of Unidentified Infrared Bands in the mid-infrared 6--9 $\\mu m$ (possibly due to PAHs) as extracted from the Midcourse Space Experiment survey (at 8, 12, 14 and 21 $\\mu m$) is compared with the radio emission. The HIRES processed IRAS maps at 12, 25, 60, and 100 $\\mu m$, have also been used for comparison. The spatial distribution of the temperature and the optical depth of the warm dust component around S 201 region, has been generated from the mid-infrared images. }  \\titlerunning{A study of the galactic star forming region S 201.......} \\authorrunning{D.K. Ojha et al.} ",
        "introduction": "The high mass star forming region S 201, corresponding to the radio source W 5A, is at a distance of 2.3 kpc from the Sun (Mampaso et al. 1987). The S 201 region is coincident with an IRAS source (IRAS 02593+6016) with increasing spectrum from 12 to 100 $\\mu m$. The IRAS fluxes lead to a luminosity of $\\sim$ 1.1$\\times$10$^4$ L$_{\\odot}$ (Zinchenko et al. 1997) for this source. The HII region S 201 is part of an extended complex of gas and dust stretched in the east-west direction and it is bordered on the west by the molecular cloud IC 1848 identified with the radio source W 5.  Earlier observations of S 201 region in the infrared wavelengths (Kleinmann et al. 1979, Thronson et al. 1984, Mampaso et al. 1989, Carpenter et al. 1993) and the radio (Felli et al. 1987, Fich 1993, Omar et al. 2002) have shown an interesting structure presumably due to the interaction between the HII region and the adjacent molecular cloud (Martin \\& Barrett 1978). The radio source has a bright, arc-shaped edge on one side and a smoothly decreasing surface brightness distribution on the opposite side as seen in the VLA observations at 2 and 6 cm (Felli et al. 1987). They modeled this configuration in terms of a three-dimensional electron distribution produced by the ionizing radiation of an early-type star located outside a spherical molecular cloud and found that an O9 zero age main sequence (ZAMS) star placed at a distance of 0.38 pc from a molecular cloud of 0.42 pc radius allows a good fit to the their data.  We have started a near-infrared (NIR) and radio continuum observational programme of study of regions of massive star formation. This is our first paper in this series. In this paper we study the Galactic star forming region S 201 in the infrared and the radio wavelengths with the aim of identifying the stellar populations in the directions of the IRAS source \\mbox{(IRAS 02593+6016)}. This paper combines new NIR observations from Telescopio Infrarosso del Gornergratm (TIRGO) \\& Telescopio Nazionale Galileo (TNG) telescopes, Italy with mid-infrared (MIR) data (8--21 $\\mu m$) from Midcourse Space Experiment (MSX) and radio continuum observations at 610 \\& 1280 MHz from Giant Metrewave Radio Telescope (GMRT), India.  In Sect. 2, we present the broad- and narrow-band NIR observations, as well as complementary MIR and radio continuum observations. Sect. 3 deals with the results and the discussion and we summarize our conclusions in Sect. 4. ",
        "conclusions": "A detailed radio and infrared study of an embedded young stellar cluster associated with the Galactic star forming region S 201 is presented here. The colour-colour and colour-magnitude diagrams have been constructed to identify young stellar objects and estimate their spectral types. The high sensitivity and high angular resolution radio continuum maps based on the GMRT observations at 610 and 1280 MHz have been generated. These radio maps show interesting morphological details, including an interesting arc-shaped structure highlighting the interaction between the HII region and the adjacent molecular cloud. Three luminous infrared sources have been identified which are located within the radio nebulosity. The spatial distribution of temperature and optical depth of the interstellar dust component in the S 201 region, has been presented based on the mid- and far-infrared measurements from the MSX and IRAS (HIRES) missions.  We have detected a compact embedded star cluster located within the ionized nebula responsible for the radio continuum emission. The cluster has richness similar to the clusters surrounding early type Herbig Be stars; the most luminous member is consistent with an O6--O8 zero age main sequence star. Our analysis confirms that this single star is responsible for most of the Lyman continuum emission required to sustain the ionized radio nebula, with all the other cluster members contributing at most 10\\% of the ionizing photons. This explains the reason for the success of the model of Felli et al. (1987) (which assumed a single massive star for excitation), in explaining the radio continuum morphology of S 201. Our observations also confirm the position of this massive star as well as the location of the molecular cloud being eroded by the ionizing radiation. The PDR at the interface between the ionized gas and the molecular cloud is traced to the west of the bright arc-shaped feature in radio continuum and Br$\\gamma$, by the diffuse H$_2$ emission detected in our narrow-band near-infrared images."
    },
    "0311/astro-ph0311326_arXiv.txt": {
        "abstract": "The finding of a Type Ic supernova connected with GRB 030329 showed a massive star origin for this burst, supporting evidence for this association in previous bursts with lightcurve bumps at the appropriate time for a supernova. Here, we explore the possibility that all long bursts have massive star progenitors, interacting with either the freely expanding wind of the progenitor or the shocked wind. We present models for the afterglows of GRB 020405 and GRB 021211, which are a challenge to wind interaction models. Considering sources for which wind interaction models are acceptable, a range of wind densities is required, from values typical of Galactic Wolf-Rayet stars to values $\\sim 10^2$ times smaller. The reason for the low densities is unclear, but may involve low progenitor masses and/or low metallicities. If mass is a factor, a low density event should be associated with a low mass supernova. The interpretation of bursts apparently interacting with constant density media as interaction with a shocked wind requires both a range of mass loss densities and a range of external pressures. The highest pressures, $p/k\\ga 10^8$ cm$^{-3}$ K, may be due to an extreme starburst environment, which would imply that the burst is superposed on an active star forming region. Although the range of observed events can be accomodated by the shocked wind theory, special circumstances are necessary to bring this about. Finally, we consider the high velocity, high ionization absorption features observed in some afterglow spectra. If the features are circumstellar, the presence of the burst in a starburst region may be important for the formation of clumps near the burst. ",
        "introduction": "The spectroscopic finding of SN 2003dh in the afterglow light from GRB 030329 (Stanek et al. 2003; Hjorth et al. 2003) showed a direct link between a cosmological GRB (gamma-ray burst) and the supernova explosion of a massive star, supporting the previous identification of the nearby SN 1998bw with GRB 980425 (Galama et al. 1998). This event also supports the view that the ``bumps'' found in the light curves of a number of afterglows are, in fact, supernova light. Sources with plausible supernova light from photometry include GRB 970228 (Reichart 1999; Galama et al. 2000), GRB 980326 (Bloom et al. 1999), GRB 011121 (Bloom et al. 2002; Garnavich et al. 2003), and GRB 020405 (Price et al. 2003b). The recent GRB 021211 shows tentative evidence for supernova light in HST (Hubble Space Telescope) observations (Fruchter et al. 2002) and evidence for a supernova spectrum like that of the normal Type Ic  SN 1994I in the light of the burst has been reported (Della Valle et al. 2003). In this case, the possible supernova is fainter than SN 1998bw and is comparable in luminosity to  SN 1994I.  The implication of the supernova light is that the progenitor object is a massive star. The existing evidence is consistent with the supernova occurring at the same time as the GRB, and that is the assumption that we make. An implication of  a massive star progenitor is that the environment for the progenitor is determined by the mass loss wind from the star (M\\'esz\\'aros, Rees, \\& Wijers 1998; Dai \\& Lu 1998; Chevalier \\& Li 1999, 2000, hereafter CL99, CL00 respectively; Livio \\& Waxman 2000). If there is a free wind from the progenitor star, the wind density is $\\rho=Ar^{-2}=\\dot M/4\\pi r^2 v_w$, where $\\dot M$ is the mass loss rate and $v_w$ is the wind velocity. If the progenitor star is a Wolf-Rayet star, characteristic mass loss parameters are $\\dot M = 10^{-5}\\ml$ and $v_w=10^3 \\kms$; the value of the density can be scaled to the corresponding value for these parameters, $A_*=A/(5\\times 10^{11} {\\rm~ gm~cm^{-1}})$ (CL99).  The existing information on the afterglows of GRB 970228 and GRB 980326 is consistent with free wind interaction, although there is insufficient information to tie down the wind density (CL99, CL00). In the case of GRB 011121, wind interaction is indicated, with $A_*\\approx 0.02$  (Price et al. 2002c). GRB 020405 presents more of a problem, because Berger et al. (2003) found that a free wind model is not indicated and they proposed a model involving interaction with a uniform medium. For GRB 021211, Kumar \\& Panaitescu (2003) consider both uniform medium and wind models, both of which have certain difficulties. For their wind model, the surrounding density must be low, $A_*=5\\times 10^{-4}$. Because these two sources are of special interest from the massive star progenitor point of view, we consider models for their afterglows in \\S~2. In \\S~3, we note that a number of sources appear to be better modeled as interacting with a constant density medium, and investigate the possibility that these objects involve interaction with a shocked stellar wind (Wijers 2001). In view of the evidence for a low density environment around some GRBs, we consider the range of wind densities around Wolf-Rayet stars in \\S~4. Another possible constraint on the environment of GRBs comes from absorption lines observed in afterglow spectra. Possible implications of the lines are discussed in \\S~5. Discussion and conclusions are in \\S~6. We do not consider the tantalizing evidence for X-ray lines, which have not yet attained a high degree of statistical significance. ",
        "conclusions": "In this paper, we have emphasized some consequences of the hypothesis that all the long duration GRBs have massive star progenitors. Because the massive stars are expected to have their close-in surroundings modified by the progenitor winds, we consider both free winds and shocked winds as possible surrounding media for the afterglow phase. The properties of a shocked wind can be affected by the pressure of the surrounding medium when the pressure is high, as occurs in a starburst region. A number of afterglows have been interpreted in terms of interaction with a constant density medium with density $\\sim 20-30$ cm$^{-3}$ (Yost et al. 2003; Frail et al. 2003). Yost et al. note that this density is compatible with that of interstellar clouds or the interclump medium of Galactic molecular clouds. A different interpretation is needed in the massive star progenitor case, and we investigated the possibility that the medium is the shocked wind bubble in a starburst region. The relatively high densities found in some afterglows require that the burst occur in an extreme starburst region. In the case where there is direct interaction of a GRB with the interstellar medium, the interclump medium in a molecular cloud in a starburst region is likely to have a considerably higher density than the densities observed ($\\ga 10^3$ cm$^{-3}$, Chevalier \\& Fransson 2001). A significant fraction of GRBs must occur in starbursts, so the fact that such high densities have not been observed may be an indication that the surroundings of GRBs have been modified by stellar winds or ionizing radiation.  The fact that there is a substantial number of bursts that require a very high pressure surrounding medium may be a drawback to this scenario, but such sources are likely to be overrepresented in the sample because they are luminous and can be observed in detail. Among the afterglows that can be interpreted as interaction with a free wind, the highest density objects are compatible with expectations for the wind from a typical Galactic Wolf-Rayet star, but the lower densities imply a wind densities that are lower by a factor $\\sim 100$. Because the density in a shocked wind is higher than that in a free wind at the same radial point, the low density requirement is not alleviated by appealing to a shocked wind. One factor contributing to a low density wind may be a high wind velocity, as appears to occur for WO stars. Other factors that can contribute are a low stellar mass or a low metallicity, although it is not clear whether these factors can provide the low densities that are needed. A related point is that some of the densities deduced around Type Ic supernovae are surprisingly small. For moderately efficient production of synchrotron emission from SN 1998bw and SN 2001ap, the surrounding wind density corresponds to $A_*\\sim 0.01-0.05$ (Li \\& Chevalier 1999; Berger, Kulkarni, \\& Chevalier 2002). It is only if the fraction of energy density going into magnetic fields ($\\epsilon_B$) is small that the value of $A_*$ comes to $\\sim 1$. At the same time, moderately high values of mass loss from Wolf-Rayet stars are expected in order to obtain compatibility with stellar evolution (e.g., Langer 1989).  We have discussed afterglow models in terms of interaction with either a free wind or a constant density medium. These cases are the simplest for modeling because the dynamical situation can be treated in terms of a self-similar solution. If both of these cases are relevant, there is also the expectation that the GRB blast wave should in some cases be observed to traverse the termination shock wave, which marks the transition between the two types of media. Wijers (2001) has described some of the basic changes that might be expected for this transition, but they have not been clearly observed in any afterglow, which is a point against the general scenario. However, there is a need for more detailed hydrodynamic simulations of this interaction, as well as calculations of the expected emission properties."
    },
    "0311/hep-th0311260_arXiv.txt": {
        "abstract": "{We calculate the quasinormal modes (QNMs) for gravitational perturbations of the Schwarzschild black hole in the five dimensional (5D) spacetime with a continued fraction method. For all the types of perturbations (scalar-gravitational, vector-gravitational, and tensor-gravitational perturbations), the QNMs associated with $l=2$, $l=3$, and $l=4$ are calculated. Our numerical results are summarized as follows: (i) The three types of gravitational perturbations associated with the same angular quantum number $l$ have a different set of the quasinormal (QN) frequencies; (ii) There is no purely imaginary frequency mode; (iii) The three types of gravitational perturbations have the same asymptotic behavior of the QNMs in the limit of the large imaginary frequencies, which are given by $\\omega T_H^{-1}\\rightarrow\\log{3}+\\,2\\pi i (n+1/2)$ as $n\\rightarrow\\infty$, where $\\omega$, $T_H$, and $n$ are the oscillation frequency, the Hawking temperature of the black hole, and the mode number, respectively. } ",
        "introduction": "The quasinormal (QN) ringing of a Schwarzschild spacetime was first observed in \\cite{vish} through numerical calculations of the gravitational wave scattering by the black hole. Since then, the quasinormal modes (QNMs) of black holes have been extensively studied. The classical motivation behind the exploration of the QNMs of black holes is twofold: One is to answer the question of whether the spacetime is stable, and the other to know what kind of oscillations will be excited in the spacetime as some perturbations are given. The latter is quite important from the observational point of view because we could determine fundamental parameters of a black hole, such as the mass or the angular momentum, through the information of the QNMs. A number of studies on QNMs of several different spacetimes containing black holes have been done. (for a review see, e.g.,  \\cite{kokkotas99,nollert99}). Recently QNMs have acquired a different status, since it was conjectured that they may be connected to black hole area quantization and quantum gravity \\cite{hod,CLY03,All}.   Most studies on QNMs of black holes were restricted to the four dimensional (4D) case, compatible with astrophysical scenarios. However, motivated by the TeV-scale gravity proposals \\cite{hamed} for instance, higher-dimensional (higher-D) theories of gravity have recently attracted much attention.  Up to recently, there were no master equations for examining the QNMs for gravitational perturbations of higher-D black holes.  Due to the absence of these master equations, only the QNMs of the simplest test fields, namely massless scalar fields, around the higher-D black holes were calculated \\cite{CDL03,IUM03,Ko03a,MN03}.   The situation has changed recently. Kodama and Ishibashi \\cite{KI03a} have derived master equations for gravitational perturbations in a higher-D Schwarzschild black hole spacetime.  Putting $D=2+n$, it was shown that for the case $n\\ge3$, the gravitational perturbations can be divided into three classes, namely, scalar-gravitational, vector-gravitational, and tensor-gravitational perturbations, according to their tensorial behavior on the $n$-sphere. The scalar-gravitational and vector-gravitational perturbations correspond to the polar and axial perturbations in the 4D spacetime. The tensor-gravitational perturbations are a new kind appearing in higher-Ds.  Kodama and Ishibashi have also shown that all the types of gravitational perturbations can be reduced to a simple Schr\\\"odinger-type wave equation like Regge-Wheeler or Zerilli equations. It is important to ask the question of whether there is a special relationship among the scalar-gravitational, vector-gravitational, and tensor-gravitational perturbations because in 4D the scalar-gravitational and vector-gravitational perturbations have a special relation yielding the same quasinormal (QN) frequencies. There seems to be no such relation between the scalar-gravitational and vector-gravitational perturbations \\cite{KI03a}.  Recently, some studies on the QNMs for gravitational perturbations of higher-D Schwarzschild black holes have appeared. Konoplya has calculated the fundamental QNMs of Kodama and Ishibashi's master equations with a JWKB approximation, and has found that three types of perturbations have different fundamental QN frequencies \\cite{Ko03b}.  Berti, Cavagli\\`a, and Gualtieri have also done similar calculations but for a wide range of angular eigenvalues $l$ of perturbations \\cite{BCG03}. Cardoso, Lemos, and Yoshida have calculated the QNMs for the vector-gravitational and tensor-gravitational perturbations up to higher-order modes \\cite{CLY03}.  As for the asymptotic behavior in the limit of highly damped modes, it has been shown that the three types of gravitational perturbations have the same behavior regardless of the spacetime dimensions or the angular quantum numbers of the perturbations \\cite{Bi03,CLY03}.   The purpose of this paper is to go a step further up and explore the QNMs in the special 5D higher-D black hole spacetime. In particular, we want to find out whether or not there is a relationship among the QNMs of the scalar-gravitational, vector-gravitational and tensor-gravitational perturbations in this 5D Schwarzschild black hole.  The 5D case can be considered representative of the other higher-D cases.  A prime purpose of the present study is to extensively calculate the QNMs for the scalar-gravitational perturbations. The scalar-gravitational perturbations are the most important because they can be excited more easily than other types of perturbations. We calculate the QN frequencies not only for low-order modes but also for relatively higher-order modes, enabling us to discuss the asymptotic behavior of the highly damped QNMs. In the present study Nollert's approach is employed numerically to obtain the QN frequencies.  The paper is organized as follows: In section 2 we present the basic equations employed for obtaining QNMs in the 5D Schwarzschild spacetime. Numerical results are given in section 3, and in section 4 we conclude. ",
        "conclusions": "We have calculated the QNMs for the three types of gravitational perturbations of the 5D Schwarzschild black hole through a continued fraction method. In order to examine the QNMs, we made use of Schr\\\"odinger-type wave equations for determining the dynamics of the gravitational perturbations \\cite{KI03a}.  To apply the continued fraction method, we expanded the eigenfunctions around the black hole horizon in terms of a Fr\\\"obenius series. It was found that scalar-gravitational perturbations obey an eight-term recurrence relation, and vector-gravitational and tensor-gravitational perturbations obey a four-term recurrence relation.  For all the types of perturbations, the QNMs associated with $l=2$, $l=3$, and $l=4$ were calculated in the present study.  Our numerical results are summarized as follows; (i) The three types of gravitational perturbations have different QNM frequencies; (ii) There is no purely imaginary frequency mode (the above two results are consistent with Kodama and Ishibashi's results that there is no function satisfying equation (\\ref{super})); (iii) The QNMs belonging to the three types of gravitational perturbations have the same asymptotic behavior in the limit of the large imaginary frequencies, given by $\\omega T_H^{-1}\\rightarrow\\log{3}+2\\pi i (n+1/2)$ as $n\\rightarrow\\infty$; This asymptotic behavior was already expected \\cite{MN03,Bi03}.   For the black hole dynamics, the lower-order QNMs play an important role after the initial non-linear effects have become insignificant. In the 4D Schwarzschild black hole, one can guess which frequency will be excited in the ringing phase since the axial and polar QNMs have the same frequencies. On the other hand, there are three different sets of QNM frequencies for the Schwarzschild black hole in $D\\ge 5$ dimensions. Our results show that the fundamental QNM for the scalar-gravitational perturbations have the smallest frequency and the longest damping-time. This means that the QNMs for the scalar-gravitational perturbations can live longer than other types of perturbations if all the types of perturbations are excited with the same initial amplitudes. Probably, in general situations things are not so simple.  What type of perturbations will be excited will be strongly dependent on the situation. In order to expect what frequency will be excited in the ringing phase, one has to model the source terms of the master equations. So far, only the case where a test particle radially falls into a black hole has been considered \\cite{BCG03}. In this case, due to the symmetry of the motion, only the scalar-gravitational perturbations are excited. The general case remains as a work to be solved.   \\vskip 2mm"
    },
    "0311/astro-ph0311461_arXiv.txt": {
        "abstract": "The scaling of velocity fluctuation $\\delta v$ as a function of spatial scale $\\ell$ in molecular clouds can be measured from size-linewidth relations, principal component analysis, or line centroid variation.  Differing values of the power law index of the scaling relation $\\langle \\delta v^{2} \\rangle^{1/2} \\propto \\ell^{\\gamma_{3D}}$ in three dimensions (3D) are given by these different methods: the first two give $\\gamma_{3D} \\simeq 0.5$, while line centroid analysis gives $\\gamma_{3D} \\simeq 0$. This discrepancy has previously not been fully appreciated, as the variation of {\\it projected} velocity line centroid fluctuations ($\\langle {\\delta}v_{lc}^{2} \\rangle^{1/2}$~$\\propto$~$\\ell^{\\gamma_{2D}}$) is indeed described, in 2D, by $\\gamma_{2D}$~$\\approx$~0.5. However, if projection smoothing is accounted for, this implies, in 3D, that $\\gamma_{3D}$~$\\approx$~0. We suggest that a resolution of this discrepancy can be achieved by accounting for the effect of density inhomogeneity on the observed $\\gamma_{2D}$ obtained from velocity line centroid analysis. Numerical simulations of compressible turbulence are used to show that the effect of density inhomogeneity statistically but not identically reverses the effect of projection smoothing in the case of {\\it driven} turbulence so that velocity line centroid analysis does indeed predict that $\\gamma_{2D}$~$\\approx$~$\\gamma_{3D}$~$\\approx$~0.5. For decaying turbulence, the effect of density inhomogeneity on the velocity line centroids diminishes with time such that at late times, $\\gamma_{2D}$~$\\approx$~$\\gamma_{3D}$~+~0.5 due to projection smoothing alone. Deprojecting the observed line centroid statistics thus requires some knowledge of the state of the flow. This information can be inferred from the spectral slope of the column density power spectrum and a measure of the standard deviation in column density relative to the mean. Using our numerical results we can restore consistency between line centroid analysis, principal component analysis and size-linewidth relations, and we derive $\\gamma_{3D}$~$\\approx$~0.5, corresponding to shock-dominated (Burgers) turbulence, which also describes the simulations of driven turbulence at scales where numerical dissipation is negligible. We find that this consistency requires that molecular clouds are continually driven on large scales or are only recently formed. ",
        "introduction": "The presence of turbulence in molecular clouds is diagnosed by the prevalence of superthermal linewidths. A second important diagnostic is the presence of macroscopic velocity fluctuations, and the exponent, $\\gamma_{3D}$, that describes the variation of root mean square velocity fluctuation ($\\langle {\\delta}v^{2} \\rangle^{1/2}$) with spatial scale, $\\ell$, in 3D : $\\langle {\\delta}v^{2}\\rangle^{1/2}$~$\\propto$~$\\ell^{\\gamma_{3D}}$. The exponent, $\\gamma_{3D}$, can vary depending on the state of turbulence in the medium; e.g., the Kolmogorov (1941) case has $\\gamma_{3D}$~$\\approx$~1$/$3, which applies to incompressible turbulence and may apply to decaying compressible turbulence (Falgarone et al. 1994); the compressible, shock-dominated case (Burgers 1974), has $\\gamma_{3D}$~$\\approx$~1$/$2. The scale dependence of velocity dispersion in molecular clouds will in general be dependent on the energy injection spectrum, in addition to the transfer of energy between scales. For a pure turbulent cascade of kinetic energy, measurement of $\\gamma_{3D}$ provides fundamental physical information on the nature of interstellar turbulence.  In this paper we describe a long-standing but largely unappreciated discrepancy between observational estimates of $\\gamma_{3D}$. Velocity line centroid analysis, operating in (projected) 2D suggests that $\\gamma_{3D}$~$\\approx$~0 when the effects of projection smoothing (von Hoerner 1951; O'Dell \\& Castaneda 1987; Stutzki et al. 1998; Miville-Desch{\\^e}nes, Levrier, \\& Falgarone 2003; Brunt et al. 2003) are accounted for. Size-linewidth relations (e.g. Solomon et al. 1987) and principal component analysis (Heyer \\& Schloerb 1997; Brunt et al. 2003) suggest that  $\\gamma_{3D}$~$\\approx$~0.5.  We argue in this paper that, for driven turbulence, density inhomogeneity induces small scale structure on the projected velocity line centroid field such that the effects of projection smoothing are statistically but not identically reversed and thus that the observed $\\gamma_{2D}$ obtained from velocity line centroid analysis does in fact closely approximate the true 3D index, $\\gamma_{3D}$. Some indication that this is true was found by Ossenkopf \\& Mac Low (2002).  We begin with an overview of the statistical properties of velocity fields in 3D, and in projected 2D as seen by velocity line centroid analysis. In {\\S}~\\ref{sec:numerical} we describe our numerical data, and in {\\S}~\\ref{sec:observational} we describe our observational data. The results of velocity line centroid analysis applied to the numerical simulations of turbulence, both with and without density weighting, are discussed in {\\S}~\\ref{sec:numericalresults}. In {\\S}~\\ref{sec:observationalresults} we estimate the observed power spectrum of line centroid velocity fluctuations in the NGC~7129 molecular cloud, and interpret the measured power spectrum exponent and previously measured line centroid structure function exponents, accounting both for projection smoothing and density inhomogeneity. ",
        "conclusions": "}}  We have examined numerical simulations of supersonic turbulence in molecular clouds in order to understand the effects of projection smoothing and density inhomogeneity on observable velocity line centroid fields ($v_{lc}(x,y)$). We find that the observed index, $\\gamma_{2D}$, describing the variation of projected line centroid velocities ($\\langle {\\delta}v_{lc}^{2} \\rangle^{1/2}$~$\\propto$~$\\ell^{\\gamma_{2D}}$) is related to the intrinsic index, $\\gamma_{3D}$, (where $\\langle {\\delta}v^{2} \\rangle^{1/2}$~$\\propto$~$\\ell^{\\gamma_{3D}}$) via the following relation : $$ \\gamma_{2D} \\;\\;\\; = \\;\\;\\; \\gamma_{3D} \\;\\;\\; + \\;\\;\\; 0.5 \\;\\;\\; + \\;\\;\\; \\delta\\kappa/2,$$ where the increase of 0.5 is a direct consequence of projection smoothing, and $\\delta\\kappa$ is an empirically determined correction for the effects of density inhomogeneity.  For driven turbulence, both with and without magnetic fields, $\\delta\\kappa$~$\\approx$~--1 (with variations $\\sim$~$\\pm$~0.22), thus implying $\\gamma_{2D}$~$\\approx$~$\\gamma_{3D}$ due to the accidental cancellation of the effects of projection smoothing and small-scale structure induced by density inhomogeneity on the line centroid field. For decaying turbulence, $\\delta\\kappa$ tends to $\\sim$0 at late times, so that $\\gamma_{2D}$~$\\approx$~$\\gamma_{3D}$~+~0.5 due entirely to projection smoothing.  The variations in $\\delta\\kappa$ observed in high Mach number driven turbulence have no obvious dependence on Mach number, $M$. As $M\\longrightarrow$0 (the incompressible case) one should also expect that $\\delta\\kappa\\longrightarrow$0, as is indeed seen at late times in the decaying runs. We have inferred that $\\delta\\kappa$ is determined mostly by the amplitude of density fluctuations. The scatter in $\\delta\\kappa$ in the supersonic regime emphasizes the need for continued high spatial dynamic range surveying of the molecular ISM in order to provide a statistically useful sample of fields.  Deprojection of observed scaling relations in velocity line centroid fields requires some knowledge of the state of the flow. This information can be inferred from the spectral slope of the column density power spectrum and a measure of the standard deviation in column density relative to the mean ($\\sigma_{N/N_{0}}$). Our numerical results show that $\\gamma_{3D}$ is now consistently estimated at $\\sim$0.5 using velocity line centroid analysis and using principal component analysis and size-linewidth relations. This consistency requires that molecular clouds are continually driven or only recently formed."
    },
    "0311/astro-ph0311182_arXiv.txt": {
        "abstract": "We have observed the QSO 3C\\,345 with VSOP at 1.6\\,GHz and 5\\,GHz and with the VLBA at 22\\,GHz and 43\\,GHz, with 24 epochs in total, from 1997.40 to 1999.69. This study, combined with previous observations, reveals periodic changes in the parsec-scale jet. This quasi-periodicity can be found in the initial position angles, the trajectories and the flux density evolution of the jet components, with a period of about 9\\,yr. This period should reflect the dynamics of the material in the central region of 3C\\,345 and it is possibly related to a putative binary black hole in the center of this object. ",
        "introduction": "The 16th magnitude quasar 3C\\,345 ($z=0.595$) can be regarded as one of the archetypical sources for studies of superluminal motion with components moving along strongly curved trajectories on pc-scales. Several components with apparent velocities of $2-20c$ have been observed in the jet of 3C\\,345 (Zensus, Cohen, \\& Unwin 1995; Lobanov 1996). Periodic outbursts occurring every $10.1 \\pm 0.8$\\,yr were found by Zhang, Xie, \\& Bal (1997) from the B-band historical light curve.  The ejection directions of the jet components from the core (the optically thick region) vary, and the component trajectories differ significantly from each other (e.g. Zensus et al. 1995). The curvature of the component trajectories increases towards the core, while the trajectories straighten at larger core separations. This makes the innermost region the most interesting for studying the kinematics of the jet components. ",
        "conclusions": "The results presented above emphasize the existence of an underlying quasi-periodic process in the pc-scale jet of 3C\\,345 with a time period of about 9\\,yr. A possible scenario, which can explain the observed quasi-periodicity, is the binary black hole (BBH) system (e.g., Begelman, Blandford, \\& Rees 1980). The BBH model was applied to describe the optical light curve and the kinematic and flux density evolution of the jet component C7 in 3C\\,345 (Lobanov \\& Roland 2002). The model yields an orbital period of 170\\,yr and a precession period of 2500\\,yr. The $2.6^\\circ \\mbox{year}^{-1}$ of the initial P.A. shown in Figure 2 may be related either to the orbital or to the precession period in this system. The observed 9\\,yr period is too short to be related to the orbital period in the BBH system. The 9\\,yr period is more likely to be caused by the rotation of the jet or the rotation of the accretion disc. Further work has to be done to determine the origin of the 9\\,yr period."
    },
    "0311/astro-ph0311527_arXiv.txt": {
        "abstract": "We present results from a temporal analysis of the longest and the most sensitive X-ray observations of TV~Columbae -- a cataclysmic variable classified as an intermediate polar. The observations were carried out with the $Rossi X-Ray Timing Explorer (RXTE)$ Proportional Counter Array, $ROSAT$ Position Sensitive Proportional Counter, and $ASCA$. Data were analyzed using a one-dimensional CLEAN and Bayesian algorithms. The presence of a nearly sinusoidal modulation due to the spin of the white dwarf is seen clearly in all the data, confirming the previous reports based on the $EXOSAT$ data. An improved period of $1909.7\\pm2.5$s is derived for the spin from the $RXTE$ data. The binary period of 5.5hr is detected unambiguously in X-rays for the first time. Several side-bands due to the interaction of these periods are observed in the power spectra, thereby suggesting contributions from both the disk-fed and the stream-fed accretion for TV~Col. The accretion disk could perhaps be precessing as side-bands due to the influence of 4 day period  on the orbital period are seen. The presence of a significant power at certain side-bands of the spin frequency indicates that the emission poles are asymmetrically located. The strong power at the orbital side-bands seen in both the $RXTE$ and $ROSAT$ data gives an indication for an absorption site fixed in the orbital frame. Both the spin and the binary modulation are found to be energy- dependent. The fact that the spin modulation amplitude in TV~Col decreases with energy is confirmed. Hardness ratio variations and the energy dependent modulation depth during the spin modulation can be explained by partially covered absorbers in the path of X-ray emission region in the accretion stream. Increased hardness ratio during a broad dip in the intensity at binary phase of 0.75--1.0 confirms the presence of a strong attenuation due to additional absorbers probably from an impact site of the accretion stream with the disk or magnetosphere. ",
        "introduction": "The cataclysmic variable (CV), TV~Columbae, is classified as an intermediate polar (IP) containing a magnetic white dwarf accreting from the Roche-Lobe of a late-type dwarf companion via an accretion disk and an accretion channel. It was identified by Charles et al. (1979) as the optical counterpart of a hard X-ray source (2A 0526-328) first discovered with the $Ariel~5$ satellite (Cooke et al. 1978). An optical star with V$\\sim$13--14 mag, TV~Col is at a distance of $368^{+17}_{-15}$ pc (McArthur et al. 2001).  A 1911$\\pm$5 s period has been detected from this source in X-rays using the $EXOSAT$ observatory (Schrijver, Brinkman, \\& van der Woerd 1987).  The X-ray period has been identified as the spin period of the white dwarf, and the spin-phased light curves have been found to be energy dependent (Norton \\& Watson 1989). X-ray observations of TV~Col have also been reported from the $ROSAT$, $ASCA$, and $Ginga$ satellites by Vrtilek et al. (1996), Ezuka \\& Ishida (1999), and Ishida \\& Fujimoto (1995), respectively, but they mainly discuss the spectral characteristics of the source rather than the timing properties. Vrtilek et al. (1996), however, reported a failure to detect the X-ray spin period in the $ROSAT$ observation analyzed by them. A period of 1938$\\pm$10s was claimed to be present in the optical bands (UBV) at a level of $\\sim$6\\% (full amplitude) by Bonnet-Bidaud, Motch, \\& Mouchet (1985) (see their note added in proof) after the discovery of 32 minute period in X-rays with the $EXOSAT$ satellite was reported by Schrijver et al. (1985).  TV~Col also exhibits additional periods of 5.2 hr, 5.5 hr and 4 days (Motch 1981; Hutchings et al. 1981; Bonnet-Bidaud et al. 1985; Barrett, O'Donoghue, \\& Warner 1988; Hellier, Mason, \\& Mittaz 1991; Hellier 1993; Augusteijn et al. 1994) observed in optical photometry and spectroscopy, which add further complexities. None of these periods have been seen convincingly in the X-ray observations reported so far, although Hellier, Garlick, \\& Mason (1993) did report some evidence for intensity dips separated by an orbital cycle in the data taken with the $EXOSAT$ observatory. The spectroscopic period of 5.5 hr is believed to represent the binary orbital modulation and shows partial eclipses.  The 4-day period is possibly due to the nodal precession of the accretion disk around the white dwarf (Barrett et al. 1988 and Hellier 1993). The 4-day period was also reported from UV observations by Mateo, Szkody, \\& Hutchings (1985). The 5.2 hr period likely represents the beat between the 5.5 hr and 4-day periods. Recently Retter et al. (2003) have reported another period at 6.3 hr in the optical data, and identified it as representing the permanent positive superhump period of the binary system. The positive superhump is usually explained as the beat between the binary period and the precession of an eccentric accretion disk in the apsidal plane (see reviews by Patterson 2001; O'Donoghue 2000).  The power spectrum of IPs in X-ray band can be very complex containing the fundamental system frequencies, their harmonics and sometimes the side-band frequencies due to beating of the fundamental periods and their harmonics. The X-rays, which are believed to originate from the shock heated plasma region near the surface of the white dwarf, show modulation at all or some of the above mentioned frequency components. Hence a long and well sampled hard X-ray data particularly taken from an instrument with large effective area and high time resolution can provide important clues about the mode of accretion and the geometry of these sources.  Here, we report on the longest and the most sensitive broad-band X-ray observations carried out in the hard X-rays with the $Rossi X-Ray Timing Explorer (RXTE)$, and in the soft X-rays with the $ROSAT$ satellites. Although these observations are not simultaneous these are the best available for carrying out temporal analysis to search for various periodicities in the X-ray emission from TV~Col and to study their energy dependence. Results from the timing analysis of simultaneous soft and hard X-ray observations with the $ASCA$ satellite are also presented. A detailed timing analysis of these observations has not been reported so far except for a preliminary report in Singh et al. (2003). The paper is organized as follows: In the next section we present details of observations and data reduction in the hard X-ray and the soft X-ray energy bands. Section 3 contains results obtained from the power spectral analysis, and the light curves folded on the spin period and the orbital period of the system, followed by a discussion and interpretation of results in \\S 4.  Finally, we present a summary of our results and the conclusions in \\S 5. ",
        "conclusions": "  \\begin{enumerate}  \\item Power density spectra of X-ray emission from TV~Col show strong peaks at frequencies $\\omega$ and $\\Omega_0$ suggesting these to be two fundamental frequencies of the system, i.e., the spin frequency and the orbital frequency. The presence of spin modulation in X-rays is thus confirmed and a significant power at $\\omega$ is detected at all energies between 0.1 to 20 keV. The $\\Omega_0$ component is clearly present in the energy bands below 10 keV, but is absent in the hardest energy band of 10--20 keV.  \\item Significant power at two fundamental periods $\\omega$ and $\\Omega_0$ and several side-band components of these two periods in the power spectra of TV~Col provides evidence that contributions from both the disk-fed and the stream-fed accretion are present. The presence of certain side-bands of $\\omega$ suggests asymmetrically placed poles. The effects of 4-day precession period can also be seen in the side-bands of $\\Omega_0$.  \\item Modulation at the spin period being a function of energy is confirmed. It is $\\sim$ 3--4 times larger in the softest X-rays when compared with the value in the hardest X-rays observed. The observed energy dependent spin modulation and its anti-correlation with the hardness ratios defined above 2 keV can be explained by a partially covered absorber system, with a small change in the covering fraction being sufficient to explain the in-phase soft X-ray modulation without affecting the X-ray intensity above 2 keV.  \\item The effect of the orbital period is seen clearly in the hard and soft X-ray light curves of TV~Col. Modulation due to the orbital period is also seen to be more prominent in the soft X-rays than in the hard X-rays. X-ray light curves show a broad dip at orbital phase between $\\phi$=0.7--1.0. The hardness ratio variations during the dip can be understood by invoking an additional absorber (or absorbers) co-rotating in the orbital frame. A significant substructure during the dip suggests that these additional absorbers may be inhomogeneous. The absorbers could be located near the region where the accretion stream interacts with the accretion disk.  \\end{enumerate}"
    },
    "0311/astro-ph0311288_arXiv.txt": {
        "abstract": "Using adaptive optics we study the binary population in the nearby OB~association Scorpius~OB2. We present the first results of our near-infrared adaptive optics survey among 199 (mainly) A- and B-type stars in Sco~OB2. In total 151 components other than the target stars are found, out of which 77 are probably background stars. Our findings are compared with data collected from literature. Out of the remaining 74 candidate physical companions 42 are new, demonstrating that many stars A/B stars have faint, close companions. ",
        "introduction": "The primordial binary population (PBP) is defined as the population of binaries as established just after the gas has been removed from the forming system, i.e. when the stars can no longer accrete gas from their surroundings \\cite{kouwenhoven2003}. Characterizing the PBP is important for our understanding of the process of star formation, the formation and evolution of OB associations, the origin of the field star population and OB runaway stars, and the production and evolution of binary systems. OB associations are ideal sites for the study of the PBP. Since OB associations are young ($5-50$~Myr) and low density ($\\approx 0.1~{\\rm M}_\\odot~{\\rm pc}^{-3}$) stellar systems, the effects of stellar and dynamical evolution are modest. OB associations are practically cleared of gas and the full stellar population (from OB stars to brown dwarfs) is present. Sco~OB2 (Figure~\\ref{figure: eps}) is the closest young OB association, its proximity facilitating the detection of close and faint companions. The membership and stellar content has been established with {\\it Hipparcos} \\cite{dezeeuw1999}, and many binary surveys have been performed in the past. ",
        "conclusions": ""
    },
    "0311/astro-ph0311241_arXiv.txt": {
        "abstract": "We present initial results from a wide field photometric survey of the Fornax dwarf spheroidal (dSph) galaxy.  The focus here is on a small overdensity of stars located approximately $17'$ ($670$ pc in projection) from the centre of Fornax.  Based on imaging data in both $V$ and $I$ bands down to $V \\approx 21$, the dimensions of the feature are $\\sim$1.7$' \\times 3.2'$ ($68 \\times 128$ pc) with an average surface brightness of $25.8$ mag/arcsec${}^2$ in $V$.  A follow-up analysis using the deep $B$ and $R$ band data obtained by \\citet{ste98} indicates that the overdensity is apparently dominated by a relatively young stellar population with an age $\\approx 2$ Gyr, though the current data do not rule out the presence of older stars.  Our preferred interpretation is that the overdensity represents shell structure in the Fornax dSph galaxy, a phenomenon previously unseen in dwarf galaxies.  The shell may be the remnant of a merger with Fornax of a smaller, gas-rich system that occured approximately 2 Gyr ago.  The comparitively recent star formation seen in some dSph and dE galaxies may then also be the result of infall of small relatively gas-rich systems. ",
        "introduction": "Dwarf spheroidal (dSph) galaxies provide important clues to our understanding of the universe. For example, the high mass-to-light ratios found in these systems make them excellent laboratories for the study of dark matter.  Further, due to the relatively large number of dwarf galaxies, they are useful tracers of matter (both luminous and dark) in the local universe. At present, the most widely accepted model describing the formation of structure in the universe is the Cold Dark Matter (CDM) paradigm, based on hypothetical non-baryonic particles \\citep{pee93,pea99}.  CDM predicts that relatively small-scale clumps of dark matter formed from fluctuations in the density spectrum of the early universe.  Large structures were then formed through a process of hierarchical merging.  According to this theory, the initial dark matter haloes had masses of $\\sim$10$^7$ M${}_{\\odot}$ \\citep{may02, moo99}, somewhat smaller than the masses of Galactic dSphs ($2 \\times 10^7$ M${}_{\\odot}$ to $7 \\times 10^7$ M${}_{\\odot}$, e.g.\\ \\citet{mat98}). Consequently, the inference from the simulations is that the current Galactic dSphs have also formed from accretion events in the early Universe.  However, high resolution simulations of individual dark matter halos (Power 2003, personal communication) indicate that late infall is negligible in typical dwarf halos.  Therefore, discounting a fleeting encounter with a dark clump in the Galactic halo, dwarf galaxies are not expected to show any significant substructure. Nevertheless, substructure is observed in at least one Galactic dSph, namely Ursa Minor (e.g.\\ \\citet{oa85,dmd01,bel02,pms03,kwg03}), though in this particular case the origin of the substructure may lie with the tidal field of the Galaxy.  Although CDM accurately reproduces the large-scale structure of the universe, its predictions do not match observations of small scale regions with a high contrast in density.  For example, the mass distribution of low surface brightness (LSB) galaxies inferred from their rotation curves do not have the steep density cusp at the core expected from CDM models (e.g.\\ \\citet{wel03}).  In addition, the number of dwarf galaxies surrounding our Galaxy is an order of magnitude smaller than the mass spectrum predicted by CDM (e.g.\\ \\citet{moo99}). Further, in the CDM picture, the Galaxy's halo was built up through mergers of low-mass systems such as dwarf galaxies.  However, the abundance patterns observed in present day dwarfs appear to be distinct from those exhibited by Galactic halo field stars \\citep{she01,tol03}.  Indeed, age, chemical and dynamical inconsistencies make currently existing dwarfs unlikely to be the major building blocks of large systems such as spiral and elliptical galaxies (e.g.\\ \\citet{tos02}).  The Fornax dwarf spheroidal galaxy is the brightest object of this type in the Galactic halo ($M_V = -13.2$, from \\citet{mat98}).  It lies at a distance of $138 \\pm 8$ kpc, and has a tidal radius of $r_t = 71' \\pm 4'$  \\citep{ih95} equivalent to $r_t = 2.85 \\pm 0.16$ kpc.  Fornax is one of nine known dwarf spheroidal companions of the Milky Way, all of which show little or no indication of HI or HII (with the possible exception of Sculptor, see \\citet{car98,bou03}).  Fornax displays a very complex star formation history, beginning at least 12 Gyr ago and extending intermittently until $\\sim 500$ Myr ago (see \\citet{ste98,sav00} and references therein). There are also indications, in the form of an excess of stars beyond the `tidal radius' of the best-fit King model, that the gravitational potential of the Galaxy is tidally heating the outer regions of this dSph \\citep{ih95}.  However, \\citet{wal03} find no such excess, albeit for a larger `tidal radius' ($r_t = 98'$) than found by \\citet{ih95}. Possible asymmetries in the distribution of stars in the inner regions of Fornax have been reported by \\citet{pwh61}, \\citet{esk88} and \\citet{ste98}.  The latter authors drew particular attention to the difference in distribution between the youngest stars and that for the bulk of the population.  A measurement of the proper motion of Fornax, made using Hubble Space Telescope data, indicates that this dwarf is currently at or near perigalacticon \\citep{pia02}.  In this paper we examine an overdensity of stars situated approximately $17'$ from the centre of the Fornax dSph.  This overdensity does not correspond to any previously reported asymmetry in the distribution of Fornax stars.  By collecting accurate photometry down to the red clump at $V=21$, we have been able to determine the surface brightness and dimensions of this feature. Subsequent analysis of the deep survey data of Fornax provided by \\citet{ste98} reveals that the overdensity is apparently dominated by a relatively young main sequence stellar population.  We propose that this clump is the first detection of shell structure in a dwarf galaxy: such structure, despite being well observed in large galaxies, was not previously known in small systems such as dwarf galaxies.  The recognition of shell structure in elliptical galaxies (e.g.\\ \\citet{mal80}) provided an important insight into understanding the interactions between a large galaxy and a smaller companion. Phase-wrapped shells in the haloes of large galaxies became obvious through the process of photographic amplification \\citep{mal78} of images of elliptical galaxies (e.g.\\ \\citet{mal83b}).  Subsequent simulations by \\citet{her88,her89} indicated that a large galaxy disrupting a small companion will produce the observed structure, with some ($\\sim10\\%$) of the material of the companion galaxy distributed in shells around the primary.  The level of ordered structure in the shells reveals the type of interaction between the two objects: a direct collision will produce highly ordered, phase-wrapped shells around the primary galaxy, while a companion on a non-radial orbit will produce a confused set of shells.  Observational evidence supports this process of shell formation.  In particular, studies of the colour of the shells shows that they are generally bluer than the colour of the primary galaxy, reflecting the stellar population of the disrupted lower mass companion \\citep{pen86,mcg90}.   In the following section we describe our Fornax observations and the data reduction techniques.  Section 3 outlines the analysis of the data and the results are discussed in section 4.  The conclusions are presented in section 5. ",
        "conclusions": "We have presented photometric data for a stellar clump located $\\sim$17$\\arcmin$ from the centre of the Fornax dwarf spheriodal galaxy. The clump is approximately 68 $\\times$ 128 pc in size and is $\\sim$40\\% more dense than the surrounding regions.  It is located on the Fornax minor axis and is elongated in the direction of the dSph's major axis.  By removing the Fornax field contamination through a CMD density function analysis, we have shown that the clump appears dominated by a stellar population approximately 2 Gyr in age.  Our interpretation of this clump is that it is a shell feature analogous to those frequently seen around more luminous E galaxies.  As such, it would represent the first detection of shell structure in a dwarf galaxy.  In this interpretation we propose that a smaller gas-rich system merged with Fornax in the relatively recent past ($\\sim$2 Gyr ago) and that the shell is a consequence of this merger.  Models of such a merger, kinematic studies of the shell stars, and further deep imaging of the region surrounding this dSph to constrain any additional shell structure, are all required to investigate this hypothesis more thoroughly.  It may be that in many cases the comparitively recent star formation seen in some dSph and dE galaxies is also the result of the infall of lower mass, relatively gas-rich systems."
    },
    "0311/astro-ph0311077_arXiv.txt": {
        "abstract": "{We present new multi-frequency EVN, MERLIN and VLA observations of the radio source 3C\\,264, sensitive to linear scales ranging from the parsec to several kiloparsecs. The observations confirm the existence of regions with different properties in the first kiloparsec of the jet.  The most remarkable feature is the transition between a well collimated narrow jet at distances from the core below 80 pc, to a conical-shaped wide jet, with an opening angle of $20^{\\circ}$. Another change of properties, consisting of an apparent deflection of the jet ridge line and a diminution of the surface brightness, occurs at a distance of $\\sim 300$ pc from the core, coincident with the radius of a ring observed at optical wavelengths. Our observations add new pieces of information on the spectrum of the radio-optical jet of 3C\\,264, with results consistent with a synchrotron emission mechanism and a spectrum break frequency in the infrared.  Brightness profiles taken perpendicularly to the jet of 3C\\,264 are consistent with a spine brightened jet at distances below 100 pc from the core, and an edge-brightened jet beyond, which can be interpreted as evidence of a transverse jet velocity structure. Our observations do not allow us to distinguish between the presence of a face--on dust and gas disk at the center of the host galaxy of 3C\\,264, or rather an evacuated bubble. However, the properties of the jet structure, the changes in the polarization angle, and the plausible jet orientation can be naturally brought into agreement in the bubble scenario. ",
        "introduction": "The radio source \\object{3C\\,264} (B1142+198) is a low luminosity radio galaxy optically identified with NGC\\,3862, an elliptical galaxy in the cluster Abell 1367 with $m_v = 13.67$ and a redshift $z=0.0217$ (Smith et al. \\cite{smith}). This radio galaxy has been observed at many different wavelengths from the radio to the X-ray, allowing a good determination of its nuclear spectral energy distribution (Capetti et al. \\cite{capetti2}). 3C\\,264 has a Fanaroff-Riley type I radio structure (Fanaroff \\& Riley \\cite{fanaroff}) with a head-tailed morphology at kiloparsec scales\\footnote{We assume throughout a Hubble constant of $H_0=72$ km s$^{-1}$ Mpc$^{-1}$ and a deceleration parameter of $q_0=-0.55$, so that $1''$ corresponds to 412 pc}, a prominent core and a wiggling jet extending toward the northeast that ends in a blob of emission at $28''$ (11.5 kpc) from the core (Gavazzi, Perola \\& Jaffe \\cite{gavazzi}; Bridle \\& Vall\\'ee \\cite{bridle}; Baum et al. \\cite{baum1}; Lara et al. \\cite{lara97}). There is evidence of counterjet emission in the southwest direction from the core (Lara et al. \\cite{lara97}). Both the jet and counterjet are embedded in a vast and diffuse region of low surface brightness emission which seems to be dragged toward the north, possibly revealing the existence of a high density intra-cluster medium.  Simultaneous EVN\\footnote{EVN: European VLBI Network} and MERLIN\\footnote{MERLIN: Multi-Element Radio Linked Interferometer Network} observations at 5.0 GHz by Lara et al. (\\cite{lara99}) showed for the first time the detailed structure of 3C\\,264 at sub-kiloparsec scales. It consists of a one-sided jet with evident variations in its morphological properties with distance: {\\em i)} the strong core and innermost jet (0 - 10 pc); {\\em ii)} a well-collimated and narrow region (10-100 pc); {\\em iii)} a region with strong widening, kinks and filaments (100-300 pc); {\\em iv)} a faint and narrow region after a jet deflection (300 - 400 pc from the core).  Besides this rich radio structure, another peculiarity of 3C\\,264 is the existence of an optical jet coincident with the radio jet, most possibly of synchrotron nature (Lara et al. \\cite{lara99}) and extending up to $2''$ (0.82 kpc) (Crane et al. \\cite{crane}; Baum et al. \\cite{baum2}). There is also evidence of a ``ring'' around the host galaxy core with a radius between $0''.75$ and $1''$ (300 - 400 pc) (Baum et al. \\cite{baum2}).  We present in this paper new observations of the radio galaxy 3C\\,264, made with MERLIN and the EVN at 1.6 GHz, and VLA\\footnote{VLA: Very Large Array} observations at 1.6, 5.0, 22.5 and 43.3 GHz. The new data allow a direct comparison with previous sub-kiloparsec scale observations at 5.0 GHz (Lara et al. \\cite{lara99}) and shed light on the properties of the inner kiloparsec of the jet in 3C\\,264, a length scale which is crucial to the understanding of the FR I vs. FR II dichotomy of radio galaxies. ",
        "conclusions": "We present in this paper new observations of the radio source 3C\\,264 which are sensitive to linear scales ranging from the parsec to several kiloparsecs. Observations at 1.6 GHz were made simultaneously with the EVN and MERLIN in 1997. Observations at 1.6, 5.0, 22.5 and 43.3 GHz were made two years later with the VLA in its most extended configuration. The main conclusions are:  \\begin{enumerate} \\item The observations confirm the existence of different regions in the jet of 3C\\,264, as previously found by Lara et al. (\\cite{lara99}). The most remarkable feature is the transition between a well collimated narrow jet at distances from the core below 80 pc, to a conical shaped wide jet, with an opening angle of $20^{\\circ}$. This transition coincides also with a transition in the transverse structure of the jet, which is consistent with a spine brightened jet in the collimated region and an edge brightened jet in the conical region.  This result is also in agreement with the possible presence of jet velocity structure (fast spine surrounded by a slower layer) as found in other FR I sources (e.g. \\object{B2\\,1144+35}; Giovannini et al. \\cite{giov99}) or Bl Lacs (Mkn\\,501; Giroletti et al. \\cite{giroletti})  \\item An apparent deflection of the jet ridge line, together with a diminution of the surface brightness and a change in the polarization properties, is present at a distance of $\\sim 300$ pc from the core, strikingly coincident with the radius of a ring observed at optical wavelengths.  \\item From observational data we estimate an initial jet velocity $\\beta=0.99$ and an orientation of the jet with respect to the observer $\\theta = 50^{\\circ}$, consistent with previous estimations of these parameters (Baum et al. \\cite{baum2}).  \\item Our data contribute to a better knowledge of the spectrum of the radio-optical jet of 3C\\,264. We find a radio spectral index of -0.58 which has to be compared with a spectral index of -1.34 at optical wavelengths. The spectrum between radio and optical wavelengths is consistent with a synchrotron emission mechanism and the existence of a spectrum break frequency in the infrared.  \\item The properties of the jet structure, the changes in the polarization angle, and the plausible jet orientation ($\\sim 50^{\\circ}$) with respect to the observer can be naturally brought into agreement with the presence of a spherical bubble in the center of NGC\\,3862, even if the existence of a face on dust and gas disk cannot be completely ruled out. \\end{enumerate}"
    },
    "0311/astro-ph0311594_arXiv.txt": {
        "abstract": "Annihilation radiation from neutralino dark matter at the Galactic center (GC) would be greatly enhanced if the dark matter were strongly clustered around the supermassive black hole (SBH). The existence of a dark-matter ``spike'' is made plausible by the observed, steeply-rising {\\it stellar} density near the GC SBH. Here the time-dependent equations describing gravitational interaction of the dark matter with the stars are solved. Scattering of dark matter particles by stars would substantially lower the dark matter density near the GC SBH over $10$ Gyr, due both to kinetic heating, and to capture of dark matter particles by the SBH. This evolution implies a decrease by several orders of magnitude in the observable flux of annihilation products compared with models that associate a steep dark matter spike with the SBH. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311307_arXiv.txt": {
        "abstract": "The large number of quasars found in the Sloan Digital Sky Survey has allowed searches for elusive, clear lines of sight towards He {\\sc ii} Ly$\\alpha$ absorption, a sensitive probe of the intergalactic medium. The few known systems indicate that He {\\sc ii} reionization occurs at $z >3$. We report the detection of a He {\\sc ii} Ly$\\alpha$ absorption edge in a quasar spectrum at $z=3.50$, the most distant such feature found to date. The candidate quasar was selected from a $z\\sim3$ sample in the SDSS spectroscopic quasar survey and confirmed as part of an {\\it HST}/STIS SNAP survey. We discuss the general characteristics of the absorption feature, as well as the probability for discovery of additional such objects. ",
        "introduction": "Much of the intergalactic hydrogen % exists in a highly ionized state as evidenced by the lack of \\h\\ Gunn-Peterson troughs in the optical spectra of high-redshift quasars \\citep{gp}. The first presence of a Gunn-Peterson trough was detected in \\he\\ absorption, in a quasar at $z = 3.28$ \\citep{jakobsen}. Follow-up observations at a higher S/N level and spectral resolution \\citep{hogan,heap} found an average \\he\\ optical depth $\\tau = 4.0 \\pm 0.5$ at $z \\sim 3$ and at least $\\tau=5$ at $z\\sim 3.2$. \\he\\ absorption displays significant structure, consisting of unabsorbed void regions and deep absorption troughs \\citep{reimers}. When contrasted with the low \\he\\ opacity measured at $z \\sim 2.5$ \\citep{dkz,gak,zheng} these findings indicate a sudden increase in optical depth. While the large opacity at $z \\sim 3.2$ is conceivably related to a fluctuation in the number of local ionizing sources, it may instead (e.g., if confirmed in multiple additional independent sightlines) more generally signal an epoch in which intergalactic helium is not fully ionized. Recently, a similar pattern for intergalactic hydrogen has been observed in the $z > 6$ quasars. \\citet{becker} and \\citet{fan} report a \\h\\ optical depth of at least $\\tau=15$ at $z\\sim 6.2$, signaling the late stages of reionization of the intergalactic medium (IGM). The recent {\\it WMAP} results \\citep{map} suggest that the IGM reionization may have started at $z \\sim 17$, and there may have been a second stage of hydrogen reionization that took place at $z\\sim 6$ \\citep{cen}.  The reionization of intergalactic hydrogen and helium are cosmic milestones. As the early universe cooled and evolved, neutral hydrogen atoms in the IGM gradually became bathed in the ionizing ultraviolet background radiation produced by the first stars and quasars. With increasing UV background radiation, the IGM eventually became fully ionized. Atomic hydrogen and helium are affected by the same ionizing UV radiation in the distant universe -- but because the \\he\\ \\lya\\ transition is related to the number of ionizing photons at 4 Rydbergs (generally smaller than the number at 1 Rydberg), the the reionization of intergalactic helium is expected to occur later than hydrogen, approximately at $z\\sim 3-5$ \\citep{haardt,madau2,loeb}.  Indirect evidence \\citep{theuns} suggests that the final stage of \\he\\ reionization may take place at $z \\sim 3.2-3.4$, but only a very small number of independent sightlines have thus far been identified that are suitable to probe \\he\\ even to $z\\sim 3$. Moreover, each such sightline is potentially susceptible to systematic bias; these sightlines are, by selection, the most unobscured of all known quasar sightlines to high redshift. It is therefore important to find additional sightlines (to average over possible systematics in global conclusions), and also to extend \\he\\ studies to higher redshift. Higher redshift sightlines are, of course, potentially the most useful as each such object may provide a measure of \\he\\ absorption spanning the full redshift range likely to be most interest. Our knowledge of the intergalactic helium has been derived from only four lines of sight, and only two are between $3.0 < z < 3.3$. More such examples are needed to establish the statistics and evolution of intergalactic helium at redshift $z>3$. The discovery of any bright, UV-unobstructed quasar in this redshift range is therefore interesting.  We report here the detection of a \\he\\ absorption feature in the spectrum of quasar SDSS J234625.67-001600.4 (hereafter abbreviated as SDSSJ2346-0016), the highest redshift \\he\\ feature found to date. This new quasar sightline potentially will allow a measure of \\he\\ absorption from $z=2.8$ to $z=3.5$ in a single object, efficiently sampling the full redshift range commonly suggested for the epoch of helium reionization. We describe the method of selection, preliminary observations from {\\it HST}/STIS, and the probability for further discovery of similar objects. ",
        "conclusions": ""
    },
    "0311/astro-ph0311125_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro}  Deuterium-bearing molecules are important as probes of the very cold phases of molecular clouds prior to star formation.  The \\hhdp\\ ion is especially important as tracer of \\hhh, the primary ion in dense molecular clouds, which does not have a dipole moment and hence no pure rotational lines.  In addition, at low temperatures ($\\ltsim10$~K), \\hhdp\\ has the ability to channel D atoms from their main reservoir, HD, into heavier species.  This process leads to abundance ratios of DCO$^+$/HCO$^+$ and N$_2$D$^+$/N$_2$H$^+$ of $\\sim$10$^{-3}-10^{-1}$ observed towards dense cores, much larger than the elemental D/H ratio of $\\sim$10$^{-5}$.  Recent observations of multiply deuterated H$_2$CO, CH$_3$OH, H$_2$S and NH$_3$ (see Ceccarelli, this volume) suggest that under extreme conditions, a significant fraction of D may be transferred to heavy molecules. We wish to quantify the role of \\hhdp\\ in this process, and compare with other processes such as grain surface reactions (Caselli, this volume).   The ground-state $1_{01}$--$0_{00}$ transition of para-\\hhdp\\ at 1370~GHz will be a prime target for GREAT on SOFIA. However, with the upper energy level 65~K above ground, this line will only be excited in relatively warm ($\\gtsim$20~K) regions, where chemical fractionation is ineffective. For colder sources, the $1_{10}$--$1_{11}$ ground-state transition of ortho-\\hhdp\\ at 372~GHz is more suitable, which can be observed from the ground under good conditions.  At the low temperatures ($\\ltsim$10~K) and high densities ($>10^5$~cm$^{-3}$) of pre-stellar cores, reactive collisions with ortho--H$_2$ keep the ortho-para ratio of \\hhdp\\ at $\\sim$1, orders of magnitude above the LTE value \\cite{pagani92,gerlich02}.  Until 2002, only two detections of \\hhdp\\ had been obtained, which indicated abundances of $\\sim$10$^{-11}$--10$^{-12}$ towards Class~0 objects \\cite{h2d+99,stark03}.  In October 2002, we~\\cite{h2d+03} observed strong \\hhdp\\ emission towards the pre-stellar core L~1544, and derived an abundance of $\\sim$1$\\times 10^{-9}$ in the central $\\approx$20$''$. Such a high abundance suggests that in this region, all CNO-bearing species are depleted onto dust grains, a situation explored in more detail by Walmsley (this volume).  Data taken in June 2003 at the CSO indicate that the same phenomenon takes place in at least five other pre-stellar cores. More observations are scheduled for December 2003.  These data will be presented in a forthcoming paper. Here we investigate the line profile of \\hhdp\\ in L~1544, and derive its abundance outside the central region. ",
        "conclusions": ""
    },
    "0311/astro-ph0311313_arXiv.txt": {
        "abstract": "{ We present new millimeter and infrared spectroscopic observations towards the radio nebula G10.0$-$0.3, which is powered by the wind of the Luminous Blue Variable star \\lbv \\, also closely associated with the soft gamma-ray repeater \\sgr, and believed to be located in the giant Galactic \\hii\\ complex W31.  Based on observations of $CO$ emission lines and $NH_3$ absorption features from molecular clouds along the line of sight to G10.0$-$0.3, as well as the radial velocity and optical extinction of the star powering the nebula, we determine its distance to be 15.1$^{+1.8}_{-1.3}$ kpc in agreement with Corbel et al. (1997).  In addition, this strengthens the association of \\sgr\\ with a massive molecular cloud at the same distance. All soft gamma-ray repeaters  with precise location are now found to be associated with a site of massive star formation or molecular cloud. We also show that W31 consists of at least two distinct components along the line of sight.  We suggest that G10.2$-$0.3 and G10.6$-$0.4 are located on the $-$30 \\kms\\ spiral arm at a distance from the Sun of 4.5 $\\pm$ 0.6 kpc and that G10.3$-$0.1 may be associated with a massive molecular cloud at the same distance as the LBV star, i.e. 15.1$^{+1.8}_{-1.3}$ kpc, implying that W31 could be decomposed into two components along the line of sight. ",
        "introduction": "\\begin{figure*}[hbt!] \\centerline{\\epsfig{figure=W31.eps, width=18cm}} \\caption{ Representation of the location of each object introduced in this paper and their relative positions on the plane of the sky. {\\bf a)} Large radio map at 408 MHz of the field of W31, adapted from Shaver \\& Goss (1970). The location of the main components of W31 are labeled, as well as the radio nebula G10.0$-$0.3 (produced by the wind of \\lbv ). {\\bf b)} Radio image of the radio nebula G10.0$-$0.3 at 1.4 GHz from Kulkarni et al. (1994). {\\bf c)} K' (2.15 $\\mu$m) near-infrared image ($\\sim$ 2 \\arcmin\\ $\\times$ 2 \\arcmin) of the core of G10.0$-$0.3. The location of \\lbv\\ and \\sgr\\ are indicated by a cross. The ISOCAM LW3 (12-18 $\\mu$m) contours are superposed on this image in order to highlight the spatial extend of the cluster of massive stars around \\sgr\\ (adapted from Fuchs et al. 1999).  } \\end{figure*}   W31 is one of the largest \\hii\\ complexes in the Galaxy, with intense star-forming regions that have been observed from radio to near-infrared wavelengths (e.g. Ghosh et al. 1989; Blum et al. 2001; Kim \\& Koo 2002). At low spatial resolution, W31 appears as three main extended \\hii\\ regions: G10.2$-$0.3, G10.3$-$0.1 and G10.6$-$0.4 (Shaver \\& Goss 1970) (Fig. 1).  The radio nebula G10.0$-$0.3 lies within W31 on the plane of the sky (Fig. 1), and has drawn considerable attention in recent years due to the intriguing objects nearby.  At one time, G10.0$-$0.3 was suggested to be a plerionic supernova remnant powered by a rare soft gamma-ray repeater, SGR~1806$-$20 (Kulkarni \\& Frail 1993; Kouveliotou et al. 1998). SGR~1806$-$20, in turn, was thought to be associated with an almost equally rare luminous blue variable (LBV) star (van Kerkwijk et al. 1995) which lies at the time-variable (in both flux and morphology) core of this nebula (Vasisht, Frail, \\& Kulkarni 1995). However, the revised Inter-Planetary Network (IPN) localization of SGR 1806$-$20 provides a position inconsistent (see Fig. 1) with that of the LBV star and radio core of G10.0$-$0.3 (Hurley et al. 1999), though the LBV position is consistent with the radio core within the uncertainties ($\\sim 2 \\arcsec$).  Recent {\\it Chandra} and infrared observations confirm that the SGR lies $\\sim 12 \\arcsec$ away from the LBV and radio core (Eikenberry et al. 2001; Kaplan et al. 2002). Furthermore, Gaensler et al. (2001) argue that G10.0$-$0.3 is not a supernova remnant at all, but is rather powered by the tremendous wind of the LBV star at its core.  Infrared observations of the field of SGR 1806$-$20 reveal that the LBV star is not alone, but appears to be part of a cluster of embedded, hot, luminous stars (\\cite{fuc99}), and the IPN position for SGR 1806$-$20 is consistent with membership in the star cluster (Eikenberry et al. 2001).  Given this somewhat confusing history, we take a moment to summarize our current understanding of G10.0-0.3: 1. G10.0-0.3 is a radio nebula (NOT supernova remnant) with emission powered by the LBV star spatially coincident with its core. 2.  The LBV star is part of a cluster of luminous stars embedded in a molecular cloud.  3. \\sgr \\ is likely another member of this cluster of stars, and is spatially distinct from the LBV star. We plot in Fig. 1 the various objects (and their relative location on the plane of the sky) that we will discuss in this paper.  Because of the proximity of these unusual objects, the distance to them can provide significant insight into their physical properties, giving this measurement particular importance.  Corbel et al. (1997) proposed a distance estimate based on observations of molecular clouds along the line of sight.  They used $CO$ spectroscopy to estimate the hydrogen column density towards G10.0$-$0.3, and from that an absorption column density as a function of distance.  Taking the measured X-ray absorption towards SGR 1806$-$20 and an estimate of the optical extinction to the LBV star, they concluded that G10.0$-$0.3, \\sgr\\ and W31 lie $14.5 \\pm 1.4$ kpc from the Sun.  \\begin{figure*}[hbt] \\centerline{\\epsfig{figure=lbv_co.eps,width=18cm}} \\caption{\\co\\ (top) and \\tco\\ (bottom) spectrum in the direction of \\lbv \\ and \\sgr.  Antenna temperatures have been converted into main beam brightness temperatures. Each molecular cloud is labeled by its name. } \\end{figure*}  However, this situation seemed complicated by newer infrared stellar spectroscopy by Blum et al. (2001) and radio/millimeter observations by Kim \\& Koo (2002).  Blum et al. (2001) present infrared spectra of members of a star cluster in the \\hii\\ region G10.2$-$0.3, also within W31 on the plane of the sky.  Based on the spectra, they derive spectral/luminosity classes and extinctions for the stars, which, combined with infrared photometry, place them and G10.2$-$0.3 unambiguously at a distance of $d \\simeq 3.4$ kpc. However, the extinction towards G10.2$-$0.3 is much smaller ($\\Delta A_V \\approx 15$ mag) than the extinction towards \\lbv \\ and the star cluster close to \\sgr\\ (Eikenberry et al. 2001) and the correlated X-ray absorption towards \\sgr\\ (Mereghetti et al. 2000).  Furthermore, 21-cm continuum and several molecular line maps of the region (Kim \\& Koo 2002) show that G10.0$-$0.3 is rather separated from the primary components of W31 (G10.2$-$0.3 and G10.3$-$0.1).  Thus, as of this writing, there is no definitive physical linkage between \\gc\\ and the major components of W31.  Together, these imply that G10.0$-$0.3 may lie at a different distance along the line of sight than G10.2$-$0.3.  It is this new, more-complicated situation which motivates us to reconsider the distance to G10.0$-$0.3 (and by extension \\lbv\\ and \\sgr ) in the light of previous work by Corbel et al. (1997) and Blum et al. (2001) as well as new observations.  In this paper, we present newer, higher velocity -resolution $CO$ spectroscopy towards \\lbv \\ and its associated radio nebula \\gc, and also towards 2 (out of 3) of the brightest \\hii\\ regions of W31. We also present the $NH_3$ absorption spectrum originally mentioned in Corbel et al. (1997).  We then add higher-resolution infrared spectroscopic observations of the LBV star at the center of G10.0$-$0.3.  In Section 2, we present the observations and data reduction as well as the main results from our observations.  In Section 3, we use these data to derive a robust distance estimate for \\lbv \\ and \\gc. We then discuss in some detail the implications of this distance for \\sgr \\ and \\lbv , as well as the structure of W31 as a whole .  Finally, in Section 4 we present our conclusions. ",
        "conclusions": "  \\begin{itemize}  \\item Based on $NH_3$ absorption from MC73 and the velocity of \\lbv \\ , we unambiguously constrain the distance to G10.0$-$0.3 and \\lbv \\ to be in the range of 13.2 - 21.5 kpc.  \\item Combining this constraint with $CO$ observations of molecular clouds along the line of sight, we further refine the distance measurement to G10.0$-$0.3 and \\lbv \\ to be 15.1$^{+1.8}_{-1.3}$ kpc.  \\item This distance estimate is confirmed by the consistency between the measured extinction towards \\lbv \\ and the extinction from the molecular clouds along the line of sight inferred from their $CO$ spectra.  \\item Based on their distinct extinctions and the newly-resolved parental molecular clouds (MC13A and MC13B), we conclude that the stellar cluster in G10.2$-$0.3 lies in the foreground to the cluster containing \\lbv \\ .  This shows that W31 consists of at least 2 discrete components along the line of sight.  \\item We suggest that G10.2$-$0.3 and G10.6$-$0.4 are located on the $-$30 \\kms\\ spiral arm at a distance from the Sun of 4.5 $\\pm$ 0.6 kpc and  \\item We also suggest that G10.3$-$0.1 may be associated with a massive molecular cloud at the same distance as \\lbv \\ (15.1$^{+1.8}_{-1.3}$ kpc).  \\item We confirm that \\sgr\\ is located at a distance from the Sun of 15.1$^{+1.8}_{-1.3}$ kpc and that it is associated with a very massive molecular cloud. All SGRs with precise location are associated with a site of massive star formation.  \\end{itemize}"
    },
    "0311/astro-ph0311580_arXiv.txt": {
        "abstract": "We discuss in detail the evolutionary properties of low mass stars $\\rm (M~\\le~1~M_\\odot)$ having metallicity lower than $\\rm Z=10^{\\rm - 6}$ from the pre- main sequence up to (almost) the end of the Asymptotic Giant Branch phase. We also discuss the possibility that the large [C,N/Fe] observed on the surface of the most Iron poor star presently known, HE0107-5240, may be attributed to the autopollution induced by the penetration of the He convective shell into the H rich mantle during the He core flash of a low mass, very low metallicity star. On the basis of a quite detailed analysis, we conclude that the autopollution scenario cannot be responsible for the observed chemical composition of HE0107-5240. ",
        "introduction": "The idea that stars born shortly after the Big Bang are still {\\it alive} has such a profound impact on the human emotionality that their evolutionary properties have been investigated in spite of the fact that their formation in an environment deprived of metals (Z=0) is at least not favored because of the lack of a strong coolant.  From a purely theoretical ground the fragmentation of primordial clouds is still largely debated and unclear, so that the current literature on the subject still ranges from papers in which essentially massive stars form \\citep{abn00} to papers in which a significant production of low mass stars $\\rm (M\\sim1~M_\\odot)$ is possible \\citep{nu02}. However the continuous discovery of progressively more metal poor low mass stars constitutes a very good excuse to trust the primordial formation of low mass stars and hence to study their evolutionary properties.  The first computation of the Main Sequence (MS) properties of primordial (Z=0) low mass stars has been performed, as far as we know, by \\cite{cp75} while the first evolution up to the He core flash has been published by \\cite{da82}. In the latter paper all the key properties of the evolution of these stars were already clearly outlined. In particular, the presence of a Red Giant Branch (RGB) powered mainly by the PP chain rather than by the CNO cycle together to its natural consequences: a dimmer core mass- luminosity relation and a more external (in mass) ignition of the He burning (due to the slower advancing of the H shell). In the same paper it was also suggested the possibility of the ingestion of protons in the \"active\" He (and C rich) convective shell and hence the possibility to produce stars showing very large surface CN abundances. In a following paper \\cite{gd83} tried to perform similar computations but the formation of a (fast growing) convective core towards the end of the central H burning phase forced them to stop the computations. Such a strong instability was not encountered in any further study. \\cite{fih90} fully reanalyzed the evolution of a primordial $\\rm 1~M_\\odot$ from the MS up the He flash and described in detail the evolution of this star in comparison with stars of the same mass but having a larger metallicity, i.e. $\\rm Z=10^{- 4}$ and $\\rm Z=2\\times10^{-2}$. Since we will not tempt to rediscover once again the \"funny\" properties of low mass zero metallicity stars, we refer the reader interested in a deep knowledge of the properties of these stars to that paper. One of the main results found by \\cite{fih90} was that the outer border of the He convective shell, that forms at the He core flash, penetrates, unambiguously, the H-rich mantle. Such a result was discussed and interpreted as due to the combined effects of the unusually external location of the He ignition and to the low entropy barrier existing at the He-H interface as well. The properties that follow the ingestion of the protons in the He convective shell were firstly discussed by the same group in a following paper \\citep{hif90}. In particular they showed that the convective shell splits in two (at the mass location where the local burning timescale becomes comparable with the mixing timescale), the He convective shell dumps out, the H convective shell eventually merges with the convective envelope an hence that the envelope of the star becomes strongly enriched in the products of both the H and He burnings. In a third paper, \\cite{fii00} recomputed several models of low mass zero (and non zero) metallicity and mainly confirmed their previous findings. By the way, tabulations of evolutionary tracks of extremely metal poor stellar models have been published in those years by \\cite{cc93} and \\cite{cct96}: the He core flash properties of zero metallicity stars were not addressed at all in those papers. \\cite{wcss00} recomputed a few low mass zero metallicity models and concluded very firmly that there isn't {\\it any evidence} of mixing between the He convective shell that develops at the He core flash and the H rich envelope. In the following year the same group \\citep{scsw01} found the mixing between the He convective shell and the H-rich mantle at the He core flash of a $\\rm 1~M_\\odot$ (Y=0.23,Z=0); they also studied the dependence of the occurrence of the mixing on several parameter (equation of state, initial He abundance, initial mass, diffusion efficiency, external pollution and opacity). Their main result was that the occurrence of this mixing is very robust with respect to all the changings but for the dependence on the initial He abundance (an increase of $\\rm Y_{\\rm ini}$ from 0.23 to 0.24 inhibits the mixing) and on the possible pollution of the envelope (that also inhibits the mixing). For the smaller masses ($\\rm 0.8~M_\\odot$) they always found the mixing to occur! In a further paper \\citep{sscw02}, they extended the tests to the dependence on the metallicity and found that the mixing occurs up to a metallicity of $\\rm Z<10^{\\rm -6}$ for masses of the order of $\\rm 0.8~M_\\odot$. However, other similarly recent computations by \\cite{mgcw01} and \\cite{sll02} did not find any mixing between the He convective shell and the H rich mantle at the onset of the He core flash.  In the last years we have continuously worked on the evolution of zero metallicity stars \\citep{cdls01,cl02,lc02,blc03,lcb03} in a wide mass range that goes from the low masses up to the more massive ones ($\\rm 100~M_\\odot$). We have discussed over the years the properties of both the primordial AGB and core collapse supernovae. In this paper we present our low mass models of extremely low metallicity and discuss the fit to the most Iron poor star presently known, HE0107-5240 \\citep{chris02}. The paper is organized as follows. In section 2 we briefly remind the current setup of the evolutionary code FRANEC (release 4.98); section 3 is devoted to the description of the properties of our \"reference\" stellar model, i.e. a $\\rm 0.8~M_\\odot$ having Y=0.23 and Z=0. Additional test models are discussed in section 4 while the fit to HE0107-5240 is discussed in section 5. ",
        "conclusions": ""
    },
    "0311/astro-ph0311063_arXiv.txt": {
        "abstract": "We describe our explicit Lorentz-invariant solution of the Einstein and null geodesic equations for the deflection experiment of 2002 September 8 when a massive moving body, Jupiter, passed within 3.7' of a line-of-sight to a distant quasar.  We develop a general relativistic framework which shows that our measurement of the retarded position of a moving light-ray deflecting body (Jupiter) by making use of the gravitational time delay of quasar's radio wave is equivalent to comparison of the relativistic laws of the Lorentz transformation for gravity and light. Because, according to Einstein, the Lorentz transformation of gravity field variables must depend on a fundamental speed $c$, its measurement through the retarded position of Jupiter in the gravitational time delay allows us to study the causal nature of gravity and to set an upper limit on the speed of propagation of gravity in the near zone of the solar system as contrasted to the speed of the radio waves. In particular, the $v/c$ term beyond of the standard Einstein's deflection, which we measured to 20\\% accuracy, is associated with the aberration of the null direction of the the gravity force (``aberration of gravity'') caused by the Lorentz transformation of the Christoffel symbols from the static frame of Jupiter to the moving frame of observer. General relativistic formulation of the experiment identifies the aberration of gravity with the retardation of gravity because the speed of gravitational waves in Einstein's theory is equal to the speed of propagation of the gravity force. We discuss the misconceptions which have inhibited the acceptance of this interpretation of the experiment.  We also comment on other interpretations of this experiment by Asada, Will, Samuel, Pascual-S\\'anchez, and Carlip and show that their ``speed of light\" interpretations confuse the Lorentz transformation for gravity with that for light, and the fundamental speed of gravity with the physical speed of light from the quasar. For this reason, the ``speed of light\" interpretations are not entirely consistent with a retarded Li\\'enard-Wiechert solution of the Einstein equations, and do not properly incorporate how the phase of the radio waves from the quasar is perturbed by the retarded gravitational field of Jupiter.  Although all of the formulations predict the same deflection to the order of $v/c$, our formulation shows that the underlying cause of this deflection term is associated with the aberration of gravity and not of light, and that the interpretations predict different deflections at higher orders of $v/c$ beyond the Shapiro delay, thus, making their measurement highly desirable for deeper testing of general relativity in future astrometric experiments like Gaia, SIM, and SKA. \\vspace{1cm} ",
        "introduction": "The recent Very Long Baseline Interferometric (VLBI) experiment of the ultra-precise measurement of the relativistic deflection of a quasar's radio waves (light) by the gravitational field of Jupiter \\cite{kop2001,fk-apj} has captivated the keen interest of relativists who were trying to deeply understand gravitational physics.  The experiment and its physical interpretation has illustrated some of the important subtleties in the understanding, the approaches, and the mathematical techniques of general relativity (GR).  The experimental goal was the measurement of the (non-radial or tangential) deflection component associated with the retarded position of Jupiter caused by its orbital motion and mathematically connected to the finite value of the fundamental speed entering in front of the time derivatives from the metric tensor in the gravity field and null geodesic equations (that is, the fundamental constant $c$ in the Christoffel symbols and the curvature tensor). This fundamental speed of gravity causes the retardation in propagation of gravity force from a moving massive object which can be measured in the light-ray deflection experiments. In this paper we use general theory of relativity and denote the fundamental speed of gravity by $c$. Label $c_g$ for the speed of gravity is understood as a parameter in the experimental fitting of observations to the Einstein theory. The symbol $c$ is frequently called as the speed of light. However, one must clearly separate the relativistic effects caused by the physical speed of light from those associated with the fundamental speed of gravity in the Eiunstein equations. We shall show in the present paper how to distinguish these effects to avoid ambiguity.  Einstein's general principle of relativity assumes that the gravity field obeys to the special principle of relativity and transforms from one inertial frame to another in accordance with the Lorentz transformation of the Minkowski space-time. Hence, it postulates that in a geometrized system of units \\cite{mtw} the speed of gravity $c=1$ in any frame of reference. We measured the impact of the retarded gravitational field of Jupiter on light deflection (as well as the instantaneous Shapiro deflection of $1100~\\mu$arcsec) of $50\\pm 9~\\mu$arcsec, in excellent agreement with the GR prediction of $51~\\mu$arcsec based on assumption that $c=1$ irrespectively of the reference frame.  We have interpreted these result as a measurement of the fundamental speed of gravity $c=1.06\\pm 0.21$ and confirmation that gravity interacts with moving particles (photons) not instantaneously.  However, at present there are four other interpretations, which do not agree among themselves, concerning the results of the experiment \\footnote{Recent publication of Frittelli \\cite{frit1} agrees with our understanding of the gravitational physics of the Jovian deflection experiment}. \\begin{itemize} \\item the experiment measured the speed of the radio waves from the quasar in the static gravitational field of Jupiter \\cite{ass}, \\item the experiment measured the aberration of the radio waves from the quasar \\cite{will-apj,carlip}, \\item the experiment measured no significant property of light or of gravity \\cite{samuel}, \\item the experiment measured the light-time delay discovered by R{\\o}mer in 1676 \\cite{pask} \\end{itemize} \\vskip 0.1cm\\noindent  In this paper, we streamline our previous theoretical calculations by working out an explicit Lorentz-invariant Li\\'enard-Wiechert solution of the Einstein gravity field equations in \\S 2.  The VLBI measurement of the phase of a radio wave, which has propagated in the field of a moving body (Jupiter), is described in \\S 3.  We demonstrate that the fundamental speed of the characteristics of the gravity field equations (which we call for brevity ``the speed of gravity\") and the aberration of gravity force are indeed associated with the experiment.  The major misconceptions about the experiment and properties of gravity in general relativity are summarized in \\S 4.  Finally, in \\S 5 we outline why these alternative interpretations of the experimental results differ with our general relativistic interpretation, even though all interpretations predict the same angular deflection result to order $v/c$ beyond the Shapiro delay. Appendix describes how light propagates in a simple bi-metric model of gravity with two metrics, and provides a clear evidence that a light-ray deflecting body deflects light from its retarded position connected to observer by a null line lying on the gravity null cone that is because gravity propagates with finite speed. ",
        "conclusions": "We have developed a Lorentz-invariant general-relativistic theory of light-deflection experiments associated with moving bodies in the solar system. The formulation is based on the retarded Li\\'enard-Wiechert solution of the Einstein equations, and explicitly calculates the effect on the phase of the radio wave from the quasar caused by the finite speed of gravity which is a primary fundamental constant in Einstein's field equations. Our post-Newtonian expansion of the retarded coordinate of the light-ray deflecting body shows that the retardation of gravity effect after making its post-Newtonian expansion, is present already in linear aberration-like terms of order $v/c$ beyond the Shapiro delay; hence, we interpret the $v/c$ term as caused by the aberration of gravity in the expression for the gravitational force (\\ref{j}) exerted by Jupiter on the photons.  We have used the concept of the gravito-magnetic field caused by translational motion of Jupiter as an aid in understanding the gravitational phenomena associated with this experiment.  The VLBI measurements confirmed that the deflection is associated with the retarded position of Jupiter, with the speed of gravity equal to that of light to 20\\% accuracy \\cite{fk-apj} and this retardation is not a preferred frame effect associated with the PPN parameter $\\alpha_1$.  Other contradictory interpretations of the experimental results are associated with: (1) an analysis of the Einstein equations which are not fully Lorentz-invariant; (2) misconceptions about the experimental method and analysis; (3) misunderstanding of the nature of the interaction of the gravitational field from moving objects with the {\\it phase} of the quasar radio signal which is a null hypersurface. According to GR the phase is influenced in {\\it all} orders of $v/c$ by the propagating gravitational field of Jupiter and this perturbation reveals itself in the linearized order of $v/c$ as the aberration of gravity force of Jupiter due to the finite speed of gravity.  We thank the Eppley Foundation for Research (award 002672) for support.  \\appendix"
    },
    "0311/astro-ph0311533_arXiv.txt": {
        "abstract": "s{ Recent results of the SLOTT-AGAPE and POINT-AGAPE collaborations on a search for microlensing events in direction of the Andromeda galaxy, by using the pixel method, are reported.  The detection of 4 microlensing events, some likely to be due to self--lensing, is discussed. One microlensing light curve is shown to be compatible with a binary lens. The present analysis still does not allow us to draw conclusions on the MACHO content of the M31 galaxy.} ",
        "introduction": "Since Paczy\\'nski's original proposal\\cite{pacz} gravitational microlensing has been proben to be a powerful tool for the detection of the dark matter component in galactic haloes in the form of MACHOs. Searches in our Galaxy towards LMC\\cite{macho,eros} show that up to 20\\% of the halo could be formed by objects of around $M \\sim 0.4\\,M_\\odot$, but these results are still debated\\cite{jetzer}.  Searches towards M31, nearby and similar to our Galaxy, have also been proposed\\cite{crotts92,agape93,jetzer94}. This allows to probe a different line of sight in our Galaxy, to globally test M31 halo and, furthermore, the high inclination of the M31 disk is expected to provide a strong signature (spatial distribution) for halo microlensing signals.  Along a different direction, results of a microlensing survey towards M87, where one can probe both the M87 and the Virgo cluster haloes, have also been presented\\cite{m87}.  For extragalactic targets, due to the distance, the sources for microlensing signals are not resolved. This claims for an original technique, the \\emph{pixel method}, the detection of flux variations of unresolved sources\\cite{agape97,agape99,tom96}, the main point being that one follows flux variations of every pixel in the image instead of single stars.  I review here the results from two different survey of M31 aimed at the detection of microlensing events, carried out by the  SLOTT-AGAPE\\cite{mdm1,mdm2} and by the POINT-AGAPE collaborations\\cite{point01,point03}. The WeCapp\\cite{wecapp,wecapp03} and the MEGA\\cite{mega} collaborations have also presented a handful of microlensing events. ",
        "conclusions": "As a general outcome of the results presented, we stress that the detection of microlensing events towards M31 is now established. The open issue to be still explored is the study of the M31 halo fraction in form of MACHOs. With respect to this analysis, the events detected so far are all compatible with stellar lenses, but the MACHO hypothesis is still open, and we recall that the analysis for the INT data is still not concluded (besides a third year data, variations with $\\Delta R_{max}>21$ have still to be studied). Once this analysis completed, it is the necessary to ``weight'' it with an efficiency study of the pipeline of detection before meaningfully comparing its results with the prediction of a Monte Carlo simulation. This should eventually allow us to draw firm conclusions on the halo content in form of MACHO of M31.   \\begin{figure}[ht] \\centerline{\\epsfxsize=4.0in \\epsfbox{fig3.eps}} \\caption{Three years data light curves for the 4 POINT-AGAPE microlensing events.} \\label{int4} \\end{figure}"
    },
    "0311/astro-ph0311019_arXiv.txt": {
        "abstract": "We review recent theoretical results on the formation of the first stars in the universe, and emphasize related open questions. In particular, we discuss the initial conditions for Population~III star formation, as given by variants of the cold dark matter cosmology. Numerical simulations have investigated the collapse and the fragmentation of metal-free gas, showing that the first stars were predominantly very massive. The exact determination of the stellar masses, and the precise form of the primordial initial mass function, is still hampered by our limited understanding of the accretion physics and the protostellar feedback effects. We address the importance of heavy elements in bringing about the transition from an early star formation mode dominated by massive stars, to the familiar mode dominated by low mass stars, at later times. We show how complementary observations, both at high redshifts and in our local cosmic neighborhood, can be utilized to probe the first epoch of star formation. ",
        "introduction": "The first stars to form out of un-enriched, pure H/He gas marked the crucial transition from a homogeneous, simple universe to a highly structured, complex one at the end of the cosmic dark ages. Extending the familiar scheme of classifying stellar populations in the local universe (Baade 1944) to the extreme case of zero metallicity, the first stars constitute the so-called Population~III (e.g., Bond 1981, Cayrel 1986, Carr 1987, 1994). The quest for Population~III stars has fascinated astronomers for many decades, going back to the first tentative ideas of Schwarzschild \\& Spitzer (1953). Recently, the subject has attracted an increased interest, both from a theoretical and observational perspective. New empirical probes of the high redshift universe have become available, and our ability to carry out sophisticated numerical simulations has improved dramatically. In this review, we attempt to summarize the current state of this rapidly evolving field.  The first generation of stars had important effects on subsequent galaxy formation (see Carr et al. 1984 for an early overview). On the one hand, Population~III stars produced copious amounts of UV photons to reionize the universe (e.g., Tumlinson \\& Shull 2000, Bromm et al. 2001b, Schaerer 2002, 2003, Tumlinson et al. 2003, Venkatesan \\& Truran 2003, Venkatesan et al. 2003). Recently, the {\\it Wilkinson Microwave Anisotropy Probe (WMAP)} has observed the large-angle polarization anisotropy of the cosmic microwave background (CMB), thus constraining the total ionizing photon production from the first stars (e.g., Cen 2003a,b, Ciardi et al. 2003, Haiman \\& Holder 2003, Holder et al. 2003, Kaplinghat et al. 2003, Kogut et al. 2003, Sokasian et al. 2003a,b, Wyithe \\& Loeb 2003a,b). The large optical depth to Thomson scattering measured by {\\it WMAP} has been interpreted as a signature of a substantial early activity of massive star formation at redshifts $z\\ga 15$ (see Section 5.1). Secondly, the supernova (SN) explosions that ended the lives of the first stars were responsible for the initial enrichment of the intergalactic medium (IGM) with heavy elements (e.g., Ostriker \\& Gnedin 1996, Gnedin \\& Ostriker 1997, Ferrara et al. 2000, Madau et al. 2001, Mori et al. 2002, Thacker et al. 2002, Bromm et al. 2003, Furlanetto \\& Loeb 2003, Mackey et al. 2003, Scannapieco et al. 2003, Wada \\& Venkatesan 2003, Yoshida et al. 2004). An intriguing possibility unique to zero-metallicity massive stars is the complete disruption of the progenitor in a pair-instability supernova (PISN), which is predicted to leave no remnant behind (e.g., Barkat et al. 1967, Ober et al. 1983, Bond et al. 1984, Fryer et al. 2001, Heger \\& Woosley 2002, Heger et al. 2003). This peculiar explosion mode could have played an important role in quickly seeding the IGM with the first metals (see Section 4.2).   To place the study of the first stars into the appropriate cosmological context, one has to ask (see Section 2): {\\it When and how did the cosmic dark ages end?} The dark ages denote the cosmic time between the emission of the CMB half a million years after the big bang, when the CMB photons shifted into the infrared, so that the universe would have appeared completely dark to a human observer, and the formation of the first sources of light (see Barkana \\& Loeb 2001, Loeb \\& Barkana 2001, Miralda-Escud\\'{e}~2003 for comprehensive reviews of this cosmic epoch). Historically, the evocative -- and now widely used -- dark age analogy was first introduced into the literature briefly after the ascent of non-baryonic, cold dark matter scenarios (e.g., Sargent 1986, Rees 1993). In the context of popular cold dark matter (CDM) models of hierarchical structure formation, the first stars are predicted to have formed in dark matter (DM) halos of mass $\\sim 10^{6}M_{\\odot}$ that collapsed at redshifts $z\\simeq 20-30$ (e.g., Tegmark et al. 1997, Yoshida et al. 2003a). The first quasars, on the other hand (e.g., Umemura et al. 1993, Loeb \\& Rasio 1994, Eisenstein \\& Loeb 1995), are likely to have formed in more massive host systems, at redshifts $z \\ga 10$ (Haiman \\& Loeb 2001, Bromm \\& Loeb 2003a), and certainly before $z\\sim 6.4$, the redshift of the most distant quasar known (Fan et al. 2003). Here, we will not discuss quasar formation in depth, and we refer the reader to the references above.   The most fundamental question about the first stars is how massive they typically were (see Section 3). Results from recent numerical simulations of the collapse and fragmentation of primordial gas clouds suggest that the first stars were predominantly very massive, with typical masses $M_{\\ast}\\ga 100 M_{\\odot}$ (Bromm et al. 1999, 2002, Nakamura \\& Umemura 2001, Abel et al. 2000, 2002). The simulations have taught us the important lesson that regardless of the detailed initial conditions, the primordial gas attains characteristic values of temperature and density. These in turn correspond to a characteristic fragmentation scale, given approximately by the Jeans mass, and are explained by the microphysics of molecular hydrogen (H$_{2}$) cooling (see Section 3.2). The minihalos which host the formation of the first stars have virial temperatures of order $\\sim 1000$~K, below the threshold for cooling due to atomic hydrogen lines ($\\sim 10^4$~K). Therefore, cooling and fragmentation inside these minihalos is possible only via rotational-vibrational transitions of H$_{2}$, which can be excited even at these low temperatures (see Section 2 for details).  Despite the progress already made, many important questions remain unanswered. A particularly important one, for example, is: {\\it What is the functional form of the primordial initial mass function (IMF)?}  Having constrained the characteristic mass scale, still leaves undetermined the overall range of stellar masses and the power-law slope which is likely to be a function of mass (e.g., Yoshii \\& Saio 1986, Larson 1998, Nakamura \\& Umemura 2001, 2002, Omukai \\& Yoshii 2003). In addition, it is presently not known whether binaries or, more generally, clusters of zero-metallicity stars, can form (see Section 3.2). This question has important implications for the star formation process. If Population~III star formation typically resulted in a binary or multiple system, much of the progenitor cloud's angular momentum could go into the orbital motion of the stars in such a system. On the other hand, if an isolated formation mode were to predominate, the classical angular momentum barrier could be much harder to breach (see Larson 2003).  How was the epoch of Population~III star formation eventually terminated (see Section 4)? Two categories of negative feedback effects are likely to be important, the first being radiative and the second chemical in nature (e.g., Ciardi et al. 2000b, Schneider et al. 2002a, Mackey et al. 2003). The radiative feedback consists in soft UV photons produced by the first stars photodissociating the rather fragile H$_{2}$ molecules in the surrounding gas, thus suppressing the corresponding H$_{2}$ cooling (e.g., Haiman et al. 1997, 2000, Omukai \\& Nishi 1999, Ciardi et al. 2000a, Nishi \\& Tashiro 2000, Glover \\& Brand 2001). Massive Population~III stars could then no longer form. There is some debate, however, whether the radiative feedback from the first stars could not have had an overall positive sign (e.g., Ferrara 1998, Machacek et al. 2001, 2003, Oh 2001, Ricotti et al. 2001, Cen 2003b). The formation of H$_{2}$ molecules is catalyzed by free electrons, and whichever process increased their presence would therefore also boost the abundance of H$_{2}$. For example, an early background of X-ray photons, either from accretion onto black holes (e.g., Madau et al. 2003), or from the remnants of Population~III SNe, could ionize the gas and thus increase the abundance of free electrons (e.g., Haiman et al. 2000, Venkatesan et al. 2001, Glover \\& Brand 2003, Machacek et al. 2003). This debate is ongoing, and it is too early for definitive conclusions (see Section 4.1 for details).  The second feedback effect that will be important in terminating the epoch of Population~III star formation is chemical in nature, or, more precisely: it is due to the enrichment of the primordial gas with the heavy elements dispersed by the first SNe (see Section 4.2). Numerical simulations of the fragmentation process have shown that lower mass stars can only form out of gas that was already pre-enriched to a level in excess of the `critical metallicity', estimated to be of order $10^{-4}-10^{-3}$ the solar value (Bromm et al. 2001a; see also Schneider et al. 2002a). Depending on the nature of the first SN explosions, and in particular on how efficiently and widespread the mixing of the metal-enriched ejecta proceeds, the cosmic star formation will at some point undergo a fundamental transition from an early high-mass (Population~III) dominated mode to one dominated by lower mass stars (Population~II).   We conclude our review with a discussion of a few select observational probes of the nature of the first stars (Section 5). Predicting the properties of the first stars is important for the design of upcoming instruments, such as the {\\it James Webb Space Telescope\\footnote{See http:// ngst.gsfc.nasa.gov.} (JWST)}, or the next generation of large ($>10$m) ground-based telescopes.  The hope is that over the upcoming decade, it will become possible to confront current theoretical predictions about the properties of the first sources of light with direct observational data. The increasing volume of new data on high redshift galaxies and quasars from existing ground- and space-based telescopes, signals the emergence of this new frontier in cosmology. ",
        "conclusions": ""
    },
    "0311/astro-ph0311196_arXiv.txt": {
        "abstract": "We have produced a catalog of 378 Galactic O stars with accurate spectral classifications which is complete for $V<8$ but includes many fainter stars. The catalog provides cross-identifications with other sources; coordinates (obtained in most cases from Tycho-2 data); astrometric distances for 24 of the nearest stars; optical (Tycho-2, Johnson, and Str\\\"omgren) and NIR photometry; group membership, runaway character, and multiplicity information; and a web-based version with links to online services. ",
        "introduction": "The advent of internet search engines and databases is producing a revolution in astronomy. Nowadays, it is possible to use a number of web-based tools to extract spectroscopic, photometric, kinematic, or astrometric data from almost any source published in the last century. Also, new large-scale surveys are making huge amounts of data available to any astronomer sitting in front of a computer screen connected to the internet. This has made possible a new kind of astronomer, the electronic one, who can work with previously unanalyzed data and combine it with information from search engines to produce interesting science without ever visiting an observatory.  However, behind the possibilities offered by this avalanche of information lurks a danger. When one compares data from widely diverse sources, the possibility exists that apples and oranges are being mixed: good quality results with poorer ones, data with different systematic errors, or measurements obtained with different calibrations (filter sets, observatories, data reduction packages\\ldots). One of the fields where this danger is more apparent is spectral classification, a subject where lower quality data or the technique of the observer can easily introduce systematic errors. It is in those cases where a critical compilation of a catalog using criteria as uniform as possible becomes important, since different sources can provide quite different classifications. As an example of the problems in this field, if one searches in Simbad for the spectral classification of $\\delta$ Ori A, one of the brightest O stars in the sky, 41 references are returned with spectral classes between O8 and O9.5 and luminosity classes between III and Iab. $\\theta^1$ Ori A, one of the stars in the Trapezium, provides an even more dramatic example: the two results obtained in Simbad are O7 and O8 V n; in reality it is an early-B star (see next section).  A number of lists and catalogs of spectral types for O stars are available in the literature. The works of \\citet{Morgetal55}; \\citet{Hilt56}; and \\citet{Lesh68} provided original accurate spectral classifications for ``northern'' ($\\delta \\gtrsim -20\\degr$) stars which, taken together, constitute a quite uniform catalog for that part of the sky. The southern equivalent can be obtained by merging the spectral classifications obtained by \\citet{Hiltetal69} and \\cite{Garretal77}. Those catalogs, though, are far from complete and most of their spectral classifications were done before the introduction of spectral types O3 \\citep{Walb71b} and O2 \\citep{Walbetal02b}. A different approach was taken by \\citet{Goy80} and by \\citet{Garmetal82}, who built catalogs of 970 and 765 stars, respectively, by compiling spectral classifications from different sources (39 in the first case, 8 in the second one). These catalogs are impressive for its number of stars but, unfortunately, its spectral classifications are far from uniform and complete. This is reflected in \\citet{Goy80} having multiple spectral classifications for many of its stars, and in \\citet{Garmetal82} with many stars not having luminosity classes and with a non-negligible number of errors (see next section for a few of them). Another alternative is to use objective prisms to perform large-scale surveys, of which the most comprehensive one is the ongoing Michigan Spectral Catalog \\citep{HoukSwif99}. The resolution of those surveys tends to be lower and the spectral classifications can have systematic errors with respect to studies performed with slits (in the Michigan catalog a number of stars classified as O are actually of early-B type), making the former less precise. Finally, a recent catalog \\citep{Reed03} takes the broad approach of compiling ``OB'' (meaning O to approximately B2) stars by adding stars on the basis of optical photometry alone. Given good-quality photometry, this method is capable of increasing the sample considerably (\\citealt{Reed03}, indeed, contains over $16\\, 000$ stars) but, given the color degeneracy existent between O and early B stars, it is not of much use to distinguish between them and, given the relative numbers of the two samples, one would expect that most of the stars in that catalog are actually of early-B type.  Our approach in this article will be to build a catalog of O stars with the following conditions. (1) Stars should be selected on the basis of their optical spectroscopy alone. (2) Sources for spectral classification should be accurate and as uniform as possible. (3) Valuable additional information (positions, photometry, distances\\ldots) should be added only if they come from uniform catalogs or if their accuracy can be checked. (4) The catalog should be made accessible through the web and be actively maintained in order for updates to reach the users as fast as possible. With those criteria in mind, we intend this article to be just the first version of the Galactic O-star (GOS) catalog. In the future, if resources are found, we intend to add data from other missions, such as UV photometry and spectroscopy from GALEX and IUE; expand the catalog to include more stars, increasing the completeness limit in magnitude; and collect more data, e.g. producing a spectroscopic library. Updates (and possible errata) will be posted at the web site {\\tt http://www.stsci.edu/\\~{}jmaiz/GOSmain.html}.  This paper is organized as follows. In section 2 we describe how the catalog was constructed. In section 3 we lay out the criteria used for the spectral classification. In sections 4, 5, and 6 we describe how the astrometry, the photometry, and other information in the GOS catalog was compiled, respectively. ",
        "conclusions": ""
    },
    "0311/astro-ph0311475_arXiv.txt": {
        "abstract": "The VIMOS VLT Deep Survey (VVDS) is underway to study the evolution of galaxies, large scale structures and AGNs, from the measurement of more than 100000 spectra of faint objects. We present here the results from the first epoch observations of more than 20000 spectra. The main challenge of the program, the redshift measurements, is described, in particular entering the ``redshift desert'' in the range $1.5<z<3$ for which only very weak features are observed in the observed wavelength range. The redshift distribution of a magnitude limited sample brighter than $I_{AB}=24$ is presented for the first time, showing a peak at a low redshift $z\\sim0.7$, and a tail extending all the way above $z=4$. The evolution of the luminosity function out to $z=1.5$ is presented, with the LF of blue star forming galaxies carrying most of the evolution, with L* changing by more than two magnitudes for this sub-sample. ",
        "introduction": "Understanding how galaxies and large scale structures formed and evolved is one of the major goals of modern cosmology. In order to identify the relative contributions of the various physical processes at play and the associated timescales, a comprehensive picture of the evolutionary properties of the constituents of the universe is needed over a large volume and a large time base. Samples of high redshift galaxies known today reach less than a few thousand at redshifts $1-3$, and statistical analysis suffer from small number statistics, small explored volumes, selection biases, which prevent detailed analysis. In the local universe, large surveys like the 2dFGRS and the SDSS, will contain from $2.5\\times10^5$ up to $10^6$ galaxies to reach a high level of accuracy in measuring the fundamental parameters of the galaxies and AGNs populations. Similarly, we need to gather large numbers of galaxies at high redshifts to accurately describe the various populations in environments ranging from low density to the dense cluster cores, and relate the properties observed at different redshifts to identify the main processes driving evolution.  This goal can only be achieved through massive observational programs assembling galaxies and QSO samples representative of the universe at the various look-back times explored. As the observed galaxies are at large distances and therefore very faint, instruments have to be conceived to combine wide field, high throughput, and high multiplex gain in order to efficiently observe large samples. Multi-object spectrographs are routinely in operation since $\\sim$15 years, and the new generation now in place on the $6-10$m telescopes, like DEIMOS on the Keck-10m or VIMOS on the VLT-8m, allows to explore large volumes of the distant universe through the observations of many tens of thousand of objects.  Because of this leap forward in instrument performances, several large deep surveys are now in progress. We are presenting here the VIMOS VLT Deep Survey (VVDS). It is combining a common observing strategy to assemble more than 100000 redshifts in 3 large galaxy datasets observed over more than 16deg$^2$ in 5 equatorial fields, selected in magnitude from $I_{AB}\\leq22.5$ to $I_{AB}\\leq25$. The general strategy of the survey has been presented in e.g. Le F\\`evre et al. (2001). ",
        "conclusions": "The VVDS is on-going. The preliminary first epoch results show:\\\\ -- a mean redshift lower than anticipated for a $I_{AB}\\leq24$ sample, despite a high redshift tail reaching $z\\sim5$. \\\\ -- a strong evolution of the global LF out to $z\\sim1.5$, with the blue star forming galaxies identified as the population carrying most of the evolution, while the LF of the early type (``redder'') galaxies does not seem to evolve much.  Many different properties of the sample are being currently being explored."
    },
    "0311/astro-ph0311369_arXiv.txt": {
        "abstract": "Recent advances in understanding of magnetohydrodynamic (MHD) turbulence call for substantial revisions in the picture of cosmic ray transport. In this paper we use recently obtained scaling laws for MHD modes to calculate the scattering frequency for cosmic rays in the ISM. We consider gyroresonance with MHD modes (Alfv\\'{e}nic, slow and fast) and transit-time damping (TTD) by fast modes. We provide calculations of cosmic ray scattering for various phases of interstellar medium with realistic interstellar turbulence driving that is consistent with the velocity dispersions observed in diffuse gas. We account for the turbulence cutoff arising from both collisional and collisionless damping. We obtain analytical expressions for diffusion coefficients that enter Fokker-Planck equation describing cosmic ray evolution. We calculate the scattering rate and parallel spatial diffusion coefficients of cosmic rays for both Alfv\\'{e}nic and fast modes. We conclude that fast modes provides the dominant contribution to cosmic ray scattering for the typical interstellar conditions in spite of the fact that fast modes are subjected to damping. We show that the efficiency of the scattering depends on the plasma $\\beta$ since it determines the damping of the fast modes. We also show that the streaming instability is modified in the presence of turbulence. ",
        "introduction": "Most astrophysical systems, e.g. accretion disks, stellar winds, the interstellar medium (ISM) and intercluster medium are turbulent with an embedded magnetic field that influences almost all of their properties. High conductivity of the astrophysical fluids makes the magnetic fields {}``frozen in'', and influence fluid motions. The coupled motion of magnetic field and conducting fluid holds the key to many astrophysical processes.  The propagation of cosmic rays (CRs) is affected by their interaction with magnetic field. This field is turbulent and therefore, the resonant interaction of cosmic rays with MHD turbulence has been discussed by many authors as the principal mechanism to scatter and isotropize cosmic rays (Schlickeiser 2002). Although cosmic ray diffusion can happen while cosmic rays follow wandering magnetic fields (Jokipii 1966), the acceleration of cosmic rays requires efficient scattering. For instance, scattering of cosmic rays back into the shock is a vital component of the first order Fermi acceleration (see Longair 1997).  While most investigations are restricted to Alfv\\'{e}n modes propagating along an external magnetic field (the so-called slab model of Alfv\\'{e}nic turbulence), obliquely propagating MHD modes have been included in Fisk et al. (1974) and later studies (Bieber et al. 1988, Pryadko \\& Petrosian 2000). The problem, however, is that the Alfv\\'{e}nic turbulence considered in their studies is isotropic turbulence, and this contradicts to the modern understanding of MHD turbulence (Goldreich \\& Shridhar 1995, see Cho, Lazarian \\& Yan 2002 for a review and references therein).  A recent study (Lerche \\& Schlickeiser 2001) found a strong dependence of scattering on turbulence anisotropy. Therefore the calculations of CR scattering must be done using a realistic MHD turbulence model. An important attempt in this direction was carried out in Chandran (2000). However, only incompressible motions were considered. On the contrary, ISM is highly compressible. Compressible MHD turbulence has been studied recently (see review by Cho \\& Lazarian 2003a and references therein). Schlickeiser \\& Miller (1998) addressed the scattering by fast modes. But they did not consider the damping, which is essential for fast modes. In this paper we discuss in detail the various damping processes which can affect the fast modes. To characterize the turbulence we use the statistics of Alfv\\'{e}nic modes obtained in Cho, Lazarian \\& Vishniac (2002, henceforth CLV02) and compressible modes obtained in Cho \\& Lazarian (2002, henceforth CL02, 2003b,c).  Yan \\& Lazarian (2002) used recent advances in understanding of MHD turbulence (CL02) to describe cosmic ray propagation in the galactic halo. In this paper we undertake a comprehensive study of cosmic ray scattering rates. In $\\S$2, we describe the statistics of Alfv\\'{e}nic and compressible turbulence in various conditions, including both high $\\beta$ and low $\\beta$ cases. In $\\S$3, we describe the resonant interactions between the MHD modes and CRs. The scattering by Alfv\\'{e}n modes and fast modes is presented in $\\S$4. We apply our results to different phases of ISM, including Galactic halo, hot ionized medium (HIM), warm ionized medium (WIM) and also partially ionized medium. In $\\S$5, we describe the streaming instability in the presence of background turbulence. Discussion of our results is provided in $\\S$6 while the summary is given in $\\S$7. ",
        "conclusions": "We have attempted to apply a realistic model recently obtained for MHD turbulence to the problem of cosmic ray scattering taking into account the turbulence cutoff arising from both collisional and collisionless damping. We considered both gyroresonance and transit-time damping (TTD),  and used QLT to obtain our estimates.  In earlier papers it was frequently assumed that Alfv\\'enic turbulence follows Kolmogorov spectrum while the problem of CR scattering was considered. This is incorrect because MHD turbulence is different from hydrodynamic turbulence if magnetic field is dynamically important. Realistic MHD turbulence can be decomposed into, Alfv\\'{e}n, slow and fast modes. They have different statistics (see CL02). The Alfv\\'{e}n and slow modes follow GS95 scalings and show scale-dependent anisotropy. The anisotropy of Alfv\\'{e}n and slow modes reduce the scattering efficiency to such a degree that their impact on cosmic ray propagation is marginal.  Fast modes, however, are promising as a means of scattering due to their isotropy. Even in spite of the damping fast modes dominate scattering if turbulent energy is injected at large scales. We provided calculations for various phases of ISM, including galactic halo, HIM, WIM and partially ionized media. As the fast modes are subjected to different damping processes, the CR scattering varies . In galactic halo and HIM, the fast modes are subjected to collisionless damping. The mean free path in WIM is much smaller, and ion viscosity dominates the damping. All these dampings decrease with $\\theta$ and marginal for parallel modes. As the result, the fast modes at small scales become anisotropic with $k_{\\parallel}\\gg k_{\\perp}$. In this sense, fast modes are in slab geometry but with less energies. These modes can efficiently scatter CRs with the resonant scale $k_{\\parallel,res}=\\Omega/v_{\\parallel}$ larger than the cutoff through gyroresonance. TTD is usually weaker for pitch angle scattering than gyroresonance, but dominant for momentum diffusion. In all the cases, the scattering efficiency decreases with plasma $\\beta$ because the damping increases with $\\beta$ and the perturbations of turbulence decreases with $\\beta$.   How good is our model of ISM turbulence? We used spectra for Alfv\\'en and fast modes which follow from theoretical constructions and which were tested numerically (CL02). Within the ISM similar spectra develop dynamically when random driving is applied at large scales (see Mac Low 2003) and therefore we expect that the distribution of energy over scales for different modes to coincide with the assumed ones. The partition of energy over modes depend, however, on driving. For random driving with $\\delta V\\sim V_A$ we expect approximately equal energy in fast and Alfv\\'{e}n modes. As the details of ISM driving are not clear and the QLT provides results correct up to a factor of the order unity, our predictions for fast mode scattering contain an uncertain factor of the order unity. We believe, however, that our results for scattering by Alfv\\'en modes that given by Eq.(\\ref{ana}) are subject to more uncertainties related to the turbulence driving. Indeed, the scattering rates that follow from Eq.(\\ref{ana}), although orders of magnitude larger than those in Chandran (2000), are still {\\it extremely} small compared to models with isotropic spectra. Because it is so small, higher order effects may become important. Does the steady state anisotropic spectrum has always enough time to develop? The dynamic establishing of anisotropic spectra of Alfv\\'{e}nic turbulence requires a few eddy turnover times. For imbalanced driving this can take somewhat longer time (see discussion in $\\S$5). Therefore one can imagine situations when the spectrum deviates of its steady state form given by Eq. (3). These deviations are most likely to be small. Therefore we do not expect that in typical ISM conditions the contribution of Alfv\\'{e}nic modes is comparable to that of the fast modes. Identification of the particular situations when the opposite may be true requires a more systematic study of particular driving mechanisms and is the subject of future research. We also note, that if energy is injected at the small scales, the anisotropy of the Alfv\\'{e}nic modes is not large and the scattering can be appreciable. The deviations of the spectrum from the steady state one are not important.  Although scattering via fast modes is isotropic at large scales and becomes slab-type at small scales at which damping gets important,  our model is radically different from earlier discussed slab and isotropic+slab models (Bieber et al. 1994, Schlickeiser \\& Miller 1998).We used theoretically motivated models of turbulence that have been tested numerically. Therefore the transitions from one regime to another as well as the values of the intensity of perturbations have good justification within our approach.   We adopted a particular set of parameters of ISM while doing the calculation. The basic conclusion, namely, fast modes dominate the cosmic ray scattering, will remain the same if the parameters are altered. However, since the $\\beta$ value in ISM remains uncertain from observation, the damping scale $k_{c}$ is somewhat uncertain. And this determines the energy limit down to which gyroresonance is dominant."
    },
    "0311/hep-ph0311189_arXiv.txt": {
        "abstract": "We calculate the cosmic microwave background (CMB) anisotropy spectrum in models with millicharged particles of electric charge $q\\sim 10^{-6}-10^{-1}$ in units of electron charge. We find that a large region of the parameter space for the millicharged particles exists where their effect on the CMB spectrum is similar to the effect of baryons. Using WMAP data on the CMB anisotropy and assuming Big Bang nucleosynthesis value for the baryon abundance we find that only a small fraction of cold dark matter, $\\Omega_{mcp}h_0^2 < 0.007$ (at 95\\%~CL), may  consists of millicharged particles with the parameters (charge and mass) from this region. This bound significantly narrows the allowed range of the parameters of millicharged particles. In models without paraphoton millicharged particles are now excluded as a dark matter candidate.  We also speculate that recent observation of 511~keV $\\gamma$-rays from the Galactic bulge may be an indication that a (small) fraction of CDM is comprised of the millicharged particles. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311308.txt": {
        "abstract": "{The  luminous unstable star (star system) $\\eta$ Carinae is surrounded by an optically bright bipolar nebula, the {\\it Homunculus} and a fainter but much larger nebula, the so-called {\\it outer ejecta}. As images from the EINSTEIN and ROSAT satellites have shown, the outer ejecta is also visible in soft X-rays, while the central source is present in the harder X-ray bands. With our CHANDRA observations we show that the morphology and properties of the X-ray nebula are the result of shocks from fast clumps in the outer ejecta moving into a pre-existing denser circumstellar medium. An additional contribution to the soft X-ray flux results from mutual interactions of clumps within the ejecta. Spectra extracted from the CHANDRA data yield gas temperatures $kT$ of $0.6-0.76$ keV. The implied pre-shock velocities  of 670-760\\,\\kms\\ are within the scatter of the velocities we measure for the majority of the clumps in the corresponding regions. Significant nitrogen enhancements over solar abundances are needed for acceptable fits in all parts of the outer ejecta, consistent with CNO processed material and non-uniform enhancement. The presence of a diffuse spot of hard X-ray emission at the S\\,condensation shows some contribution of the highest velocity clumps and further underlines the multicomponent, non-equilibrium nature of the X-ray nebula. The detection of an X-ray ``bridge'' between the northern and southern part of the X-ray nebula and an X-ray shadow at the position of the NN\\,bow can be attributed to a large expanding disk, which would appear as an extension of the equatorial disk.  No soft emission is seen from the Homunculus, or from the NN\\,bow or the ``strings''. % ",
        "introduction": "\\subsection{The object: $\\eta$ Carinae\\label{sec:eta}}  With luminosity $L\\sim 10^{6.7}\\,L_{\\sun}$, $\\eta$ Car marks the upper boundary of the empirical {\\it Hertsprung-Russell diagram\\/} (HRD). It is one of the most massive stars known ($M > 120\\,M_{\\sun}$), somewhat evolved, and unstable; see e.g.\\ Humphreys \\& Davidson (1994), Davidson \\& Humphreys (1997) and Hillier et al. (2001).  In the optical, $\\eta$ Carinae is known to be an irregularly variable source. \\ The historical lightcurve covers several centuries and shows the star's most sudden change around  1843 when it brightened to $-$1$^{\\rm m}$ (Herschel 1847; Innes 1903; Viotti 1995; Humphreys et al.\\ 1999) and drastically decreased its brightness by more than  7$^{\\rm m}$ during the following 20 years, an event known as the ``Great Eruption''.  $\\eta$ Carinae is an evolved massive star, now classified a {\\it Luminous Blue Variable\\/} (LBV). LBVs are very massive stars in a transitional phase between the main-sequence and the Wolf-Rayet state. Only stars above $\\sim$  50\\,M$_{\\sun}$ enter this unstable phase, which occurs at an age of roughly 3 10$^6$ years (e.g.\\ Langer et al.\\ 1994). After the main-sequence phase such stars first evolve towards the red, cooler regime in the HRD, but as  they enter the LBV phase they reverse their evolution before becoming red supergiants. The temperature-dependent luminosity boundary in the HRD near which  LBVs are situated is known as the {\\it Humphreys-Davidson Limit\\/} (HD-Limit; Humphreys \\& Davidson 1979, 1994).  A  periodicity first found in high excitation lines (Damineli 1996, Damineli et al.\\ 1997) and in $2-10$ keV X-ray emission (Corcoran et al.\\ 1995, 2001a; Ishibashi et al.\\ 1999) from the star raised the possibility that  $\\eta$ Carinae is a binary system. In such a system a wind eclipse of a shock produced by the the colliding winds from two hot stars would be responsible for the observed X-ray variability. In this scenario, $\\eta$ Carinae, the more massive component, has a mass in excess of  60\\,M$_{\\sun}$ while the companion would have a mass of about $30-60$\\,M$_{\\sun}$  (Damineli et al.\\ 1997, Corcoran et al. 2001a). For further information concerning the high-energy X-ray emission in connection with the central source we refer the reader to, for e.g.,\\ Ishibashi et al.\\ (1999), Corcoran et al.\\ (2001a,b),  Pittard \\& Corcoran (2002) and references therein.  \\subsection{The Homunculus nebula: optical and X-ray emission}    One of the main characteristics of the LBV phase is the star's optical variability,  which occurs on different timescales and amplitudes: from a few tenths of a magnitude within several months, to $1-2$ magnitude variations on timescales of decades, to variations of several magnitudes as seen in the giant eruptions which occur probably once or twice during the LBV phase. Whether all LBVs undergo giant eruptions is not clear. Therefore LBVs are sometimes divided into the ``S Dor'' variables and the ``eruptive'' variables depending on whether or not an eruption has taken place (Humphreys 1999). $\\eta$ Carinae is the canonical eruptive-type LBV. Together with the high increase in mass loss which occurs during the LBV phase (several\\,$10^{-4}$\\,M$_{\\sun}\\,{\\rm yr}^{-1}$ at least) these eruptions -- in which the stars peel of large parts of their outer envelopes -- are responsible for the formation of {\\it LBV nebulae\\/} (Nota et al.\\ 1995; Weis 2001a).   The bright ejecta-nebula around $\\eta$ Carinae was first clearly identified in the 1940s (Gaviola 1946, 1950; Thackeray 1949, 1950) when it was about 30\\% smaller than today.  Gaviola called it the {\\it Homunculus} because it resembled a little man in his photographs.  At high spatial resolution, it has become one of the most dramatic examples of bipolar morphology, having two roughly symmetric polar lobes and an {\\it equatorial disk} of material (Duschl et al.\\ 1995) which is sometimes called the skirt (e.g. Morse et al.\\ 1998).  Its polar diameter is currently about 17\\arcsec\\ or 0.2\\,pc. Fainter condensations outside the Homunculus were also detected more than 50 years ago (Thackeray 1949; Walborn 1976), and manifest a much larger outer nebula, seen in deep HST images (e.g.\\ Weis 2001a) and images  in the soft X-ray energy  band. This {\\it outer ejecta}, which consists of a large collection of ``bullets'', ``knots'' and ``filaments'', has a diameter of about 60\\arcsec\\ or 0.67\\,pc. Figure \\ref{fig:hst} displays an HST image taken in July 1997 in the F658N filter, which shows a mixture of H$_{\\alpha}$ and [N\\,{\\sc ii}] emission.   Here the intensity levels are chosen to show the faint outer material, and we have indicated the location and size of the bipolar Homunculus by contours.    \\begin{figure} {\\resizebox{\\hsize}{!}{\\includegraphics{aa0306_f1.eps}}} \\caption{An HST F658N image which shows the optical emission from $\\eta$ Carinae's Homunculus and the outer ejecta. The contrast has been optimized to show the fainter outer emission.  The bright central region including the bipolar Homunculus is displayed using  contours. } \\label{fig:hst} \\end{figure}   \\begin{figure} {\\resizebox{\\hsize}{!}{\\includegraphics{aa0306_f2.eps}}} \\caption{Previously identified features in the Homunculus and the outer ejecta are marked. Part of the nomenclature is according to Walborn et al.\\ (1978) } \\label{fig:name} \\end{figure}   Both parts of the nebula -- the Homunculus and the outer ejecta -- are expanding. The expansion velocity of the  Homunculus is about 650\\,\\kms\\ (Thackeray 1961; Davidson \\& Humphreys 1997) while in the outer ejecta observed velocities average at roughly 600\\,\\kms\\ -- slow enough to suggest either an earlier ejection date or else deceleration  -- and reach maximum expansion velocities up to 2000\\,\\kms\\ or more (Weis 2001a,b; Weis et al.\\ 2001). We will show in an upcoming paper that it is indeed possible that the outer ejecta was  created 1843 and has been slowed down (Weis \\& Duschl, in prep.).   The outer ejecta produces soft X-rays. The EINSTEIN satellite separated the nebular from the stellar X-ray emission (Chlebowski et al.\\ 1984), while images with spatial resolution $\\sim$ 5\\arcsec\\ were obtained later with the High Resolution Imager (HRI) aboard the {\\it R\\\"ontgensatellite}, ROSAT. Weis et al.\\ (2001) presented a detailed analysis of the ROSAT data and a comparison with optical images and kinematic data.  The X-ray nebula is hook-shaped, with X-ray bright knots near the locations of the prominent {\\it S\\,ridge} and  {\\it W\\,arc} (features named by Walborn 1978,  see also  Figs. \\ref{fig:name}, \\ref{fig:nameconv}). Considering also the kinematic data, it thus became clear that shocks associated with fast-moving visible knots produce much of the observed diffuse soft X-ray flux.  Since the outer knots and filaments have characteristic sizes ranging from a small fraction of an arcsecond to several arcseconds, the shock structures were not well resolved in the ROSAT images.  Seward et al.\\ (2001) reported the first CHANDRA X-ray images in the $0.2-10$ keV band using the Advanced CCD Imaging Spectroscopy (ACIS) imaging array with resolution $\\approx 1\\arcsec$ and compared them to HST images.  However, the energy resolution of this early CHANDRA observation was  seriously degraded by radiation damage which severely hampered the search for spectral variations in the nebulosity.  Here we describe a spatially resolved  analysis of the X-ray nebulosity around $\\eta$ Car, based on a the zeroth-order image from a  deep observation obtained using the CHANDRA High Energy Transmission Grating and the ACIS spectroscopic array.     \\subsection{Aim and advances of the new CHANDRA observations}  We used a dataset\\footnote{Dataset No.: 200057 ACIS-S plus HETG; P.I.: M.F. Corcoran} which was obtained by the CHANDRA X-ray Observatory primarily to obtaine a grating spectrum of the central source. CHANDRA images have about 10 times better spatial resolution than ROSAT HRI, allowing us to better define the X-ray source regions and possible coincidences with optical counterparts.  Even though the X-ray resolution is still a factor of 5 worse than the best optical images, we can now constrain the identifications fairly well.  In addition to morphological comparisons of X-ray, optical and kinematic data (see also Weis et al.\\ 2001), we can also explore the X-ray spectral properties of selected regions in the ejecta.  Our new data and results are presented in the following way. Sect. 2 describes the new X-ray observations and their reduction; Sect. 3 briefly summarizes earlier work and data sets which are necessary and useful for this paper.  In Section 4 we discuss new results from the CHANDRA data. Section 5 contains a comparison between the X-ray data and the kinematics of the optical gas,  while in Sect. 6 we present a discussion of our results, and we summarize the main results in Sect. 7. ",
        "conclusions": "\\subsection{Origin of the X-ray nebula's shape}  The observed properties of the X-ray emission is a combination of  the speed of the outflowing material and the density of the ambient medium into which the material is moving.  The density of the ambient medium has to be sufficiently high to produce detectable emission.  Velocities of the warm gas are derived from our Echelle data supplemented by transverse velocities from proper motion measurements (Walborn et al.\\ 1978).  Densities of the ambient medium are derived from direct measurements using the [S\\,{\\sc ii}] lines (Dufour et al.\\ 1997, Meaburn et al.\\ 1993), but for most of the gas we have to rely on emission measure derived from HST H$_{\\alpha}$ images, which can be contaminated by red- or blueshifted [N\\,{\\sc ii}] as previously discussed. While the absolute values for n$_e^2$ is somewhat compromised by such contamination, the relative surface brightnesses still trace the relative overall column densities of the H\\,{\\sc ii} gas.  With this  information we have to understand why there is no strict correlation between the warm gas seen in the optical and the hot X-ray gas. There are two alternatives which might explain why the density is higher at the places which are  X-ray bright.  The first possibility is an enhancement of  density produced by an older stellar wind. During its main sequence phase, $\\eta$ Carinae would have possessed a fast, low density wind which would  have formed an interstellar  bubble (Weaver et al.\\ 1977). The interior of this bubble would be nearly evacuated  since all material inside the bubble would be swept up to a  radius of the order of 70\\,pc . As $\\eta$ Carinae evolved into a blue supergiant and LBV, however,  the wind velocity would decrease (down to several 100\\,\\kms)  and the mass loss rate would increase to 10$^{-3}M_{\\sun}yr^{-1}$ (e.g. Langer et al.\\ 1994, Garc{\\'\\i}a-Segura et al.\\ 1996), which would increasingly fill the bubble with denser gas. After the ``Great Eruption'' the expansion of the outer ejecta  into the remnants of this older, denser LBV wind (which at 0.35\\,pc has a density of roughly 150 times that of the ISM) would be able to slow down the ejecta to its current expansion velocity (Weis \\& Duschl, in prep) and could produce sufficient X-ray emission in the outer region of the nebula. Had this stellar wind been non-spherical, the ejecta would encounter the  denser wind at different distances from the star.  Alternatively, higher densities may manifest localized regions in which the outer ejecta consists of many turbulent clumps which have different velocities.  These clumps would collide into and over take each other, so the density is locally enhanced. This seems a reasonable explanation especially in the S\\,condensation where a large number of clumps with different velocities are detected. X-ray emission is formed in regions with many clumps interacting with high velocities. In the E\\,rim and outer regions of the S\\,ridge fewer knots are detected and one might speculate that here the outer ejecta collides into the older stellar wind.  The Homunculus nebula itself is expanding at a velocity of 650\\,\\kms\\ on average but is not visible in soft X-rays, though velocities this large should produce soft X-ray emission. The lack of detectable soft emission from the Homunculus is likely due to the low density (as judged from the low surface brightness in the H$_{\\alpha}$ and N\\,[{\\sc ii}] HST images) in the medium into which the Homunculus is  expanding, which has been cleared out by earlier episondes of mass loss. Therefore the Homunculus expands practically freely and no shocks are formed.  \\subsection{Velocity fields and plasma temperatures\\label{sec:plasma}}  The ACIS-S spectra of the outer ejecta are decently fitted with single temperature coronal equilibrium models as we showed in sect. \\ref{sec:spect}.  Still, looking at the $\\chi^2$ deviations, residual spikes are visible which indicate poorly fitted lines/line complexes.  Indeed, several of these spikes correspond to lines complexes of e.g. Fe\\,{\\sc xvii} and Fe\\,{\\sc xviii} around 0.8 keV, Ne\\,{\\sc x} and Ne\\,{\\sc ix} at about 1.0 keV, Mg\\,{\\sc xi} and  Mg\\,{\\sc xii} at about 1.3 keV and finally Si\\,{\\sc xiii} and Si\\,{\\sc xiv} at around 1.8 keV. While Doppler line velocity displacements are big enough to create noticeable shifts in several places in the outer ejecta , the highest velocity regions only represent a minor part of the total emission of the outer ejecta as seen in our Echelle spectra.  Additionally, significant residuals are also visible in the X-ray spectra of regions without extremely fast moving knots. Therefore velocity offsets cannot fully explain the residuals. As traced by the Echelle data, we see not one typical expansion velocity but a distribution of velocities in all regions of the outer ejecta. Therfore we have different velocities which imply that the X-ray emission is composed of a distribution of emissivities, so that our single temperature fits are too simplistic.  The plasma properties we derive are therefore not the properties of the whole region, but rather the density weighted properties. We see the dominant emission component of a complex plasma, which is most probably also the reason why our single temperature fits are relatively good despite the multi-component and, due to the short timescales, non-equilibrium nature of the X-ray plasma in the outer ejecta (Weis et al.\\ 2001). Our interpretation of the optical gas velocities as shock velocities rest on the details of the mass loss history of the LBV. The ejecta do not expand into a fast freely expanding wind as assumed by Leutenegger et al.\\ (2003), but a dense very slow moving wind as explained above.  The temperatures of the X-ray emitting gas are in good agreement with the velocity of the dominant emission component (weighted by surface brightness, therefore column density).  What was missing is the very hot gas corresponding to the extremes of the velocities.  This we may have found now with the ACIS  data.  \\subsection{Missing X-ray emission from the NN\\,bow}  We detected soft X-ray emission from all the optical line emitting regions outside the Homunculus and the presumably low density envelope around it. Besides this region, there is one notable structure without X-ray emission: the NN\\,bow. This roughly trapezoidal region is regarded as part of the outer ejecta and appears as a possible radial extension of the equatorial disk (Duschl et al.\\ 1995). Surprisingly, the region is devoid of all soft emission, as can best be seen in the overlay of the soft X-ray emission on the HST WFPC2 F658N image (Fig. \\ref{fig:overlay}, middle panel). Since the radial velocities of the ionized gas belonging to the NN\\,bow are typically around $-650$\\,\\kms\\ (Meaburn et al.\\ 1993,  Weis et al.\\ 2001) (similar to other regions of the outer ejecta), one would expect to see similar X-ray emission. The apparent minimum in the X-ray emission around the NN\\,bow could be explained as a particularly deep X-ray shadow. The soft X-ray flux limit is as low as the one observed for the Homunculus itself.  Assuming again shocks as the creation mechanism of the X-ray emission we would expect gas with a plasma temperature around 0.56\\,keV, and the total absence would point at an H\\,{\\sc i} foreground density of about 10$^{22}$ cm$^{-1}$, similar to the column density through the Homunculus. Still, since the gas in the NN\\,bow is blueshifted, one would expect to see this gas on the front side of the dense gas. It appears therefore that the NN\\,bow is expanding into a low density region. Nothing is stopping the expansion in line of sight, the density of the ambient medium is too low to create soft X-rays, and there are also -- in contrast to the S\\,condensation -- no faster knots or clumps overtaking the slower gas in this region. At the northern end, however, shocks might form, and here at the optically bright NN\\,condensation X-ray emission is present. It is further possible that X-ray emission behind the NN\\,bow is  absorbed by the NN\\,bow itself.  The only measurements of the density of the gas in the NN\\,bow is  a lower limit from [S\\,{\\sc ii}] line ratios which as mentioned above are close to the high density limit (Dufour et al.\\ 1997, Weis 2002).  The density of the outer ejecta (outside the bright clumps) appears rather homogeneous (Weis et al., in prep.) with at least n$_e \\sim 10^4$ e$^-$ cm$^{-3}$, consistent with the NN\\,bow lower limit.  For an assumed cylindrical geometry and size based on the observed dimensions in the HST images (diameter of the cylinder d=0.065\\,pc) we get therefore an absorbing column of at least $\\sim 2 \\times 10^{22}$ cm$^{-2}$ through the NN\\,bow. This is large enough to completely hide any $0.6$ keV plasma at the sensitivity of our data.   Another interpretation is that an X-ray shadow could be formed if  the equatorial disk (Duschl et al.\\ 1995) or its extension shadows the soft X-ray emission in the center and western part of the NN\\,bow. In this scenario the expansion of the disk produces X-ray emitting shocks only at the outer edges, forming the bright X-ray rim at the NN\\,condensation and the ``bridge'' (see also next section). Unfortunately, with the present data set details of the absorption and possible changes of the absorption across the NN\\,bow cannot be derived due to a lack of X-ray photons. It is of interest to note that there is an X-ray bright spur just on the opposite side of the NN\\,bow (the S\\,condensation) which is spatially consistent with the plane of the equatorial disk.  Thus, both self absorption and/or expansion into low density environment can explain the lack of soft X-rays from the NN\\,bow.  Much deeper X-ray observations are needed to sort out the individual contributions of the processes.   \\subsection{The X-ray bridge across the Homunculus\\label{sec:bridge}}  In addition to the general lack of soft X-ray emission from the Homunculus region we can see in the CHANDRA data a ``bridge'' of soft emission crossing the Homunculus from the south-east to the north-west (see Fig.\\ \\ref{fig:colorchandra} and Fig. \\ref{fig:overlay}). We extracted a spectrum from this structure and find that it is rather similar to the X-ray spectra of the outer ejecta. A single temperature Mewe-Kaastra-Liedahl model as fitted to the other spectra (solar abundance except for nitrogen) gives an adequate fit to the data and yields a temperature of 0.67\\,keV and a nitrogen overabundance of a factor of 4. This is in good agreement with values found in other regions of the outer ejecta. The bridge has a somewhat lower X-ray surface brightness than most of the other regions in the outer ejecta (middle panel Fig.\\ \\ref{fig:images}), but shows similar small scale variations in the surface brightness.  It appears from all the properties as another part of the outer X-ray nebula.  The most straightforward interpretation for the soft X-ray bridge across the Homunculus is therefore a clump or a group of clumps located physically in front of the Homunculus, which was not seen in previous observations due to insufficient spatial and spectral resolution. The bridge is slightly below the optical disk and might represent shocks at the very outer ends of the disk, which in projection would appear offset from the bright optical emission of the disk. Unfortunately, the strong scattered light from the Homunculus prevents us from detecting related H$_{\\alpha}$ emitting gas in our Echelle spectra and therefore determining the velocity field of this gas. In the HST images, any optical extension of the disk is also difficult to detect due to the very bright optical emission of the Homunculus.    \\subsection{Hard X-ray emission from the nebula around $\\eta$ Carinae\\label{sec:hard}}   We have shown in section \\ref{sec:images} that some of the hard X-ray emission coinciding with the S\\,condensation and the SS\\,ridge are untouched by the PSF subtraction and therefore real. This implies that the S\\,condensation and some of the SS\\,ridge show hard emission, consistent with plasma temperatures comparable to the central source.  The implied shock velocities are well in excess of 1500\\,\\kms. Unfortunately the number of photons in both features are too small to generate a meaningful spectrum, but the result lends some support for the notion that not all of the hard tail and the Fe\\,{\\sc xxv} visible in the total nebular spectrum (Fig.\\, \\ref{fig:shellstar}) are due to wings of the central source's PSF, but intrinsic to the highest velocity knots in the nebula.\\\\  \\subsection{No X-rays from the strings} Some of the most puzzling features of the outer nebula are several extremely narrow optical emission filaments, dubbed ``strings'' (Weis et al. 1999), ``whiskers'' (Morse et al.\\ 1998) or ``spikes'' (Meaburn et al.\\ 1996). Because of their high velocity -- nearly 1000\\,\\kms\\ (Weis et al. 1999) -- and large density -- n$_e \\sim 10^4$ e$^-$ cm$^{-3}$ (Weis 2002) it is worthwhile to search for X-ray emission from these structures. Detection of such emission would constrain the properties of the interaction of the ``strings'' with the surrounding medium and would  further limit the possible creation mechanisms of these enigmatic features. The northwestern strings (\\#3 and  \\#4 as designated in Weis et al. 1999) are located in the complex region of the W\\,arc. Plenty of diffuse X-ray emission is present, but no emission can be convincingly associated with the strings. The southeastern  strings are located in a much less crowded region. Unfortunately no diffuse X-ray emission can be seen either at the tip or along the body of the strings (\\#1, \\#2 and \\#5). The reason for this non-detection is most probably in the very small working surface of the strings (Weis et al., in prep.).  Much deeper CHANDRA data will be needed to search for the interaction of the strings with the surrounding material."
    },
    "0311/astro-ph0311565_arXiv.txt": {
        "abstract": "{ In this paper we present surveys of two open clusters using photometry and accurate astrometry from the SuperCOSMOS microdensitometer. These use plates taken by the Palomar Oschin Schmidt Telescope giving a wide field ($5^{\\circ}$ from the cluster centre in both cases), accurate positions and a long time baseline for the proper motions. Distribution functions are fitted to  proper motion vector point diagrams yeilding formal membership probabilities. Luminosity and mass functions are then produced along with a catalogue of high probability members. Background star contamination limited the depth of the study of Alpha Per to $R=18$. Due to this the mass function found for this cluster could only be fitted with a power law ($\\xi(m) = m^{-\\alpha}$) with $\\alpha=0.86^{+0.14}_{-0.19}$. However with the better seperation of the Pleiades' cluster proper motion from the field population results were obtained down to $R=21$. As the mass function produced for this cluster extends to lower masses it is possible to see the gradient becoming increasingly shallow. This mass function is well fitted by a log normal distribution.\\\\ ",
        "introduction": "The young open cluster Alpha Per lies at a distance of 183 pc ($(M-m)_{o}=6.31$, van Leeuwen 1999) with an age of 90 Myr (Stauffer, Barrado \\& Bouvier 1999). We assume in this paper that  $A_{R}=0.23$(O'Dell, Hendry \\& Collier Cameron 1994). It has been reasonably well studied in the past using various techniques. Prosser's catalogue (Prosser 1993) contains a comprehensive list of candidate member stars identified by photometry, spectroscopy and radial velocity measurements. The faintest of these has $V \\approx 21.8$($M_{V} \\approx 15.5$). These candidate stars have been supplemented by those identified in similar studies by Zapatero~Osorio et al.~(1996), Stauffer et al.~(1999) and Rebolo et al.~(1992). In addition Prosser, Randich \\& Simon~(1998) identified candidate stars which were optical counterparts to X-ray sources within the cluster.\\\\ An early proper motion study by Heckmann, Dieckvoss \\& Kox~(1956) identified stars with proper motions consistent with cluster membership. However this study only went down to $V \\approx 12$ and did not calculate formal membership probabilities. Prosser (1992) also used stellar proper motions as an initial selection technique. Although this study was deeper than the earlier one ($V=18.8$) it also did not calculate formal membership probabilities. \\\\ Most recently Barrado~y~Navascu\\'{e}s et al~(2001) produced a deep CCD survey of the cluster. This identified several candidate low mass stars and Brown Dwarfs down to $I_{c}=22$. They then used this data to produce a mass function. The mass function was fitted with a power law whose slope was found to be $\\alpha = 0.59$.\\\\ The Pleiades is the best studied open cluster, here we assume $(M-m)_{o}=5.4$ (van Leeuwen 1999) and $A_{R}=0.15$ (O'Dell, Hendry \\& Collier Cameron 1994). In recent years several deep CCD surveys of the cluster such as Dobbie et al (2002) and Moraux et al (2003) have produced candidates down to $0.03 M_{\\odot}$. These have been complimented by proper motion studies such as those of Hambly et al~(1999) and Moraux, Bouvier and Stauffer (2002). However the only wide field, high precision proper motion survey is that of Hambly, Hawkins and Jameson~(1991). Adams et al (2001) used 2MASS data along with proper motions to produce a wide field study. This contained many candidate stars outside the tidal radius of the cluster showing background star contamination is a problem for many techniques. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{applot.ps}} \\caption{A plot showing the sky coverage of the Alpha Per study. The solid line represent the later POSSII fields while the dotted lines are the earlier POSSI fields. The dashed circles marks out a five degree radius from the cluster centre (marked by an x).} \\label{applot}  \\end{figure} \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{plplot.ps}} \\caption{A plot showing the sky coverage of the Pleiades study. The solid line represent the later POSSII fields while the dotted lines are the earlier POSSI fields. The dashed circles marks out a five degree radius from the cluster centre (marked by an x).} \\label{plplot}  \\end{figure} ",
        "conclusions": "We have presented the first proper motion survey of the cluster Alpha Persei to calculate formal membership probabilities. A  Luminosity Function for the range $R=10$ to $R=18$ has been produced and a Mass Fuction derived from this. The Mass Function was fitted with a power law distribution with $\\alpha=0.86^{+0.14}_{-0.19}$. A catalogue of 339 high probability members has been created. This  has been cross-correlated with previous studies of the cluster showing that a large number of stars previously thought to be cluster members have high calculated membership probabilities. A similar study has also been undertaken on the Pleiades and while not the first of it's type it has produced several low mass star candidates and yielded a better constrained Mass Function at the faint end."
    },
    "0311/astro-ph0311086_arXiv.txt": {
        "abstract": "{\\par \\ifpreprintsty % \\vskip2.5pc \\begin{center}% {\\large \\abstractname\\par}% \\end{center}% \\vskip.5pc \\fi \\bgroup \\ifdim\\prevdepth=-1000pt \\prevdepth0pt\\fi \\hsize2\\columnwidth \\if@twocolumn\\else\\leftskip=0.10753\\textwidth \\rightskip\\leftskip\\fi \\dimen0=-\\prevdepth \\advance\\dimen0 by17.5pt \\nointerlineskip \\small\\vrule width 0pt height\\dimen0 \\relax }   \\begin{document}  \\title{Astrophysical tau neutrinos and their detection by large neutrino telescopes} \\author{ Edgar Bugaev,$^1$ Teresa Montaruli,$^{2,3}$ and Igor Sokalski$^{3,1}$\\\\[1mm] {\\it $^1$ INR, 60th October Anniversary Prospect 7a, 117312, Moscow, Russia\\\\ $^2$ Physics Department, Bari University, Via Amendola 173, 70126 Bari, Italy\\\\ $^3$ INFN/Bari, Via Amendola 173, 70126 Bari, Italy\\\\} } \\date{ \\today } \\twocolumn[ \\begin{@twocolumnfalse}   We present results of the detailed Monte Carlo calculation of the rates of double-bang events in 1 km$^3$ underwater neutrino telescope with taking into account the effects of $\\tau$-neutrino propagation through the Earth. As an input, the moderately optimistic theoretical predictions for diffuse neutrino spectra of AGN jets are used. ",
        "introduction": "\\label{sec:1} During these years there have been many discussions \\cite{1,2} on the possibilities of very high energy $\\tau$-neutrino detection by large neutrino telescopes, which would prove oscillations of muon neutrinos of astrophysical origin (\"astrophysical long baseline experiment\" \\cite{3}).  It is known now that neutrino oscillations definitely exist. Nevertheless, the  detection of astrophysical $\\tau$-neutrinos is still very interesting and useful, e.g., for a better understanding of neutrino properties as well as properties of astrophysical sources \\cite{4}.  Intrinsic flux of $\\tau$-neutrinos from a typical source in which there is acceleration of protons and production of neutrinos through $pp$- and $p\\gamma$-reactions is very small. The dominant channel of $\\tau$-neutrino birth is through the inclusive production of charmed mesons \\cite{5}, \\begin{equation} pp,p\\gamma \\longrightarrow D^{+}_{S} + X,\\;\\;\\;\\;\\;\\; D^{+}_{S} \\longrightarrow \\tau\\nu_{\\tau}\\;\\;,\\nonumber \\end{equation} and the smallness of $\\nu_{\\tau}$-flux  follows from the relations \\begin{equation} \\frac{\\sigma_{\\gamma p\\to DX}}{\\sigma_{\\gamma p \\to \\pi X}} \\alt 10^{-3}\\;,\\;\\; \\frac{Br(D_{S}^{+}\\to\\nu_{\\tau})}{Br(\\pi\\to\\nu_{\\mu})}\\sim 7\\cdot 10^{-2}.\\nonumber \\end{equation} Therefore, the production of $\\nu_{\\tau}$ at source is negligible, and if all muons can decay, it can be expected for the 3 flavors the following proportion: \\begin{equation} \\Phi^{0}_{std}=\\{\\phi^0_{e},\\phi^0_{\\mu},\\phi^0_{\\tau}\\}=\\{\\frac{1}{3},\\frac{2}{3},0\\}.\\nonumber \\end{equation}  Assuming maximal atmospheric mixing and $U_{e\\;3}=0$, that is: \\begin{equation} \\sin \\theta_{1\\;3}=0\\;,\\;\\;\\; \\sin 2 \\theta_{2\\;3}=1\\;\\;\\;\\nonumber \\end{equation} one obtains, after averaging over many oscillations, the ideal equipartition between neutrino flavors in the astrophysical neutrino flux at Earth, \\begin{equation} \\Phi_{std}=\\{\\frac{1}{3}, \\frac{1}{3}, \\frac{1}{3}\\},\\nonumber \\end{equation} independently on the value of the solar mixing angle $\\theta_{1\\;2}$. Moreover, in the case when there is maximal atmospheric mixing and $U_{e \\;3}=0$, for any proportions between flavors at the source one has, after propagation, the equality $\\phi_{\\mu}=\\phi_{\\tau}$ \\cite{4}.  The experimental verification of the relation $\\phi_{\\mu}=\\phi_{\\tau}$ by detection of astrophysical diffuse fluxes of $\\nu_{\\tau}$ and $\\nu_{\\mu}$ with flavor identification capabilities is very important because the inequality $\\phi_{\\mu}\\ne\\phi_{\\tau}$ is predicted in some exotic scenarios, e.g. \\cite{4}: i) if CPT invariance is violated in the neutrino sector and the neutrino and antineutrino mixing matrices do not coincide, $\\phi_{\\mu}$, in general, is not equal to $\\phi_{\\tau}$, ii) if neutrinos can decay \\cite{6} {\\it en route } from the sources to the Earth, the neutrino fluxes and flavor ratios are very sensitive to uncertanties in the neutrino mixing matrix and strongly depend on the hierarchy of neutrino masses. So, the inequality $\\phi_{\\mu}\\ne\\phi_{\\tau}$ would reveal unconventional neutrino physics.  Inside of some astrophysical and cosmological objects there are processes in which very energetic quarks  and gluons fragmenting into jets of hadrons are expected to be produced. Even the usual inclusive spectrum of high energy $pp$- or $p\\gamma$-collisions always contains the jet component (described well by perturbative QCD). Jets can appear in decays of (hypothetical) supermassive particles which, in turn, are produced in decays of topological defects \\cite{7}, in processes of superheavy dark matter annihilation inside of stars \\cite{8}, in a process of Hawking evaporation of primordial black holes \\cite{9}. It was shown in \\cite{10} that decays of top quarks contained in the  jets lead to production of $\\tau$-neutrino flux from the jets, which is of the same order as $\\nu_{\\mu},\\nu_{e}$ fluxes (at large neutrino energies, close to the value of the jet mass). This means that the intrinsic $\\tau$-neutrino flux from the sources in which the jet phenomena are essential is, in general, not suppressed (in comparison with intrinsic $\\nu_{\\mu},\\nu_{e}$ fluxes).  We studied in this paper the possibility of detection of extragalactic diffuse $\\tau$-neutrino fluxes by large ($\\sim1$km$^{3}$) underwater neutrino telescopes. In Sec.\\ref{sec:2} we discuss the general problems connected with the detection of very high energy $\\tau$-neutrinos ($E_{\\nu}\\agt 10^{6}\\text{GeV}$). In Sec.\\ref{sec:3} we present the diffuse neutrino spectra for the AGN jets chosen for the numerical calculations. Details of the calculations, results and conclusions are given in Sec.\\ref{sec:4}. ",
        "conclusions": "\\label{sec:4} First estimates of double-bang event rates in km$^3$ neutrino telescope were done in \\cite{26}, for down-going $\\tau$-neutrinos only. Detailed calculations of $\\tau$-neutrino propagation through the Earth carried out by Dutta et al. \\cite{27} (using power law incident neutrino spectra) showed that, in spite of the absence of neutrino absorption, the contribution of lower hemisphere events to the total event rate in a telescope cannot be too large. This is due to a strong decrease of neutrino energies during propagation. It was concluded in \\cite{27} that the most promising way to detect $\\tau$-neutrino astrophysical fluxes is an analysis of shower events in the telescope.  In this work we present results on Monte Carlo calculations of rates of double-bang events in a 1 km$^3$ telescope for both hemispheres using the diffuse neutrino spectra shown on Fig.1. We simulate both DIS CC and NC neutrino interactions for all flavors in the energy range $10^3\\text{GeV}<E_{\\nu}<10^9\\text{GeV}$ using the CTEQ3\\_~DIS structure functions. For $\\tau$-lepton propagation we used the MUM code \\cite{28} that was updated accounting for the corrections for photonuclear interactions of $\\tau$-leptons \\cite{29}. Some details of the calculations and the resulting neutrino fluxes of all flavors after propagation through the Earth were presented in our previous work \\cite{30}.  The rate of totally contained double-bang events is given by the formula \\begin{eqnarray} N=2\\pi\\rho N_A \\int_{-1}^{1}\\int_{E_{min}}^{\\infty}  V_{eff}(E_{\\nu_{\\tau}},\\theta) I(E_{\\nu_{\\tau}}, \\theta)\\times\\nonumber\\\\ \\sigma^{CC}(E_{\\nu_{\\tau}}) d E_{\\nu_{\\tau}}d(\\cos\\theta).\\nonumber \\end{eqnarray} Here, $N_A$ is the Avogadro number, $\\rho$ is the density of the medium (we use $\\rho=1\\text{g cm}^{-3}$ which is close to sea water/ice density), $I(E_{\\nu_{\\tau}}, \\theta)$ is the differential $\\nu_{\\tau}$ flux entering the detector at zenith angle $\\theta$, $\\sigma^{CC}$ is the total deep inelastic neutrino cross section (charge current part). $V_{eff}$ is the effective volume given by the relation: \\begin{eqnarray} V_{eff}(E_{\\nu_{\\tau}},\\theta)=S_{p}(\\theta) (L-R_{\\tau}(E_{\\nu_{\\tau}})),\\nonumber \\end{eqnarray} where $S_p$ is the projected area for tracks generated isotropically in azimuth at the fixed $\\theta$ directions on a detector of parallelepiped form, $L$ is the geometrical distance between entry and exit point, $R_{\\tau}$ is the $\\tau$ lepton range. The minimum value of the neutrino energy, $E_{min}=2\\cdot 10^6\\text{GeV}$,corresponds to a $\\tau$-lepton range of $R_{\\tau}^{min}=70\\text{m}$.  The double-bang topology can be considered as a background free signature of $\\tau$-neutrino appearance \\cite{1} since there should be no atmospheric muons which could produce in the detector 2 showers with comparable amounts of photons. Nevertheless, this reasonable  assumption needs to be verified through a full simulation.  Concrete calculations were done for two geometries of the neutrino detector: \\begin{eqnarray} 1\\times 1\\times 1\\;\\;\\; \\text{km}^3 \\;\\;\\;\\;\\;\\;\\; (\\text{IceCube-like \\cite{31}}),\\nonumber\\\\ 1.4\\times 1.4\\times 0.6\\;\\;\\;\\text{km}^3\\;\\;\\;\\;\\;\\;\\;(\\text{NEMO-like \\cite{32}}).\\nonumber \\end{eqnarray}  In Table~1 the corresponding double-bang rates for lower, upper and both hemispheres are shown. \\begin{tabular}{c} ~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~   ~\\\\ ~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~   ~\\\\ TABLE~1. ~~~~~~~~ Number of totally contained double-\\\\ bang events in km$^3$ detector per year.~~~~~~~~~~~~~\\\\ \\\\ \\hline \\hline ~~Spectrum~~~~~~~  ~~~~Rate~$(N_{-2\\pi} / N_{2\\pi} / N_{4\\pi})$ \\\\ \\hline ~~~~~~~~~~~~~~~~~~~~~~~~~  ~~~~~~~~~~~~~~~NEMO-like~~~~~~~~~~~~~ \\\\ ~~~~~~~\\cite{21}~~~~~~~~~  ~~~~~~~~~~1.0~~~~~~2.1~~~~~~3.1~~~~~ \\\\ ~~~~~~~\\cite{18}~~~~~~~~~  ~~~~~~~~~~1.4~~~~~~3.1~~~~~~4.5~~~~~ \\\\ \\hline ~~~~~~~~~~~~~~~~~~~~~~~~~  ~~~~~~~~~~~~~~IceCube-like~~~~~~~~~~ \\\\ ~~~~~~~\\cite{21}~~~~~~~~~  ~~~~~~~~~~0.7~~~~~~1.6~~~~~~2.3~~~~~ \\\\ ~~~~~~~\\cite{18}~~~~~~~~~  ~~~~~~~~~~1.0~~~~~~2.3~~~~~~3.3~~~~~ \\\\ \\hline ~~~~~~~~~~~~~~~~~~~~~ \\end{tabular}  Calculated values for upper hemisphere are $3-6$ times lower compared to \\cite{26} since more optimistic predictions for diffuse neutrino fluxes were used there.  Finally, we found using the theoretical predictions in \\cite{18,21} for AGN diffuse neutrino fluxes that in a 1~km$^3$ neutrino telescope one can expect a marginally observable rate of double-bang events from $\\tau$-neutrinos."
    },
    "0311/astro-ph0311279_arXiv.txt": {
        "abstract": "We present the detection of 34 \\lya\\ emission-line galaxy candidates in a 80$\\times$ 80 $\\times$ 60 co-moving Mpc region surrounding the known $z=2.38$ galaxy cluster J2143-4423. The space density of \\lya\\ emitters is comparable to that found by Steidel et~al. when targeting a cluster at redshift 3.09, but is a factor of $5.8 \\pm 2.5$ greater than that found by field samples at similar redshifts.  The distribution of these galaxy candidates contains several 5-10 Mpc scale voids. We compare our observations with mock catalogs derived from the VIRGO consortium $\\lambda$CDM n-body simulations. Fewer than 1\\% of the mock catalogues contain voids as large as we observe. Our observations thus tentatively suggest that the galaxy distribution at redshift 2.38 contains larger voids than predicted by current models.  Three of the candidate galaxies and one previously discovered galaxy have the large luminosities and extended morphologies of ``\\lya\\ blobs''. ",
        "introduction": "The two-point correlation coefficient of high redshift ($z>2$) galaxies is quite similar to that of galaxies today (e.g. Steidel et~al.\\ 1996, 2000, Giavalisco et~al. 1998). The {\\em topology} of the distribution of high redshift galaxies, however, is not yet clear. The two-point correlation coefficient alone is not very sensitive to this topology. In the local universe, a large fraction of galaxies lie in filaments and sheets, such as the Great Wall (Geller \\& Huchra 1989). These filaments may be several hundred Mpc in length, and are separated by voids which can be tens of Mpc across.  Our understanding of how and when filaments and voids were established depends on the geometry of the universe and biasing of galaxies (e.g. White \\& Frenk 1991; Kauffmann, Nusser, \\& Steinmetz 1997; Mo, Mao, \\& White 1999).  Mapping such large scale structure at high redshifts requires a careful balance of field of view (to encompass the structures), redshift coverage (to accept enough test objects in the structure while avoiding confusion by overlapping structures in distance), and photometric sensitivity (to find enough test objects). Photometric redshifts (see Hogg et~al. 1998 for a review) provide an efficient estimate of redshift for large imaging surveys, but they are not accurate enough to avoid confusion from multiple overlapping structures, and cannot delineate the structures directly.  Bright QSOs may be markers of dense regions (Ellingson, Yee \\& Green, 1991), but are too sparse to provide multiple samples within individual structure elements.  Narrow-band imaging to find \\lya\\ emitting galaxies has many advantages as a technique for mapping the topology of the galaxy distribution at high redshifts. The technique picks out galaxies in a narrow-range of redshifts, so the two-dimensional distribution of their positions on the sky can be used to constrain the topology alone, without the need for expensive follow-up spectroscopy. At sufficiently low flux limits the space density of detectable \\lya\\ emitters becomes comparable to that of Lyman-break galaxies (eg. Hu \\& McMahon 1996).  They should be dense tracers of large scale structure. With large-format prime-focus imagers, it is relatively straightforward to survey very large volumes of the early universe. The fast beam required for large field coverage puts a lower limit the width of the narrow-band filter. Typically this limits the line-of-sight depth covered in an exposure to be roughly equal to or greater than the width.  A disadvantage of \\lya\\ searches is that the \\lya\\ emission from a galaxy is very hard to predict theoretically, due to the very high optical depth in this line, its dependency on metalicity and star formation rate, and the ease with which it can be obscured by dust.  Two studies find some evidence for filamentary structure in the distribution of \\lya\\ emitting galaxies on scales of around 5Mpc (Campos et~al. 1999, M{\\o}ller \\& Fynbo 2001). Other surveys have found evidence for clustering in \\lya\\ emitting galaxies, but not for filamentary structure (Steidel et~al. 2000, Ouchi et~al. 2002).  In this paper, we present the detection of 34 candidate \\lya\\ emitting galaxies in a region containing J2143-4423, a z=2.38 cluster of galaxies and damped \\lya\\ absorbers (Francis et~al., 1996, 1997, 2000). Our survey covers a region 80$\\times$80$\\times$60 co-moving Mpc in size, and hence is sensitive to much larger structures than previous studies at comparable redshifts.  The narrow-band imaging also detects four extended nebulae of \\lya\\ emission, often dubbed ``blobs'' (Steidel et~al., 1999, Keel et~al. 1999, Francis et~al., 2000).  Throughout the paper, we assume a $\\Lambda$-dominated flat Universe ($H_0 = 65$ km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_M=0.3, \\Omega_{\\Lambda}=0.7$). ",
        "conclusions": "\\subsection{The Space Density of \\lya\\ Emitters}  We are sensitive to \\lya\\ emitting galaxies over a 45\\arcmin$\\times$45\\arcmin\\ field minus 9\\arcmin$\\times$9\\arcmin in the South-East corner, in the redshift range $2.356 < z < 2.401$.  For our adopted cosmology, this corresponds to a volume 6100 co-moving Mpc$^2$ in area and 59 co-moving Mpc deep. Figure~\\ref{fig_lyacont} shows the density contours of \\lya\\ candidates estimated using an adaptive kernel routine (Beers, Flynn \\& Gebhardt 1990).  The initial smoothing scale is 12 arcminutes. The highest density contour is $5.7 \\times 10^{-2}$ sources per square arcminute corresponding to a spatial density of $3.0 \\times 10^{-4}$ sources per cubic Mpc. The average density over the whole field is $1.9 \\times 10^{-2}$ sources per square arcminute corresponding to a spatial density of $1 \\times 10^{-4}$ sources per cubic Mpc.  \\begin{figure} \\plotone{f5.eps} \\caption{Density contours for the distribution of \\lya\\ candidates estimated with the adaptive kernel routine. The lowest contour is one galaxy per resolution element $3.3\\times10^{-3}$ galaxies arcmin$^{-2}$. The highest contour is $5.7\\times10^{-2}$ galaxies arcmin$^{-2}$.\\label{fig_lyacont}} \\end{figure}  Steidel et~al. (2000) surveyed 20,000 cubic co-moving Mpc at redshift 3.09, targeting a cluster of Ly-dropout galaxies. Their survey was sensitive to lower \\lya\\ luminosities than ours, and they used a lower rest-frame equivalent width threshold. They find 9 candidates meeting our luminosity and equivalent width thresholds corrected for redshift, yielding a density of $4.5 \\times 10^{-4}$ sources per cubic Mpc. Steidel et~al. calculate that this space density is a factor of $6.0 \\pm 2.4$ greater than that of smaller survey of Cowie \\& Hu (1998). This implies that peak density in our field is a factor of 4 overdense.     \\subsection{Star Formation Rates}  To estimate the star formation rate we first assume that star formation powers the observed \\lya\\ flux.  This is almost certainly a great underestimate: even tiny amounts of dust will greatly reduce the observed \\lya\\ flux due to its high optical depth.  We use an un-reddened \\lya/\\ha\\ ratio of 8:1, and the conversion factor between H$\\alpha$ flux and star formation rate of Kennicutt (1983). Our faintest candidates (NB=23.25) have inferred star formation rates of $\\sim 9 M_{\\sun}{\\rm yr}^{-1}$.  Summing over all galaxies down to our flux limit, and making a crude correction for incompleteness amongst the faintest, the integrated star formation rate density (assuming no dust obscuration) is $2.4 \\times 10^{-3} M_{\\sun}{\\rm Mpc}^{-3}{\\rm yr}^{-1}$: well below the value ($\\sim 0.1 M_{\\sun}{\\rm Mpc}^{-3}{\\rm yr}^{-1}$) derived from integrated blue light by Madau, Pozzetti \\& Dickinson (1998).  We also estimate star formation rate from the 1600\\AA\\ UV continuum using the Madau et~al. formula. At z=2.38 the 1600\\AA\\ continuum falls in the $V$-band. We detect 23 of our candidates in the V-band.  The star-formation density implied by these candidates is $5 \\times 10^{-3} M_{\\sun}{\\rm Mpc}^{-3}{\\rm yr}^{-1}$.  The star formation rate for a candidate at the V-band limit of 25.3mag is $\\sim 18 M_{\\sun}{\\rm yr}^{-1}$.  The 14 galaxies not detected in the $V$-band could add, as an upper limit, $0.6 \\times 10^{-3} M_{\\sun}{\\rm Mpc}^{-3}{\\rm yr}^{-1}$ to the star formation rate density.  Our \\lya\\ candidates therefore contribute only a small fraction ($\\sim 5\\%$) of the overall star formation rate density at z=2.38.  In a detailed study of the spectra of Lyman-break galaxies, Shapley et~al. (2003) conclude that galaxies with the strongest \\lya\\ emission have bluer UV continua, weaker interstellar absorption and smaller star-formation rates than galaxies with weak \\lya\\ emission or \\lya\\ absorption. \\lya\\ emission from galaxies with the highest star-formation rates is absorbed by dust.  \\subsection{The Distribution of \\lya\\ Emission Candidates}  The distribution of our candidate resolved \\lya\\ emitting sources is shown in Fig~\\ref{fig: lya positions}. There is no obvious overdensity surrounding the previously known cluster at the field center. Instead, most of the candidates lie in a broad band extending from the north-east of the field to the south.  There is a significant is the lack of galaxies in the north and west. These regions have long integration times. There are three close pairs or triplets of candidates, but otherwise the candidates are well spaced.  The pairs have a spacing of about 20 arcsec or a minimum separation of 200 proper kpc.  The spacing of the galaxies in the triplet is about 50 arcsec or a minimum separation of 440 proper kpc.   \\subsubsection{Comparison to a random distribution}  We used two statistics to compare our data to a random distribution. We compare these statistics to results from Monte Carlo simulations generated with the same exposure-time mask. In this section we show that the distribution of our candidates has a significant excess of close ($< 1$\\arcmin ) pairs, and of large (6\\arcmin\\ -- 8\\arcmin ) voids.  The first statistic is essentially the angular two-point correlation function. For a series of angular separations, we take each data point in turn and count the number of other data points within that angular radius.  The second statistic is a two-dimensional analogue of the void probability function (VPF). For each angular scale, we randomly place 1000 circles with that radius on our field, so that the circles do not extend beyond the edge of our data. We count the fraction of these circles which did not contain any data points. The requirement that the circles not extend beyond the edge of the data makes this statistic most sensitive to voids near the center of the field. Void probability functions are sensitive to structures such as voids and filaments, to which the two-point correlation function is relatively insensitive.  Our data were compared to random Monte Carlo simulations generated as follows.  At each location in our image, the total exposure time was calculated. Using these total exposure times, the relative flux limits in every part of our field were computed. The number/magnitude distribution of our brighter candidates can be reasonably approximated as a power-law, with the number per unit magnitude increasing by a factor of roughly 5. Combining this number-magnitude relation with the flux limits, we calculated the relative probability of finding a candidate per unit area in any given part of our image. We then used this relative probability map to generate 2000 fake data sets, each containing the same number of sources as in the real data.  Our data have a significant excess of galaxies with separations less than 1\\arcmin . Only 0.1\\% of the Monte Carlo simulations produced this many close galaxies. On larger scales, our data still show an excess of galaxies, but this is no longer significant at the 5\\% level.  The void probability function (Fig~\\ref{fig: vpf}) is more interesting.  There is a clear excess of voids on scales of 5 arcminutes or greater. On scales of 5\\arcmin\\ -- 8\\arcmin , this excess is significant at the 99.9\\% confidence level (ie. fewer than 0.1\\% of our randomly generated data sets had VPFs as large as our data).  \\begin{figure} \\plotone{f6.eps} \\caption{ The void probability function of our data (solid triangles) compared to the average void probability functions of randomly distributed points (solid line) and void probability functions derived from the $\\lambda$CDM n-body simulations (open squares). \\label{fig: vpf} } \\end{figure}  We conclude that there is a significant excess of voids, and of small scale clustering, in our data. This conclusion is insensitive to the details of how we generated the random data sets: changes in (for example) the assumed number/magnitude relation made little difference to our results.  \\subsubsection{Comparison with $\\lambda$CDM n-body Simulations}  The distribution of our \\lya\\ candidates is thus inconsistent with a random distribution, with 99.5\\% confidence. Assuming that they are indeed galaxies at redshift 2.38, is their distribution consistent with theoretical predictions? In this section, we show that while the two-point correlation coefficient of our galaxy candidates is quite consistent with the predictions of a $\\lambda$CDM simulation, the distribution of our candidates shows more voids on 6\\arcmin\\ -- 8\\arcmin\\ scales, with 97\\% confidence.  To evaluate the apparent excess of voids, we compared our measured distribution of galaxies against the $\\lambda$CDM n-body simulations performed by the VIRGO consortium (Kauffmann et~al. 1999). We used their galaxy catalog, generated for redshift 2.12. For each galaxy, they list its position, stellar mass, gaseous mass and star formation rate.  Mock catalogs were generated as follows. Our observations were targeted at the source B1, a probable pair of giant elliptical galaxies surrounded by an extensive neutral hydrogen cloud (Francis et~al. 2000). B1 has an estimated stellar mass of $10^{10.9} M_{\\sun}$ and gaseous mass of $10^{9.7} M_{\\sun}$. We picked five galaxies in the $\\lambda$CDM simulations that had stellar and gaseous masses at least this large, and which lay far enough from the edge of the simulation volume.  The regions around these ``B1-analogue'' galaxies should be good matches to our data. For comparison, we also picked eight random locations within the $\\lambda$CDM simulation volume. We then extracted the galaxies that lay within a volume 70$\\times$70$\\times$46 co-moving Mpc (for our adopted cosmology) centered at each of these locations. There were $180 \\pm 30$ (1 $\\sigma$) galaxies in their catalogs in each of these volumes.  We then had to sparsely sample the galaxies within each of these volumes to generate mock catalogs with (a) 37 galaxies, and (b) a probability of being detected consistent with our exposure-time map. We did this in three different ways.  \\begin{enumerate}  \\item Randomly. As \\lya\\ emission can be generated in many ways (by star formation, shocks, cooling flows, AGN etc), and its escape in measurable quantities depends on accidents of the geometry and dust distribution inside galaxies, we first assumed that all the galaxies in the $\\lambda$CDM simulations had the same probability (37/180 $\\sim$ 20\\%) of being detected by our survey.  \\item Star Formation Rate. If the \\lya\\ flux is generated by star formation, only galaxies with star formation rates in excess of $8 M_{\\sun}{\\rm yr}^{-1}$ would have been detected. We therefore took only galaxies with star formation rates at least this large, and generated random sub-samples with the correct size and probability of detection as a function of position in our field. Only 25\\% of the galaxies within the $\\lambda$CDM cubes had star formation rates this large.  \\item Mass. We took only galaxies with masses greater than $10^{10.3} M_{\\sun}$, once again picking random sub-samples of the appropriate size and with probabilities of detection scaling appropriately with exposure time in each part of our image.  \\end{enumerate}  For each data cube, and for each of these three sub-sampling techniques, we repeated the random sub-sampling 2000 times. We then calculated the angular two-point correlation function and the void probability function for each sub-sampling of each data cube, and compared our results with the observations.  Fig~\\ref{fig: skyplot} shows our data, and five of the mock $\\lambda$CDM data cubes (centered on the five B1-analogue galaxies). The cubes shown have been randomly sub-sampled: those sampled on the basis of star formation rate or mass look very similar. Fig~\\ref{fig: vpf} shows the average void probability function of the randomly sub-sampled data cubes (centered on B1-analogue galaxies).   \\begin{figure} \\plotone{f7.eps} \\caption{ The distribution on the sky or our candidate \\lya emitting galaxies (top left) compared to five simulated data sets drawn from the $\\lambda$CDM n-body simulations. These simulations use random sub-sampling and are centered on B1-analogue galaxies. \\label{fig: skyplot} } \\end{figure}  The angular two-point correlation coefficient of our data, and of the $\\lambda$CDM simulations are in excellent agreement. As Fig~\\ref{fig: vpf} shows, however, the agreement is not as good for the void probability function. Our data have a void probability function considerably higher than that of the simulations on large scales.    Is this difference significant? The VPFs of our 13 simulated data cubes (five centered on B1-analogue galaxies and eight randomly located) actually showed rather little scatter, when averaged over all the 2000 random sub-samplings of each. The dominant source of scatter was the selection of the 37 galaxies in each simulated volume, which was typically twice as large.  Our uncertainty is thus dominated by the small sample size rather than by cosmic variance.  If we take the five data cubes centered on B1-analogues, and generate 2000 random sub-samplings of each, fewer that 0.1\\% of all these random sub-samplings have void probability functions as large as we observe on scales of 5\\arcmin\\ -- 8\\arcmin . If we sub-sample on the basis of star formation rate or galaxy mass, this fraction rises to around 1\\%. We would expect this, as more massive galaxies, and those with larger star formation rates, are expected to be more strongly biased. If we repeat these calculations for the randomly centered data cubes, the void probability function also rises slightly. This is presumably because centering our data cube on a B1 analogue precludes the possibility of a void in the center of the field.  Thus the observed distribution of galaxies has an excess of voids over the $\\lambda$CDM simulations, significant at roughly the 99\\% level. A qualitative impression of this can be gained from Fig~\\ref{fig: skyplot}. While all the $\\lambda$CDM simulations show empty regions near the edge of the field, this is an artifact of the shorter exposure times in these regions. The simulations do not show large voids in the central $\\pm 13$ arcminutes. Due to the restriction that we only place voids so that they do not overlap the edges of our field, our VPF statistic is mostly sensitive to the distribution of galaxies in this central region.  \\subsubsection{Discussion of the Voids}  We thus have tentative evidence from our survey of larger voids in the galaxy distribution at redshift 2.38 than predicted by one particular $\\lambda$CDM simulation. With only 97\\% confidence, a sample size of 34 and no spectral confirmation for most of our candidate galaxies, this evidence must be regarded as tentative at best. But if it is confirmed by larger samples, what is it telling us?  Steidel et~al. (2000) did not see voids, but they covered too small a region on the sky to be sensitive to the scales of void we are detecting. Ouchi et~al. (2003) imaged a region $\\sim 25$\\% the size of ours, searching for \\lya\\ candidates at $z=4.86$. Voids can be seen in their data if only brightest objects are selected so that their number per unit area matches ours.  Voids might also explain the enormous field-to-field variance in the space density of \\lya\\ sources found by Pascarelle, Windhorst \\& Keel (1998).  The spiky distribution of the redshifts of Lyman-dropout galaxies seen by Steidel et~al. (1998) and Adelberger et~al. (1998) would also be consistent with large voids in the high redshift galaxy distribution.  If the distribution of galaxies really contains more 10 co-moving Mpc scale voids than predicted, either $\\lambda$CDM is predicting too uniform a distribution of galaxies, or the voids contain dark matter halos which for some reason do not contain detectable galaxies.  There is some evidence for the latter hypothesis. The sight-line to background QSO 2138$-$4427 passes through one of the voids in our data, but the QSO spectrum shows a metal enriched damped \\lya\\ absorption-system at our redshift (Francis \\& Hewett 1993, Francis, Wilson \\& Woodgate 2001). So there is clearly at least one dense concentration of gas and some stars within one of our voids.  \\subsection{Blobs}  Some high redshift galaxy clusters contain extended \\lya\\ emission nebulae, or ``blobs'' (Steidel et~al.\\ 1999).  The blobs are radio quiet ($< 140~\\mu$Jy at 1.4 GHz), but otherwise share morphological similarity with the nebulae around radio galaxies (e.g. Kurk et~al.\\ 2001).  \\lya\\ blobs are large, bright gas clouds ($\\sim 100 h^{-1}$~kpc, $\\Omega_m =0.3, \\Omega_{\\Lambda}=0.7$; \\lya\\ flux up to $1.8\\times \\sim 10^{-15}$ erg cm$^{-2}$ s$^{-1}$). They appear to be common in regions of significant galaxy overdensity ($\\sim 10$~times that of the field) at high redshift.  Most blobs are associated with a host galaxy, though not symmetrically centered on it.  Blobs may break into smaller knots of \\lya\\ and continuum emission, with velocity differences $\\Delta V\\simlt 2000$ \\kms\\ (Steidel et~al. 2000).   Blobs can be among the brightest high-$z$ sub-mm sources (Chapman et~al.  2001).  Some blobs (Francis et~al. 2001) show CIV emission (and may host AGN) but others do not.  We detect four \\lya\\ blobs in the narrow-band image.  One of these, ``B1'', has previously been detected (Francis et~al.\\ 2001).  Figure \\ref{fig: blob images}\\ shows the broad- and narrow-band images of each blob, together with the continuum subtracted emission-line image. Each blob conforms to the expected size and brightness of this class of object.  Table \\ref{tab: blobs}\\ summarizes the characteristics of the blobs.  \\begin{figure} \\plotone{f8.eps} \\caption{REDUCED RESOLUTION. Images of the blobs.  The top row shows the broad band image of each blob; the middle row the narrow-band image; and the bottom row the difference of the two, scaled to give an image of the \\lya\\ emission.  Each image is $30\\arcs \\times 30\\arcs$. \\label{fig: blob images} } \\end{figure}   \\begin{deluxetable*}{llllll} \\tablecaption{\\lya\\ Blobs \\label{tab: blobs}} \\tablewidth{0pc} \\tablehead{ \\colhead{Name} & \\colhead{Redshift} & \\colhead{major axis} & \\colhead{\\lya\\ flux} & \\colhead{CIV emission} & \\colhead{Ref.\\tablenotemark{a}} \\\\ \\colhead{} & \\colhead{($z$)} & \\colhead{(\\arcs)} & \\colhead{(\\ecs)} & \\colhead{} & \\colhead{} } \\startdata  SB-1 & 3.09 & 17           & $1.4\\times10^{-15}$    & no         & 1,2          \\\\ SB-2 & 3.09 & 15           & $1.2\\times10^{-15}$    & no         & 1,2          \\\\ KB-A & 2.39 & $\\sim$8      & $1.1\\times10^{-15}$    & yes        & 3,4          \\\\ KB-B & 2.39 & $\\sim$8      & $0.45\\times10^{-15}$   & yes        & 3,4          \\\\ B1   & 2.38 & 17           & $1.8\\times10^{-15}$    & yes        & 5, this work \\\\ B5   & 2.38 & 8            & $1.5\\times10^{-15}$    & \\nodata    & this work    \\\\ B6   & 2.38 & 8            & $1.4\\times10^{-15}$    & \\nodata    & this work    \\\\ B7   & 2.38 & 7            & $0.7\\times10^{-15}$    & \\nodata    & this work  \\enddata \\tablenotetext{a}{ 1 -- Steidel et~al. 2000; 2 -- Bautz \\& Garmire 2001; 3 -- Windhorst et~al. 1998; 4 -- Keel et~al. 1999; 5 -- Francis et~al. 2001}  \\end{deluxetable*}   These blobs may be pre-galactic clouds of gas, in the process of initial collapse or destined to be stripped away into the intracluster medium (ICM).  In the blobs, we may witness a stage of structure formation which is illuminated by an unusual or transitory phenomenon. The excitation mechanism for the \\lya\\ emission is not known, but three primary models have been suggested:  In the early stages of a galaxy's formation, gas collapses into the dark matter potential well and cools via radiation (e.g. Fabian et~al. 1986).  Haiman, Spaans \\& Quataert (2000) calculate that cooling flows (with gas temperatures of a few $\\times 10^4$ K) could produce the observed flux range, size, surface brightness profiles and number density observed from \\lya\\ blobs.  While Fardal et~al. (2000) doubt that cooling flows would have sufficient luminosity, Francis et~al.  (2001) suggest that shocks within the infalling gas can trigger the blob emission in the cooling flow.  Confirmation that the blobs are high redshift cooling flows would have profound implications for the interpretation of semianalytic galaxy formation models (e.g. Fall \\& Efstathiou 1980) since we would be observing the initial collapse of gas into dark matter potential wells.  Taniguchi et~al. (2001 and references therein) suggest that starburst-powered winds trigger \\lya\\ emission in the expelled gas. Dust could obscure the host galaxy, with the \\lya\\ emission arising from a superwind extending $\\sim 100$ kpc from the host.  The dusty host galaxies would eventually evolve into ordinary elliptical galaxies, and the blobs would presumably dissipate into the intracluster medium.  The SCUBA detections support this model, implying rapid star formation (Chapman et~al. 2001) and a probability of hosting an AGN (Barger et~al.  2001).  Superwinds would indicate that the blobs are the next phase of galaxy evolution after the initial collapse.  Taniguchi \\& Shioya (2000) estimate this phase would last only 0.1 Gyr, occurring after the first 0.5 Gyr of star formation.  The expelled gas would eventually be stripped into the ICM, though it would be metal-enriched having been processed by supernovae in the host galaxy.  An AGN may photo-ionize the gas, similar to radio galaxy nebulae and Seyfert 2 galaxies with extended emission line regions (e.g. Villar-Martin, Tadhunter \\& Clark 1997).  For example, a \\lya\\ nebula has been found around the HDF-S QSO which is radio-quiet (Bergeron et~al. 1999; Palunas et~al. 2000a).  For the blobs, we do not see a strong UV continuum source; however, there may be an AGN surrounded by an obscuring torus (e.g. Krolik \\& Begelman 1986). In fact, most of the Chandra point sources do not show optical evidence for AGN (Mushotzky et~al. 2000). Assuming an ionization parameter of 10 and a density of 1 cm$^{-3}$\\ at a distance of 10 kpc, a central source with luminosity $9\\times 10^{45}$ ergs s$^{-1}$\\ is required to ionize a \\lya\\ blob with $L_x\\sim 10^{44}$~ergs s$^{-1}$.  The gas in AGN-supported blobs would contribute to the ICM, but with potentially lower metallicity than that from a starburst, thus diluting rather than enriching the ICM."
    },
    "0311/astro-ph0311423_arXiv.txt": {
        "abstract": "By means of high-resolution N-body simulations in a $\\Lambda$CDM cosmology, we verify that scaling relations similar to those observed for nearby galaxy clusters are also defined by their dark matter hosts; the slopes, however, are not the same. We then show that the scaling relations of galaxy clusters can be explained as the result of the cosmological collapse of density fluctuations at the appropriate scales, plus a systematic trend of the $M/L$ ratio with cluster mass. The empirical fact that the exponent of the Faber-Jackson relation of elliptical galaxies is significantly different (higher) than that of clusters, force us to conclude that the galaxy scaling laws might derive from the cosmological collapse of density fluctuations at the epoch when galactic scales became non-linear, plus modifications afterward due to early-time dissipative merging. ",
        "introduction": "It has been shown \\cite{S93,A98} that nearby galaxy clusters define scaling relations involving their optical luminosity $L$, effective radius $R_{\\mathrm e}$ and velocity dispersion $\\sigma$. Reanalyzing the sample of \\cite{S93}, we obtain $L\\propto \\sigma^{2.18\\pm0.52}$ for the Faber-Jackson (FJ) relation, $L\\propto R_{\\mathrm e}^{1.55\\pm0.19}$ for the Kormendy relation, and $L\\propto R_{\\mathrm e}^{0.9\\pm 0.15} \\,\\sigma^{1.31\\pm 0.22}$ for a nearly edge-on view of the Fundamental Plane, where $L$ is in $10^{12}L_\\odot/h^2$, $\\sigma$ in 1000 km/s, $R_{\\mathrm e}$ in Mpc$/h$.  Similar scaling laws are also followed by early-type galaxies, with the only significant difference of the FJ relation: in fact, $L\\propto \\sigma^4$ for luminous ellipticals \\cite{Bern03a}.  For a sample of 13 massive ($M>10^{14}{\\mathrm M}_\\odot/h$) DM halos obtained from N-body cosmological simulations \\cite{num_proceed}, we have computed the projected radius containing half the total mass ($ R_{\\mathrm h}$) and the mean projected velocity dispersion within it ($\\sigma_{\\mathrm h}$). Well defined scaling relations relating $ R_{\\mathrm h}$, $\\sigma_{\\mathrm h}$ and the total mass $M$ are found, and are interpreted as the result of the cosmological collapse of density fluctuations and the weak homology of DM halos \\cite{NFW}. The slopes of these relations, however, do not coincide with the observed ones: the simplest way to satisfy at once the theoretical and observational scaling laws is to assume that $M/L\\propto L^{0.3}$, in good agreement with what inferred observationally \\cite{Girardi02}.  \\begin{figure} \\begin{center} \\includegraphics[width=\\textwidth]{lanzoniF1.eps} \\end{center} \\caption[]{Nearly edge-on view of the FP of galaxy clusters, and of their DM halos under the assumption $M/L\\propto L^{0.3}$} \\label{fig1} \\end{figure}  Since a similar trend of $M/L$ with $L$ is also found for early-type galaxies, and the cosmological predictions for the DM halo scaling relations are the same at galactic and cluster scales, the higher slope of the galactic FJ relation implies that some additional mechanism played a role in the formation/evolution of ellipticals: indeed, the combined effect of dissipationless merging (which tends to increase the FJ exponent) and gas dissipation (which acts in the opposite direction) might be the solution \\cite{fpcmer}. ",
        "conclusions": "We conclude that, while the scaling relations of galaxy clusters can be explained by the cosmological collapse of density fluctuations at the appropriate scale, plus a systematic trend the total $M/L$ ratio with cluster luminosity, the scaling laws of elliptical galaxies seem to suggest that additional processes (like early-time dissipational merging) probably played an important role after the cosmological collapse of the density fluctuations at these scales. For a discussion of how $M/L$ for clusters could increase with $L$, we refer to \\cite{fpcmer}."
    },
    "0311/astro-ph0311109_arXiv.txt": {
        "abstract": "We report the earliest detection of an extremely bright optical afterglow of the gamma-ray burst (GRB) 030329 using a 30cm-telescope at Tokyo Institute of Technology (Tokyo, JAPAN). Our observation started 67 minutes after the burst, and continued for succeeding two nights until the afterglow faded below the sensitivity limit of the telescope (approximately 18 mag). Combining our data with those reported in GCN Circulars, we find that the early afterglow light curve of the first half day is described by a broken power-law ($\\propto$ $t^{-\\alpha}$) function with indices $\\alpha_{1} = 0.88 \\pm 0.01$ (0.047$<t<$ $t_{\\rm \\rm b1}$ days), $\\alpha_{2} = 1.18 \\pm 0.01$ ($t_{\\rm \\rm b1}$ $<t<$ $t_{\\rm \\rm b2}$ days), and $\\alpha_{3} = 1.81 \\pm 0.04$ ($t_{\\rm \\rm b2}$ $<t<$ 1.2 days), where $t_{\\rm \\rm b1}$ $\\sim 0.26$ days and $t_{\\rm \\rm b2}$ $\\sim 0.54$ days, respectively. The change of the power-law index at the first break at $t\\sim0.26$ days is consistent with that expected from a ``cooling-break'' when the cooling frequency crossed the optical band. If the interpretation is correct, the decay index before the cooling-break implies a uniform ISM environment. ",
        "introduction": "The overall behavior of GRB afterglows has been successfully explained by the standard fireball model (e.g. Piran et al. 1999). It is expected that the deviation from a simple power-law in the afterglow light curve will provide wealth of information on the environment and physical parameters of GRBs. For example, observations of early GRB afterglows confirmed that ``breaks'' exist in the light curves of a number of GRBs. Such breaks may be understood in the framework of the standard fireball model either as a ``jet-break'' where the bulk Lorentz factor of the relativistic jet decrease to the inverse of jet opening angle (Sari, Piran $\\&$ Halpern 1999, Rhoads 1999), or the ``cooling-break'' where the high energy electrons start to lose most of their energy rapidly by synchrotron emission which emit the observed optical photons (Sari, Piran $\\&$ Nakar 1998). While there are number of convincing cases for jet-breaks, identification of a cooling-break in the optical light curve has been difficult since it requires detection of a subtle change in the power-law index ($\\Delta \\alpha \\sim 0.25$). In order to study afterglow light curves in detail, continuous coverage and a high signal-to-noise ratio are required.  The situation has been dramatically improved since the advent of High Energy Transient Explorer-2 (HETE-2). HETE-2 can determine the positions of GRBs onboard, and notify ground-based observers of the GRB coordinates within 1$-$10 minutes after the burst (Ricker et al. 2002). For example, the locations of GRB 021004 (e.g., Fox 2002) and GRB 021211 (e.g., Fox $\\&$ Price 2002) were disseminated within less than a minute after the bursts, which prompted detailed studies of GRB afterglows in the very initial phases while they are bright.  GRB 030329 was detected by the HETE-2 satellite on 29 March 2003 at 11:37:14.7 UT. Position was determined by the ground analysis, and the location was reported to GCN at 73 minutes after the burst (Vanderspek et al. 2003). A very bright($R \\sim 13$ mag) optical transient (OT) was reported at $\\alpha = 10^{h} 44^{m} 50^{s} .0$, $\\delta = + 21^{\\circ} 31^{\\prime} 17.8'' (J2000.0)$ (Peterson $\\&$ Price 2003, Torii 2003) inside the SXC error circle. This is the brightest GRB ever detected by HETE-2 with the 30$-$400 keV fluence of $1.2 \\times 10^{-4}$ erg cm$^{-2}$, and precise and continuous follow-up observations were carried out by dozens of telescopes distributed around the world. Optical spectroscopic observations have determined its redshift as $z$ = 0.1685 (Greiner et al. 2003), which is one of the closest ever known and is possibly related to the exceptional brightness of this afterglow. Moreover, spectra taken after several days reveal the evolution of broad peaks in the spectra characteristic of a supernova. The spectral similarity to SN 1998bw (e.g., Galama et al. 1998, Iwamoto et al. 1998) and other energetic supernovae such as 1997ef provides strong evidence that GRB 030329 is associated with core-collapse supernovae (Dado, Dar $\\&$ Rujula 2003, Hjorth et al. 2003, Stanek et al. 2003, Kawabata et al. 2003). In order to investigate the kinetic energy of GRB and the immediate vicinity of its progenitor, the early light curve is important. In this letter, we report the earliest detection of the optical afterglow of GRB 030329 starting 67 minutes after the burst. ",
        "conclusions": "\\subsection{Light curve at $0.05 < t < 0.26$ days}  We have presented a light curve of the early phase of the optical afterglow of GRB 030329 starting 67 minutes after the burst. This is the earliest detection of GRB 030329 ever reported (Peterson $\\&$ Price 2003, Torii et al. 2003, Uemura et al. 2003).  Burenin et al. (2003) reported follow-up observations of GRB 030329 as early as 6 hours after the burst, using BVRI filters. In each of the filters, they observed a gradual flux decay which can be accurately described as a power-law $F_{\\nu}$ $\\propto$ $t^{-1.19}$. They also found a characteristic break in the light curve $t_{\\rm brk}$ $\\sim$ 0.57 day, after which the afterglow flux started to decline faster. The power-law slopes of the light curves changed from $-$1.19 to $-$1.9 for $t \\ge t_{\\rm brk}$. Notably, this break is nicely consistent with our second break ($t_{\\rm b2}$) within error bars, where the power-law slopes changes from $-1.18 \\pm 0.01$ to $-1.81 \\pm 0.04$ after the break of $t_{\\rm b2} \\sim 0.54$ days (see above). The first break ($t_{\\rm b1}$) is not discussed in Burenin et al, since they started their observations just around this break ($t_{\\rm b1} \\sim 0.26$ days).  Price et al. (2003), Burenin et al. (2003), and Tiengo et al. (2003) found that the results of their observations are consistent with the model where the afterglow emission is generated during the deceleration of the ultra-relativistic collimated jet. They found that the break in the power-law light curve, at $t\\sim0.5-0.6$ days, can be interpreted as the jet-break, i.e., the break which occurs when the angular structure of the ultra-relativistic collimated jet becomes observable (Sari, Piran $\\&$ Halpern 1999, Rhoads 1999). This interpretation is supported by the fact that the break occurred simultaneously in different colors. Furthermore, a change in power-law slope from $-$1.19 to $-$1.9, is approximately consistent with that generally observed in jet-break. Therefore, our major concern in this paper is to understand the nature of the first break, $t_{\\rm b1}$, and examine the consistency between the above scenario in the frame work of standard GRB fireball theories (Piran et al. 1999). \\vspace*{-0.3cm}  \\subsection{Break at $t \\sim 0.26$ days}  There are two possible break frequencies in the spectra, $\\nu_{\\rm m}$ and $\\nu_{\\rm c}$, where $\\nu_{\\rm m}$ is the synchrotron frequency, $\\nu_{\\rm c}$ is the cooling frequency above which electrons lose their energy rapidly by synchrotron radiation (Sari, Piran $\\&$ Narayan 1998). Since $\\nu_{\\rm m}$  and $\\nu_{\\rm c}$ are the functions of time, a break in the light curve could be observed when $\\nu_{\\rm c}$ and/or $\\nu_{\\rm m}$ crossed over the observed frequency $\\nu_{\\rm R}$. Therefore, we examined six possible cases to understand the first break ($t_{\\rm b1}$), according to the relation between  $\\nu_{\\rm m}$, $\\nu_{\\rm c}$ and $\\nu_{\\rm R}$. In the standard GRB scenario, $\\nu_{\\rm c}$ $\\le$ $\\nu_{\\rm m}$ is often called ``fast cooling'' since all electrons cool rapidly, whereas $\\nu_{\\rm c}$ $\\ge$ $\\nu_{\\rm m}$   is referred to ``slow cooling'' since only the high energy population of electrons cool efficiently. We will also extend our discussion to discriminate between ``a homogeneous interstellar medium (ISM) model'' (e.g., Sari $\\&$ Piran 1999)  and  ``a pre-existing stellar wind model'' (e.g., Chevalier $\\&$ Li 1999) for the GRB environment. The relationships between observed spectral index and model predictions are compared in Table 1.  We first consider the case where both $\\nu_{\\rm c}$ and $\\nu_{\\rm m}$ are above the observed optical frequencies  ($\\nu_{\\rm R} < \\nu_{\\rm m} < \\nu_{\\rm c}$ or $\\nu_{\\rm R} < \\nu_{\\rm c} < \\nu_{\\rm m}$: case (3) and (6) in Table 1). In these two situations, observed flux at $\\nu_{\\rm R}$ should increase with time, which strongly conflicts with the observed declining light curve. On the contrary, if the both cooling frequency $\\nu_{\\rm c}$ and the minimum frequency $\\nu_{\\rm m}$ are below the optical band ($\\nu_{\\rm c} < \\nu_{\\rm m} < \\nu_{\\rm R}$ or $\\nu_{\\rm m} < \\nu_{\\rm c} < \\nu_{\\rm R}$: case (1) and (4) in Table 1), the predicted optical spectral index would be $\\beta = p/2$, where $p$ is the electron spectral index. Since the photon spectral index of this afterglow was $\\beta = 0.66$ at $t = 0.26$ days (Burenin et al. 2003), we expect $p = 1.32$, which is unusually flat for an electron population accelerated in a GRB. Furthermore, power-law index in the light curve should be $\\alpha = \\frac{2-3p}{4} \\sim 0.49$, which is again inconsistent with $\\alpha_{1} \\sim 0.88$ obtained with our data.  Case (5) $\\nu_{\\rm c} < \\nu_{\\rm R} < \\nu_{\\rm m}$ in Table 1 is also ruled out since the predicted power-law index $\\alpha = 0.25$ (Sari, Piran $\\&$ Narayan, 1998) is too flat compared to the observed value of $\\alpha_{1} \\sim 0.88$. Therefore, we argue that the possible solution is $\\nu_{\\rm m} < \\nu_{\\rm R} < \\nu_{\\rm c}$ for the time region of $t \\leq t_{\\rm b1}$. In this case, however, if the burst occurred in pre-existing stellar wind, the optical decay slope is predicted to be $\\alpha = \\frac{3\\beta}{2} - \\frac{\\delta}{8-2\\delta} \\sim 1.49$, with $\\delta = 2$ for a wind model (Panaitescu, Meszaros $\\&$ Rees 1998), which is quite steeper than that observed, and hence we can rule out wind-interaction model. In summary, $\\nu_{\\rm m} < \\nu_{\\rm R} < \\nu_{\\rm c}$ and ISM model (case (2) in Table 1) is the only possible solution to reproduce both the temporal/spectral index of the optical afterglow of GRB 030329 at $0.05 < t < 0.26$ days.  In such a slow cooling case, time variation of afterglow flux is given by $F \\propto t^{-3(p-1)/4}$ for $\\nu < \\nu_{\\rm c}$ and $F \\propto t^{-(3p-2)/4}$  for $\\nu_{\\rm c} < \\nu$, respectively (Sari, Piran $\\&$ Narayan 1998). By assuming $\\alpha_{1}$ = 0.88, the electron spectral index is estimated as $p = 2.17$. Note that, this electron spectral index agrees well with those of electrons accelerated in relativistic shock waves (e.g., Dado, Dar $\\&$ Rujula 2003). Furthermore, we expect that power-law slope of the light curve would change from 0.88 to 1.13 for $\\nu_{\\rm c} < \\nu$. Again, this is approximately consistent with the observed spectral index after $t_{\\rm b1}$, where $\\alpha_{2} = 1.18 \\pm 0.01$. Therefore, we conclude that the first break in the optical afterglow light curve at $t_{\\rm b1}$ is the most probably cooling-break where the cooling frequency crosses down the observed optical frequency.  Under this assumption, we can determine important physical parameters for the GRB emission. For example, we can estimate $\\epsilon_{B}$ and  $\\epsilon_{e}$, where $\\epsilon_{B}$ and $\\epsilon_{e}$ are the fractions of the shock energy given to magnetic field and electrons at the shock (Sari, Piran $\\&$ Narayan 1998). In case of slow cooling, $t_{m}<t_{R}<t_{c}$ would be expected. Since we started our observation 0.047 days after the burst, we can limit the range of $t_{m}$ as $t_{m} < t_{obs} = 0.047$ days. For $t_{c} = t_{\\rm b1} = 0.26$ days, $t_{m} < 0.047$ days, $E = 10^{52}$ erg, $n = 1$ cm$^{-3}$, $\\nu = 0.5 \\times 10^{15}$ Hz, we obtain  \\begin{equation} \\epsilon_{B} \\sim  0.05 \\biggl( \\frac{t_c}{0.26} \\biggr)^{-\\frac{1}{3}} \\biggl( \\frac{E_{52}}{1.0} \\biggr)^{-\\frac{1}{3}} \\biggl( \\frac{n_1}{1.0} \\biggr)^{-\\frac{2}{3}} \\biggl( \\frac{\\nu_{15}}{0.5} \\biggr)^{-\\frac{2}{3}} \\end{equation}  \\begin{equation} \\epsilon_{e} < 0.20 \\biggl( \\frac{\\epsilon_B}{0.05} \\biggr)^{-\\frac{1}{4}} \\biggl( \\frac{t_m}{0.047} \\biggr)^{\\frac{3}{4}} \\biggl( \\frac{E_{52}}{1.0} \\biggr)^{-\\frac{1}{4}} \\biggl( \\frac{\\nu_{15}}{0.5} \\biggr)^{\\frac{1}{2}} \\end{equation}  We can also constrain the peak time of reverse shock $t_{A}$ days after the burst (Sari $\\&$ Piran 1999).  \\begin{eqnarray} t_A \\sim 0.03 \\biggl( \\frac{\\epsilon_B}{0.05} \\times \\frac{1}{0.1} \\biggr)^{-3} \\biggl( \\frac{E_{52}}{1.0} \\biggr)^{-1} \\biggl( \\frac{n_{1}}{1.0} \\times \\frac{\\nu_{c,15}}{0.5} \\biggr)^{-2} \\end{eqnarray}   The Lorentz factor depends on time, $\\gamma(t) \\sim (3E/256 \\pi n m_{p}c^{5}t^{3})^{1/8}$ (Piran 1999). And the magnetic field is calculated using $\\epsilon_{B}$ and the Lorentz factor by $B = \\gamma c \\sqrt{32 \\pi \\epsilon_{B} n m_{p}}$. Assuming $E = 10^{52}$erg, the Lorentz factor and the magnetic field strength at two characteristic breaks time are $\\gamma = 9.7, B = 0.86$ gauss at $t_{\\rm b1} = 0.26$ days and $\\gamma = 7.4, B = 0.64$ gauss at $t_{\\rm b2} = 0.54$ days, respectively.  The values of $\\epsilon_{B}$ and $\\epsilon_{e}$ estimated in the preceding chapter are in good agreement with the averages of these parameters for GRBs calculated by Panaitescu $\\&$ Kumar (2001), which are $\\log \\epsilon_{B} = -2.4 \\pm 1.2$ and $\\epsilon_{e} = 0.062 \\pm 0.045$.  \\subsection{Bump at $t \\sim 0.08-0.09$ days}  Finally, we comment on a small ``bump'' of the light curves at $t \\sim 0.08-0.09$ days ($t_{\\rm bump}$) with an amplitude of $\\sim0.1$ mag. Uemura et al. (2003) reported a change of slopes from 0.74 to 0.95 days at $t = 0.085$ days. However our earliest data at $t < 0.08$ days has a slope steeper than 0.74. The light curve at $t > 0.09$ days lies on the extrapolation of this earliest segment. We consider this feature as a bump rather than a break.  Short time variabilities, i.e., ``bumps and wiggles'', may be associated with the forward/reverse shock structures along the afterglow emitting regions (Kobayashi $\\&$ Zhang 2003), repeated energy injection from the central engine, or fluctuation in the density of the interstellar medium (Nakar, Piran $\\&$ Granot 2003).  First, we can rule out a case with the forward/reverse shock structure, since it predicts the light curve should not have the same power-law index before and after the bump. A case with repeated energy injection is also ruled out since after the injection the light curve after the bump should have the same power-law decay slope, but with a larger normalization. Therefore, we conclude that the bumps in the light curve is likely due to the fluctuation in the external density of the interstellar medium (Nakar, Piran $\\&$ Granot 2003).  We can estimate the distance from the central engine $R(t)$ (Piran 1999) and the density variation (Nakar, Piran $\\&$ Granot 2003) at $t_{\\rm bump}$.  \\begin{equation} R(t) \\sim 2.2  \\times 10^{17} \\biggl[ 3 \\biggl( \\frac{E_{52}}{1.0} \\biggr) \\biggl( \\frac{t (\\rm sec)}{7300} \\biggr) \\bigg/ \\pi m_{p} c \\biggl( \\frac{n_1}{1.0} \\biggr) \\biggr]^{\\frac{1}{4}} {\\rm cm} \\end{equation}  \\begin{equation} (n / n_{0}) \\sim 1.1 \\times [ (F_{\\nu}/F_{0}) / 1.1 ] ^{4/(1+p)} \\end{equation}   We find that density is enhanced about 10$\\%$ at the distance of $2.2 \\times 10^{17}$ cm."
    },
    "0311/astro-ph0311615_arXiv.txt": {
        "abstract": "We present new spectropolarimetry of a sample of nearby Compton-thin Seyfert 2 galaxies (ie those with $N_H<10^{23}$cm$^{-2}$).  We show that the detection rate of scattered broad H$\\alpha$ in this sample is considerably higher than in Seyfert 2 galaxies as a whole.  Our results also show that in this low obscuration set it is possible to find scattered broad H$\\alpha$ even when the global properties of the galaxy are largely dominated by the host galaxy and not the active galactic nucleus.  These results argue against the existence of a population of `pure' Seyfert 2 galaxies. ",
        "introduction": "It is now well established that many Seyfert 2 galaxies show characteristics consistent with an embedded Seyfert 1 core (ie the presence of a broad line region) that is hidden from our direct line of sight by surrounding dust.  This is in good agreement with the basic unified model for Seyfert galaxies (Antonucci 1993) which suggests that an optically thick circumnuclear torus, with a scale height at least as large as the broad line region itself, surrounds the active nucleus.  The orientation of this torus relative to our line of sight then determines whether the galaxy is classified as a Seyfert 1 or 2 galaxy.  The detection of scattered broad H$\\alpha$ lines provides the firmest evidence for a hidden broad line region (HBLR) in Seyfert 2 galaxies, as has been demonstrated by, for example, Antonucci and Miller (1985), Miller and Goodrich (1990), Young et al.\\ (1996a), Heisler, Lumsden \\& Bailey (1997), Moran et al.\\ (2000), Lumsden et al.\\ (2001) and Tran (2001, 2003).  A significant fraction of all Seyfert 2 galaxies show evidence for an HBLR (eg.\\ Tran 2001, Moran et al.\\ 2000, Lumsden et al.\\ 2001).  This fraction increases as the luminosity of the active galactic nucleus (AGN) increases (Lumsden \\& Alexander 2001, Gu \\& Huang 2002, Martocchia \\& Matt 2002, Tran 2003).  It is also clear that the galaxies known to have HBLRs tend to have warmer mid-far infrared colours than those galaxies without (Heisler et al.\\ 1997, Lumsden et al.\\ 2001 and Tran 2001, 2003).  Hard x-ray spectroscopy also allows us to investigate the cores of these galaxies even at levels of obscuration that would completely hide the broad line region in the optical (eg Maiolino et al.\\ 1998, Risaliti, Maiolino \\& Salvati 1999). Alexander (2001) pointed out that the combination of the x-ray data and the spectropolarimetry ruled out the model proposed in Heisler et al.\\ (1997) in order to explain the warmth of the infrared colours in the HBLRs.  Heisler et al.\\ suggested that the colour was simply due to the orientation of the torus and hence the obscuration to the infrared emitting zone, so that more obscured systems appeared colder.  By contrast, Alexander showed that the level of x-ray obscuration did not vary on average between galaxies both with and without HBLRs.  It has also been shown that the Seyfert 2s which do not show HBLRs tend to have global properties that are more in keeping with those of their host galaxies (eg.\\ Alexander 2001, Lumsden et al.\\ 2001, Tran 2001, 2003).  However, radically different interpretations have been reached from these data. Tran (2001, 2003) suggests that at least some of the non-HBLRs genuinely lack a Seyfert 1 core.  Nicastro, Martocchia \\& Matt (2003) have suggested that the main difference between those Seyfert 2s showing scattered broad H$\\alpha$ and those without is the accretion rate of the central AGN.  They suggest that the rate is simply too low in the Seyfert 2s without scattered broad H$\\alpha$ to support an extensive broad line region.  They note that at least some of these galaxies have low inferred hard x-ray absorption, suggesting that we should be able to see the broad line region if it exists.  The Nicastro et al.\\ model is clearly compatible with Tran's observations, though the fraction of galaxies without an intrinsic broad line region should be less for suitable luminosity selected samples in the Nicastro et al.\\ model than Tran claims to find.  By contrast, Lumsden \\& Alexander (2001) showed that the luminosity dependence of HBLR detectability was dominant (ie suggesting more luminous active galaxies supported a larger scattering volume).  Lumsden et al.\\ (2001) further suggested that the level of obscuration was still a weak contribution to the appearance of AGN once the effects of the intrinsic AGN luminosity were allowed for, although other studies have not confirmed this result (Alexander 2001, Tran 2001, Gu et al.\\ 2001, Gu \\& Huang 2002).  These observations are compatible with the alternative version of the unified model given by Lawrence (1991).  He suggested that the inner surface of the torus would be pushed outwards from the AGN as the luminosity increased purely because the dust sublimation radius would increase.  In practice an accretion disk has a self shielding effect so the result may not be as dramtic as Lawrence suggested, but it is still a potential factor.  It should also be noted that the findings of Lumsden \\& Alexander are not inconsistent with the Nicastro et al.\\ result, since the accretion rate may be the fundamental parameter determining the AGN luminosity, and the presence of suitable scatterers must scale with the luminosity of the source.  The aim of this paper is to test which of these possibilities is a better match to the actual results for Seyfert 2s with known low obscuration from hard x-ray data.  In these sources obscuration to the scattering sites should be effectively irrelevant if the extinction is reasonably compact.  This should allow us to test the other models outlined above more easily than in previous samples which span the full range of obscuration.  In order to test these ideas we have acquired spectropolarimetry at both high signal-to-noise and higher spectral resolution than in our previous observations.  The latter allows a better separation of possible multiple polarising components. ",
        "conclusions": "Our new results on Compton thin Seyfert 2 galaxies indicate that most show evidence for an HBLR.  The exceptions are NGC~5506, which has very extended obscuration but is known from infrared spectroscopy to have a narrow line Seyfert 1 core, and NGC~6300 for which we do not have sufficient signal-to-noise to definitely state whether an HBLR is present or not.  At least one factor in our success rate in detecting HBLRs is the higher spectral resolution data we have acquired.  This makes it much simpler to detect weak broad scattered lines masked by polarised narrow lines as shown in Section 3. Our results also show counter examples of HBLRs in the `pure' Seyfert 2 region of radio/infrared colour space proposed by Tran (2001, 2003), and for the first time we find HBLRs in galaxies with cool far-infrared colours ($F_{60}/F_{25}>4$).  This strongly suggests that the split proposed by Tran is not as simple as appears from his data.  It is still worth considering briefly the evidence for a class of Seyfert 2s lacking a broad line region.  Nicastro et al.\\ (2003) have shown that at least some of the `pure' Seyfert 2s identified by Tran must have very low accretion rates.  However, some caution must be applied.  X-ray and spectropolarimetric data are rarely contemporaneous, and it is known that there can be wide variations in the observed level of x-ray flux with time (eg Turner et al.\\ 1997).  A good example to consider here is the case of NGC~2992, which has varied in x-ray luminosity by over an order of magnitude in the last twenty years (Gilli et al.\\ 2000), reaching a maximum again recently.  Our spectropolarimetry coincides with this rise to maximum.  At minimum, NGC~2992 would have fallen well below the minimum accretion cut-off suggested by Nicastro et al.  The narrow line region responds much less to such variability in the central engine since it is spatially much larger, and hence the lag between the inner and outer edge tends to smooth out variations.  Longer timescale monitoring of the x-ray properties of the low accretion rate Seyfert 2s would be useful in determining if their lack of an HBLR is due to some intrinsic difference in the Seyfert 2 population as a whole (which we would argue against) or simply that they are in a low state.  Once accretion stops completely such objects would rapidly lose their narrow line regions as well, and cease to classed as Seyfert 2s."
    },
    "0311/astro-ph0311173_arXiv.txt": {
        "abstract": "We report a detailed abundance analysis for {\\he}, a halo giant with $\\mbox{[Fe/H]}_{\\mbox{\\scriptsize NLTE}}=-5.3$. This star was discovered in the course of follow-up medium-resolution spectroscopy of extremely metal-poor candidates selected from the digitized Hamburg/ESO objective-prism survey.  On the basis of high-resolution VLT/UVES spectra, we derive abundances for 8 elements (C, N, Na, Mg, Ca, Ti, Fe, and Ni), and upper limits for another 12 elements. A plane-parallel LTE model atmosphere has been specifically tailored for the chemical composition of {\\he}. Scenarios for the origin of the abundance pattern observed in the star are discussed. We argue that {\\he} is most likely not a post-AGB star, and that the extremely low abundances of the iron-peak, and other elements, are not due to selective dust depletion. The abundance pattern of {\\he} can be explained by pre-enrichment from a zero-metallicity type-II supernova of $20$--$25$\\,M$_{\\odot}$, plus either self-enrichment with C and N, or production of these elements in the AGB phase of a formerly more massive companion, which is now a white dwarf. However, significant radial velocity variations have not been detected within the 52 days covered by our moderate- and high-resolution spectra. Alternatively, the abundance pattern can be explained by enrichment of the gas cloud from which {\\he} formed by a $25\\mbox{M}_{\\odot}$ first-generation star exploding as a subluminous SN~II, as proposed by Umeda \\& Nomoto (2003). We discuss consequences of the existence of {\\he} for low-mass star formation in extremely metal-poor environments, and for currently ongoing and future searches for the most metal-poor stars in the Galaxy. ",
        "introduction": "It has become clear, in recent years, that in order to understand the history of galaxy formation, and indeed, the early evolution of the universe as a whole, it is necessary to understand the nature of star formation at the earliest epochs. A number of recent studies have suggested that the very first stars that formed may have been quite massive objects, several hundred to perhaps one thousand solar masses \\citep{Brommetal:2001,Schneideretal:2002}, while others have suggested that the First Mass Function (FMF) may have included stars with masses as low as roughly the solar mass \\citep{Yoshii/Saio:1986,Nakamura/Umemura:2001}. The most massive stars would have disappeared after some tens of million years, and in this process may have contributed the first heavy element production in the early universe. The less massive stars, especially if these included stars of sufficiently low mass that their main-sequence lifetimes exceed a Hubble time, should still be observable at present. If they were indeed present at the earliest times, these stars should, to a large extent, have preserved in their atmospheres the fossil record of the element production of the most massive stars that were their immediate precursors.  The fundamental role that early-formed, low-mass stars play as the ``scribes'' of stellar generations from long ago has inspired a number of systematic searches for their presence in the Galaxy today. The HK survey of Beers, Preston and Shectman \\citep{BPSII,TimTSS} was the first survey to detect significant numbers of extremely metal-poor (EMP) stars, those which, for the purpose of this paper, we define to be stars of metallicity\\footnote{We use the common notation $\\mbox{[X/H]}=\\log_{10}\\mbox{[$N$(X)/$N$(H)]}_{\\ast}- \\log_{10}\\mbox{[$N$(X)/$N$(H)]}_{\\odot}$, and analogously for [X/Fe]. Abundances are on a scale where $\\log\\epsilon(\\mbox{H})=12$, i.e., $\\log\\epsilon(\\mbox{X}) = \\log\\left\\{N(\\mbox{X})/N(\\mbox{H})\\right\\}+12$.} $\\mbox{[Fe/H]}<-3.0$.  However, even this effort, spanning some two decades, has to date only identified about 100 stars of such extremely low metallicity.  The ESO Large Programme on ``Galaxy Formation, Early Nucleosynthesis, and the First Stars'' of Cayrel et al. \\citep{Cayreletal:2001,Hilletal:2002,Depagneetal:2002,Cayreletal:2003,Francoisetal:2003} has obtained high-resolution, high signal-to-noise ratio ($S/N$) spectra of many of the HK survey stars with $\\mbox{[Fe/H]}<-3.0$. However, it is clear that even after this sample of stars has been analysed, there will remain many open questions, and that new questions will arise.  In order to increase the number of known EMP stars suitable as targets for high-resolution, high $S/N$ spectroscopy with currently existing telescopes, the data base of digital spectra of the Hamburg/ESO objective-prism survey \\citep[HES;][]{hespaperIII} is currently being exploited by means of quantitative selection criteria \\citep{HESstarsI,HESStarsII,MiningTheSky,HESautoclass,Christlieb:2003}. Medium-resolution ($\\sim 2$\\,{\\AA}) spectroscopic follow-up observations of the $8\\,715$ candidate metal-poor stars identified to date (main-sequence turnoff stars and subgiants as well as giants) are being obtained with 1.5m--6.5\\,m telescopes at AAO, CTIO, ESO, KPNO, LCO, Palomar Observatory and SSO. To date, follow-up spectra of $\\sim 3\\,300$ HES metal-poor candidates have been obtained, and $\\sim 200$ stars with $\\mbox{[Fe/H]}<-3.0$ have been identified \\citep{Christlieb:2003}, which triples the number of known EMP stars, from the $\\sim 100$ found in the HK survey, to a total of $\\sim 300$ stars.  Results from high-resolution spectroscopy of HES EMP stars obtained with VLT-UT2/UVES and Keck/HIRES have been reported in \\citet{Depagneetal:2000}, \\citet{KeckpaperI}, \\citet{KeckpaperII}, \\citet{KeckpaperIII}, and \\citet{Cohenetal:2003}.  The lowest metallicity stars in the HK survey reach to $\\mbox{[Fe/H]}=-4.0$, equal to, but not lower than, the metallicity of the lowest metallicity star known prior to the HK survey, {\\cd} \\citep{Bessell/Norris:1984}. Hence, it was widely assumed that a physical low-metallicity limit for Galactic halo stars was reached at around $\\mbox{[Fe/H]}=-4.0$.  However, in a previous paper \\citep[][ hereafter Paper~I]{HE0107_Nature}, we reported the discovery of {\\he}, which is a factor of 20 times lower in [Fe/H] than {\\cd}. In this paper, we describe the derivation of the stellar parameters for {\\he}, and the abundance analysis in more detail (\\S \\ref{Sect:StellarParameters} and \\S \\ref{Sect:AbundAnalysis}, respectively).  In \\S \\ref{Sect:AbundancePattern} we present possible scenarios for the origin of the abundance pattern observed in {\\he}. We conclude our paper with a discussion of some consequences of these scenarios, and consider the impact of the existence of {\\he} on theories of low-mass star formation in the early Universe. ",
        "conclusions": "\\label{Sect:Conclusions}  We have found an overall metal abundance of {\\he}, disregarding the CNO elements, of less than $10^{-5}$ solar, or more than one order of magnitude less than any other known extremely metal-poor giant star. It is tempting to speculate that the chemical composition of this star might reflect nucleosynthesis processes that occured during the very first star formation period of the universe, only about 200\\,Myr from the Big Bang, according to recent WMAP data \\citep{Bennettetal:2003b, Kogutetal:2003}. To what extent this is the case, and for which elements, remains to be explored.  Obviously, an important issue to be investigated in the future is the similarity of the abundance pattern of {\\he} with other stars with $\\mbox{[Fe/H]}<-4.0$, if such stars exist.  From the discussion above we conclude that most elements heavier than Na in {\\he} may be the result of one or a few supernova explosions. We cannot disprove the alternative hypothesis that these elements were gradually accreted e.g. from the interstellar medium in the Galactic disc during passages of the star through it.  Concerning the origin of carbon and nitrogen, one hypothesis is that the star has made them through helium and hydrogen burning and mixed them to the surface as the result of a He and H flash. Also, Na may have been formed this way.  If this self-enrichment scenario is true, a problem remains -- why are the observed amounts of C and N significantly smaller than those predicted by the models? One way around this would be to assume that {\\he} was instead polluted by a binary companion that went through this development.  In this case there would still be a problem with the C/N ratio, which is predicted by current models to be one order of magnitude greater than the observed one, as well as the $^{12}$C/$^{13}$C ratio, which is predicted to be much smaller. A hypothetical companion probably also went through a number of He flashes which might have produced some s-process elements. More precise measurements of the abundance of the s-process elements, as well as oxygen, and radial velocity measurements spanning a longer period of time, will put stronger constraints on possible scenarios for the origin of C and N in the atmosphere of {\\he}.  As already discussed in \\citet{HE0107_Nature}, it has been thought, until very recently, that with the exception of the metal-depleted post-AGB stars (and G77--61, for which no accurate Fe abundance measurement is available), stars of $\\mbox{[Fe/H]}<-4.0$ did not exist. This observational limit is based on the failure of even very extensive wide-angle spectroscopic surveys, like the HK survey of Beers, Preston and Shectman \\citep{BPSII,TimTSS}, to reveal a star more metal-poor than {\\cd}. This failure is easily explainable theoretically, e.g., by the fact that a massive star of, say, $15\\,\\mbox{M}_{\\odot}$ has an evolutionary timescale that is about one order of magnitude shorter than the contraction timescale of a $0.8\\,\\mbox{M}_{\\odot}$ star. It was therefore thought that the massive members of the first generation of stars in our Galaxy, ending up as supernovae of type II (SN~II), might have enriched the ISM to $\\mbox{[Fe/H]}>-4.0$ {\\it before} low-mass stars formed. Alternatively, low-mass star formation might have been suppressed by less efficient cooling due to the absence of metals \\citep[e.g.,][]{Bond:1981}. The discovery of {\\he} can be a challenge to these scenarios.  However, if the first generation of stars was predominantly located in the mass range $20$--$130$\\,M$_{\\odot}$, and exploded as low-explosion-energy supernovae as suggested by \\citet{Umeda/Nomoto:2003}, the influence of the presence of C and O in nearly primordial gas clouds on cooling needs to be taken into account. Umeda \\& Nomoto report that ``the cooling efficiency is large enough to form small mass stars'' in this case.   One prediction of the model of Umeda \\& Nomoto is that {\\he} should have an oxygen abundance of $\\mbox{[O/Fe]}\\sim 2.8$\\,dex. We shall confront this prediction with more recent observations in a forthcoming paper.  A consequence of the scenario of Umeda \\& Nomoto would be that, if there is a critical overall metallicity, $Z_{\\mbox{\\scriptsize crit}}$, which needs to be reached in order to make the formation of low-mass stars possible, \\emph{all} low-mass stars below $\\mbox{[Fe/H]} < \\log\\left(Z_{\\mbox{\\scriptsize crit}}/Z_{\\odot}\\right)$ should be enhanced in CNO relative to Fe and other elements, giving rise to more efficient cooling of the gas, and the subsequent enhanced formation of such stars.  This scenario could be tested by the assemblage of a significant sample of metal-poor stars, extending to $\\mbox{[Fe/H]}<-4.0$, and determination of the fraction of carbon-enhanced stars in that sample.  We note that the Umeda \\& Nomoto scenario is also attractive in that it does not require the existence of a binary companion for the carbon-enhanced stars. In fact, there are at least three unevolved carbon-enhanced metal-poor stars known that do not show any radial velocity variations larger than 0.4\\,km\\,s$^{-1}$ over a period of 8 years \\citep{Preston/Sneden:2001}, ruling out, at least for these three stars, the scenario in which a formerly more massive star went through its AGB phase and transfered dredged-up material onto the surface of the less-massive companion.  It should be realized, when considering the attractive features of the Umeda \\& Nomoto model, that it, unlike the other models discussed, was proposed \\emph{a posteriori}, with knowledge about the peculiarities of {\\he}. At present, we consider it premature to judge whether this model is to be preferred to the SN plus red giant alternative.  {\\he} was found in a sample of roughly $300$ stars with $\\mbox{[Fe/H]}<-3.0$ identified in the Hamburg/ESO and HK surveys. Assuming that {\\he} is not unique, and considering that up to now only roughly $1/3$ of the HES metal-poor candidates have been followed-up on, it is not excluded that one or more additional stars with $\\mbox{[Fe/H]}<-4.0$ will be found in the HES, or other proposed surveys in the near future."
    },
    "0311/astro-ph0311345_arXiv.txt": {
        "abstract": "{ Characteristic scale lengths of nonthermal X-rays from the SN~1006 NE rim, which are observed by {\\it Chandra}, are interpreted in the context of the diffusive shock acceleration on the assumption that the observed spatial profile of nonthermal X-rays corresponds to that of accelerated electrons with energies of a few tens of TeV. To explain the observed scale lengths, we construct two simple models with a test particle approximation, where the maximum energy of accelerated electrons is determined by the age of SN~1006 ({\\it age-limited model}) or the energy loss ({\\it energy loss-limited model}), and constrain the magnetic field configuration and the diffusion coefficients of accelerated electrons. When the magnetic field is nearly parallel to the shock normal, the magnetic field should be in the range of 20--85~$\\MG$ and highly turbulent both in upstream and downstream, which means that the mean free path of accelerated electrons is on the order of their gyro-radius (Bohm limit). This situation can be realized both in the age-limited and energy loss-limited model. On the other hand, when the magnetic field is nearly perpendicular to the shock normal, which can exist only in the age-limited case, the  magnetic field is several $\\MG$ in the upstream and 14--20 $\\MG$ in the downstream, and the upstream magnetic field is less turbulent than the downstream. ",
        "introduction": "Galactic cosmic rays with energies of less than $10^{15.5}$~eV (the ``knee'' energy) are commonly believed to be generated by supernova remnants (SNRs). SN~1006 is one of the SNRs which is thought to be an accelerator of such high energy particles. Koyama et al. (1995) discovered synchrotron X-rays from the rims of this SNR, indicating the existence of accelerated electrons with an energy more than a few tens of TeV. The detection of TeV $\\gamma$-rays from the northeastern (NE) rim of SN~1006 (Tanimori et al. 1998) implies further evidence for the presence of high energy particles, since TeV $\\gamma$-rays arise from the Inverse Compton (IC) process, in which cosmic microwave background (CMB) photons are up-scattered by high energy electrons (Tanimori et al. 2001), or the hadronic process, in which $\\pi^0$ particles made by collisions between accelerated and interstellar protons decay into $\\gamma$-ray photons (Berezhko, Ksenofontov, \\& V\\\"{o}lk 2002; Aharonian \\& Atoyan 1999).  The mechanism for cosmic ray acceleration has also been studied for a long time and the most plausible process is a diffusive shock acceleration (DSA) (Bell 1978; Blandford \\& Ostriker 1978; Drury 1983; Blandford \\& Eicher 1987; Jones \\& Ellison 1991; Malkov \\& Drury 2001). Many authors have explained the observed properties of SN~1006 in the context of the DSA but the conclusions are different due to the arbitrary assumptions on unknown physical parameters, such as the magnetic field configuration, the diffusion coefficient, the injection rate, and the electron to proton ratio (Ellison, Berezhko, \\& Baring 2000; Berezhko, Ksenofontov, \\& V\\\"{o}lk 2002; V\\\"{o}lk, Berezhko, \\& Ksenofontov 2003; Aharonian \\& Atoyan 1999; Reynolds 1998; Dyer et al. 2001; Allen, Petre, \\& Gotthelf 2001; Achterberg, Blandford, \\& Reynold 1998). For example, at present the origin of TeV $\\gamma$-rays can be explained by both leptonic and hadronic models. This comes from an insufficient theoretical understanding; apart from a globally successful picture of the DSA, detailed but important processes, such as the injection or the reflection of accelerated particles that determines the above unknown quantities, are not well understood. Worse yet, previous observations in the hard X-ray band had insufficient spatial resolutions to resolve small-scale structures near the shock front, and could not strongly constrain the theoretical parameters.  Recently, Bamba et al. (2003a,b) reported on the results for not only spectral but spatial studies of thermal and non-thermal shock structure in the NE rim of SN~1006 with the excellent spatial resolution of {\\it Chandra}. Similar results are also reported by Long et al. (2003) with {\\it Chandra} data. Bamba et al. (2003a) estimated the scale length of thermal and non-thermal X-rays in both upstream and  downstream of the shock front, which means that the direct measurement of the diffusion coefficients has become possible.  In this paper, we show that the important physical parameters for the magnetic field in the acceleration site can be constrained by the spatial distribution of observed nonthermal X-rays. \\S~\\ref{section:obs} summarized results of data analyses by Bamba et al. (2003a). We construct two models in \\S~\\ref{sec:model}; one is assumed that the maximum energy of accelerated electrons are determined by the age of SNR, while the other is by the energy loss process such as synchrotron or IC cooling. Finally, \\S~\\ref{section:dis} is devoted to the  discussion on validity of our estimation. Throughout the paper, indices ``u'' and ``d'' represent upstream and downstream, respectively. ",
        "conclusions": "\\label{section:dis} We have argued the limitations on the model of DSA from the recently observed scale length of nonthermal X-ray emitting region of NE shell of SN~1006 based on the assumption that the observed spatial structure of nonthermal X-ray filaments reflects that of accelerated electrons to the maximum energy. A test particle approximation has been adopted, where the back reactions of accelerated particles are neglected, although one might have to consider corrections due to nonlinear effects in order to explain the wide-band spectrum from  radio to  TeV $\\gamma$-rays. Two models have been discussed; age-limited and energy loss-limited models. Note that in each model, the value and configuration of a magnetic field in the acceleration site can be discussed using only the spatial distribution of the synchrotron X-rays, which are emitted by accelerated electrons with energy several tens of TeV, and the roll-off frequency, which is derived by the fitting of the wide band spectrum from the radio to X-rays.  When a magnetic field is nearly parallel to the shock normal, $\\eta_\\u$ and $\\eta_\\d$ should be nearly unity, which means that the magnetic field is highly turbulent (near the Bohm limit), and the magnetic field is in the range of 20--85~$\\MG$. Relatively strong parallel magnetic fields are considered by several authors (Berezhko et al. 2002; Ellison et al. 2000). Derived value of the magnetic field seems to be higher than the usual interstellar value of a few micro Gauss. However, the mechanism pointed by Lucek \\& Bell (2000) may be able to amplify the magnetic field. These situation can be realized both in the age-limited and energy loss-limited model.  When the magnetic field is nearly perpendicular to the shock normal, relatively small values of $B_\\d=4B_\\u=14$--20~$\\MG$ are allowed, which can exist only in the age-limited model. This value is consistent with the interstellar magnetic field and assumed value in Bamba et al. (2003a) or other papers (Allen et al. 2001; Reynolds 1998) but slightly larger than that derived by CANGAROO observation (Tanimori et al. 1998, 2001), which assumes TeV $\\gamma$-rays are emitted by the IC process. This discrepancy may be solved if the filaments~1--6 are not the main sites of TeV $\\gamma$-rays generated by the leptonic process. The observed flux of TeV $\\gamma$-rays which are up-scattered CMB photons by synchrotron emitting electrons is written as \\begin{equation} {\\cal F}_{\\rm IC}=2.5\\times10^{-2} \\left(\\f{B_\\d}{20\\ \\mu {\\rm G}}\\right)^{-2} {\\cal F}_{\\rm synch} \\ , \\end{equation} where ${\\cal F}_{\\rm synch}$ is the observed flux of synchrotron X-rays and estimated as $1.8\\times10^{-12}$~ergs~s$^{-1}$~cm$^{-2}$ in the 0.5--10.0~keV band by \\cite{bamba2003}. Then, the contribution of the IC $\\gamma$-rays generated in the filaments~1--6 is only $\\sim0.5$\\% of the whole flux coming from SN~1006. This shows that the IC $\\gamma$-ray emitting sites have smaller magnetic field,  hence radiate less synchrotron X-rays than in the filaments~1--6. Or the TeV $\\gamma$-rays might be emitted by the hadronic process. In fact the most plausible position of the $\\gamma$-ray emission determined by CANGAROO is located north (and outside) of the filament~1 (Tanimori et al. 1998, 2001). However, the region where the significance is higher than half the maximum value extends over $\\sim 0.2^{\\circ}$, which is almost the same as the standard deviation of the point-spread function of the CANGAROO telescope. We need stereoscopic observations of TeV gamma-ray emissions by imaging atmospheric Cherenkov telescopes (for example, CANGAROO-III, HEGRA, and H.E.S.S.) to establish the position and extent of the emission region.   We briefly discuss how our results may change if the test particle approximation is dropped. Although we should consider the spatial structure modified, for example, in the upstream precursor region, we simply argue the cases of $r>4$ that are thought to be caused by nonlinear effects. Indeed, Berezhko, Ksenofontov, \\& V\\\"{o}lk (2002) showed that the present value of total compression ratio is  about 6. In the age-limited case, we observe that the inferred magnitude of the upstream magnetic field becomes smaller than that of the $r=4$ case as $r$ is made larger (up to 7).  However, the overall shapes and total areas of the allowed regions are changed only slightly for any case of $\\theta$. Thus the nonlinear effect would not be essential. On the other hand, in the energy loss-limited case, the allowed region in the $\\eta_\\u$--$\\eta_\\d$ plane becomes narrow and only the cases of $\\theta=0^\\circ$--10$^\\circ$ can be exist for $4<r\\la6$ while no allowed region for $r\\sim7$ case. This comes from the fact that another restriction, Eq.~(\\ref{eq_coolFG}), emerges. The larger $r$, the stronger the constraint because of the less efficient advection for fixed shock velocity $u_\\s$. Furthermore, it can be shown that there exists no case in which $w_\\d$ is given by $(K_\\d t_\\co)^\\Half$, since the condition $t_\\ac=t_\\co$ is incompatible with the condition $u_\\d t_\\co<(K_\\d t_\\co)^\\Half$ in the parameter range of our interest.   It is important to determine the magnetic field configurations in order to discuss the acceleration and/or injection efficiency. In this paper, using the spatial distribution of nonthermal X-rays, we have shown  that roughly two cases can exist; high and parallel, and low and perpendicular magnetic field. In the latter case, the back reaction of accelerated particles is small and thus a test-particle treatment is a good approximation. The magnetic field amplification process  pointed by Lucek \\& Bell (2000) does not work well. While in the former case, nonlinear effects are so efficient that the magnetic field can be large. The difference between these cases probably comes from the fact that the (proton) injection rate depends strongly on the shock obliquity and diminishes as $\\theta$ increases (V\\\"{o}lk, Berezhko, \\& Ksenofontov 2003). In addition to our result, radio polarization data with high spatial resolution may become further informations about the magnetic field configuration.  In this paper, we have adopted the plane shock approximation. For further details, it might be important to consider the curvature effect as discussed in Berezhko, Ksenofontov, \\& V\\\"{o}lk (2003) in order to produce more realistic scenario."
    },
    "0311/astro-ph0311035_arXiv.txt": {
        "abstract": "Observations of accreting systems often show significant variability (10--20 percent of accretion luminosity) on timescales much longer than expected for the disc regions releasing most of the luminosity. We propose an explicit physical model for disc variability, consistent with Lyubarskii's (1997) general scheme for solving this problem. We suggest that local dynamo processes can affect the evolution of an accretion disc by driving angular momentum loss in the form of an outflow (a wind or jet). We model the dynamo as a small--scale stochastic phenomenon, operating on roughly the local dynamical timescale. We argue that large--scale outflow can only occur when the small--scale random processes in neighbouring disc annuli give rise by chance to a coherent large--scale magnetic field. This occurs on much longer timescales, and causes a bright large--amplitude flare as a wide range of disc radii evolve in a coherent fashion. Most of the time, dynamo action instead produces small--amplitude flickering. We reproduce power spectra similar to those observed, including a $1/f$ power spectrum below a break frequency given by the magnetic alignment timescale at the inner disc edge. However the relation between the black hole mass and the value of the break frequency is less straightforward than often assumed in the literature. The effect of an outer disc edge is to flatten the spectrum below the magnetic alignment frequency there. We also find a correlation between the variability amplitude and luminosity, similar to that found in some AGN. ",
        "introduction": "Both Galactic and extragalactic accretion--powered X--ray sources display significant aperiodic variability on a wide range of timescales. Similarities in behaviour between the black hole candidate X--ray binaries (BHCs) and the nuclei of Seyfert galaxies (that presumably also harbour central black holes) have been used to suggest that the physical processes occuring in accretion flows around the black holes in these objects are essentially the same, but scaled in some appropriate manner. There are indeed suggestions that such scaling relationships might be used as an estimate for black hole masses (e.g.\\ Vaughan, Fabian \\& Nandra, 2002; Markowitz et~al.\\ 2003). Although X--ray binaries containing neutron stars as the accreting compact objects also exhibit interesting aperiodic, and quasi-periodic, behaviour (see for example the review by van der Klis 1989), we restrict our attention in this paper to black hole accretion discs, partly in the hope that understanding the physics of such accretion flows might be somewhat less complicated, but also because of the interest in scaling between stellar and supermassive objects.  We stress, however, that the processes we discuss in this paper are fully relevant not only to binary X--ray sources and AGN, but to all objects containing accretion discs, such as cataclysmic variables and protostellar objects.  Two main diagnostics are relevant for the variability of these sources. These are the shape of the power spectrum and the amplitude of the variability. Of course, the latter can be deduced from the normalisation of the former (e.g.\\ Miyamoto et~al.\\ 1992) and so strictly speaking, these are just two aspects of the same diagnostic.  The best observed black-hole candidate X--ray binary (BHC) is Cygnus X--1. In the low/hard state the normalised rms variability, $\\bar{\\sigma}$, is of order a few tens of per cent and the power spectrum can be approximated as a broken power law of the form $p(f) \\propto f^{-\\alpha}$ where $\\alpha \\simeq 0$ for $10^{-3}$~Hz $< f < 0.1$~Hz, $\\alpha \\simeq 1$ for $0.1$~Hz $< f < f_{\\rm br} \\simeq 3$~Hz, and $\\alpha \\simeq 2$ for $ f_{\\rm br} < f < 10^3$~Hz (Nowak et~al.\\ 1999; Revnivtsev, Gilfanov \\& Churazov 2000; Churazov, Gilfanov \\& Revnivtsev 2001). This implies that the flickering variability occurs predominantly on timescales in the range 0.3~s -- 10~s. In the high/soft state the normalised rms variability, $\\bar{\\sigma}$, is of the order of a few per cent, and the spectrum is a similar, broken power law, except that now $f_{\\rm br} \\simeq 10$~Hz, and $\\alpha \\simeq 1$ for $f < f_{\\rm br}$. In this state, therefore, the flickering occurs predominantly on timescales $\\ga 10$~s. Other BHCs show similar behaviour. The power spectra of GX339-4 and GS2033+338 in the low/hard state are given by Miyamoto et~al.\\ (1992) who note their strong similarity to Cyg X--1. The power spectrum of LMC X--1 is shown by Nowak et~al.\\ (2001) to be similar to Cyg X--1 in the high/soft state, and to have a normalised rms variability of $\\bar{\\sigma} \\simeq 7$ per cent.  X--ray power spectra of Seyfert 1 galaxies also typically display a broken power law, with the break occurring at frequency $f_{\\rm br} \\simeq 10^{-6}$~Hz, corresponding to a timescale of the order of a few days (Markowitz et al. 2003). For $f < f_{\\rm br}$ the slope of the spectrum has typically $\\alpha \\simeq 1$ and for $f > f_{\\rm br}$ the slope of the spectrum has typically $\\alpha \\simeq 2$ (see also Vaughan \\& Fabian 2003; Vaughan et al.\\ 2003). The normalised rms variability amplitudes range typically up to $\\bar{\\sigma} \\simeq 10$ per cent (Markowitz et~al.\\ 2003).  Uttley and McHardy (2001) have drawn attention to a linear correlation between the actual rms variability, $\\sigma$, and the flux, $F$, in both BHCs and Seyfert~1 galaxies, implying that the normalised rms variability $\\bar{\\sigma}=\\sigma/F$ is approximately independent of source brightness, at least over a small range in flux. From the data for the BHCs Cyg X--1 and SAX J1808.4-3658, given in the paper by Uttley \\& McHardy (2001), we find that the relationship is approximately of the form $\\bar{\\sigma} \\propto F^q$ where $q \\simeq 0.3$.  The earliest models for the variability were in terms of shot noise (Terrell 1972); the light curves were reconstructed from a series of independent overlapping shots. Such a picture, or variations thereof (e.g.\\ Lehto 1989), can be made to model the data reasonably well, in particular the shape of the power spectrum. However, this scenario still lacks a physical picture for the variability. Since then, various ideas have been put forward, including accretion disc turbulence (Nowak \\& Wagoner, 1995), magnetic flares in the disc corona (e.g.\\ Poutanen \\& Fabian, 1999), and MHD instabilities in the plunging region of the inner disc (Hawley \\& Krolik, 2001). These models all have a common problem in accounting simultaneously for the range of timescales and the normalised rms variability amplitudes. The latter are typically around 10 per cent, and imply emission from the inner few disc annuli, where most of the accretion energy is released (see for example the discussion by Bruch, 1992; Fritz \\& Bruch, 1998). But these models also imply that characteristic variability takes place on local dynamical timescales. At the inner radii of accretion discs around $\\sim 10 M_{\\sun}$ black holes, these relevant timescales are far too short, typically of the order of milliseconds.  An insightful paper by Lyubarskii (1997) offers a way out of these difficulties. Lyubarskii considers small amplitude local fluctuations in the accretion rate at each radius, caused by small amplitude variations in the viscosity, and then considers the effect of these fluctuations on the accretion rate {\\it at the inner disc edge}. A linear calculation shows that if the characteristic timescale of the viscosity variations is everywhere comparable to the viscous (inflow) timescale, and if the amplitude of the variations is independent of radius, then the power spectrum of luminosity fluctuations has the form $p(f) \\propto 1/f$. If the amplitude of the variations increases with radius, the slope of the power spectrum of the luminosity variations is steeper than $-1$. Lyubarskii (1997) notes that he has no physical model for the cause of such fluctuations. In particular, although the obvious candidate cause is the magnetic dynamo, the characteristic timescales for the dynamo are much shorter than the local viscous timescale.  In this paper we use the idea (Livio, Pringle \\& King 2003) that although the local dynamo timescale at each point in the disc is indeed short, the timescale for the dynamo processes in a sufficient number of independent neighbouring disc annuli to produce a poloidal magnetic field coherent enough to affect the accretion rate, by generating a disc wind or jet, can be quite long. That a disc wind/jet can be a significant driver of accretion in such objects is in line with the recent ideas about `jet--dominated' states in black hole candidate binaries proposed by Fender, Gallo \\& Jonker (2003). We propose a specific, but highly simplified, model of an accretion disc. A magnetic dynamo process generates the viscosity, but also, from time to time, produces a sufficiently well--ordered poloidal field. Therefore, the angular momentum is transferred outwards by two separate processes --- the usual magnetic viscosity and a disc wind or jet. Numerical simulation of all the physical processes involved in such an intricate system is, of course, unrealistic at the current time, and we make no attempt to describe the full MHD problem in its three--dimensional complexity (e.g.\\ Hujeirat, Camenzind \\& Livio 2002).  Instead, we make a number of simplifying assumptions about the physical processes involved, and use these to construct a numerical description of the time-dependence of such an accretion disc. We show below that this straightforward approach, based on simple physical ideas, can describe the flickering behaviour observed in accreting black hole systems. ",
        "conclusions": "In this paper we have proposed a simple model to describe how the local dynamo process in an accretion disc, which produces the effective disc viscosity, can also affect the disc evolution by driving angular momentum loss in the form of a wind or jet. While the disc viscosity produces the usual smooth inflow, the accretion rate driven by the wind/jet is highly stochastic. This comes about because of our assumption that efficient angular momentum loss to an outflow can only occur when the local poloidal field produced by the disc dynamo is sufficiently coherent. We assume that each disc annulus (of size comparable to the local disc thickness $H$) produces a poloidal field whose size and direction varies as a random walk with some characteristic timescale (Markoff process). Only when a number $\\sim R/H$ of such neighbouring annuli produce sufficiently aligned poloidal fields can outflow and angular momentum loss occur.  We have implemented these ideas in a simple set of equations that determine the evolution of the disc, and hence the disc and jet(outflow) luminosities. We have chosen the parameters of the model ($B_s$ and $f_{\\rm LPP}$) such that the poloidal field strength produced by the disc, while smaller than the tangled field within the body of the disc, is locally large enough to give a significant outflow. Thus, typically, in our models, $L_{\\rm jet}$ and $L_{\\rm disc}$ are comparable in magnitude. From an observational point of view this is not unreasonable (e.g.\\ Fender et~al.\\ 2003; D'Elia, Padovani \\& Landt 2003).  We have chosen simple parameters for our disc models, with constant ratio $H/R$ and constant viscosity parameter $\\alpha$. In addition, we have assumed that local dynamo timescale $\\tau_d$ is a constant $\\times$~the local dynamical timescale $\\Omega^{-1}$ in the disc. These assumptions imply that the disc is self--similar and that the local viscous timescale $\\tau_\\nu$ is proportional to the timescale $\\tau_{\\rm mag}$ on which the local poloidal field alignment occurs.   We find that in general the disc luminosity shows irregular behaviour. Small amplitude flickering is superposed on variations in the average luminosity, together with occasional large amplitude flares. We have investigated the properties of the flickering with the goal of obtaining an understanding of what the flickering might be able to tell us about the properties of the disc. Lyubarskii (1997) showed that the power spectrum has the form $p(f) \\propto 1/f$ if the magnitude of viscosity varies on a timescale everywhere the same as the local viscous time ($\\tau_\\nu = 1/f_\\nu)$. In our case the first timescale is that for angular momentum loss in the outflow $\\tau_{\\rm mag} = 1/f_{\\rm mag}$. This is the case shown in Figure 1, where $f_\\nu = f_{\\rm mag}$. The $1/f$ spectrum continues up to a break frequency $f_{\\rm br}$ corresponding to the inner disc edge, beyond which the spectrum steepens. For cases in which the local magnetic alignment timescale was longer than the local viscous time (i.e.\\ $f_{\\rm mag} \\la f_\\nu$) we find a similar behaviour, with the frequency $f_{\\rm br}$ corresponding to the local magnetic alignment timescale at the inner disc edge. We do not agree with the assumption found in the literature (e.g.\\ Vaughan \\& Fabian 2003) that the break frequency corresponds to the viscous timescale at a particular radius. If the local magnetic alignment timescale is everywhere shorter than the local viscous timescale, the spectrum is steeper than $1/f$ throughout (i.e.\\ red noise). We have also shown that the effect of an outer disc edge is to produce a flattening of the spectrum for frequencies $f < f_o$, where $f_o$ is the magnetic alignment frequency at that edge.  We have investigated how the amplitude of the flickering varies with luminosity and with $\\alpha$. We find that in any one run, i.e.\\ at fixed $\\alpha$, the normalised rms amplitude is approximately constant, in line with the findings of Uttley \\& McHardy (2001) from observation. We also show that the normalised amplitude of the flickering increases with increasing $\\alpha$, varying in our models from about 3~per cent for $\\alpha = 0.006$ to almost 30 per cent for $\\alpha = 0.6$. This probably comes about because the larger the viscosity, the more rapidly stochastic behaviour can be swept inwards from the outer disc regions. Thus, it may be that large amplitude flickering is a symptom of a large disc viscosity.  We have also demonstrated the possibility of angular momentum loss to an outflow giving rise to avalanche--like behaviour resulting in a large flare in luminosity. This is because angular momentum loss directly to an outflow changes the characteristic nature of the equation governing the evolution of the disc from diffusive to wave--like. We note that much of the flickering behaviour also comes about because of the wave--like nature of the response. This not only enhances, but also keeps coherent, any perturbations to the disc.  Finally, we stress again that, although we have applied our findings to accretion discs around black holes (BHCs and AGN), the whole range of phenomena found here, and much more (for example, interactions between the disc and the central object's surface and/or its magnetosphere), are likely to occur in all systems in which accretion discs are found."
    },
    "0311/astro-ph0311203_arXiv.txt": {
        "abstract": "We use gravitational lenses from the Cosmic Lens All-Sky Survey (CLASS) to constrain the evolution of galaxies since redshift $z \\sim 1$ in the current $\\LCDM$ cosmology. This constraint is unique as it is based on a mass-selected lens sample of galaxies. Our method of statistical analysis is the same as in Chae (2003). We parametrise the early-type number density evolution in the form of $(1+z)^{\\nu_n}$ and the velocity dispersion as $(1+z)^{\\nu_v}$. We find that $\\nu_n=-0.11^{+0.82}_{-0.89}$ ($1\\sigma$) if we assume $\\nu_v =0$, implying that the number density of early-type galaxies is within 50\\% to 164\\% of the present-day value at redshift $z=1$. Allowing the velocity dispersion to evolve, we find that $\\nu_v=-0.4^{+0.5}_{-0.4}$ ($1\\sigma$), indicating that the velocity dispersion must be within $57\\%$ and $107\\%$ of the present-day value at $z=1$. These results are consistent with the early formation and passive evolution of early-type galaxies.  More stringent limits from lensing can be obtained from future large lens surveys and by using very high-redshift quasars ($z \\ga 5$) such as those found from the Sloan Digital Sky Survey. ",
        "introduction": "Currently there are about 70 multiply-imaged systems due to galactic mass scale gravitational lenses. The statistics of gravitational lenses depend  on three key ingredients, namely, the cosmology, the number density of potential lenses as a function of redshift and the dynamical properties of galaxies (e.g., velocity dispersions and the surface mass densities). Gravitational lenses hence encode information of the cosmology, the galaxy mass profiles and the evolution history of galaxies. At present, the lens sample is too small to constrain all the ingredients simultaneously. Most previous studies concentrated on constraining the cosmological constant assuming non-evolving populations of lenses (e.g., Fukugita et al.\\ 1992; Kochanek 1996; Helbig et al.\\ 1999). Under this assumption, the most recent lens statistics study of Chae et al.\\ (2002) finds that the lens statistics are best-fitted by a present-day matter density of $\\omega0 \\approx 0.3$ and a cosmological constant of $\\lambda0 \\approx 0.7$. This result is consistent with results from a variety of other studies, including the cosmic microwave background radiation (e.g., de~Bernardis et al.\\ 2000; Spergel et al. 2003), Type~Ia supernovae at cosmological distances (e.g., Riess et al. 1998; Perlmutter et al. 1999), and the large-scale structures in the universe (e.g., Peacock et al. 2001).  In light of the convergence of the cosmological model, it becomes important to use gravitational lenses for a different purpose: to study the evolution of galaxies (their number density and dynamical properties) in the $\\omega0=0.3, \\lambda0=0.7$ cosmology (hereafter $\\Lambda$CDM). Lensing limits on galaxy evolution have been explored by Mao (1991), Mao \\& Kochanek (1994), Rix et al.\\ (1994), and Jain et al.\\ (2000), all of which used optically-selected lenses. Lensing is sensitive to the evolution of galaxy properties as the lensing probability is $\\propto n\\sigma^4$, while separations are $\\propto \\sigma^2$; here $n$ is the number density and $\\sigma$ is the velocity dispersion of typical lenses. For example, (for a fixed cosmological model) a decreasing number density of galaxies with redshift($z$) lowers the lensing rate and the mean redshift of lenses while a decreasing velocity dispersion with $z$ lowers the lensing rate, the mean lens redshift and the mean angular size of image separations. Therefore, through a careful analysis of the lens redshifts, image separations and lensing probability we can constrain the evolution of the number density and dynamical properties of galaxies.  Most lensing galaxies are massive early-type galaxies as they dominate the lensing cross-sections due to their larger central mass concentrations. Gravitational lenses therefore provide a unique mass-selected sample to study the evolution of early-type galaxies, independent of and complementary to the traditional redshift surveys of galaxies (e.g., Fried et al. 2001; Im et al. 2002). This is a much debated research area. There exist two different views on the formation and evolution of early-type galaxies, namely a monolithic collapse model (Eggen, Lynden-Bell, \\& Sandage 1962) and a merger hypothesis (Toomre \\& Toomre 1972). In the monolithic collapse model, early-type galaxies are thought to have formed rapidly at high redshift and then evolve passively to the present-day. The merger model is a natural consequence of the hierarchical structure formation theory. Semi-analytic implementations of this model predict a continuous formation of ellipticals and hence a certain fraction of massive early-type galaxies must have formed since $z \\sim 1$ (e.g.\\ Kauffmann 1996; Baugh et al.\\ 1996; Kauffmann et at.\\ 1999); the fraction depends on the assumed cosmology and other assumptions and is typically one third or more (for more, see \\S4). An observational way of probing galaxy evolution is using redshift survey of galaxies. However, no consensus has been reached either with this method. For example, Kauffmann, Charlot, \\& White (1996) found rapid evolutions in the number density of ellipticals while Schade et al.\\ (1999) advocated the opposite conclusion using similar samples (see \\S4 for more details).  The purpose of this work is to use data from the Cosmic Lens All-Sky Survey (CLASS) to provide independent constraints on galaxy evolution. As we were completing this work, a complementary study has been carried out by Ofek, Rix, \\& Maoz (2003); their results are compared with our results in \\S2. ",
        "conclusions": "We have used the statistical properties of the CLASS strong lens sample (i.e., the rate of multiple-imaging and the image separations as a function of redshift) to constrain the evolution of galaxies. The lens sample is unique as it is mass-selected and hence the constraints obtained from it are independent of those from redshift surveys. The method we use is based on Chae et al.\\ (2002) and Chae (2003) and has some uncertainties, such as the adopted luminosity function of early-type galaxies and the redshift distribution of the source population in the CLASS survey (see Chae 2003 for details). However, these uncertainties are smaller than those arising from the moderate-size CLASS sample of lenses.  We find that the (comoving) number density of lensing galaxies at redshift $z=1$ must be within 50\\% to 164\\% of the present-day number density; their characteristic velocity dispersion also must be within $57\\%$ and $107\\%$ of the present value. The lensing statistics are therefore consistent with a slow evolution of galaxies in both their number density and their dynamical properties. These results are inconsistent with very fast evolution of early-type galaxies, where the number of early-types at $z = 1$ is only 20\\%-40\\% or less of the present-day values (e.g., Lin et al. 1999; Fried et al. 2001; Wolf et al.\\ 2003). Our results are, however, consistent with several other studies, in particular the study based on the Hubble Space Telescope observations of the Groth Strip (Im et al. 2002) where there find a number density evolution of $n(z) \\propto (1+z)^{-0.86\\pm 0.68}$ (see their Table 6)  in the same underlying cosmology. In the Standard Cold Dark Matter model with $\\omega0=1$, the number density of bright ellipticals at $z=1$ is predicted to be a factor of 2-3 smaller than the local value (Kauffmann 1996; Baugh, Cole \\& Frenk 1996), inconsistent with our results. However, the evolution of galaxies is expected to be slower in the $\\LCDM$ cosmology. Indeed, our results are consistent with the modest evolution out to redshift $\\sim 1$ predicted by Kauffmann et al.\\ (1999; see their {Fig.}~9) for the same cosmology. Quantitatively, for ellipticals with a stellar mass $\\ga 10^{10} M_\\odot$, the space density at $z=1$ is about 30\\% lower than the present day value in the fiducial model of Cole et al.\\ (2000). This result is robust to changes in the amount of star formation occurring in bursts at high redshift (C. Baugh, private communication; Baugh et al.\\ 1996). Our results are consistent with these predictions.  The effects of galaxy evolution on the lens statistics will be more dramatic for very high-redshift quasars ($z \\ga 5$, Fan et al. 2003). For a source at $z=5$, half of the multiple-imaging cross-section (with separations between 0.3 to 5 arcseconds) is contributed by galaxies with redshifts greater than 1.2 if there was no evolution of galaxies. Hence the lensing probability will be reduced significantly if the comoving number density of early-type galaxies is much smaller at redshifts $z \\ga 1$ compared with the local number density. A large sample of very high-redshift quasars is therefore an independent and effective way of probing galaxy evolution.  The lensing constraints on the galaxy evolution are already quite competitive compared with those from other methods. However, they still suffer from the small number of lenses ($13$) in our CLASS statistical sample. With planned upgrades in major radio instruments (such as eVLA and e-Merlin), it is possible to obtain a radio lens sample that is an oder of magnitude larger. When such a sample becomes available, the lensing constraint will become more stringent and more physical evolution (rather than our toy) models can be tested."
    },
    "0311/astro-ph0311217_arXiv.txt": {
        "abstract": "New 21cm observations with the Green Bank Telescope (GBT) show that a significant fraction of the HI in the inner Galaxy's halo $\\sim 1$ kpc from the midplane exists in the form of discrete clouds. Some look very much like a Spitzer (1968) ``standard'' diffuse cloud but with their HI in two phases. They mark the transition between the neutral disk and the highly ionized halo. The dominant motion of the clouds is Galactic rotation, but some have  random velocities of as much as 50 km~s$^{-1}$.  They are part of the Galaxy and are not related to high-velocity clouds, yet their origin is obscure. ",
        "introduction": "In the inner parts of the Milky Way most of the hydrogen is confined to a narrow disk which establishes a plane of fundamental physical significance (Blaauw et al. 1960).   But the more we look out of the disk, the more HI we see at curious locations. Many mechanisms  can lift HI from below, and the Milky Way, like other galaxies,  is probably accreting gas from above, so the halo is likely an untidy place, containing HI  that  1) has been pushed outward by mechanical and photon energy from young stars and supernovae, 2)  has condensed from a halo of very hot gas and now returns in a cool phase to the plane, 3) is falling into the Milky Way for the first time, having been stripped from neighboring galaxies like the Magellanic Stream, or coming in as  high-velocity clouds (Shapiro \\& Field 1976; Heiles 1984; Norman \\& Ikeuchi 1989; Houck \\& Bregman 1990; Wakker et al. 1999; Sancisi et al. 2001; de Avillez \\& Berry 2001; Konz et al. 2002; Putman et al. 2003).  This paper discusses new observations of the HI at some distance from the Galactic plane which show the organization of halo HI for the first time, but at present raise more questions than answers.    \\begin{figure} \\plotfiddle{lockman_fig1.ps}{6.45in}{0}{80}{80}{-280}{-75} \\caption{HI emission above the Galactic plane in a region around $\\ell = 25\\deg$ as observed with the GBT at $9'$ angular resolution.  The data cover  $V_{LSR} = 134 \\pm 1$ km~s$^{-1}$, which is at the tangent point in this direction, where one degree in either coordinate corresponds to a linear distance of 135 pc.  The clouds at $b \\approx +7\\fdg6$ are 1 kpc above the plane. } \\end{figure} ",
        "conclusions": ""
    },
    "0311/astro-ph0311021_arXiv.txt": {
        "abstract": "{We have used the Westerbork Synthesis Radio Telescope  to detect HI 21cm absorption at $z \\sim 0.7645$ in the gravitational lens system towards \\pmn. The 21cm profile has two broad components, with peak optical depths of $0.047 \\pm 0.007$ and $0.039 \\pm 0.007$, at heliocentric redshifts $0.76470 \\pm 0.00006$ and $0.76348 \\pm 0.00006$, respectively. The redshift of the stronger component matches that of CaII H and K absorption detected earlier. The absorption has a total velocity width of $\\sim 500$~\\kms~(between nulls) and an equivalent width of $7.1 \\pm 0.08$~\\kms. This would imply a total HI column density of $2.6 \\pm 0.3 \\times 10^{21}$~\\cm, for a spin temperature of 200~K and a covering factor of unity. The high estimated HI column density is consistent with the presence of large amounts of dust at the lens redshift; the intervening dust could be responsible for the extremely red colour of the background quasar. ",
        "introduction": "The $z \\sim 2.2$ radio-loud quasar \\pmn~is an exceedingly interesting object. Its radio continuum consists of at least 6 components, with a maximum image separation of $\\sim 700$~mas (\\cite{winn02} (W02); \\cite{winn03} (W03)). Five of these images (A -- E) have the same spectral index and a curved arc of emission is seen between A and B. These properties caused W02 to argue that the object is a gravitationally lensed system. Similarly, Gregg et al.~(2002) used the multiplicity of components in a near-IR image as well as the high inferred source luminosity as evidence for lensing. Interestingly enough, \\pmn~is only the second lens of galactic mass to show an image multiplicity greater than four; the other such lens is the six-component system towards CLASS~B1359+154 (\\cite{rusin01}).  CaII H and K lines have been detected towards \\pmn, at $z = 0.76451 \\pm 0.00016$ (\\cite{hall02}); the lens is thus likely to be at this redshift. The CaII absorption is very strong, with rest frame equivalent width $\\sim 5.3$ \\AA, similar to that seen in the lens B0218+357 (\\cite{browne93}). Further, the quasar is highly reddened in optical wavebands, with $B - K \\ge 11$ (\\cite{gregg02}); this suggests the presence of large amounts of dust, either in the host galaxy of the quasar or in the lensing galaxy. While the differential extinction of the quasar images suggests that at least some of the reddening is due to dust in the absorber, it has so far not been possible to conclusively determine the location of the dust. Assuming the dust is at $z = 0.7645$ implies $E(B-V) = 1.315$, i.e. a visual extinction $A_{\\rm V} = 3.85$ (\\cite{hall02}). In conjunction with the strong CaII lines, this would suggest that the lens is a gas-rich, late-type galaxy.  More recently, Keeton \\& Winn (2003) have come up with the first quantitative model of the lensing system that explains most of its observed properties. The background source is assumed to have two components and the lensing system to consist of two galaxies. The observed 5 components A -- E are then the images of one of the background sources and component F is an image of the second, with the remaining two images of the latter source too faint for detection. The model is supported by new high-resolution radio and optical observations (W03), which found that F has a different spectral index from images A -- E. These observations also showed that the observed large sizes of components C and E are not intrinsic but instead arise due to scatter broadening at the lower frequencies.  Finally, W03 find marginal evidence for a detection of the two lens galaxies in Hubble Space Telescope (HST) images of the system; Keeton \\& Winn (2003) predict these to be spiral galaxies, each with a velocity dispersion of about $\\sim 120$~\\kms, and a projected separation of $\\sim 0.4''$.  We present, in this Letter, Westerbork Synthesis Radio Telescope (WSRT) observations of HI 21cm absorption at $z \\sim 0.7645$ towards \\pmn. The usual association of dust with high column densities of neutral gas confirms that the dust is likely to be concentrated in the intervening lens system. We describe the WSRT observations and data analysis in section~\\ref{sec:obs}; the spectra, results and models are presented and discussed in section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{The 21cm spectrum and derived quantities} \\label{sec:21cm}  \\begin{figure} \\centering \\epsfig{file=fig1.eps,width=3.5in} \\caption{Final WSRT HI 21cm spectrum towards \\pmn. The x and y axes are heliocentric frequency (in MHz) and flux density (in mJy) respectively; the top axis is heliocentric velocity (in \\kms), relative to $z = 0.7645$. The spectrum has been boxcar smoothed by 3 channels and has a velocity resolution of $\\sim 10.8$~\\kms~and an RMS noise of $\\sim 4.1$~mJy. } \\label{fig:lowres} \\end{figure}  The final WSRT HI 21cm spectrum towards \\pmn~is shown in figure~\\ref{fig:lowres}; this has been boxcar smoothed by three channels (and re-sampled) and has a velocity resolution of $\\sim 10.8$~\\kms~and an RMS noise of $\\sim 4.1$~mJy. The redshift of the CaII lines (\\cite{hall02}) is shown below the spectrum, with error bars.  Two deep features can be immediately identified in the spectrum of Fig.~\\ref{fig:lowres}, with their peak optical depths at heliocentric frequencies of $804.898 \\pm 0.03$ and $805.455 \\pm 0.03$~MHz, i.e. heliocentric redshifts of $0.76470 \\pm 0.00006$ and $0.76348 \\pm 0.00006$ respectively. The redshift of the first of these features agrees with that of the CaII lines, within the errors. This is also the deeper of the two HI components, with a peak optical depth of $0.047 \\pm 0.007$; the second feature has a peak optical depth of $0.039 \\pm 0.007$.  The HI absorption is extremely wide, with a total velocity spread (between nulls) of $\\sim 500$~\\kms. This is similar to the width of the CaII lines, which have $\\sigma_{\\rm v} = 220 \\pm 40$~\\kms~(\\cite{hall02}).  The equivalent width of the 21cm profile is $\\int \\tau_{\\rm 21} \\mathrm{d}V = 7.06 \\pm 0.08$~\\kms. This is is higher than that measured in either of the other gravitational lenses in which HI absorption has so far been detected, B0218+357 ($EW = 2.94 \\pm 0.02$~\\kms; Kanekar et al. 2003) and PKS~1830$-$21 ($EW = 5.8 \\pm 0.4$~\\kms; Chengalur et al. 1999). The large equivalent width implies a very HI column density $\\NHI = 2.6 \\pm 0.3 \\times 10^{21} \\lb T_{\\rm s} / 200 {\\rm K} \\rb \\lb 1 / f \\rb$~\\cm~(e.g. Kanekar \\& Chengalur 2003), where $f$ is the covering factor of the absorber and $T_{\\rm s}$, its spin temperature; this is well into the range observed in damped Lyman-$\\alpha$ systems (e.g. Wolfe et al. 1986). We note that a spin temperature of $\\sim 200$~K is typical of spiral galaxies, both locally (\\cite{braun92}) and at low redshift (\\cite{kanekar03}). Of course, the above estimate of $\\NHI$ gives the sum of HI column densities towards {\\it all} the different source components giving rise to the 21cm absorption; it is not possible to measure the column density towards individual source components with the available data.  The extremely high estimated HI column density implies a high dust content in the $z \\sim 0.7645$ lens and explains the observed reddening in the optical images (\\cite{gregg02}; W03). Assuming that the molecular hydrogen column density is roughly equal to $\\NHI$, the total hydrogen column density $N_{\\rm H} = 2 N_{\\rm H_2} + \\NHI \\sim 7.8 \\times 10^{21}$~\\cm~(again assuming $T_{\\rm s} = 200$~K and a covering factor of unity). This implies $E(B-V)\\sim 1.35$, using a Galactic dust to gas ratio; further assuming a Galactic extinction law ($R_{\\rm V} = 3.1$; \\cite{binney98}) results in a visual extinction $A_{\\rm V} \\sim 4.15$. Of course, it is likely that the 21cm absorption profile (and hence the estimated column density) contains contributions from more than one line of sight, with variable optical depth and kinematics across the complex lensing image; it is thus not possible to estimate the visual extinction towards individual source components.   \\subsection{Modelling the absorption} \\label{sec:models}  \\begin{figure*} \\epsfig{file=fig2a.ps,width=2.4in,height=2.4in} \\vskip -2.29in \\hskip 2.38in \\epsfig{file=fig2b.ps,width=2.37in} \\vskip -2.1in \\hskip 4.8in \\epsfig{file=fig2c.ps,width=2.4in,height=2.4in} \\caption{Disk galaxy models for the absorption profile. Central Panel [B]: The six principal image components (A--F) of the 1.7~GHz VLBA image of W03 (grey-scale), with two sets of ellipses drawn to indicate the orientation and inclination of the intervening model disk galaxies. The galaxy centre is marked by a $+$ and a $\\times$ for Models~I and II, respectively; the two filled squares indicate the mass centroids of the two mass lens model of Keeton \\& Winn (2003). Left Panel [A]: 21cm spectrum derived from Model I (heavy line) superimposed on the observed 21cm profile (dashed line). Right Panel [C]: 21cm spectrum derived from Model II (heavy line) superimposed on the observed 21cm profile (dashed line). In both cases, the zero of the velocity scale corresponds to 805 MHz ($z=0.76448$, heliocentric) . See text for discussion. } \\label{fig:models} \\end{figure*}  The lens towards \\pmn~creates a complex radio image of six image components (A -- F), with a maximum angular separation of $\\sim 681$~mas (W02). Components A and B are seen to connect in a section of lensing arc at faint levels in the 1.7~GHz VLBA image (W03). Components C and E appear to be strongly scatter broadened at frequencies below 5~GHz and are almost undetected in VLBA images at 1.7~GHz (W02,W03), as they are over-resolved by even the shortest baselines of the VLBA. Note, however, that their flux must be present in the integral flux density measured in our spectrum at 805~MHz; further, they provide additional lines of sight through the $z \\sim 0.7645$ absorber. Next, while component F has a significantly steeper spectrum than the other five components, its total flux density is only $11 \\pm 1$~mJy at 1.7~GHz; it is thus quite unlikely to have sufficient flux density at 805~MHz to give rise to either of the two main 21cm components. Further, while it is possible that the some of the wide, weak absorption arises against this component, this would require a fairly high optical depth ($\\tau \\ga 0.5$). The 21cm absorption is thus most likely to arise against the brighter components A, B, C and D (and perhaps the arc of radio emission between A and B), each of which have flux densities $\\ga 100$~mJy at 1.7~GHz. We note that component A also shows evidence for scatter broadening at frequencies below 5~GHz while D is the most strongly reddened of the source components; on the other hand, B is the bluest of all components and shows no evidence for scatter broadening in the VLBA images (W03). This suggests that there is relatively less dust and gas along the line of sight to component B; the two main 21cm features thus appear most likely to arise from absorption against components A and D.  We have explored these possibilities with a simple kinematic model using a differentially rotating disk absorber. The modelling does not produce a unique result, but rather succeeds in demonstrating that a single, massive galaxy could indeed be responsible for the 21cm absorption profile. If a more complex, multi-galaxy system were to be invoked, there would be ample free parameters to fit the integral spectrum reported here. Further observational progress will require radio interferometry to separate the absorption along different lines of sight.  The philosophy of modelling the 21cm profile shape is based on the assumption that narrow absorption features occur in the spectrum when a compact, high-surface brightness image of the background radio source ``selects'' a small patch from an extended, smooth velocity field of a rotating galaxy.  Thus, the two main narrow 21cm features of Fig.~\\ref{fig:lowres} are expected to be associated with specific source components of the background source, once VLBI observations are conducted in the 21~cm line. The modelling uses the source structure obtained from the 1.7~GHz VLBA maps of W03 for the compact radio components, with a scaling in flux to both match the relative flux densities of the components in the higher frequency bands, and still produce the integral flux density measured at 805~MHz by our WSRT observations. The more extended components C and E had to be supplemented with diffuse gaussian components to account for the flux missing from the VLBA map.  The two models discussed here (and shown in Fig.~\\ref{fig:models}[B]) are representative of two broad classes which can reproduce the observed 21cm profile. It should be emphasized, however, that neither class is able to account for {\\it all} the observations simultaneously, including the scatter broadening and reddening of different source components. Both models are axially symmetric, planar systems with rotation curves that are flat at $V_{\\rm rot}$ in the outer regions, with a linear rise from the centre to $R_{\\rm c}$. The systemic velocity of the galaxy is offset from 805~MHz ($z=0.76448$) by $V_{\\rm off}$, and the galaxy centroid is specified by the projected positional offsets $\\Delta X$ (west), $\\Delta Y$ (north) relative to image D. The HI optical depth profile for each pixel is broadened by a velocity dispersion of $\\sigma_{\\rm v}=7$~km~s$^{-1}$ before computing that pixel's contribution to the integral absorption spectrum.  \\subsubsection{Model I}  Models belonging to class I are motivated by the need to inflict heavy reddening on image D without reddening B, coupled with accounting for the scatter broadening of A and C. However, this model does not account for the scatter broadening of image E. The model parameters are: position angle of the receding major axis $p=-50^{\\circ}$ ($+p$, measured east from north), inclination to sky plane $i=85^{\\circ}$, $V_{\\rm rot}=230$~km~s$^{-1}$, $R_{\\rm c}=1.4$~kpc, exponential gas scale height $h=0.7$~kpc, $V_{\\rm off}= 35$~km~s$^{-1}$, and galaxy position at $\\Delta X=490$~mas, $\\Delta Y=370$~mas. In this highly inclined model, the HI column density measured perpendicular to the galaxy plane is only $3{\\times}10^{19}(T_{\\rm s}/100)$~cm$^{-2}$ in the region that covers components A and C, but it rises to $3{\\times}10^{20}(T_{\\rm s}/100)$~cm$^{-2}$ for B and D. The high disk inclination increases the line of sight column densities by an order of magnitude. The 21cm spectrum resulting from this model is shown in fig.~\\ref{fig:models}[A]. Here, the two narrow features arise from absorption against  A and D; the extended weak absorption wing out to $\\sim +300$~\\kms~is achieved by tipping the intervening disk to a high inclination and thickening the HI layer with a scale height of $\\sim 700$~pc (in the spirit of Briggs et al.~(1985) and Prochaska and Wolfe~(1997)).   \\subsubsection{Model II}  The second, somewhat orthogonal, class of models (of which Model II is an example) has an infinitesimally thin disk, oriented to cast a shadow on images A, B, C, and E, thereby accounting for the observed scatter broadening of A, C, and E, but implying that all would be vulnerable to reddening in contradiction with B being the bluest and D being the reddest of the images.  The model parameters are: $p=+52^{\\circ}$, $i=72.5^{\\circ}$, $V_{\\rm rot}=250$~km~s$^{-1}$, $R_{\\rm c}=0.8$~kpc, (gas scale height $h=0$~kpc), $V_{\\rm off}=-15$~km~s$^{-1}$, and galaxy position at $\\Delta X=535$~mas, $\\Delta Y=290$~mas. This model has an exponential radial decline in HI column density, with scale length 1.5~kpc, from a central peak of $1.5{\\times}10^{21}(T_{\\rm s}/100)$~cm$^{-2}$ measured perpendicular to the disk plane. The spectrum resulting from the above orientation is shown in fig.~\\ref{fig:models}[C]. Here, the two narrow features arise against images A and B, while the broad shoulder at positive velocities stems from absorption against image C and the radio arc connecting A and B."
    },
    "0311/astro-ph0311398_arXiv.txt": {
        "abstract": "Material accreted onto a neutron star can stably burn in steady state only when the accretion rate is high (typically super-Eddington) or if a large flux from the neutron star crust permeates the outer atmosphere. For such situations we have analyzed the stability of nonradial oscillations, finding one unstable mode for pure helium accretion. This is a shallow surface wave which resides in the helium atmosphere above the heavier ashes of the ocean. It is excited by the increase in the nuclear reaction rate during the oscillations, and it grows on the timescale of a second. For a slowly rotating star, this mode has a frequency $\\omega/(2\\pi)\\approx (20-30\\textrm{ Hz})(l(l+1)/2)^{1/2}$, and we calculate the full spectrum that a rapidly rotating ($\\gg30\\textrm{ Hz}$) neutron star would support. The short period X-ray binary 4U 1820--30 is accreting helium rich material and is the system most likely to show this unstable mode, especially when it is not exhibiting X-ray bursts. Our discovery of an unstable mode in a thermally stable atmosphere shows that nonradial perturbations have a different stability criterion than the spherically symmetric thermal perturbations that generate type I X-ray bursts. ",
        "introduction": "Type I X-ray bursts are caused by unstable nuclear burning on accreting neutron stars (NSs) (see recent reviews by Bildsten 1998; Strohmayer \\& Bildsten 2003) in low mass X-ray binaries (LMXBs). They have rise times of seconds with decay times ranging from tens to hundreds of seconds depending on the thickness and composition of the accumulated material. These bursts repeat every few hours -- the timescale to accumulate an unstable amount of fuel on a NS. Evidence that nuclear burning powers the bursts is the observed ratio of the time averaged accretion luminosity to the time averaged burst luminosity ($\\approx 40$ for hydrogen rich accretion and $\\approx 125$ for pure helium accretion). The thermonuclear instability which triggers these bursts is usually due to the extreme temperature sensitivity of \\ta reactions (Hansen \\& Van Horn 1975; Woosley \\& Taam 1976; Maraschi \\& Cavaliere 1977; Joss 1977; Lamb \\& Lamb 1978).  Ever since the original work of Hansen \\& Van Horn (1975), the approach to assessing the stability of burning on a NS has been to apply a spherically symmetric thermal perturbation to a spherical model (Schwarzschild \\& H\\\"{a}rm 1965). This is either done numerically through a time-dependent code (e.g. Woosley et al. 2003), analytically with a one-zone model (e.g. Fujimoto, Hanawa \\& Miyaji 1981), or using an eigenvalue approach (e.g. Narayan \\& Heyl 2003). These results roughly agree and find a thermal runaway once a critical amount of fuel has accumulated. These calculations also find the accretion rate above which burning is stable.  We are now taking an alternative approach to this problem. Using a NS envelope that is stable to spherically symmetric thermal perturbations (i.e. not bursting), we ask whether the envelope is unstable to a non-axisymmetric perturbation. Since the immediate response to such a perturbation is a wave, the proper way to proceed is through a study of the nonradial oscillation modes. The criterion for thermal versus mode stability is mathematically different (as we will explain in \\S2.2 and \\S3.3), but it is not yet clear how this contrast will manifest itself observationally. Nonradial oscillations in burning NS envelopes have been studied before (for example McDermott \\& Taam 1987; Strohmayer \\& Lee 1996), but these are for atmospheres that are thermally unstable, leaving our question unanswered.  To simplify this initial work, we assume that the NS is accreting pure helium and burning this fuel at the same rate it accretes. This steady state burning proceeds if the NS envelope is thermally stable (and no type I X-ray bursts are occuring). It is stable when either the temperature sensitivity of \\ta reactions has weakened due to a high envelope temperature ($\\gtrsim5\\times10^8$ K), or the thermal balance of the burning layer is fixed by a high flux from deeper parts of the NS (Paczy\\'{n}ski 1983a). We then find the eigenfrequencies and test the stability of g-modes in such an envelope. The one unstable mode we find is a shallow surface wave in the helium layer that is supported by the buoyancy of the helium/carbon interface. It has a non-rotating frequency of $\\omega/(2\\pi)\\approx (20-30\\textrm{ Hz})(l(l+1)/2)^{1/2}$, depending on the accretion rate. This mode is driven unstable by the $\\epsilon$-mechanism (i.e. pumped by the perturbation in nuclear burning) and grows in amplitude on the timescale of a second. This is very different from other observed stellar pulsations, such as Cepheid variables, which are driven by the $\\kappa$-mechanism (pumped by changes in opacity). Unlike the damped higher radial order g-modes, the unstable mode suffers little radiative damping in the outer atmosphere. Rapid rotation like that seen from many accreting NSs (200-700 Hz) will yield multiple oscillation frequencies with a well defined pattern. If such a frequency pattern were seen from a NS it would give strong evidence that a mode has been excited.  In \\S 2 we describe the physics of a NS star envelope accreting and burning helium in steady state. The conditions required for steady state accretion are reviewed. Nonradial adiabatic oscillations in a plane-parallel envelope are discussed in \\S 3, and though much of this has been described in previous papers, we include additional details and strategies that are relevant for NSs. The oscillatory equations are derived in \\S 3.1, and in \\S 3.2 we calculate the radial eigenfunctions and find the g-mode frequencies for a non-rotating NS. The stability of these modes is investigated in \\S 3.3, including a look at excitation and damping times. In \\S 3.4 we find the unstable mode's frequency on a rapidly rotating neutron star. The LMXB 4U 1820--30 is discussed in \\S 4 along with other candidates that might show the excited mode we have found. We conclude in \\S 5 by summarizing our findings and discussing future directions for this work. ",
        "conclusions": "This study makes clear the distinction between thermal and mode perturbations of the thin envelope of a NS. We find one excited mode on a thermally stable helium accreting NS which is a shallow surface wave riding in the helium atmosphere above the ash ocean. It has a frequency of $(20-30\\textrm{ Hz})(l(l+1)/2)^{1/2}$ (depending on the accretion rate) and has an excitation time of about a second. This mode is modified by rotation to give multiple possible observable frequencies. Seeing such a frequency pattern from a NS could be used to constrain properties of the NS star (e.g. its spin or radius). The rotationally modified eigenfunctions will have different levels of visibility depending on both their final amplitudes and their angular eigenvalues. The LMXB 4U 1820--30 is accreting helium rich material and is the best candidate for having an envelope like the ones we consider. No such oscillations have yet been reported from it, but perhaps now that predictions have been made, a search for modes would be easier, especially when bursting is absent. Making comparisons with just this one system is limited, and it would be useful to do a similar mode stability analysis on NSs that are accreting a solar mix of hydrogen and helium, especially during times following a superburst when type I bursts are not occuring.   A further study could also address the eventual nonlinear evolution of the excited mode. Such a study might predict observational effects that this mode may exhibit. If this mode were to saturate at a low amplitude (as is common in most other stars) then the most observable feature would be the periodic modulation of the surface luminosity. If the mode saturation is the result of mode coupling, then other frequencies might also be visible. If the amplitude becomes exceedingly nonlinear, the resulting large temperature excursions may well lead to a more rapid consumption of the fuel than the rate of supply, possibly triggering limit cycle behavior.  This study highlights the importance a shallow surface wave may play on NSs, and we plan to continue an investigation to understand better what other NS environments might show this mode. An interesting study would be the time before or after type I X-ray bursts when the NS envelope should have the mean molecular weight discontinuity necessary to exhibit this mode. In the latter case the frequency of the mode would be changing due to the cooling of the atmosphere. This implies that shifting frequency searches may be the correct way to find these modes. A future paper will consider such effects in more detail, including the size and sign of the expected shifts.  Above all, this study has shown that spherically symmetric studies are not complete and can easily miss important physics that may be relevant for correctly understanding the properties of NSs. We hope to extend our work to type I X-ray burst ignition, which has not been fully investigated in the non-axisymmetric case. It is still somewhat unsatisfactory to think that a NS ignites spherically all at once, and our study makes clear that nonradial perturbations may give different stability criterion for ignition than what has commonly been used.  We thank Andrew Cumming for some of the original inspiration for beginning this research and Phil Arras for discussions concerning g-modes and rotational modifications. We have benefited greatly from conversations with Philip Chang, Chris Deloye, Aristotle Socrates, Anatoly Spitkovsky, Ronald Taam, Dean Townsley, and Greg Ushomirsky. We thank the anonymous referee for suggestions which have added clarity to this manuscript. This work was supported by the National Science Foundation under grants PHY99-07949 and AST02-05956 and by NASA through grant NAG 5-8658.  L. B. is a Cottrell Scholar of the Research Corporation."
    },
    "0311/astro-ph0311484_arXiv.txt": {
        "abstract": "Gamma-ray bursts can appear to be a hundred times as luminous as supernovae, but their underlying energy source(s) have remained a mystery.  However, there has been evidence for some time now of an association of gamma-ray bursts with supernovae of Type Ib and Ic, a fact which has been exploited by a number of models, to explain the gamma-ray burst phenomenon. Here we interpret the results of basic observations of SN 1987A and of pulsars in globular clusters, to propose the energy source, which powers at least some long-duration gamma-ray bursts, as core-collapse following the merger of two white dwarfs, either as stars or stellar cores.  The beaming and intrinsic differences among gamma-ray bursts arise, at least in part, from differing amounts and composition of the gas in the merged stellar common envelopes, with the more energetic bursts resulting from mergers within less massive envelopes.  In order for the beams/jets associated with gamma-ray bursts to form in mergers within massive common envelopes (as with SN 1987A), much of the intervening stellar material in the polar directions must be cleared out by the time of core-collapse, {\\it or} the beams/jets themselves must clear their own path.  The core-collapse produces supernovae of Type Ib, Ic, or II (as with SN 1987A, a SNa IIp), leaving a weakly magnetized neutron star remnant with a spin period near 2 milliseconds.  There is no compelling reason to invoke any {\\it other} model to account for gamma-ray bursts.  Far from being an unusual event, SN 1987A is typical, having the same merger source of initiation as 95\\% of all supernovae, the rare exceptions being Ia's induced via gradual accretion from a binary companion, and Fe catastrophe II's. ",
        "introduction": "Gamma-ray bursts (GRBs) were discovered over three decades ago (Klebesadel, Strong, \\& Olson 1973), with two distinct temporal classes emerging: short and long, with distributions centered near 0.3 and 20 s, being nearly separated at $t = 2 s$, and having larger and smaller hardness ratios respectively \\citep{pac99}, though with nearly 50\\% overlap \\citep{kou93}.  That longer duration GRBs at cosmological distances should be associated with fading afterglows was predicted a decade ago \\citep{paczyn93,kat93}, and confirmed more recently in the optical and X-ray bands, as for GRB 970508 \\citep{cos97,vpa97}, and later in the radio, as for GRB 970703 \\citep{fra97}. The identification and redshift determination of the host galaxy for GRB 970508 definitively confirmed its extragalactic nature \\citep{met97}.  The next year Galama et al.~(1998) found the Ic SN 1998bw within the error circle of GRB 980425 \\citep{pia99}.  Just afterward Bloom et al.~(1999) measured a possible SNa Ic component in the fading afterglow of GRB 980326, which led to other such possible associations, as for GRB 970228. That GRB 980425 was four orders of magnitude less luminous than typical cosmological bursts and the low redshift of SN 1998bw (0.0085), led to widespread skepticism of this association \\citep{whe00}.  Recently, however, the case for a SNa-GRB association has been bolstered by the association of GRB 030329 with SN 2003bh \\citep{hjo03}.  Supernovae are generally classified according to their spectra, Type II for those with hydrogen lines and Type I for those without \\citep{fil97}.  These are each further divided into three subclasses, with Ia's showing strong Si lines in their early spectra, indicating a thermonuclear runaway of a white dwarf star, while Ib's show helium, and Ic's neither hydrogen nor helium, but intermediate mass elements instead.  Ib's and Ic's are thought to share core-collapse as a common cause with all Type II's, except that their progenitors have been stripped of hydrogen (Ib), or both hydrogen and helium (Ic). Ia's also differ from Ib's, Ic's, and II's, in their luminosities, in some cases producing, via C-O burning, nearly a solar mass of $^{56}$Ni -- ten times the amount typical of all other types of SNe.  A number of models have been proposed to explain GRBs.  Some earlier ones, which may not survive a GRB-SNa association, invoke neutron star (NS) collisions with other NSs or black holes (BH) (Lattimer \\& Schramm 1974; Eichler et al.~1989; Narayan, Paczy\\'nski, \\& Piran 1992).  Other, more recent models produce SNa-like explosions by invoking collapses of massive stars, into magnetars \\citep{whe2000}, and black holes: with strong magnetic fields via ``hypernovae'' and ``supranovae'' \\citep{iwa98,vie99}, and without, via ``collapsars'' \\citep{mac99}.  This paper shows that a conceptually very simple model involving the merger of two white dwarfs (WDs), as cores of massive stars or not (``double degenerate''), satisfies the observational constraints of long duration GRBs and most SNe well, while others fall short of this goal. ",
        "conclusions": "We have shown that a very simple model for binary WD-WD merger explains the details of the axisymmetric remnant expanding from, and the ringed gas structure surrounding, SN 1987A \\citep{che94}, as well as the apparent 2.14 ms pulsar signal for the compact remnant (M2000). By extending this picture to other Type II, as well as Type Ib and Ic supernovae, which comprise at least 90\\% of all SNe, we have accounted for the anisotropies recently discovered in many other SNe, as well as all of the fundamental details of long-duration GRBs. At this point, other models to explain this class of GRBs appear to be superfluous.  The nature of the energy source behind GRBs can {\\it definitively} be confirmed only, perhaps, for pulsars, from magnetars to MSPSRs, by detecting an underlying periodicity. Thus, taking ``imposter'' SNe into account, high time-resolution observations of {\\it all} nearby ($< 50$ Mpc) SNe should be made, using the largest and most sensitive optical and radio telescopes, to search for pulsar remnants. Unlike SN 1987A, opacity will not be an issue for many of these.  Observations of only SNe which happen to be associated with GRBs might not be fruitful, as these are usually much farther away, and the proximity of the line of sight to the rotation axis of the binary merger could reduce the amplitude of the pulsar signal.  Although a number of extragalactic SNe (but no Ib's or Ic's) have already been searched \\citep{mid84}, and this has continued sporadically, without success, only scant resources have been devoted to this effort so far. Continuing indefinitely {\\it without} making further, more sensitive searches, is to risk wasting time and resources.  The results from SN 1987A, now less controversial due to strong evidence for precession in two other ``normal'' pulsars, B1828-11 and B1642-03 (Stairs, Lyne, \\& Shemar 2000; Shabanova, Lyne, \\& Urama 2001), imply a pulsed optical signal of 10 or more solar luminosities, at 5.0 -- 6.5 years of age (M2000). This is hardly surprising, as the Crab pulsar's (nearly entirely pulsed) optical output is 4 solar, and nanosecond radio bursts in pulsars are bright enough to be seen in neighboring galaxies \\citep{han03}.  This new understanding of GRBs and SNe may represent the end of an era, begun four decades ago, of exploration and discovery in high-energy astrophysics.  The core-collapse events, and early stages of the 2 ms pulsar remnants from SNe in the Local Group, may be detectable by gravitational wave detectors such as LIGO (Santostasi, Johnson, \\& Frank 2003), provided an early, rudimentary spin frequency history can be obtained by detecting the pulsars in other bands, such as the optical and/or radio. Thus this new understanding also {\\it begins} an era, which will span the first part of our own 21$^{st}$ Century, of extragalactic pulsar astronomy."
    },
    "0311/astro-ph0311167_arXiv.txt": {
        "abstract": "We study the evolution of perturbations during the domination and decay of a massive particle species whose mass and decay rate are allowed to depend on the expectation value of a light scalar field. We specialize in the case where the light field is slow-rolling, showing that during a phase of inhomogeneous mass-domination and decay the isocurvature perturbation of the light field is converted into a curvature perturbation with an efficiency which is nine times larger than when the mass is fixed. We derive a condition on the annihilation cross section and on the decay rate for the domination of the massive particles and we show that standard model particles cannot dominate the universe before nucleosynthesis. We also compare this mechanism with the curvaton model. Finally, observational signatures are discussed. A cold dark matter isocurvature mode can be generated if the dark matter is produced out of equilibrium by both the inflaton and the massive particle species decay. Non-Gaussianities are present: they are chi-square deviations. However, they might be too small to be observable. ",
        "introduction": "Strongly supported by the recent observations of the Wilkinson Microwave Anisotropy Probe (WMAP) satellite \\cite{WMAP1,WMAP2,Peiris,WMAP4,WMAP5}, inflation \\cite{guth81} has now become the dominant contender for generating adiabatic density perturbations with an almost flat spectrum. In the standard picture, the observed large scale density perturbations are due to the fluctuations of the inflaton field, which are created during a period of accelerated expansion by amplification of quantum vacuum fluctuations \\cite{MFB}.  Recently, alternative mechanisms for generating density perturbations after inflation have been proposed. They all assume that the early universe is filled with at least one light scalar field $\\phi$ whose energy density is negligible during inflation. Fluctuations in the field are amplified during the inflationary phase with a quasi scale-invariant spectrum and their amplitude fixed by the energy scale of inflation, $\\delta \\phi \\sim \\frac{H_*}{2\\pi}$, where $H_*$ is the Hubble parameter at the horizon crossing. If the vacuum expectation value (VEV) of the light field $\\phi$ during inflation is small, its density perturbation $\\delta \\rho_\\phi /\\rho_\\phi \\sim \\delta \\phi/\\phi$ is larger than the perturbations generated during inflation. However, since the light field is sub-dominant, its perturbation is initially of the isocurvature type. Later, it can be converted into a curvature perturbation. For a successful conversion one of these two ingredients are necessary:  (1) the field $\\phi$ must come (close) to dominate the universe, or (2) it must induce fluctuations in a second component which eventually comes to dominate the universe.   In case (1) the scalar field $\\phi$ is called curvaton. In the curvaton scenario \\cite{curvaton1,curvaton2,curvaton3}, the isocurvature perturbation stored in the curvaton field is transformed into a curvature perturbation during a phase in which $\\phi$ oscillates at the bottom of its potential behaving as a dust-like component and thus dominating over the radiation.   Case (2) has been inspired by the idea that coupling constants and masses of particles during  the early universe may depend on the VEV of some light scalar field. This idea is motivated by supersymmetric theories and theories inspired by superstrings where coupling `constants' and masses of particles are usually functions of the scalar fields of the theory. An interesting proposal is that large scale perturbations could be generated from the fluctuations of the inflaton coupling to ordinary matter \\cite{gamma1,gamma2}. These can be converted into curvature perturbations during the reheating phase \\cite{Turner} when the inflaton decays into radiation and reheats different patches of the universe with different temperatures and energy densities (see however \\cite{Enqvist}). This mechanism of conversion has been called inhomogeneous reheating in \\cite{gamma1}. An extension and generalization of this mechanism is that the masses of particles produced during the reheating are allowed to vary \\cite{gamma3}. If sufficiently long-lived, the massive particles can eventually dominate the universe. Due to the fluctuations in the masses the mass-domination process becomes inhomogeneous and can convert the density perturbation of $\\phi$ into a curvature perturbation.  Some authors \\cite{Riotto,Mazum,Tsuji} (see also \\cite{Zalda}) have studied the evolution of large scale perturbations during the decay of the inflaton in the case of a fluctuating decay rate. Two of these groups \\cite{Riotto,Mazum} make extensive use of the formalism introduced in Ref.~\\cite{wmu} (see however also \\cite{KS,hama}) for interacting fluids. Indeed, on several occasions the light scalar field can be treated as a fluid and the formalism of Ref.~\\cite{wmu} consistently applied. In this paper we extend the analysis of \\cite{Riotto,Mazum} and we include the case where the inhomogeneous reheating is due to the decay of massive particles whose mass and decay rate depend on the VEV of a light field. In the limit of a constant mass the results of \\cite{Riotto,Mazum} are recovered.  We also derive the condition for the domination of the massive particles. Since at early times they are in thermal equilibrium, their abundance depends on their annihilation cross section. As we shall see the condition for the domination does not depend on their mass but only on their annihilation cross section and decay rate. Our results are then compared with those obtained for the curvaton model and some of the observational consequences are discussed. If the inflaton and/or the massive particles decay out of equilibrium isocurvature perturbations can be generated. The inhomogeneous inflaton decay and mass-domination can also lead to non-Gaussianities in the adiabatic spectrum of perturbations. These can lead to observational signatures in the cosmic microwave background (CMB) that can be possibly observed in future experiments.   The paper is organized as follows. In the next section we model the inhomogeneous mass-domination mechanism by writing down the conservation laws for the radiation, the massive particles and the light field, and we derive the coupling between these species. In Section III we write the background and perturbation equations obtained from the conservation laws. These equations are completely general. They can be used for the inhomogeneous inflaton decay as for the inhomogeneous mass-domination model. In Section IV we concentrate on the limit that the light field is in slow-roll and we discuss the efficiency of the inhomogeneous inflaton decay and mass-domination models. We also compare them to the curvaton model. Finally, in Section V we discuss the observable signatures of the inhomogeneous mass-domination i.e., the possible presence of isocurvature perturbations and non-Gaussianities. In the final section we draw our conclusions. ",
        "conclusions": "In this paper we have studied the evolution of perturbations during a phase dominated by massive particles whose mass and decay rate can fluctuate in space and time. If the fluctuations are set by the VEV of a light scalar field overdamped during inflation, the isocurvature perturbation in the scalar field can be converted into a curvature perturbation, and this can be the main mechanism of generation of large scale perturbations for structure formation. We have derived a set of perturbation equations that can be used in full generality for any kind of dependence of the mass and decay rate on the light field. Making use of these perturbation equations we have recovered the results of \\cite{gamma1,Riotto,Mazum} for the inhomogeneous reheating with varying decay rate. We have also discussed the condition for the massive particles to dominate the universe before decaying. This condition does not depend on their mass, but depends on the annihilation cross section and decay rate. Standard model massive particles cannot dominate the universe. Furthermore, we have shown that when the mass of the massive particles is allowed to vary, the mechanism of conversion can be nine times more efficient. The final total curvature perturbation is $\\zeta = - (1/6) \\delta \\Gamma/\\Gamma - (4/3) \\delta M/M $. This is our main result. Finally, we have compared this with the curvaton model discussing differences and similarities.   On the observational side we have discussed two possible signatures of the mass-domination mechanism: correlated adiabatic and isocurvature perturbations and non-Gaussianities. If present, a cold dark matter isocurvature perturbation provides some important information on the mechanism of generation of the dark matter and on the vacuum expectation values of the inflaton and light field during inflation. There are non-Gaussianities generated by this mechanism, which are $\\chi^2$. In order to precisely compute them one has to study the evolution of second order perturbations. In the limit where $f_{\\rm NL}$ is large, by simply using Eq.~(\\ref{eq:ngalpha}) the non-linear parameter is $f_{\\rm NL} =- 5/(6 r \\alpha_\\phi)$. When both the mass and the decay rate of the massive particles fluctuate, due to the high efficiency the non-Gaussianities can be much smaller than what is possibly observable, $f_{\\rm NL} \\ll -5$. \\\\   {\\bf Acknowledgment}  It is a pleasure to thank Roberto Trotta for very fruitful discussions and suggestions. I also acknowledge Ruth Durrer, David Langlois, Karim Malik, and Jean-Philippe Uzan for very helpful suggestions and comments, Robert Brandenberger for carefully reading the manuscript, and Nicola Bartolo, Antonio Riotto and Sabino Matarrese for drawing my attention on the importance of the evolution of second order perturbations. Part of this work was done at the Kavli Institute of Theoretic Physics. This research was supported in part by the National Science Foundation under Grant No. PHY99-07949 and by the Swiss National Science Foundation."
    },
    "0311/astro-ph0311092_arXiv.txt": {
        "abstract": "{ Fifteen Be/X-ray binaries and candidates in the SMC were observed serendipitously with the EPIC instruments of XMM-Newton during two observations of SNR\\,0047-73.5 and SNR\\,0103-72.6 in October 2000. A total of twelve of those sources are detected. For eleven of them an accurate position and in part detection of X-ray pulsations support the proposed identification as Be/X-ray binaries. In one case the improved X-ray position excludes the previously suggested identification with an \\Halp\\ emission line star found within the ROSAT error circle. The detection of pulsations (172.2 s, 320.1 s and 751 s) from three hard X-ray sources with periods known from ASCA observations confirm their proposed identifications with ROSAT sources and their optical Be star counterparts. In addition, pulsations with a period of 263.6 s were found from XMMU\\,J004723.7$-$731226=RX\\,J0047.3$-$7312. For \\sax\\ a pulse period of 341.2$\\pm$0.5 s was determined, continuing the large spin-up seen with ASCA, BeppoSAX and Chandra between 1996 and 1999 with a period derivative of $-$1.6 s yr$^{-1}$ covering now 4.5 years. The 0.3-10.0 keV EPIC spectra of all eleven Be/X-ray binaries and candidates are consistent with power-law energy distributions with derived photon indices strongly peaked at 1.00 with a standard deviation of 0.16. No pulsations are detected from RX\\,J0049.2$-$7311 and RX\\,J0049.5$-$7310 (both near the 9 s pulsar AX\\,J0049$-$732) and  RX\\,J0105.1$-$7211 (near AX\\,J0105$-$722, which may pulsate with 3.3 s), leaving the identification of the ASCA sources with ROSAT and corresponding XMM-Newton objects still unclear. We present an updated list of high mass X-ray binaries (HMXBs) and candidates in the SMC incorporating improved X-ray positions obtained from Chandra and XMM-Newton observations. Including the results from this work and recent publications the SMC HMXB catalogue comprises 65 objects with at least 37 showing X-ray pulsations. ",
        "introduction": "X-ray observations of the Small Magellanic Cloud (SMC) with recent observatories like ASCA, BeppoSAX, ROSAT and RXTE revealed a large number of Be/X-ray binary systems. Most of them were discovered through the detection of X-ray pulsations \\citep[e.g. see summary of ASCA X-ray sources in the SMC][]{2003PASJ...55..161Y}, indicating the spin period of the neutron star which orbits an early type star. Optical follow-up observations are then required to determine the spectral type of the high mass star and in the SMC only Be stars were identified (with the only secure exception of the supergiant system SMC\\,X-1), making the SMC very different to the Milky Way and the Large Magellanic Cloud in absolute numbers of high mass X-ray binaries with respect to the galaxy mass and also in their ratio of Be to supergiant OB systems.  \\citet[][ hereafter HS00]{2000A&A...359..573H} searched for Be/X-ray binaries in the ROSAT catalogues of \\citet{2000A&AS..142...41H} and \\citet{2000A&AS..147...75S} by correlating them with catalogues of optical \\Halp\\ emission line objects \\citep{1993A&AS..102..451M,2000MNRAS.311..741M}. HS00 presented a list of 47 Be/X-ray binaries and candidates with X-ray positions typically better than 10\\arcsec. Since the launch of Chandra and XMM-Newton new X-ray pulsars and Be/X-ray binary candidates were discovered \\citep{2001A&A...369L..29S,2003A&A...403..901S,2003ApJ...584L..79M}. Together with eleven pulsars detected by RXTE with very uncertain X-ray positions we know about sixty high mass X-ray binaries (HMXBs) and candidates in the SMC, from which more than thirty are known as X-ray pulsars \\citep{2003PASJ...55..161Y,2001foap.conf..239M}.  The position uncertainties of X-ray sources in the SMC obtained by previous missions of up to tens of arc-seconds (including the most precise ROSAT positions) made it in many cases impossible to uniquely correlate sources from the different X-ray missions, nor allowed exact optical identifications. We therefore started to systematically analyse archival XMM-Newton data of HMXBs and candidates in the SMC in order to confirm their proposed nature. First results were presented in \\citet{2003A&A...403..901S}. Here we continue this work analyzing the XMM-Newton EPIC data of two fields in the SMC which contain in total fifteen known HMXBs and candidates. ",
        "conclusions": "We performed temporal and spectral analyses of eleven Be/X-ray binaries and candidates observed serendipitously by the EPIC instruments on board XMM-Newton in two observations of the supernova remnants \\snra\\ and \\snrb. From five of the objects pulsations in their X-ray flux are detected, four of them were previously known pulsars and in one case, XMMU\\,J004723.7$-$731227, the XMM-Newton data revealed 263 s pulsations for the first time. XMMU\\,J004723.7$-$731227 coincides in position with RX\\,J0047.3$-$7312 which was suggested to be a Be/X-ray binary due to its likely identification with an \\Halp\\ emission line object (HS00). The detection of the pulse period strongly supports this identification.  The detection of pulsations, together with the good X-ray positions allowed to securely identify sources detected by various X-ray space missions, like e.g. the 172 s pulsar AX\\,J0051.6$-$7311 with XMMU\\,J005152.2$-$731033 (via pulsations) and RX\\,J0051.9$-$7311 (via accurate position). The latter was previously also suggested to be the counterpart of XTE\\,J0050$-$732\\#1, a 16.6 s pulsar. The 345 s pulsar \\sax\\ was found to continue its large spin down with a rate of $-$1.6 s yr$^{-1}$ extending now over 4.5 years. The 323 s pulsar AX\\,J0051$-$733 was found to exhibit very similar properties as \\sax\\ with respect to spin period, period derivative and X-ray luminosity.  The large number of Be/X-ray binaries in the SMC allows first statistical studies of their spectral and temporal properties. XMM-Newton and Chandra are ideally suited to detect pulsars among the Be-/X-ray binary candidates at low source fluxes and future observations are required to further increase the sample and allow to analyse it in a homogeneous way."
    },
    "0311/astro-ph0311571_arXiv.txt": {
        "abstract": "Green Bank Telescope (GBT) $\\lambda$21cm observations have revealed a faint, yet extensive \\ion{H}{1} cloud population surrounding the Andromeda Galaxy (M31).  The newfound objects are likely analogs to the high-velocity \\ion{H}{1} clouds (HVCs) seen around the Milky Way. At least 20 discrete features are detected within 50~kpc of the M31 disk, with radial velocities that are comparable to those of outer disk rotation. In addition, a filamentary ``halo'' component of at least 30~kpc extent is concentrated at the M31 systemic velocity. Some of the discrete features are organized into elongated systems with velocity continuity, suggestive of tidal streams. The discrete population can be characterized by a steep power-law distribution of number versus \\ion{H}{1} mass in the range between 10$^5$ and 10$^7$~M$_\\odot$. The velocity line-width of discrete clouds is correlated with the cloud \\ion{H}{1} mass: such that if the clouds are gravitationally bound this implies a dark- to \\ion{H}{1} mass ratio of $\\sim$ 100:1. Possible origins for the discrete and ``halo'' M31 features include: a Local Group ``cooling flow'', tidal debris from recent mergers or interactions, and the gaseous counterparts of low mass dark-matter halos. ",
        "introduction": "Much of the sky is covered with neutral hydrogen clouds whose velocities, though within several hundreds of km s$^{-1}$ of zero, are nonetheless anomalous and cannot be explained by normal rotation of the Galaxy \\citep{mulleretal63, wvw97}.  It is attractive to consider that these high-velocity HI clouds (HVCs) might be material left over from the formation of the Milky Way, now in the process of being accreted \\citep{oort70}.  If this is the case, other galaxies may be expected to have similar gaseous components.  A key advantage of studying such extragalactic clouds is that their distance can be comparatively well-constrained. In practice, it has proven difficult to achieve the combination of high spatial resolution, mass sensitivity, and large field-of-view necessary to detect HVC analogs in external systems.  Many nearby galaxies have a few peculiar \\ion{H}{1} features with \\ion{H}{1} mass as low as $\\sim$ 10$^7$~M$_\\odot$ (eg. NGC~628, \\citet{kamphuisbriggs92}), but only recently have instruments had sufficient sensitivity to begin revealing complete systems of peculiar velocity \\ion{H}{1} (eg.  the slowly rotating ``halo'' component of NGC~2403, \\citet{fraternalietal02}).  In this Letter, we report on deep, wide-field observations conducted with the Green Bank Telescope (GBT) which have resulted in the detection of a complex system of \\ion{H}{1} gas extending to at least 50~kpc radius around M31. ",
        "conclusions": "At a distance of 770 kpc, most clouds in our sample have \\ion{H}{1} masses in the range $0.15-1.3\\times10^6$~M$_{\\odot}$.  Such clouds would not have been detected as discrete objects in previous surveys of more distant galaxies, which generally have probed only the mass range above $\\sim10^7$~M$_{\\odot}$ \\citep{braunburton01}. An indication for the more general existence of low mass circum-galactic populations comes from the recent detection of a 15\\% excess of \\ion{H}{1} mass found on scales of several hundred kpc relative to that seen at tens of kpc in the environment of apparently isolated galaxies \\citep{braunetal03}.  Placing accurate limits on the total HI mass of the M31 halo cloud population is difficult because we lack a well-measured radial scale length for the population, must cope with potential incompleteness due to Galactic confusion, and have no confirmed minimum cloud mass.  To the degree that the observed cloud mass function is consistent with a sensitivity-bounded power law given by $dN(\\rm{M}_{HI})~\\alpha~\\rm{M}_{HI}^{-2}~d\\rm{M}_{HI}$, our estimate of HI mass for the halo cloud population within the GBT field is $\\sim 3-4\\times10^7$~M$_{\\odot}$.  The upper limit of this range is calculated assuming the true population is bounded by a minimum HI cloud mass of $10^5$~M$_{\\odot}$ and that there are fewer than $\\sim 100$ of such objects, the vast majority of which are not yet detected. We also included an extra 25\\% to account for clouds confused with Galactic emission, but even so, the \\ion{H}{1} mass of Andromeda's discrete halo cloud population in our data amounts to only $1\\%$ of the mass of M31's HI disk.  Below, we explore the idea that the halo clouds may trace a more substantial amount of ionized gas and possibly also dark matter.  An obvious source of high velocity \\ion{H}{1} is tidal stripping from recent or ongoing mergers.  We detect a gaseous feature which is partially coextensive with the stellar Andromeda Stream \\citep{mcconnachieetal03}, and of comparable spatial extent, as well as a component of possibly stripped gas adjacent to NGC~205.  However, most of the discrete clouds detected in our GBT survey are rather isolated, lacking any apparent relation to known M31 companion galaxies, and do not have obvious indications if internal tidal distortions.  Nevertheless, the faintest dwarf galaxies and stellar streams are notoriously hard to detect (Armandroff et al. 1999). One might argue that some dwarfs have yet to be discovered, while others have been entirely consumed by M31.  Indeed, the comparatively young (6-8 Gyr) halo stars found by Brown et al. (2003) in the ACS Andromeda deep field do suggest that M31 has seen a major merger (with a massive star forming galaxy) or several minor mergers.  Cooling of a tenuous inter-galactic medium \\citep{oort70} is a second viable source of clouds.  In this view, halo clouds would condense from, and remain confined by, coronal gas which is located around M31 or perhaps pervades the Local Group.  Sembach et al. (2002) presented evidence for an extended and highly ionized Galactic corona or Local Group medium traced by high-velocity \\ion{O}{6} absorption.  Likewise, Lockman et al. (2002) conducted a sensitive HI survey of 860 sightlines at $\\delta > -43\\arcdeg$, which suggested that a low-column density ``mist'' of high-velocity neutral gas surrounds the Milky Way ($\\rm{N}_{HI} \\ga 8\\times10^{17}$~cm$^{-2}$ along $40\\%$ of sightlines).  Together, these studies show that the Galactic HVC phenomenon extends to much lower column densities than traditionally appreciated and that the classical Galactic HVCs are only part of a more ubiquitous multi-phase medium. The filamentary ``\\ion{H}{1} halo'' we detect concentrated on the M31 systemic velocity may be a manifestation of an M31 ``cooling flow''. In such a scenario, the detected \\ion{H}{1} may represent only the tip of the iceberg, in terms of baryonic mass, since the gas may well be only of order 1\\% neutral at the relevant low volume densities \\citep{sembachetal99}. The primary source of such coronal gas may actually be the action of a large-scale ``galactic fountain'' \\citep{shapiro, bregman80, savageetal03} in the recent or distant past.  A third component of high velocity \\ion{H}{1} near M31 might be the gas associated with a putative population of low mass dark-matter halos. Current simulations of the Local Group in a $\\Lambda$CDM cosmology predict a large population of low mass DM halos at the present epoch \\citep{klypinetal99,mooreetal99} which dramatically outnumber known dwarf galaxies. If the M31 halo clouds are tracers of substantial dark-matter concentrations, this should be reflected in their internal line-widths. We plot the observed \\ion{H}{1} mass versus FWHM line-width for discrete clouds in Figure 3.  Line-width measurements were determined using Gaussian fits to cloud spectra. The line-width distribution has an approximate lower bound roughly consistent with thermal broadening for gas at 10$^4$ K (FWHM = 24 \\kms). A systematic increase in FWHM line-width is observed with increasing \\ion{H}{1} mass, reaching some 70 \\kms\\ for the clouds near 10$^6$~M$_\\odot$.  Although Davies' Cloud is significantly offset from the remainder of the distribution in Fig. 3, high-resolution imaging has revealed much higher internal line-widths in that object as well \\citep{deheijetal02a}. To demonstrate the expected distribution of line-width with mass, in Fig.~3 we show a curve corresponding to $V^{2} = 100 G M_{HI} R^{-1}$, where the characteristic discrete cloud radius, R, has been held constant at 500 pc, based on our interferometric WSRT detections of some of the cloud cores (Braun et al. 2003, in prep.).  This curve, corresponding to a dark to \\ion{H}{1} mass ratio of about 100:1, is not intended to fit to the data, but merely to provide a basis for comparison. The hypothesis of a kinematically dominant dark-matter component appears to be consistent with the observed line-widths of the discrete M31 halo clouds.  Another aspect of the dark-matter mini-halo scenario that can be checked is the expected number of such objects in the appropriate mass range in the vicinity of M31. Sternberg et al. (2002) predict $\\sim 25$ dark matter mini-halos associated with gravitationally confined HI within a radius of 40 kpc around M31, based on the simulations of Klypin et al. (1999) and Moore et al. (1999). This is fully consistent with the 20 discrete M31 halo clouds we have detected.  The Sternberg et al. calculations also suggest that circum-galactic objects of such low mass and peak \\ion{H}{1} column density should be only $\\sim10$ \\% neutral, implying $M_{HI+HII} \\sim 10 M_{HI}$ and $M_{DM} \\sim 10 M_{HI+HII}$. Associated ionized gas could perhaps be detected via deep Fabry-Perot imaging \\citep{blandhawthornmaloney99,tufteetal02} or absorption line spectroscopy of background quasars \\citep{murphyetal00,trippetal03,sembachetal03}.  In summary, our GBT observations of the Andromeda Galaxy have revealed the first extensive extragalactic counterpart to the Galactic HVC population. Both discrete and diffuse components are detected. We find some supporting evidence for at least three different possible origins of the high velocity gas, namely: tidal disruption, halo condensation and association with low-mass DM halos."
    },
    "0311/nucl-th0311103_arXiv.txt": {
        "abstract": "The contribution from hyperons to the bulk viscosity of neutron star matter is calculated. Compared to previous works we use for the weak interaction the one-pion exchange model rather than a current-current interaction, and include the neutral current $nn\\leftrightarrow n\\Lambda$ process. Also the sensitivity to details of the equation of state is examined. Compared to previous works we find that the contribution from hyperons to the bulk viscosity is about two orders of magnitude smaller. ",
        "introduction": "\\label{chap:hypvis;sec:int} The bulk viscosity of matter in neutron stars has recently received considerable attention in connection with damping of neutron star pulsations and gravitational radiation driven instabilities, especially in the damping of $r$-modes~\\cite{And2001}. If the $r$-modes are unstable, i.e. if the damping time-scales due to viscous processes in neutron star matter are longer than the gravitational radiation driving time-scale, a rapidly rotating neutron star could emit a significant fraction of its rotational energy and angular momentum as gravitational waves, which could be detectable. \\\\ \\indent It has been shown~\\cite{Hae2000,Hae2001a,Hae2002,Jon2001,Lin2002} that the bulk viscosity in a neutron star is caused by energy dissipation associated with  nonequilibrium weak interaction reactions in a pulsating dense matter. Strong interaction processes do not play a role, because the strong interaction equilibrium is reached so fast that these processes can be considered to be in thermal equilibrium compared to the typical pulsation time-scales of $10^{-4}-10^{-3} \\ \\textrm{s}$. The relevant nonequilibrium   weak interaction reaction rates and the magnitude of the bulk viscosity depend on the density and the composition of neutron star matter. \\\\ \\indent In a relatively low density neutron star composed mainly of neutrons $n$ with admixtures of protons $p$, electrons $e$, and muons $\\mu$, the bulk viscosity is mainly determined by  the reactions of the nonequilibrium modified Urca process, \\begin{eqnarray} N \\ + \\ n \\ \\rightarrow \\ N \\ + \\ p \\ + \\ l \\ + \\ \\overline{\\nu}_l, \\ \\ N \\ + \\ p \\ + l \\ \\rightarrow \\ N \\ + \\ n + \\ \\nu _l, \\label{eq:modurca} \\end{eqnarray} with $N=n, p$  and  $l=e, \\mu$. The bulk viscosity was studied by Sawyer~\\cite{Saw1989} for $npe$ matter and by Haensel et al.~\\cite{Hae2001a} for $npe\\mu$ matter. For the process in Eq.~(\\ref{eq:modurca}) the relaxation time, i.e. the time it takes to restore equilibrium in case of a perturbation, is strongly temperature dependent, namely $\\tau ^{-1} \\propto T^6$. \\\\ \\indent At densities $n_B$ of a few times saturation density, $n _0$ ($n _0 \\approx 0.15 \\ {\\rm fm}^{-3}$), the direct Urca process, \\begin{eqnarray} n \\ \\rightarrow \\ p \\ + \\ l \\ + \\ \\overline{\\nu}_l, \\ \\ p \\ + \\ l \\ \\rightarrow \\ n \\ + \\ \\nu _l, \\label{eq:directurca} \\end{eqnarray} may also be allowed depending on whether or not the proton fraction exceeds the Urca limit of $x_p \\approx 0.11$ \\cite{Lat1991a}. The contribution of this process to the bulk viscosity was computed by Haensel and Schaeffer~\\cite{Hae1992} for $npe$ matter and by Haensel \\textit{et al.}~\\cite{Hae2000} for $npe\\mu$ matter. Compared to the modified Urca process a smaller number of particles is involved % which leads to a weaker temperature dependence, $\\tau ^{-1} \\propto T^4$. As a consequence at typical neutron star temperatures, $ T \\sim 10^9-10^{10} \\ \\textrm{K}$, its contribution to the bulk viscosity is typically $4-6$ orders of magnitude larger than that from the modified Urca process. The largest difference compared with that from the modified Urca process is reached at the low temperatures. \\\\ \\indent At about the same densities hyperons may appear in the neutron star core, first the $\\Sigma ^{-}$ and $\\Lambda $ hyperons, followed by $\\Xi ^0$, $\\Xi ^{-}$, and $\\Sigma ^{+}$ at higher densities. Here we will restrict ourselves to the $\\Sigma ^{-}$ and $\\Lambda$ hyperons. In addition to the semileptonic hyperon processes, weak nonleptonic hyperon processes also occur, specifically the processes \\begin{equation} n \\ + \\ n \\ \\leftrightarrow \\ p \\ + \\ \\Sigma ^{-}, \\label{nnpS} \\end{equation} \\begin{equation} p \\ + \\ n \\ \\leftrightarrow \\ p \\ + \\ \\Lambda, \\label{pnpL} \\end{equation} and % \\begin{equation} n \\ + \\ n \\ \\leftrightarrow \\ n \\ +  \\ \\Lambda. \\label{nnnL} \\end{equation} At low temperatures these processes contribute more efficiently to the bulk viscosity than the direct Urca process and the semileptonic hyperon ones, because they contain no neutrino phase space factor; at typical neutron star temperatures, $T< 10^{10} \\ \\textrm{K}$, the phase space of neutrinos is almost negligible compared to that of baryons. Hence, for the weak nonleptonic hyperon processes of Eqs.~(\\ref{nnpS}-\\ref{nnnL}) the temperature dependence of the inverse relaxation time  is $\\tau ^{-1} \\propto T^2$. \\\\ \\indent Historically the first semi-quantitative calculation of bulk viscosity in neutron matter was carried out by Jones \\cite{Jon1971}. In this and all later works the weak nonleptonic process was calculated using a baryon current-current interaction, i.e. a contact $W$ exchange.  More recently the contribution to the bulk viscosity from the various  weak nonleptonic hyperon processes has been reconsidered by several authors using a modern equation of state (EoS). Haensel et al.~\\cite{Hae2002} studied the bulk viscosity for the $nnp\\Sigma ^{-}$ process % within the nonrelativistic limit, and they found the bulk viscosity to be several orders of magnitude larger than that of the direct and the modified Urca processes. Also applying the contact interaction of $W$ exchange, Lindblom and Owen~\\cite{Lin2002} calculated the contribution  of the $pnp\\Lambda$ process in addition to that of the $nnp\\Sigma ^{-}$ process. \\\\  \\indent One may question the validity of the $W$ exchange process. First there is no $W$ exchange contribution to the $nnn\\Lambda$ process and therefore it has not been considered in the above approaches~\\cite{Jon1971,Hae2002,Lin2002}. On the other hand it is well known that in nuclear physics experiments on weak $\\Lambda$ decay in large hypernuclei the $nnn\\Lambda$ process is found equally important as the $pnp\\Lambda$ process. Secondly, Jones noticed already several orders of magnitude difference between $\\tau _{nnn\\Lambda}$ and the theoretical $W$ exchange based value of $\\tau _{nnp\\Sigma ^-}$~\\cite{Jon2001}. This is interesting, because the rate for the $nnn\\Lambda$ process must depend on a very wide class of possible weak hadronic processes. Because most of these processes also contribute to $nnp\\Sigma ^-$ process, the $\\tau _{nnp\\Sigma ^-}$ in reality is probably of the same order of magnitude as $\\tau _{nnn\\Lambda}$. \\\\ \\indent However, in nuclear physics it is customary to model the weak nonleptonic $NY \\to NN$ process in terms of meson exchange (pion and kaon) with one phenomenological weak $\\pi NY$ or $KNN$ vertex. In this approach one can describe both the  observed weak decay rate of the  $\\Lambda$ as well as branching ratios in hypernuclei quite well~\\cite{Alb2002}. Therefore, in the following we will  calculate the bulk viscosity using the meson exchange picture to describe the hyperon processes in Eqs.~(\\ref{nnpS}-\\ref{nnnL}). \\\\ \\indent In practice the rates also depend on the details of the equation of state, e.g. the hyperon fractions as a function of the density. A variety of model equations of state have been constructed with widely varying properties during the last decades, some of these based upon a microscopic free space $NN$ and NY interactions, while others are phenomenological parametrizations of the energy density to higher densities. In particular the relativistic mean field approaches, as used in Ref.~\\cite{Hae2002}, and Ref.~\\cite{Lin2002}, although microscopic in nature, do not start from a realistic nucleon-nucleon interaction. % Here we employ the effective EoS based on the work of Balberg and Gal~\\cite{Bal1997}. To get some idea about the sensitivity  of the bulk viscosity to the details of the EoS we consider two different parameter sets for the density dependence of the multi-body potential energy. \\\\ \\indent Superfluidity can also influence the bulk viscosity. Below the critical temperature, superfluidity~\\cite{Hae2002,Lin2002} suppresses the reaction rates by roughly a factor $\\exp(-\\Delta /T)$ with gap energy $\\Delta$. It generally leads to a smaller bulk viscosity. Above the critical temperature, it has no influence on the reaction rates. Many properties of superfluid matter such as the gaps and the critical temperatures are still known with large uncertainties, e.g. $T_c \\sim 10^{8}-10^{10} K$. Hence, the bulk viscosity is only considered here for nonsuperfluid matter. For recent papers about the effects of superfluidity we refer to Haensel \\textit{et al.}~\\cite{Hae2002} and Lindblom and Owen~\\cite{Lin2002}.   \\\\ \\indent The main goal in this study is threefold (i) to compute the viscosity using the  more realistic one-pion exchange (OPE) instead of the contact $W$ exchange description in the nonleptonic weak hyperon processes, (ii) to include the weak neutral current  $nnn\\Lambda$ process, and (iii) to examine the sensitivity of the bulk viscosity to the EoS. The starting point of our derivation of the bulk viscosity for the $nnp\\Sigma ^{-}$, $pnp\\Lambda$, and $nnn\\Lambda$ processes is the finite temperature Green's function formalism in the quasi-particle approximation (QPA). The coupling constants for these processes using OPE can be verified from pionic hyperon decay experiments. The application of these coupling constants gives the correct order of magnitude for the rate of nonmesonic hyperon decay in large hypernuclei. Covariant expressions for the matrix elements are derived. \\\\ \\indent The plan of this paper is as follows. The EoS is discussed in Section~\\ref{chap:chap:vis;sec:EoS}. The bulk viscosity for neutron star matter is discussed in section~\\ref{chap:vis;sec:gen}. Section~\\ref{chap:vis;sec:colrate} is devoted to the collision rate of the weak nonleptonic hyperon process as in Eqs.~(\\ref{nnpS}-\\ref{nnnL}). The results are presented and discussed in section~\\ref{chap:vis;sec:resdis}. Finally the conclusion is given in section~\\ref{chap:vis;sec:con}. ",
        "conclusions": "\\label{chap:vis;sec:con} In this paper the bulk viscosity due to weak nonleptonic hyperon processes has been studied. This viscosity is relevant in connection with damping of neutron star pulsations, especially in the damping of r-modes. In particular we considered pion exchange, which is based upon empirical input, in Born approximation to describe the weak nonleptonic hyperon processes instead of the $W$ exchange~\\cite{Hae2002,Jon2001,Lin2002}. The conclusions are: i) The bulk viscosity in the OPE picture is about 1-2 orders of magnitude smaller than that in the $W$ exchange picture. ii) The rates of the $nnp\\Sigma$ and $pnp\\Lambda$ process are of the same order of magnitude using OPE. This result is in contrast with the case of $W$ exchange, because the $pnp\\Lambda$ process is one order of magnitude larger than $nnp\\Sigma$ process. iii) The $nnn\\Lambda$ process can be included in the calculation for the bulk viscosity using OPE. iv) The bulk viscosity is rather insensitive for the EoS used except with respect to the appearance of the $\\Lambda$. \\\\     \\indent Therefore, the contact $W$ exchange interaction is probably to naive for quantitative calculations. In a more realistic OPE approach the contributions of the hyperons to the bulk viscosity are less pronounced."
    },
    "0311/astro-ph0311236_arXiv.txt": {
        "abstract": "The halos of elliptical galaxies are faint and difficult to explore, but they contain vital clues to both structure and formation. We present the results of an imaging and spectroscopic survey for planetary nebulae (PNe) in the nearby elliptical \\cena.  We extend the work of Hui \\etal (1995) well into the halo of the galaxy---out to distances of 100 and 50~kpc along the major and minor axes. We now know of 1141 PNe in \\cena, 780 of which are confirmed.  Of these 780 PNe, 349 are new from this survey, and 148 are at radii beyond 20~kpc.  PNe exist at distances up to 80~kpc ($\\sim15r_e$), showing that the stellar halo extends to the limit of our data. This study represents by far the largest kinematic study of an elliptical galaxy to date, both in the number of velocity tracers and in radial extent. We confirm the large rotation of the PNe along the major axis, and show that it extends in a disk-like feature into the halo.  The rotation curve of the stars flattens at $\\sim100$~\\kms\\ with $V/\\sigma$ between 1 and 1.5, and with the velocity dispersion of the PNe falling gradually at larger radii. The two-dimensional velocity field exhibits a zero-velocity contour with a pronounced twist, showing that the galaxy potential is likely triaxial in shape, tending toward prolate.  The total dynamical mass of the galaxy within 80~kpc is $\\sim5\\times10^{11} M_{\\sun}$, with $M/L_B\\sim13$. This mass-to-light ratio is much lower than what is typically expected for elliptical galaxies. ",
        "introduction": "\\begin{deluxetable*}{llll} \\tablewidth{0pt} \\tablecaption{Observing Runs: Imaging \\label{table:obsrun_im}} \\tablehead{ \\colhead{Data} & \\colhead{Acquired AT/With} & \\colhead{Date} } \\startdata $\\!$ [\\ion{O}{3}] on-/off-band, inner halo & CTIO, Prime Focus CCD & 1993 Jun 17--19 \\\\ $\\!$ [\\ion{O}{3}] on-/off-band, outer halo & CTIO, Prime Focus CCD & 1995 Mar 24--28 \\\\ $\\!$ [\\ion{O}{3}] and {\\it V} imaging & CTIO, Mosaic II CCD camera & 2000 Jun 4 \\\\ \\enddata \\end{deluxetable*}  Extragalactic planetary nebulae (PNe) are powerful tools for studying the dynamical states of early-type galaxies.  Using their bright [\\ion{O}{3}] emission lines, we can identify PNe and typically measure their radial velocities to an accuracy of $\\sim15\\,{\\rm km \\,s^{-1}}$.  Traditionally, the kinematics of old stellar populations are derived from absorption line profiles in integrated spectroscopy (e.g.\\ Bender, Saglia, \\& Gerhard 1994; Kronawitter \\etal 2000). This method, however, is limited to the relatively high surface brightness central regions of galaxies ($r < 2r_e$), where the mass density is dominated by baryons.  Unlike in spirals where the rotation curves can be measured out to a large radii using \\ion{H}{1} observations, very few ellipticals have large scale \\ion{H}{1} disks or polar rings that can be used to probe the potential in the halo.  So, while current theories of structure formation predict that elliptical galaxies should reside in massive dark matter halos, the dynamical evidence for these halos is not nearly as complete or as conclusive as it is for spiral galaxies (e.g.\\ Gerhard \\etal 2001, Romanowsky \\etal 2003).  Using PNe, we can now obtain kinematical information in the halos of early-type galaxies at radii up to 80~kpc, allowing us to study the two-dimensional stellar kinematics and the mass distribution in these galaxies. Because PNe are evolved asymptotic giant branch (AGB) stars, their kinematics must trace that of the underlying stellar population, making them ideal test particles for dynamical studies. PNe have been used successfully as velocity test particles in many elliptical galaxies.  As early as a decade ago, Ciardullo \\etal (1993) measured velocities for 29 PNe within $r\\sim10$~kpc for NGC~3379. Other examples include Arnaboldi \\etal (1998), who discovered 43 PNe within $r\\sim16$~kpc for NGC~1316 (Fornax A), and M{\\'e}ndez \\etal (2001), who measured velocities for 535 PNe within $r\\sim14$~kpc for NGC~4697. More recently, Romanowsky \\etal (2003) presented results from a PN spectrograph study in which $\\sim100$~PNe were observed in each of three galaxies out to radii of 4--6~$r_e$.  The work most relevant to this paper is the study of Hui \\etal (1995, hereafter H95), who first investigated the PN system of \\cena. At a distance of 3.5 Mpc (Hui \\etal 1993a, hereafter H93a), NGC 5128 is the closest large elliptical galaxy to the Milky Way and is a prime candidate for PNe studies.  At the distance of \\cena, $1\\arcmin=1.02 \\ {\\rm kpc}$. Hui \\etal 1993b (H93b) identified 785 PNe in \\cena, and H95 measured velocities for 431 PNe.  This previous study probed the halo potential out to projected distances of $\\sim20\\,{\\rm kpc}$ along the major axis and $\\sim10\\,{\\rm kpc}$ along the minor axis, and produced strong evidence for: 1)~A dark matter halo, 2)~A significant quantity of angular momentum in the halo, and 3)~Intrinsic triaxiality in the true shape of the galaxy.  In this study, we present the results of an extended survey in which we have identified and measured velocities for PNe in \\cena\\ out to projected distances of 80 and 40 kpc along the major and minor axes, respectively. We identified 356 new PN candidates, many of which are beyond 20~kpc, bringing the total number of PNe in \\cena\\ to 1141. In total, we were able to measure 349 new PN radial velocities, which increases the kinematic sample to 780 PNe. Similar studies in other galaxies have typically used 2--10 times fewer test particles (whether PNe or globular clusters), and have mostly been conducted in regions where $r\\lesssim 5 r_e$.  The data set we present in this paper represents by far the largest kinematic survey of an elliptical galaxy to date, both in the number of velocity tracers, and the radial extent in the galaxy halo ($r\\sim15 r_e$). ",
        "conclusions": "\\subsection{Comparison to Merger Simulations}  N-body simulations of disk-disk mergers suggest that spheroidal stellar populations may be the result of one or more mergers. Moreover, Barnes (1992) showed that many simulated merger remnants have misaligned kinematic axes. Recent simulations are producing two-dimensional velocity fields that are useful comparisons for our data.  In particular, we refer to the work of Bendo \\& Barnes (2000) and Naab \\& Burkert (2001).  Both of these groups simulated the mergers of two stellar disk/bulge systems with different mass ratios and angular momentum vectors at varying angles.  It is important to note that these studies did not include gas dynamics or star formation, so cannot reproduce a remnant such as \\cena\\ in detail.  However, the general trends in stellar kinematics are useful baselines for comparison. They tested two families of mergers---``equal-mass'' (1:1) and ``unequal-mass'' (3:1).  All merger remnants have masses of 3--5$\\times10^{11} M_{\\sun}$, values which are comparable to the total mass of \\cena. Equal-mass mergers tend to produce a wide variety of kinematic properties, often disrupting many of the initial characteristics of the progenitor galaxies.  Unequal-mass mergers tend to preserve the disk of the larger galaxy, and hence will more consistently have rotational support.  All kinds of mergers can produce the types of velocity field morphology we see in \\cena.  These simulations show that the projected velocity contours can kink, much like those in Statler's triaxial galaxies, and in our observed velocity field.  The remnant's final kinematics will be the result of the initial angular momentum (both rotational and orbital) of the progenitor galaxies, and the shape of the final potential.  The properties of \\cena\\ are similar to some aspects of both types of mergers.  It is important to keep in mind that even if we were certain that \\cena\\ was the product of a disk-disk merger, it would be extremely difficult to uniquely constrain the exact parameters of the merger that occurred. Given the diskiness of the outer isophotes in \\cena, and the rapid rotation along the major axis, it is possible that \\cena\\ is the result of an unequal-mass disk-disk merger.  On the other hand, the severe misalignment of the kinematic axis in the halo is more typical of 1:1 mergers.  Naab \\& Burkert (2001) use simulations to show that even these 1:1 remnants may contain a disk component, although these simulated disks are much smaller in size than the one observed in \\cena.  An interesting note is that the 3:1 simulated merger remnants in Bendo \\& Barnes (2000) with kinematic properties most like the 1:1 merger remnants were also the most prolate.  These prolate remnants also showed a relatively large population of X-tube orbits, and these may be responsible for the twist in the ZVC.  Thus, in the simplest comparison that we can make with these simulations, the morphology and kinematics of \\cena\\ are broadly consistent with the remnant of an unequal-mass merger.  Two-dimensional velocity fields have the potential to tell us much more about the structure and formation of ellipticals if the corresponding models and simulations are available.  Simulations that also include gas dissipation and star formation would be especially useful. Ongoing surveys such as SAURON and the PN spectrograph promise to accumulate a wealth of data in the coming years.  \\subsection{Mass-to-Light}  By assuming isotropy, we derive a mass within 80~kpc of $5\\times10^{11} M_{\\sun}$, which gives a mass-to-light ratio of $M/L_B\\sim13$.  We can also compare this mass to the total K-band luminosity, which should be a better measure of the old stellar population of the galaxy.  Recently, the Two Micron All Sky Survey (2MASS) imaged and measured photometry for \\cena\\ in their Large Galaxy Atlas (Jarrett \\etal 2003). They measured an apparent total magnitude of $K_s = 3.942\\pm0.016$, which translates to $M_{K_s} = -23.778$, a value slightly brighter than $M^{*}_K$ for the $K$--band galaxy luminosity function (Kochanek \\etal 2001).  This gives a $K_s$-band luminosity of  $L_{K_s} = 6.97\\times10^{10} L_{\\sun}$ and a mass-to-light ratio  of $M/L_{K_s}\\sim7$.  The reader must be warned that mass estimates are fraught with systematic uncertainties due to the many assumptions made.  For instance, we have assumed that the surface radial distribution of stars continues as either a power law, de Vaucouleurs law, or Hernquist profile even at very large radii.  Given the extremely low surface brightness of the integrated stellar light, and the small number of PNe at these distances (we detect 4 PNe beyond 60~kpc), it is currently extremely difficult to accurately constrain the density profile at these radii.  However, the amount of light in these regions is negligible for the purposes of determining the total luminosity of the galaxy. Likewise, the abundance of two-dimensional velocity data behooves us to eventually model the mass distribution with non-spherical geometries.  In this case, we are fortunate that the unique, nearly edge-on disk-like component in \\cena\\ works in our favor by simplifying the geometry of the tracer PNe in the halo.  What we can see is that the stars in \\cena\\ do exist in detectable numbers at distances as great as 80~kpc (at least in one direction), and that their velocities are fairly well-behaved, extending the flat rotation profile seen in the inner halo.  Although dark matter is required to explain the observed stellar kinematics in \\cena, our value for $M/L_B$ is noticeably low when compared to those determined for ellipticals using X-ray halos and satellite kinematics.  This is interesting, especially in light of recent results from the PN spectrograph group that report low mass-to-light ratios for ellipticals within 5~$r_e$ (Romanowsky \\etal 2003).  Bahcall, Lubin, \\& Dorman (1995) compiled the values of $M/L_B$ for galaxies in the literature and found that they are closer to $200\\pm50\\ R/(100{\\rm kpc})$ for ellipticals and $60\\pm10\\ (R /100{\\rm kpc})$ for spirals.  If we assume $H_0 = 70$ and calculate $M/L_B$ for $R=80$~kpc, then these are lowered somewhat to values of $112\\pm28$ and $34\\pm6$, but both are still larger than what we are measuring by over $3\\sigma$. H95 also remarked that the dynamical mass they measured within 25~kpc was systematically lower than that measure by Forman, Jones, \\& Tucker (1985) from ROSAT data. van Gorkom \\etal (1990) reached a similar conclusion when comparing \\ion{H}{1} dynamical mass estimates with the X-ray estimates. However, Kraft \\etal (2003) recently used Chandra and XMM data to measure the total gravitating mass with 15~kpc and determined a value of $\\sim2\\times10^{11} M_{\\sun}$, which is statistically indistinguishable from the values we obtain from our mass models within the same radius.  This independent mass determination implies that our assumptions of orbital isotropy and a high inclination for the stellar disk are not likely to be grossly wrong, at least within 15~kpc.  Could it be possible that \\cena\\ appears brighter than expected, either because of an error in measurement or because it is young?  Using our value of the gravitating mass, \\cena\\ would have to be over two magnitudes fainter than what Dufour \\etal (1979) measured in order for $M/L_B$ to be at the level of 80--100.  One potential culprit is that the obscuring dust lane makes extrapolation toward the central regions essential for obtaining a magnitude that one can compare with other galaxies.  Hence, there will always be a larger than normal uncertainty in the galaxy's measured luminosity, although this is less true for infrared wavelengths. However, we think it is unlikely that such a large error was made. Could \\cena\\ be intrinsically brighter than most ellipticals because it is young?  Using the 2001 release of the Bruzual \\& Charlot (1993) evolutionary synthesis models, we find that for solar metallicity, the mean age of the ``old'' component of \\cena\\ (excluding the dust lane region) would need to have an age of 2~Gyr in order to be two magnitudes brighter than a similarly enriched 12~Gyr old population. While there is growing evidence that some fraction of the stars in \\cena\\ formed in a secondary event (e.g.\\ Rejkuba 2003), it is unlikely the entire galaxy could be this young, especially since some globular clusters in the galaxy that are at least 10~Gyr old (discussed in Peng, Ford, \\& Freeman 2004c).  Additionally, color-magnitude diagram studies of the red giant population in the halo show no evidence for a dominant young component (Harris \\& Harris 2002).  \\subsection{Comparison with Other Ellipticals}  While some of the most massive ellipticals can have central velocity dispersions of $\\sigma\\sim300$~\\kms, \\cena\\ has a more intermediate value for its central velocity dispersion with $\\sigma\\sim150$~\\kms. Likewise, with absolute magnitudes of $M_B = -20.8$ and $M_{K_s} = -23.778$, \\cena\\ has an typical luminosity for early-type galaxies. The measured amount of total luminous mass is $1.4\\times10^{11} M_{\\sun}$, and its stellar mass-to-light ratio is $M_l/L_B\\sim3.5$. These properties of \\cena's stellar component all are what one would expect for a typical elliptical galaxy, so there is little reason to suspect that its dark halo should be less massive than in other galaxies.  While it is possible that our assumption of isotropy for the stellar orbits is incorrect, fully radial orbits (which maximizes the discrepancy) might increase the total $M/L$ by a factor of two, but not a factor of ten.  Could the low value for $M/L$ be explained by the luminous component rather than the total mass (i.e.\\ by the stellar population rather than the dark matter)?  If so, then this should be evident in the colors of the galaxy. How do the mean integrated optical colors of \\cena\\ compare to those of other, perhaps less morphologically disturbed ellipticals? We compare the \\uv\\ color of \\cena\\ to those of Virgo and Coma elliptical galaxies using photometry from Bower, Lucey, \\& Ellis (1992).  The colors of ellipticals are known to correlate with galaxy properties such as velocity dispersion, surface brightness, and luminosity so it is important to compare physically similar galaxies.  We plot the color-velocity dispersion relation for these galaxies and for \\cena, because the velocity dispersion for \\cena\\ is well-measured with a value of 150~\\kms\\ (Wilkinson~\\etal 1986, H95). We have obtained our own measure of the colors of \\cena\\ using CCD photometry of the galaxy introduced in Peng \\etal (2002), and further described in Peng, Ford, \\& Freeman (2004b). We deredden our measured colors assuming $E(B-V) = 0.115$ (Schlegel, Finkbeiner, \\& Davis 1998) to obtain $(U-V)_0 = 1.28$.  These results are plotted in Figure~\\ref{figure:colvdisp}.  We also plot the significantly bluer measurement for \\cena\\ of van den Bergh (1976, vdB76).  \\begin{figure} \\epsscale{1.0} \\plotone{f13} \\caption[(\\uv) Color-Velocity Dispersion Relation for Elliptical Galaxies with \\cena] {\\small (\\uv) Color-Velocity Dispersion Relation for Elliptical Galaxies with \\cena.  We compare our measurement of the mean (\\uv) color and velocity dispersion of \\cena\\ with those measured for Virgo and Coma early-type galaxies from Bower, Lucey, \\& Ellis (1992).  The value from van den Bergh (1976) is bluer than ours, probably because that color is for a more central region.  \\cena\\ appears significantly bluer than the cluster ellipticals for its velocity dispersion, although this may due to metallicity gradients.\\label{figure:colvdisp}} \\end{figure}  The difference between our measurement and the value from vdB76 underline the importance of aperture effects.  vdB76 measured the color of the galaxy out to a radius of $2\\farcm5$, a region that is still affected by the young stars in the dust lane.  vdB76 notes that the color of the galaxy becomes bluer as one measures regions closer to the nucleus.  We chose to generously mask the dust lane to avoid any light from obviously recent star formation.  Our measurement is for a region 2--10 arc minutes along the major axis, and 5--10 arc minutes along the minor axis.  It is not surprising that our measured color is redder than that measured by vdB76.  In Figure~\\ref{figure:colvdisp}, we also see that \\cena\\ is fairly blue for its velocity dispersion compared to Virgo and Coma ellipticals. Even our redder measurement falls below the color-velocity dispersion relation.  One could argue that this may also be an aperture effect. Elliptical galaxies are known to have metallicity (and hence, color) gradients.  The aperture through which Bower \\etal measured their photometry has a radius of $2.5 h^{-1}$~kpc, which is 1.75~kpc for $H_0 = 70$.  Measurements through an aperture of this size for \\cena\\ would be dominated by the dust lane and its associated star formation. If we assume, however, that \\cena\\ does have an obscured, redder stellar population comparable to those in the cluster ellipticals, we can make some assumptions on the nature of the color gradient and estimate what its color would be.  Franx \\& Illingworth (1990) measured the mean $(U$--$R)$ color gradient in ellipticals to be $-0.23$~mag per decade of radius.  Since most of this is due to the change in (\\ub), we assume this value as an upper limit on the gradient in (\\uv).  The mean radius at which we measure our color is $\\sim6$~kpc.  If we assume that the mean radius of the Bower \\etal measurements are $\\sim1$~kpc, then the maximum expected color difference between our values would be 0.18~mag. Adding this value to our (\\uv)$_0$ color, for a value of 1.46, does in fact bring \\cena\\ to the mean of the color-velocity dispersion relation.  The major problem with this extrapolation is that we do not see color gradients in \\cena\\ of this magnitude, although our measurements are still preliminary.  So, while we measure the color of the ``old'' stellar light in \\cena\\ to be significantly bluer than the central regions of local ellipticals at the same velocity dispersion, we will need to do further work on color gradients to determine whether this is because these galaxies are intrinsically different, or because the apertures of the measurements are different.  If the colors are intrinsically different, this could be another clue that there exists an intermediate-age component in the galaxy.  A recent star-formation event of only 10\\% by mass can significantly change the optical colors of otherwise old stellar populations.  However, a difference of only 0.12~mag between \\cena\\ and the mean Virgo-Coma relation cannot account for the difference in age of 10~Gyr needed to make up the $M/L_B$ ``deficit''.  We can compare this mass-to-light value with the work of Romanowsky \\etal (2003), who recently reported that the three ellipticals in their study are not dominated by dark matter within 5~$r_e$.  For NGC~3379, they find that the ratio of dark matter to total mass within 5~$r_e$ (9~kpc) is less than 0.32.  The situation in \\cena\\ is not very different.  For \\cena, this ratio within 9~kpc is 0.26, within 5~$r_e$ is 0.55, and within 80~kpc is 0.73.  So, we find that the transition from luminous- to dark matter-dominated enclosed mass occurs at approximately 5~$r_e$.  It is possible that much of the dark matter is beyond 80~kpc, and thus not detected by our kinematic tracers.  \\cena\\ looks to be the product of an unusual merger --- systems with such extended disks are somewhat unique.  It may be that the merger parameters which led to such an extended disk also led to the dark halo being even more extended.  If that is true, the mass measured by the PNe and GC, and also by the X-ray data, would not include much of the dark matter.  However, what is perhaps a bit puzzling is that the dark matter models that we fit to the current data do not require a low concentration, which might be expected if the dark halo was very diffuse."
    },
    "0311/astro-ph0311146_arXiv.txt": {
        "abstract": "{ We present here a new determination of the ratio of energies radiated by active galactic nuclei and by stars and discuss the reasons for the apparently conflicting results found in previous studies. We conclude that the energy radiated by accretion processes onto super massive black holes is about 1 to 5\\% of the energy radiated by stars. We also estimate that the total mass accreted on average by a super massive black hole at the centre of a typical $10^{11}\\,\\Msol$ galaxy is of about $7\\,10^{7}\\,\\Msol$. ",
        "introduction": "The universe is composed of various objects with a wide variety of emission processes. However, photons have three main origins: the big bang as the source of the cosmic microwave background (CMB), nuclear reactions in stars and gravity through accretion onto compact objects, especially super massive black holes (SMBH). Currently (at $z\\simeq 0$), the ratio of the number of photons of CMB and stellar origin in the universe is of the order of 400. It is more difficult to estimate the ratio of the nuclear energy radiated by stars and the gravitational energy radiated by active galactic nuclei (AGN). Several recent studies have shown that the energy released by nucleosynthesis is evolving with cosmic time \\citep{Madaual96,Madaual98,Trenthamal99,Somervilleal01}. It further seems that the star formation history roughly matches the evolution of energy release through accretion onto black holes \\citep{Dunlop98,Franceschinial99}, since both processes peak around $z\\sim 2$. It is therefore possible to relate both phenomena and to estimate the energy ratio radiated by AGN and by stars over the history of the universe.  The determination of this ratio has been addressed several times in the literature. \\citet{Dunlop98} constructed two models of the radio luminosity function of a sample of radio-loud quasars (RLQ), the first only considering luminosity evolution and the second combining luminosity and density evolution. He then related this luminosity to the mass accreted onto the central black hole and compared these curves with the star formation history. The two curves appear to be correlated, suggesting that when 1 $\\Msol$ is accreted by a SMBH in a radio loud quasar, $10^{7}\\,\\Msol$ are used in the star formation process. Based on this relation, \\citet{Csier01} derived a ratio of the energy radiated by RLQ and stars of $E_{\\tty{RLQ}}$/$E_{\\tty{stars}}\\sim 10^{-5}$ by assuming an accretion efficiency of 10\\,\\% for the RLQ and a stellar population made of 10 solar mass stars each radiating $4\\,10^{52}$ ergs over their lifetime. Rather than considering only radio loud AGN, \\citet{Franceschinial99} use the $0.5-2$ keV X-ray emission as a measure of the energy radiated by AGN. They find that when 1 $\\Msol\\,\\unit{yr}{-1}\\,\\unit{Mpc}{-3}$ is absorbed for star formation, the $0.5-2$ keV volume emissivity from AGN is $2.4\\,10^{40}\\ \\tr{ergs}\\,\\unit{s}{-1}\\,\\unit{Mpc}{-3}$. Considering that type I and II AGN bolometric luminosity is 250 times the $0.5-2$ keV luminosity and a stellar radiative efficiency of 0.001, \\citet{Franceschinial99} finally obtain $\\Lb\\tss{(AGN)}$/$\\Lb\\tss{(SFR)}\\sim 0.1$.  A different approach was followed by \\citet{FabIwa99}. Their estimate of this ratio is based on relations linking the bulge mass of a galaxy to both the mass of its central black hole and of its stars. They obtained a value of $\\sim$\\,0.2 for the ratio of the energy radiated by accretion processes and by stars, when assuming an accretion efficiency of 10\\,\\% for the AGN and the fact that one tenth of the stellar mass is used for nuclear fusion with an efficiency of 0.6\\,\\%.  Finally, \\citet{Elvisal02} estimated that AGN contribute by at least 7\\,\\% to the total luminosity of the universe as derived from the diffuse background from submillimeter to ultraviolet wavelengths. They obtained this result by first deriving a lower limit of the AGN X-ray emission from the X-ray background (XRB) and by applying a bolometric correction determined with the average spectral energy distribution of quasars.  The aim of this study is to compare the various studies described above and to derive a new estimate of the ratio of gravitational energy released around SMBH to the energy released by nuclear fusion in stars. We first derive the radiation energy density due to stars in Sect.~\\ref{EnStars}, then the corresponding value for accretion by SMBH based on the XRB in Sect.~\\ref{EnAGNs}. The obtained ratio is compared with previous studies in Sect.~\\ref{discussion}, where we try to identify the origin of the discrepancies. ",
        "conclusions": "We derived the relative contribution of AGN and stars to the radiation energy density of the universe. The results are given for different IMF and metallicities of the stellar population. The obtained values for the energy ratio released by AGN and stars are all between 0.01 and 0.05. In the case of a Salpeter IMF and solar metallicity we obtain a ratio of 0.04. This result cannot be compared directly with previous studies because the approaches as well as the values of some parameters used in the calculation differ from one study to the other. Actually, when using similar parameters and appropriate correction factors it seems that all previous studies do converge to values between 0.02 and 0.07 for the ratio of energy radiated by AGN and stars.  Since the CMB energy density is about $\\sim 0.2\\ \\tr{eV}\\,\\unit{cm}{-3}$, the energy density due to stars is of the order of $10^{60}\\ \\tr{ergs}\\,\\unit{Mpc}{-3}\\sim 0.02\\ \\tr{eV}\\,\\unit{cm}{-3}$ and the AGN related radiation energy density is of about $8\\,10^{-4}\\ \\tr{eV}\\,\\unit{cm}{-3}$, the general picture resulting from this work is that the CMB contributes 10 times more than stars and 250 times more than AGN to the local energy density of background photons.  We also estimate that RLQ contribute about 3400 times less than the whole AGN family to the total accretion power in the universe and that the cumulated mass accreted on average by a SMBH is of about $6.5\\,10^{7}\\,\\Msol$.  \\begin{acknowledgement} We thank G. Meynet for useful discussions on star formation related issues. \\end{acknowledgement}"
    },
    "0311/astro-ph0311370_arXiv.txt": {
        "abstract": "Interstellar dust grains exhibit complex dynamics which is essential for understanding many key interstellar processes that involve dust, including grain alignment, grain growth, grain shattering etc. Grain rotational and translational motions are affected not only by gaseous collisions, but also by interactions with ions, photons, magnetic fields etc. Some of those interactions, e.g. interactions of ions with the dipole electric moment of dust grains,  require the quantum nature of the process to be accounted for. Similarly, coupling of rotational and vibrational degrees of freedom in a grain happens due to  relaxation processes, among which the process related to nuclear spins frequently is the dominant one. This coupling modifies substantially both the dynamics of rotational and translational motions by inducing grain flips. The flipping averages our certain systematic torques and forces that act on grains. As the rate of flipping is larger for smaller grains, these grains can be ``thermally trapped'', i.e. rotate at a thermal rate in spite of the presence of systematic torques. Moreover, a subset of small grains with high dipole moments may rotate at subthermal rates due to high damping arising from grain emission in microwave range of frequencies. Translational and rotational dynamics of grains is interrelated. For instance, both rotation and gas grain relative speeds affect grain alignment as well as determine grain size distribution and structure. Translational dynamics of grains is mostly dominated by grain interactions with magnetohydrodynamic turbulence. Efficient turbulent mixing of dust grains limits the degree to which grains of different sizes may be segregated in space. ",
        "introduction": "Dust is an important constituent of the interstellar medium (ISM). It interferes with observations in the optical range, but provides an insight to star-formation activity through far-infrared radiation. As dust gets aligned in the external magnetic field (Hall 1949; Hiltner 1949) it traces the magnetic field via emission and extinction polarization (see reviews by Hildebrand et al. 2000; Lazarian 2003). Grain alignment is a property of rotating grains. Fast rotation makes grains less susceptible to disorientation by collisions with gas atoms, but it also limits the efficiency of some alignment mechanisms. In addition, fast rotating grains are the sources of microwave emission that interferes with measurements of Cosmic Microwave Background (CMB) intensity and polarization ( Draine \\& Lazarian 1998a; see also recent review by Lazarian \\& Finkbeiner 2003).  The basic properties of dust (optical extinction, dust chemistry, heating of the ISM etc) strongly depend on its size distribution (Biermann \\& Harwit 1980; O'donnell \\& Mathis 1997). The latter evolves As a result of grain collisions, whose frequency and consequences depend on grain relative velocities (Draine 1985).   Rotational and translational motions are not independent. For instance, the difference in the number of active sites of H$_2$ formation over the parts of the grain perpendicular to the axis of grain rotation can result in uncompensated force accelerating the grain (Purcell 1979, Lazarian \\& Yan 2002, henceforth LY02). However, if a grain is undergoing frequent flips due to thermal fluctuations (Lazarian \\& Draine 1999a, henceforth LD99a and \\S 2) the thrust will change direction and be averaged out. On the other hand, grain alignment presents a case when translational motion affects grain rotation. For instance, mechanical alignment of grains (see Lazarian 2003 and references therein) induces higher rotational rates if grains rotate thermally initially.   Damping of rotational and translational motions have many similarities. For instance, the interaction of grains having dipole moments with passing ions is important for damping of both types of motions. For the case of damping of translational motion the treatment of the process is given in Yan, Lazarian \\& Draine (2003, henceforth YLD03). Quantum cut-off discussed in \\S 4 is applicable to this case as well.  In what follows we discuss how dust grains rotate (\\S 2), at what rate they rotate (\\S 3), and whether quantum effects are important for grain rotation (\\S 4). We describe the translational motion of grains in \\S 5. In \\S 6 we discuss astrophysical implications of grain dynamics and provide the summary in \\S 7. ",
        "conclusions": "{\\bf Implications of Rotational Dynamics}  \\emph{Grain Alignment}  Grain alignment critically depends on rotational dynamics of grains. Incorporation of thermal fluctuations within theories of grain alignment (Lazarian 1994, 1997ab, Lazarian \\& Roberge 1997, Lazarian \\& Draine 1997, Roberge \\& Lazarian 1999, Weingartner \\& Draine 2003) introduced substantial changes in our understanding how grains get aligned. Thermal flips (LD99a,b) and thermal trapping (LD99a) that they entail have radically changed some of the contemporary paradigms of grain alignment. Indeed, if grains get thermally trapped as discussed earlier, collisions with gaseous atoms randomize vectors of angular momenta to high degree. Due to efficiency of nuclear relaxation (LD99b) most of the grains do not rotate suprathermally in the absence of radiative torques. This explains why small grains for which radiative torques are not important are only marginally aligned. An alternative explanation of this fact is given by Mathis (1986) and is related to the preferential presence of ferromagnetic inclusions within large grains. However, it is claimed in Lazarian (2003) that data in Whittet et al. (2001) is consistent only with the change of the alignment efficiency with the extinction rather than the grain size. More discussion of grain alignment can be found in the paper by Roberge (2004) in this volume.    \\emph{Structure of Grains}  While grain formation may favor loosely bound conglomerates of low fractal dimension, rapid rotation of grains destroys grain that are not sufficiently compact. The tensile stress calculated in Draine \\& Lazarian (1998b) did not place severe constrains on small grain considered unless the grains consist of loosely connected parts.   {\\bf Implications of Translational Motion}  \\emph{Shattering and Coagulation}  With the grain relative velocities known, we can make predictions for grain shattering and coagulation. For shattering, we adopt the Jones' et al. (1996) results, namely, the shattering threshold is 2.7km/s for silicate grain and 1.2km/s for carbonaceous grain. The critical sticking velocity were calculated in Chokshi et al. (1993) (see also Dominik \\& Tielens 1997).% \\footnote{There are apparent misprints in the numerical coefficient of Eq.(7) in Chokshi et al.(1993) and the power index of Young's modulus in Eq.(28) of Dominik \\& Tielens (1997). % } However, experimental work by Blum (2000) shows that the critical velocity is an order of magnitude larger than the theoretical calculation. Comparing these critical velocities with the velocity curve we obtained for various media, we can get the corresponding critical size for each of them (see Table 1). However, given the uncertainties with the parameters as discussed above, we are cautious about these numbers.    \\begin{table}  \\begin{tabular}{|c|cc|c|c|c|c|} \\hline ISM& CNM& CNM& WNM& WNM& WIM& WIM\\tabularnewline \\hline Material& Si& C& Si& C& Si& C\\tabularnewline \\hline Shattering size ($\\mu$m)& NA& NA& $>0.2$& $>0.2$& $>0.003$& $>0.001$\\tabularnewline \\hline Coagulation size ($\\mu$m)& $<0.01$& $<0.02$& $<0.02$& $<0.05$& NA& NA \\tabularnewline \\hline \\end{tabular} \\caption{The critical shattering and coagulation size in different medium. NA=not applicable.} \\end{table}  \\emph{Correlation between turbulence and grain sizes}  Change in the intensity of turbulence should entail variations in grain sizes. The grain velocities are strongly dependent on the maximal velocity of turbulence $V$ at the injection scale, which is highly uncertain. Thus the critical coagulation and shattering sizes would also depend on the amplitude of the turbulence accordingly.  \\emph{Composition of cosmic ray}  It has been shown that the composition of the galactic cosmic ray seems to be better correlated with volatility of elements (Ellison, Drury \\& Meyer 1997). The more refractory elements are systematically overabundant relative to the more volatile ones. This suggests that the material locked in grains must be accelerated more efficiently than gas-phase ions. This leads to the speculation of acceleration of grain erosion products in shocks (Epstein 1980; Cesarsky \\& Bibring 1981; Bibring \\& Cesarsky 1981). If the sputtering happens upstreams, the sputtered products will be carried downstream, where they can be further accelerated to cosmic ray energies with higher efficiencies than gas-phase ions (Epstein 1980). The destruction by shocks may also be a prodigious source of PAHs, HACs and small grains (Jones, Tielens \\& Hollenbach 1996, Tielens et al. 1994). The stochastic acceleration is more efficient for low speed grains and thus can serve as a preacceleration mechanism for shock acceleration.  \\emph{Vacuum cleaning and mechanical alignment}  Our results indicate that grains can get supersonic through interaction with fast modes. Grains moving supersonically can efficiently vacuum-clean heavy elements as suggested by observations (Wakker \\& Mathis 2000). The supersonic grains can also be aligned (see a review by Lazarian 2003 and references therein). As pointed out earlier, the scattering is not efficient for slowly moving grains so that we may ignore the effect of scattering to the angular distribution of the grains. Since the acceleration of grains increases with the pitch angle of the grain, the supersonic grain motions will result in grain alignment with long axes perpendicular to the magnetic field.  \\emph{Grain Segregation and Turbulent Mixing}  Our results are also relevant to grain segregation. Grains are the major carrier of heavy elements in the ISM. The issue of grain segregation may have significant influence on the ISM metallicity. Subjected to external forcing, e.g., due to radiation pressure\\footnote{Even in the absence of radiation pressure grains can move along magnetic field lines due to the uncompensated forces, e.g. due to an unequal number of active sites of H$_2$ formation (see P79). Those forces would be mitigated in molecular clouds, which would induce inflow of dust into molecular cloud. The latter would affect metallicity of the newborn stars.}, grains gain size-dependent velocities with respect to gas. WD01 have considered the forces on dust grains exposed to anisotropic interstellar radiation fields. They included photoelectric emission, photodesorption as well as radiation pressure, and calculated the drift velocity for grains of different sizes. The velocities they got for silicate grains in the CNM range from $0.1$cm/s to $10^{3}$cm/s. Fig.~5a shows that the turbulence produces larger velocity dispersions\\footnote{Our calculation show that for the chosen set of parameters the effects of H$_{2}$ thrust are also limited.}. Those velocities are preferentially perpendicular to magnetic field, but in many cases the dispersion of velocities parallel\\footnote{This dispersion stems from both the fact that the transpositions of matter by fast modes are not exactly perpendicular to magnetic field (see plot in Lazarian \\& Yan 2002b) and due to randomization of directions of grain velocities by magnetized turbulence (Yan \\& Lazarian 2003).} to magnetic field will be comparable or greater than the regular velocities above.  More important is that if reconnection is fast (see Lazarian \\& Vishniac 1999), the mixing of grains over large scales is provided by turbulent diffusivity $\\sim VL$. Usually it was assumed that the magnetic fields strongly suppress the diffusion of charged species perpendicular to their directions. However, this assumption is questionable if we notice that motions perpendicular to the local magnetic field are hydrodynamic to high order as suggested by Cho, Lazarian \\& Vishniac (2002). In fact, recent work by Cho et al. (2003) shows that the diffusion processes in MHD turbulence are as efficient as in hydrodynamic case if the mean magnetic field is weak or moderately strong, i.e., $\\mathbf{B}_{0}$ is $\\sim$ equipartition value. This means that from the theoretical perspective, we do expect that grains can be mixed by the MHD turbulence. Therefore we believe that the segregation of very small and large grains speculated in de Oliveira-Costa et al. (2002) is unlikely to happen for typical interstellar conditions.   {\\bf Interdependence of Rotational and Translational Motions}  Coupling of different types of motions is very important for understanding grain dynamics. As we discussed earlier, the coupling of vibrational and rotational motions results in flips that change dynamics of grains.  Similarly translational and rotational motions are interdependent. For instance, frequent flips average out not only uncompensated torques, but also the uncompensated force. As we mentioned earlier, radiative torques can rescue grains from thermal trapping and therefore we show in Fig~5 that grains larger than $\\sim 5\\times 10^{-6}$~cm experience uncompensated forces due to H$_2$ formation.  Fast rotation makes atoms less susceptible to mechanical alignment through gaseous collisions. However, as grains undergo crossovers, they can still be aligned via processes that were termed in Lazarian (1995a) crossover and cross sectional alignment (see also Lazarian \\& Efroimsky 1996; Efroimsky 2002).  Streaming grains get rotation rates that are faster than that of their thermally rotating counterparts. This may allow thermally trapped grain get ``untrapped'' and therefore acquire suprathermal rotation through uncompensated Purcell's torques (P79). If grains are ``helical'' in terms of collisions with gaseous atoms the effect of grain spinning up and getting untrapped should be even move pronounced.  In addition, some characteristics of dust grains are affected both by rotational and translational motions. For instance, both fast rotation and collisions limit the structure of grains and prevent very loose aggregates to dominate the extinction."
    },
    "0311/astro-ph0311416_arXiv.txt": {
        "abstract": "We describe a technique for improving the response of a telescope chopping secondary mirror assembly by using a signal processing method based on the Lucy deconvolution technique. This technique is general and could be used for any systems, linear or non-linear, where the transfer function(s) can be measured with sufficient precision. We demonstrate how the method was implemented and show results obtained at the Caltech Submillimeter Observatory using different chop throw amplitudes and frequencies. No intervention from the telescope user is needed besides the selection of the chop throw amplitude and frequency. All the calculations are done automatically once the appropriate command is issued from the user interface of the observatory's main computer. ",
        "introduction": "Chopping scans are widely used in radioastronomy as they provide an efficient way to reduce the adverse effects that any instabilities present either in the sky signal or some telescope equipment can have on the detection of weak signals. A chopping scan is defined as a mode of observation where the telescope's secondary mirror is rotated back and forth through some angle and where the signals from both {}``end'' positions are integrated separately and subtracted from each other. This mode is to be compared with the so-called ON-OFF position (beam switching) scan where the telescope actually moves back and forth from one end position to the other. Because of the much greater speed at which the secondary mirror can move compared to the telescope, the signal subtraction happens much faster and thus an increase in the ability to detect weak signals. By moving or chopping the secondary mirror even at a relatively low frequency (e.g., 1 Hz) one can obtain a significant improvement in baseline quality when compared to a typical beam switch. In what will follow, the secondary mirror displacements are in units of arcseconds as measured on the sky.  At the Caltech Submillimeter Observatory (CSO) a chopping secondary mirror assembly was installed in 1994 and has since been used both for heterodyne receivers and bolometer cameras (e.g., SHARC and HERTZ) observations. It is composed, in part, of a carbon fiber mirror mounted on a DC brushless motor along with a system of counterweights which greatly reduces the amount of vibration noise transmitted to the observatory's receivers or cameras. The huge advantage that this vibration suppression technique brings, for the detection of weak signals comes, however, at the cost of an increase in the inertia of the chopper assembly which causes a reduction in the speed and an increase in the settling time in the response of the system.  We show in Figure \\ref{cap:pid} a block diagram of the chopping secondary system as it has been used since its installation at the CSO. Once the user of the telescope has selected a chop throw amplitude (in arcseconds) and frequency, a square wave is sent to the input of a typical Proportional-Integral-Derivative (PID) electronic controller \\citep{Ogata} where it is compared to the position signal of the mirror (obtained through a Linear Variable Differential Transformer or LVDT). The processed error signal is then sent to a power amplifier which feeds the motor and thus continuously re-positions the mirror while the PID controller acts to minimize the error signal.  Because of the relatively slow response of the chopper assembly, and the non-linearities inherent to the system (see section \\ref{sec:Implementation}), the parameters of the PID controller cannot be held fixed at a given set of values but have to be adjusted by the user of the telescope for different chop throw amplitudes and frequencies. Although this does not present a problem in principle it has been the experience that the tuning of the controller's parameters can sometimes be a time consuming effort that reduces the efficiency of the observatory. Also, since, as will be shown later, the response time of the assembly is of the order of the chopping period (or more) it is often quite difficult to find the appropriate set of parameters that will give optimum results. Too often, the outcome of such exercise is a reduction in the performance of the chopping assembly; both in its settling time and positioning accuracy.  In the following sections of this paper we will demonstrate how a signal processing method based on the Lucy deconvolution technique \\citep{Lucy} was implemented at the CSO to solve this problem and provide a system that requires no intervention from the telescope user, while keeping hardware changes to a minimum. We will start in the next section with a brief exposition of the set of equations that govern the Lucy deconvolution technique followed by a presentation of the new chopping secondary assembly (section \\ref{sec:new-assembly}). We will finish by showing how the deconvolution technique was implemented, along with the needed modifications, and by presenting some results obtained so far. ",
        "conclusions": ""
    },
    "0311/astro-ph0311620_arXiv.txt": {
        "abstract": "{ The metallicity evolution and ionization history of the universe must leave its imprint on the Cosmic Microwave Background through resonant scattering of CMB photons by atoms, ions and molecules. These transitions partially erase original temperature anisotropies of the CMB, and also generate new fluctuations. In this paper we propose a method to determine the abundance of these heavy species in low density (over-densities less than $10^4-10^5$) optically thin regions of the universe by using the unprecedented sensitivity of current and future CMB experiments. In particular, we focus our analysis on the sensitivity of the PLANCK HFI detectors in four spectral bands. We also present results for l=220 and 810 which are of interest for balloon and ground-based instruments, like ACT, APEX and SPT. We use the fine-structure transitions of atoms and ions as a source of frequency dependent optical depth ($\\tau_{\\nu}$). These transitions give different contributions to the power spectrum of  CMB in different observing channels. By comparing results from those channels, it is possible to {\\it avoid} the limit imposed by the cosmic variance and to extract information about the abundance of corresponding species at the redshift of scattering.  For PLANCK HFI we will be able to get strong constraints ($10^{-4}-10^{-2}$ solar fraction) on the abundances of neutral atoms like C, O, Si, S, and Fe in the redshift range 1-50. Fine-structure transitions of ions like CII, NII or OIII  set similar limits in the very important redshift range 3-25 and can be used to probe the ionization history of the universe. Foregrounds and other frequency dependent contaminants may set a serious limitation for this method. ",
        "introduction": "High precision observations of CMB anisotropies are giving us unique information about the angular distribution of CMB fluctuations, as well as their spectral dependence in a very broad frequency range. HFI and LFI detectors of PLANCK\\footnote{PLANCK URL site: http://astro.estec.esa.nl/Planck/} spacecraft will provide unprecedented sensitivity in 9 broad band $( \\Delta \\nu / \\nu \\sim 20-30 \\% )$ channels, uniformly distributed in the spectral region of the CMB where contribution of different foregrounds are expected to be at a minimum. CMBpol and other proposed missions are expected to reach noise levels 20 - 100 times lower than that of PLANCK HFI with technology already available (Church 2002). The WMAP\\footnote{WMAP URL site: http://map.gsfc.nasa.gov/} satellite was designed to provide measurements of the CMB temperature anisotropies in the whole sky with an average sensitivity of 35$\\mu $K per $0.3^{\\circ} \\times 0.3^{\\circ}$ pixel at the end of the mission (Bennett et al. 2002, Page et al. 2002). Once such sensitivity limit is reached, the WMAP data will be the first real-life test for all those attempts to estimate, with extremely high precision, the key parameters of our universe using this sensitivity (Bond, Efstathiou \\& Tegmark 1997, Einsestein, Hu \\& Tegmark 1999, Prunet, Sethi \\& Bouchet, 2000). Indeed, after the first 12 months of operation, the WMAP team is recovering the first multipoles of the CMB power spectrum with an accuracy of a few percent (Hinshaw et al. 2003).   In this paper we are presenting an additional use of the tremendous sensitivity of PLANCK and CMBpol, and ground-based experiments like ACT, APEX and SPT. We propose to look for or to place upper limits on the abundances of heavy elements present in the inter galactic medium and/or in optically thin clouds of gas everywhere in the redshift range $[1,500]$. We shall focus on the fine-structure lines of neutral (CI, OI, SiI, FeI, ...) and ionized (CII, NII, OIII, ...) atoms, which might provide information about the epoch of first star formation and ionization history of the universe. Limits on abundances of heavy atoms and ions can be obtained by utilizing the frequency dependent opacity that will be generated by scattering of background photons by these species; giving rise to different $\\delta C_l$-s in in different observing bands of the experiments. Although Planck HFI is expected to detect the combined signal from the strongest lines, like OI $63.2 \\mu$, CII $157.7 \\mu$ and OIII $88.4 \\mu$, future multichannel broadband CMB anisotropy experiments like CMBpol might permit to detect contributions from these lines separately from the epoch when dark ages were terminating, universe became partially ionized and heavy element production begun. It is important that the polarization signal arising due to resonance scattering depends strongly on the properties of the transition (Sazonov et al. 2002), and together with the temperature signal will permit us to separate contributions of different species.  The observed primordial acoustic peaks and angular fluctuations should not depend on the frequency at all. This is connected with the nature of Thomson scattering which produces these fluctuations both in the time of recombination of hydrogen in the universe (Peebles \\& Yu 1970, Sunyaev \\& Zel'dovich 1970) and during the secondary ionization of the universe (Sunyaev 1977, Ostriker \\& Vishniac 1986, Vishniac 1987, and more recently Gruzinov \\& Hu 1998, da Silva et al. 2000, Seljak, Burwell \\& Pen 2000, Springel, White \\& Hernquist 2001, and Gnedin \\& Jaffe 2001). In this context, WMAP polarization measurements have recently shown strong evidence favouring an early reionization scenario, with $z_r = 20^{+10}_{-9}$ (Kogut et al. 2003, 95\\% confidence). But if any amount of chemical elements are present during the dark ages, then these species will be able to scatter the CMB in their fine-structure lines. This scattering would not only partially smooth out primordial CMB anisotropies, but in addition will generate new fluctuations  through the Doppler shift of the line associated with the motion of matter connected with the growth of density perturbations. The main difference with Thompson scattering is that the latter is giving us equal contribution over the whole CMB spectrum, whereas the discussed line scattering would give different contribution to different observing channels placed at different parts of the spectrum.  Likewise, the contribution in this case would be restricted to a very {\\em thin} slice in the universe. Hence there is a possibility to detect contribution from the lines with high transition probabilities, even though the typical optical depth we would find is very small ($\\tau_{\\nu}< 10^{-4}$).  Since every line is able to work only in a given range of redshift, having observations on different wavelengths can give an upper limit to abundances of different species at different epochs. This method relies critically in the fact that intrinsic CMB temperature fluctuations are frequency independent. Therefore, in the absence of other frequency-dependent components, the difference of two CMB maps obtained at different frequencies must be sensitive to the difference in abundance of resonant species at those redshifts probed by the frequency channels. For this reason, this method is particularly sensitive to the possible presence of any frequency-dependent signal coming from foregrounds.  The idea to use the line transitions of atoms and molecules for modifying the CMB power spectrum is not new. Dubrovich (1977, 1993) proposed the use of rotational lines of primordial molecules (LiH, HD etc.) as a source of creating new angular fluctuations in the CMB, and there have been attempts to observe these molecules at high redshifts (de Bernardis et al. 1993). Varshalovich, Khersonskii and Sunyaev (1978) were trying to estimate the spectral distortions of CMB due to absorption in fine-structure lines of oxygen and carbon at redshifts around 150-300 where the temperature of electrons should be lower than the temperature of CMB due to different adiabatic indices of radiation and matter (see Zel'dovich, Kurt \\& Sunyaev 1968). Suginohara et al.(1999) probed the possibility of detecting excess flux due to emission in these lines coming from very over-dense regions in the universe. Recently Loeb (2001) and Zaldarriaga \\& Loeb (2002) (hereafter ZL02) computed the distortions connected with the recombination of primordial Lithium and scattering in the Lithium resonant doublet line. Unfortunately the wavelength of fine-structure 2P-2S transition of Lithium atoms is too short and will be unobservable by PLANCK and balloon instruments. On the very low frequency domain, future experiments like SKA and LOFAR might detect signal from neutral hydrogen 21 cm line, which also carry important ioformation from high redshifts (Sunyaev \\& Zel'dovich 1975, Madau, Meiksin \\& Rees 1997). We consider below the ground-state fine-structure lines of heavier elements, with wavelengths of the order of $50-200 \\mu$, which will in principle be observable with PLANCK and CMBpol if they are present in the redshift range $[1,500]$.  The problem of overcoming their extremely small optical depth comprises the main idea of this paper.   We will not discuss in detail the origin and ways of enrichment of the inter-galactic gas by heavy elements. This certainly requires existence of massive stars, supernova explosions, stellar and galactic winds, and even jets from disks around young stars with cold molecular gas (Yu, Billawala \\& Bally 1999). The main goal of this paper is to show that the announced sensitivity of PLANCK detectors might permit us to set very strong upper limits to the time of enrichment of inter-galactic gas by heavy elements, the time of reionization, and maybe even to detect the heavy elements in the inter-cluster medium. The census of baryons in the local universe (Fukugita et al. 1998) shows that most of the baryon remains unobserved, and the proposed method might set way to detect its existence at high redshifts, when it had moderate or low temperature. These missing baryons are centainly out of stars, interstellar gas and intergalactic gas in clusters and groups of galaxies. However, we know that such baryons should exist because they have been detected by WMAP at the last scattering surface at $z \\simeq 1100$, and are also necessary to justify the observed abundance of deuterium and $^6$Li in the early universe. Due to the wide range of redshifts under study, and the uncertain degree of mixture and clumpiness of our species in the interstellar and inter-galactic medium, we shall assume that all elements are smoothly distributed in the sky. In other words, in the present paper we shall address exclusively the homogeneous low density optically thin case, where effects related to strong over-density of gas and collisional excitation are excluded and left as subject of an upcoming work. Resonant scattering effects will produce the discussed signal even if all the gas in the enriched plasma has over-density up to $10^3$. Under  this {\\it smooth} approximation, we will find that the effect of resonant species on the CMB power spectrum will be particularly simple, especially in the high multipole range. Indeed, in the small angular scales, we shall show that the induced change in the $C_l$'s is given by $\\delta C_l \\simeq -2\\;\\tau_{X_i} C_l$, where $\\tau_{X_i}$ is the optical depth induced by the resonant scattering of species $X$ and $C_l$ is the primordial CMB power spectrum, which is currently the main target of most CMB experiments.  In section 2 we describe the basic approach for computing optical depths and corresponding deviations in the CMB power spectra due to scattering, and then show how we can put constraints on abundances using the sensitivity of present and future CMB experiments. Section 3 presents our main results for atoms and ions, and in section 4 we discuss various possible ionization and enrichment scenarios of the universe and the nature of predicted signal for each case. We will demonstrate in the Appendix B that optically thin over-dense regions with $\\langle n_e \\rangle < 10-200$ cm$^{-3}$ are contributing to the effect of resonant scattering. The effect of foreground emissions on our analysis is discussed in section 5, and we conclude in section 6. Throughout this paper we have used the parameters for standard $\\Lambda \\mbox{CDM}$ cosmology (Spergel et al. 2003),  with $\\mbox{H}_0 = 71 \\ km \\ s^{-1} \\ Mpc^{-1}, \\ \\ \\Omega_b = 0.044, \\ \\ \\Omega_m = 0.27, \\ \\Omega_{\\Lambda} = 0.73, \\ \\mbox{and} \\ \\tau_T = 0.17$ for secondary reionization. $\\mbox{T}_0$ is the present-day value of mean CMB temperature, given by $\\mbox{T}_0 = 2.726 \\mbox{K}$ (Mather et al. 1994). \\\\    \\section[]{Basic Approach}    We now discuss the basic method of obtaining the deviations in the CMB power spectrum by resonance scattering of atoms and ions, and using this deviation to constrain their abundance during the dark ages. The interaction of CMB photons with atoms and ions will mostly consist of resonant scattering with either an atomic or a rotational/vibrational transition, depending on the species under study. In the Appendix A we detail how  we introduce the optical depth due to the resonant transition ($\\tau_{X_i}$) in the CMBFAST  code (Seljak \\& Zaldarriaga, 1996). We consider a particular resonant transition $i$ for a given species $X$, with rest-frame resonant frequency $\\nu_i$. The total optical depth encountered by CMB photons on their way from the last scattering surface to us is then obtained by adding the contributions from all lines to the standard Thomson opacity: \\begin{equation} \\tau = \\tau_{T} + \\sum_i \\tau_{X_i} \\label{eq:tautot} \\end{equation}  In order to calculate $\\tau_{X_i}$, we shall recur to the formula which gives the optical depth of a resonant transition in an expanding medium (Sobolev 1946),  \\begin{equation} \\tau_{X_i}(z) = f_i \\ {\\frac{\\pi e^2}{m_e c}} {\\frac{\\lambda_i n_{X_i}(z)}{H(z)}}, \\label{basic-eqn} \\end{equation}  where  $f_i$ is the absorption oscillator strength of the resonant transition, $\\lambda_i$ is the corresponding wavelength (in rest frame), $n_{X_i}(z)$ is number density of $X$ species at redshift $z$, and $H(z)$ is the Hubble parameter at that epoch. The oscillator strength depends on $\\lambda_i$, the Einstein coefficient of the transition $A_{ul}$, and the degeneracy of the levels involved in it: \\begin{equation} f_{i} = \\frac{m_e c}{8 \\pi^2 e^2} \\frac{g_u}{g_l} \\   \\lambda_i^2  A_{ul} \\label{oscillator} \\end{equation}   A simple and elegant treatment following Gunn \\& Peterson (1965) using $\\delta$-function line profile, instead of thermally broadened gaussian for the same line, gives same value of optical depth. Further modifications are made to this formula to take into account the finite populations in the upper transition levels, which is important for atomic fine-structure transitions, where excitation temperature for them is comparable to the radiation temperature at high redshifts. Also the effect of non-zero cosmological constant can not be neglected in the low redshift universe of interest. Considering these facts, we obtain the following formula for optical depths:  \\[ \\tau_{X_i}(z)  \\ = \\ 1.7 \\times 10^{-6} \\ \\left( \\frac{{\\cal X}_{X_i}(z)}{10^{-6}} \\right) \\ \\ \\times \\] \\begin{equation} \\phantom{xxxxxxxxxxxx} \\times \\ \\left( {\\frac{{\\cal S}(z)}{{\\cal S}(z=10)}} \\right) \\ \\left( \\frac{\\lambda_i}{100 \\ \\mu} \\right) \\ \\left( \\frac{f_i}{10^{-9}} \\right) \\ {\\cal B}. \\label{sobolev} \\end{equation}  Here ${\\cal X}_{X_i}(z)$ is the ratio between the number density $N_{X_i}(z)$ of the atomic or ionic species $X$ under consideration, with respect to the baryon number density at the same redshift: $N_b(z) = 2 \\times 10^{-7}\\; (1+z)^3$ cm$^{-3}$. i.e. ${\\cal X}_{X_i}(z) = N_{X_i} / N_b$ presents the evolution of abundance of the given species due to element production, ionization and recombination processes. We propose to constrain the minimum abundance that can be detected at that redshift in {\\it solar units}, or $[X]_{min} \\equiv {\\cal X}_{X_i}(z)\\; / \\;{\\cal X}_{\\odot}$, where, for example, ${\\cal X}_{\\odot} = 3.7 \\times 10^{-4}$ for carbon. ${\\cal S}(z)$ gives the redshift dependence of optical depth in a $\\Lambda$CDM universe: ${\\cal S}(z) = (1+z)^3 \\ [(1+z)^2 (1+\\Omega_m z) - z (2+z) \\Omega_{\\Lambda}]^{- \\frac{1}{2}} $ (see, e.g., Bergstrom 1998) , and $\\lambda_i$ is the wavelength in micron. The final term ${\\cal B}$   accounts for the actual fraction of atom/ion present in their ground state, and is governed solely by the temperature of background radiation at the redshift of scattering. For a two-level system, this fraction is simply \\begin{equation} n_l = [1+(g_u/g_l)\\ \\exp(-h \\nu_i/kT_0(1+z))]^{-1}. \\label{terror} \\end{equation} We also include here the correcting term for the induced emission in the presence of the CMB and finite population of the upper level: \\begin{equation} {\\cal B} = n_l \\ \\times \\ \\left[ 1 - \\exp{\\left( -\\frac{h\\nu_i}{k_B T_0 (1+z)} \\right)} \\right]. \\label{eq:Boltz} \\end{equation}  The resonance scattering on ions and atoms in thermal equilibrium with black body radiation does not change its intensity. However, the observed CMB also has finite primordial angular fluctuations. The effect of resonant scattering is to decrease these angular fluctuations, and to bring the system more close to thermodynamic equilibrium. Therefore resulting fluctuations observed on the frequency of the line should differ from the situation on other frequencies which are far from the resonance.  At low multipoles peculiar motions arising due to the growth of large scale density perturbations become important. All ions or atoms are moving in the same direction and change the frequency of CMB photons during resonant scattering. This leads to generation of new anisotropies of background.  \\medskip We next give a simple description of the modification of the power spectrum of the CMB when it encounters a resonant species. As mentioned in the Introduction, our first hypothesis will be that the species responsible for the scattering are homogeneous, isotropic and smoothly  (i.e., not clumped) distributed in the Universe, at least during the epoch of interaction with the CMB. One may argue that heavy species are located in halos where the first stars are produced and that their distribution in the sky can non be regarded as {\\it smooth}. In this case, the final effect would depend on the typical angular size of the patches in which the species have been spread and on their total sky coverage. This paper will observe the case where those scattering patches percolate in the sky, giving rise to a smooth, homogeneous picture. The case of patchy distribution of emitting sources will be addressed in a forthcoming work, where collissional processes are studied in an extremely dense optically thick environment.   In the conformal Newtonian gauge, (also known as the longitudinal gauge), the {\\it k}-mode of the photon temperature fluctuation at current epoch is given by (Hu \\& Sugiyama 1994) \\[ \\Delta_T (k,\\eta_0,\\mu ) \\  = \\ \\int_0^{\\eta_0}d\\eta \\;e^{ ik\\mu(\\eta-\\eta_0) } \\bigl[ \\Upsilon (\\eta ) \\left( \\Delta_{T0} - i\\mu v_b \\right) \\] \\begin{equation} \\phantom{xxxxxxxxxxxxxxxxxxxxxxxxxxxxxxxx} + \\ \\dot{\\phi} + \\psi  - ik\\mu\\psi \\bigr]. \\label{eq:dT1} \\end{equation}  Here we have neglected the polarization term (which contributes at most with a few percent of the temperature amplitude). $\\eta $ denotes the conformal time, $\\Upsilon \\equiv {\\dot \\tau}(\\eta)\\exp{-\\tau (\\eta )} $ is the visibility function  and $\\tau(\\eta) = \\int_{\\eta}^{\\eta_0} d\\eta' \\dot{\\tau} (\\eta')$ is the optical depth. $\\Delta_{T0}$ accounts for the intrinsic temperature fluctuations, $v_b$ for the velocity of baryons and $\\phi$ and $\\psi$ are the scalar perturbations of the metric in this gauge. We have also neglected all tensor perturbations. If we now introduce the optical depth associated to a resonant transition as a Dirac delta of amplitude $\\tau_{X_i}$ placed at $\\eta=\\eta_{X_i}$, \\begin{equation} d\\tau(\\eta) = \\sigma_T n_e(\\eta) d\\eta + \\tau_{X_i} \\delta(\\eta-\\eta_{X_i}) d\\eta, \\label{eq:smplapp} \\end{equation} we readily obtain that the original anisotropies have been erased by a factor $e^{-\\tau_{X_i}}$, whereas new anisotropies have been generated at the same place: \\begin{equation} \\Delta_T (k,\\eta_0,\\mu )\\;  = \\;e^{-\\tau_{X_i}}\\Delta_{T_{orig}} (k,\\eta_{X_i},\t\\mu )\\; + \\; \\Delta_{T_{new}} (k,\\eta_{X_i},\\mu ) \\label{eq:newdT} \\end{equation}  In real space, if we assume that $\\tau_{X_i}\\ll 1$ , the Dirac delta approximation for the resonant transition (eq.(\\ref{eq:smplapp})) translates into \\[ \\frac{\\delta T}{T_0} ({\\bf n},z=0) = \\left(1-\\tau_{X_i}\\right)\\; \\frac{\\delta T}{T_0} ({\\bf n},z_{X_i} ) \\] \\begin{equation} \\phantom{xxxxxxxxxxx} + \\tau_{X_i}\\; \\frac{\\delta T}{T_0}\\biggr|_{new}^{lin}({\\bf n},z_{X_i})\\; +  \\; {\\cal O}\\left[\\tau_{X_i}^2\\right]. \\label{eq:newdT2} \\end{equation} Here, ${\\bf n}$ is the unitary vector denoting the observing direction, and $\\delta T/T_0|_{new}^{lin}({\\bf n},z_{X_i})$ is the linear in $\\tau_{X_i}$ term of the newly generated anisotropies\\footnote{ In what follows, our temporal coordinate will be denoted by redshift or conformal time $\\eta$ indiferently.}. We remark again that we are assuming that the species is homogeneously distributed, so that $\\tau_{X_i}$ does not depend on position. Let us now consider two different observing frequencies: the first one corresponds to the scattering redshift $z_{X_i}$, and the another corresponds to a scattering redshift too high to expect any significant presence of the species $X$. If we compute the correlation function, $\\bigl< \\frac{\\delta T}{T_0} ({\\bf n_1},z=0) \\cdot  \\frac{\\delta T}{T_0} ({\\bf n_2},z=0)\\bigr>$, in both cases, we find that the difference of these quantities will be equal to:\\footnote{This type of manipulations of the correlation function/power spectrum which avoids the Cosmic Variance limit is possible as long as the weak signal (signal induced by scattering in this case) has different spectral dependence than the CMB.} \\[ \\Delta \\left( \\bigl< \\frac{\\delta T}{T_0} ({\\bf n_1},z=0) \\cdot  \\frac{\\delta T}{T_0} ({\\bf n_2},z=0)\\bigr>\\right) = \\] \\[ \\phantom{xxxxxxx} \\tau_{X_i} \\biggl( -\\;2\\;\\bigl< \\frac{\\delta T}{T_0} ({\\bf n_1},z_{X_i}) \\cdot \\frac{\\delta T}{T_0} ({\\bf n_2},z_{X_i}) \\bigr> \\; \\] \\begin{equation} \\phantom{xxxxxxx} +\\; 2\\; \\bigl< \\frac{\\delta T}{T_0} ({\\bf n_1},z_{X_i}) \\cdot \\frac{\\delta T}{T_0}\\biggr|_{new}^{lin}({\\bf n_2},z_{X_i})\\bigr>\\biggr) + {\\cal O}\\left[\\tau_{X_i}^2\\right]. \\label{eq:dCf} \\end{equation} That is, the term linear in $\\tau_{X_i}$ of $\\Delta \\left( \\bigl< \\frac{\\delta T}{T_0} ({\\bf n_1},z=0) \\cdot  \\frac{\\delta T}{T_0} ({\\bf n_2},z=0)\\bigr> \\right)$ is the sum of two different contributions: the suppression of original fluctuations (which, as we shall see, dominates at small angular scales), and the cross-correlation of the newly generated anisotropies with the intrinsic CMB field. The first term (suppression) does not depend on the potential or velocity fields during scattering, but only on the intrinsic CMB anisotropy field. Moreover, we must remark that those contributions are evaluated {\\em at scattering epoch}\\footnote{The dependence on cosmic epoch of eqs. (\\ref{eq:newdT2},\\ref{eq:dCf}) has been simplified. We again refer to Appendix A for a formal derivation.}. Hence, if the resonant scattering takes place {\\em before} reionization, the changes in the correlation function (or in the power spectrum, as we show below) will be sensitive to the CMB anisotropy field {\\em before} being processed by the {\\it re-}ionized medium. In the Appendix A we give a detailed computation of the change in the power spectrum due to a resonant line. As we concentrate in the optically thin limit, it is possible to make a power expansion on the optical depth $\\tau_{X_i}$: \\begin{equation} \\delta C_l \\equiv C_l^{X_i} - C_l = \\tau_{X_i} \\cdot {\\cal C}_1 + \\tau_{X_i}^2 \\cdot {\\cal C}_2 + {\\cal O}( \\tau_{X_i}^3 ). \\label{eq:deltaCl} \\end{equation}  In this equation, $C_l^{X_i}$ and $C_l$ refer to the angular power spectrum multipoles in the presence and in the absence of the $X_i$ resonant transition, respectively. For the limit of very small $\\tau_{X_i}$ one can retain only the linear order, and a direct identification of $\\delta C_l$ with the abundance $[X](z)$ is possible by means of eq.(\\ref{sobolev}): the change induced in the CMB power spectrum will be proportional to the abundance of the species responsible for the resonant scattering. In agreement with what has been established when studying $\\Delta \\left( \\bigl< \\frac{\\delta T}{T_0} ({\\bf n_1},z=0) \\cdot\\frac{\\delta T}{T_0} ({\\bf n_2},z=0)\\bigr> \\right)$, we find that ${\\cal C}_1$ changes sign at some multipole $l_{cross}$: positive values of $\\delta C_l$ imply generation of new anisotropies ($l<l_{cross}$), whereas anisotropies are suppressed for $l>l_{cross}$ (see figure \\ref{oxygen2} at very low $l$). In the Appendix A we note that positive values of $\\delta C_l$ (generation) are due to the non-zero correlation existing between anisotropies generated during recombination and those generated during the resonant scattering. This correlation is due to the coupling of $k$ modes of the initial metric perturbation field in scales of the order or larger than the distance separating the two events (recombination at redshift $1100$ and resonant scattering at redshift $5$-$25$), (Hern\\'andez--Monteagudo \\& Sunyaev, in preparation).  \\begin{figure} \\centering \\psfrag{123456789l}{Multipole, $l$} \\psfrag{123456789123456789cl}{$ l(l+1)/2 \\pi \\ C_l \\ \\ (\\mu K^2)$} \\psfrag{123456789log}{log($[\\mbox{O}]_{min} / [\\mbox{O}]_{\\odot}$)} \\includegraphics[width=10.5cm, height=7cm]{dev-hfi-O63.eps} \\caption{The nature of temperature anisotropy that will be marginally detectable with Planck HFI 143 GHz channel, using $63 \\mu$ OI line. Abundance is taken from the table(\\ref{hfitable}), as $5.3 \\times 10^{-4}$ solar at redshift 32. Here {\\bf Primary} denotes  the measured temperature anisotropy ({\\it upper line}), $C_l^{pri} + \\delta C_l \\approx C_l^{pri}$, and {\\bf Line} denotes the the newly generated anistropy, $|\\delta C_l|$, arising from line scattering ({\\it bottom, with filled squares}). This $|\\delta C_l|$ is obtained by taking the difference from the 100 GHz channel, and changes sign around  $l=3$. The corresponding noise level is denoted by {\\bf HFI noise}. Also shown is the cosmic variance limit ({\\it short-dashed line}) for comparison.} \\label{oxygen2} \\end{figure}   \\subsection{$\\delta C_l$'s at small angular scales}  With respect to the high-l range, we find that the change induced in the power spectrum takes a very simple form:  \\begin{equation} \\delta C_l \\simeq - 2\\; \\tau_{X_i}\\; C_l \\label{simple} \\end{equation}  where $C_l$ is the intrinsic power spectrum. This dependence is identical to the effect of reionization on the power spectrum at small angular scales. Indeed, in that scenario, if the optical depth due to electron scattering during this epoch is given by $\\tau_{reio}$, then we have that for $l\\gg 1$ the intrinsic CMB power spectrum generated at recombination is suppressed by a factor $1-\\exp({-2\\tau_{reio}})$, or $\\approx -2 \\tau_{reio}$ if $\\tau_{reio} < 1$. So in both (resonant and Thompson) scatterings, the shape of the induced change in the power spectrum is particularly simple and equal to $\\delta C_l \\simeq -2\\tau\\;C_l$, for a given optical depth $\\tau$ and intrinsic power spectrum $C_l$. This simplifies considerably the effect of reionization on the $\\delta C_l$'s induced by resonant transitions. Indeed, if the symbol $\\Delta^{reio}$ denotes the difference of a given quantity evaluated in the presence and in the absence of reionization, then we have that, for high $l$, $\\Delta^{reio}( C_l) = -2\\tau_{reio}\\;C_l$ and $\\Delta^{reio} \\left(\\delta C_l \\right) = -2\\tau_{reio}\\; \\left( \\delta C_l \\right)$. This is shown in figure (\\ref{fig:dlta_reio}), where we plot the quantities $V \\equiv \\left| C_l^{reio} / C_l  -1 \\right|$ and $W \\equiv \\left| \\left| \\delta C_l^{reio} / \\delta C_l \\right| -1 \\right|$, that is, the relative change of $C_l$ and $\\delta C_l$ due to the presence of reionization. We have taken $\\tau_{reio}=0.17$. We see that, for high $l$, both $V$ and $W$ approach the limit $V \\simeq W \\simeq 2\\;\\tau_{reio}$, i.e., the effect of reionization is identical in the two cases.   \\begin{figure} \\centering \\includegraphics[width=7.5cm, height=6cm]{DCl_reio.eps}  \\caption[columnwidth]{For two lines placed at different redshifts, ($z=20$ in the top, $z=10$ in the bottom), we plot the quantities $V$ (solid lines) and $W$ (dotted lines), defined as $V \\equiv \\left|  C_l^{reio} / C_l  -1 \\right|$ and $W \\equiv \\left| \\left| \\delta C_l^{reio} / \\delta C_l \\right| -1 \\right|$, respectively. In both cases, in the limit of high $l$, these quantities tend to $2\\; \\tau_{reio}$, where $\\tau_{reio}$ was taken to be $0.17$. In both panels, $\\tau_{X_i}=10^{-3}$.}  \\label{fig:dlta_reio} \\end{figure}   \\subsection{Measuring $\\delta C_l$'s and abundances}  Next we investigate the limits of detectability of the optical depth of atoms and ionss in current and future CMB missions. Our starting point is the expected uncertainty in the obtained $C_l$'s from any CMB experiment (Knox 1995): \\begin{equation} \\sigma_{C_l}^2 = \\frac{2}{\\left(2l+1\\right)f_{sky}} \\left( C_l + w^{-1}B_l^{-2} \\right)^2. \\label{eq:senst1} \\end{equation} In this equation, $C_l$ is the underlying CMB power spectrum, $f_{sky}$ is the fraction of the sky covered by the experiment, and $w$ is the pixel weight given by $w^{-1}=\\sigma_N^2 \\Omega_{pix}$, with $\\sigma_N$ the noise amplitude for pixels of solid angle $\\Omega_{pix}$.  $B_l$ is the beam window function, and we shall approximate it by a gaussian. We are assuming that the noise is gaussian, uniform and uncorrelated. The first term in parentheses of eq.(\\ref{eq:senst1}) reflects the uncertainty associated to the cosmic variance. In the ideal case of an experiment with no noise, this would be the unavoidable uncertainty when identifying the {\\it observed} $C_l$'s to the $C_l$'s corresponding to a particular cosmological model. However, our interest focuses on the comparison of power spectra measured in different frequency channels with respect to a {\\it measured} reference power spectrum, (which is supposed to be free of any {\\it contamining} species of atoms and/or ions). Under the assumption that the reference power spectrum contains only CMB and noise components, we can write the following expression for the uncertainty in the measured power spectrum difference between a {\\it probe} channel and the reference channel: \\[ \\sigma_{ \\delta C_l }^2 = \\frac{2}{(2l+1)f_{sky}} \\bigl[ \\left( \\delta C_l^{prob} + w^{-1}_{prob} B_{l,prob}^{-2} \\right)^2 \\] \\begin{equation} \\phantom{xxxxxxxxxxxxxxxxxxxxxxxxxxxxx} + \\left( w^{-1}_{ref} B_{l,ref}^{-2} \\right)^2 \\bigr]. \\label{eq:newsgcl} \\end{equation} The indexes {\\it prob} and {\\it ref} refer to the {\\it probe} and {\\it reference} channel respectively, and the first term $\\delta C_l^{prob}$ refers to the cosmic variance associated to the temperature anisotropies generated by the resonant species.  Let us now assume that a CMB experiment is observing at a frequency for which it is expected to see the effects of a resonant transition of a species $X$ at redshift $1+z_{X_i}= \\nu_{X_i}/\\nu_{obs}$. One can then make use of eq.(\\ref{eq:deltaCl}) to relate the minimum $\\tau_{X_i}$ observable with the sensitivity of the experiment: \\begin{equation} \\left( \\tau_{X_i} \\right)_{min} \\simeq \\;n \\;\\frac{ \\sigma_{\\delta C_l} }{{\\cal C}_1}. \\label{tauminn} \\end{equation} The factor $n$ expresses the $\\sigma$-level necessary to claim a detection. For all the results presented in this paper, we have set the detection threshold at 3$\\sigma$, i.e. $n$ = 3. Fig.\\ref{oxygen2} shows how a very low optical depth suffices to generate changes in the power spectrum (filled squares) which are big enough to overcome the combined noise of the two experiment channels being compared, (dotted line). Once this minimum optical depth is obtained, we can easily find the corresponding minimum abundance through eq.(\\ref{sobolev}), which is one of the main goals of our paper. In fact, this limit can be further improved by a factor of $\\sim \\sqrt{\\Delta l}$ if one computes band power spectra on some multipole range $\\Delta l$. We shall address the details of these issues in the next section, where we show that for PLANCK's HFI channels, values for the optical depth as low as $10^{-5}-10^{-7}$ can be detected. \\\\   So far we are neglecting the problematic associated to the calibration between channels and the possible presence of foregrounds. Given the frequency dependent nature of the latter, the amplitudes and their characterization of the different galactic and extragalactic contaminants are critical for our purposes. We return to these issues in section 5, were we estimate the effect of foreground contamination on our analysis.\\\\ ",
        "conclusions": "In this paper, we have computed the effect of heavy elements on the Cosmic Microwave Background. We have selected those species which might leave a greater signature in the CMB power spectrum, and analyzed the possibility of constraining their abundances by the use of forthcoming CMB data.  Under the approximation of negligible drag force, we have computed what are the changes in the CMB angular power spectra introduced by resonant species placed at redshifts which are dependent on the observing and resonant frequencies. The overall effect can be decomposed into a damping or suppression of the original CMB temperature fluctuations, and a generation of new anisotropies. For the optically thin limit to which we have restricted ourselves, damping dominates over generation of new anisotropies in the intermediate and high multipole range ($l \\gtrsim 20$).  These changes in the CMB are of very small amplitude ($0.01-0.3\\mu K$), and could be distinguished from the CMB component by means of their frequency dependence. Indeed, a comparison with respect to a {\\it reference} channel containing only non-frequency dependent temperature fluctuations could be used to quantify the amount of angular power introduced by the resonant species. However, this would only be feasible if either the other frequency-dependent components are negligible or characterized with extreme accuracy. The possible presence of foregrounds, galactic of extragalactic, whose amplitude and spectral behavior is still under characterization, is a serious aspect to be taken into account. Other technical challenges, such as the calibration of different channels and the enormous sensitivity required, should be accessible from the next generation of detector technology.  We have obtained limits for abundances of heavy elements when complete foreground removal is possible, but also have shown values when all foregrounds are present in the sky map. Our analyses have particularly focused on PLANCK HFI channels, whose very low noise levels give extremely strong limits on abundances. It is easiest to observe the proposed effect with HFI at angular scales of $\\theta \\sim 20^{\\circ}$ or $l \\sim 10$ because of the very low noise of the first three HFI channels at these multipoles. Using the 143 \\& 217 GHz channels of HFI, with the 100 GHz channel as reference, limits as low as $10^{-3}-10^{-4}$ solar abundance were obtained for atoms and ions of the most important elements like carbon, nitrogen and oxygen in the redshift range $[5,30]$. At higher multipoles ($l > 200$), we have shown that future balloon and ground-based experiments like ACT, APEX and SPT can put similar constraint on abundances, where the predicted signal is stronger because of the higher amplitude of the primordial CMB signal, and effect will be directly proportinal to the optical depths in lines. The presence of foregrounds makes all these limits about a factor of $10^2$ worse, but that must be considered as the most pessimistic scenario.  In this paper we have considered scattering of CMB photons by low density gas everywhere in the universe, where optical depths in the relevant atomic or molecular lines are very low and resonant scattering effect dominates over line emission. The effect connected with scattering from atomic and ionic fine-structure transitions depends on the amount of of atoms/ions in the ground state, which is determined by the excitation temperature. In Appendix B we showed that even up to very high over-density of the plasma, the change in excitation temperature from the CMB temperature is negligible, and our simple approximation is valid. The clustering of heavy atoms and molecules in clouds should change the predictions on the modifications of the CMB power spectrum at smaller angular scales, and will be discussed in a subsequent work.   \\medskip   \\medskip"
    },
    "0311/astro-ph0311599_arXiv.txt": {
        "abstract": "Using well-defined selection criteria applied to the LEDA galaxy catalogue we have constructed a sample of elliptical galaxies that can be taken to lie in the field. Such criteria can easily be applied to theoretical simulations for direct comparison with observations. The variation of the number of `isolated' ellipticals with selection criteria is also investigated. A preliminary study of the environment of the field ellipticals shows that, in the mean, they are surrounded by a population of dwarf galaxies, out to projected radii of at least 500 kpc, with a radial density profile of $r^{-0.6\\pm 0.2}$ and a luminosity function slope of $\\alpha \\sim -1.8$. The results are compared and contrasted to the satellite population around isolated spiral galaxies. ",
        "introduction": "From the many observations of the two-point correlation function, both in its angular (e.g. Lidman and Peterson 1996, Maddox et al. 1990) and spatial (e.g. Ratcliffe et al. 1998, Giuricin et al. 2001) form, it is now well established that galaxies are clustered out to distances of at least 10 Mpc. In addition, numerous studies of the distribution of bright galaxies (e.g. Dressler 1980) show that elliptical galaxies preferentially occur in regions of high galactic density whilst spiral galaxies dominate in the field. These basic observations have led to a multitude of theoretical simulations to explain the observed clustering properties and morphological segregation (e.g. Baugh et al. 1996). A major implication of the majority of these models is that elliptical galaxies are formed from the merger of many sub-clumps during the early stages of the evolution of the Universe. Although not ruling out the presence of elliptical galaxies in low density environments, the hierarchical models suggest that on average they are very different from the cluster ellipticals, with likely evidence of recent star-formation and/or merger events. There have been many studies of elliptical galaxies in low-density environments with somewhat inconclusive results, with some studies suggesting only minor star formation at low redshifts (e.g. Silva and Bothun 1998, Bernardi et al. 1998) whilst other studies have shown strong evidence for recent merger/star formation activity (e.g. Treu et al., 1999, 2001, Kuntschner et al. 2002).  A major problem with the current comparisons between theory and observation is the lack of a consistent definition of a field galaxy. Several of the studies (e.g. Treu et al. 1999, 2001, 2002, Aars et al. 2001) use redshift surveys of the brighter galaxies to derive a sample of isolated ellipticals. However, the incompleteness of the redshift catalogues may lead to the inclusion of several ellipticals that have close neighbours. This has led many to a final visual inspection to confirm their isolated nature, destroying the objectiveness of the selection criteria. Only using an extensive redshift survey can an objective sample be constructed (e.g. Kuntschner et al. 2002) but even then incompleteness in the catalogue can lead to erroneous selection of non-field galaxies.  The studies have also concentrated on the properties of the galaxies themselves, with little, if any, analysis of the local environment of the elliptical which, from the morphology-density relationship, is likely to also have a very significant effect on the properties of the elliptical galaxy.   In this paper, we use an objective definition of a field galaxy applied to an all-sky galaxy catalogue that can also be applied to the theoretical simulations. We also investigate the variation of the number of field ellipticals with selection criteria. Using all-sky photographic surveys we have made a preliminary study of the environment of these galaxies in a search for a surrounding faint dwarf galaxy population. ",
        "conclusions": "We have identified a sample of elliptical galaxies that lie in regions where the local density of bright ($M_B < -17$) galaxies is very low, indicating that such objects are not exclusively associated with groups or clusters of galaxies. A study of the local environment around them shows an excess of faint galaxies, presumably satellites, out to a projected distance of at least 500 kpc and with a projected density varying as $r^{-0.6\\pm 0.2}$. The numbers of these dwarfs suggests a steep faint end of the luminosity function, in contradiction to that found for the field but in good agreement with that found for the outer regions of clusters. A considerable number of questions remain which can only be answered through a more detailed study of these objects and their environment."
    },
    "0311/astro-ph0311285_arXiv.txt": {
        "abstract": "We discuss recent advances in the study of dusty, massive galaxies at $z>1$ arising from the first extensive spectroscopic surveys. ",
        "introduction": "There is a growing evidence for the existence of a population of massive galaxies at comparatively early epochs in the Universe, $z\\gs 1$--2 (e.g.\\ Franx et al.\\ 2003; Glazebrook et al.\\ 2004). As extreme examples of the process of galaxy formation these systems provide a strong test of theoretical models and an excellent opportunity to observe the interplay of the physical processes involved in galaxy formation.  The first question that this population poses is: When and how were these massive galaxies formed?  Are they assembled on a timescale comparable to their star-formation in a single massive burst, or from a more prolonged build-up from sub-components which had already undergone substantial star formation?  Observations from the local Universe provide some information on the formation history of massive galaxies.  The most luminous, and by implication most massive, galaxies in the local Universe are the giant ellipticals (gE's) which frequently reside in high-density regions such as clusters.  The stellar population of these systems are dominated by old, metal-rich stars suggesting that the bulk of their star-formation activity occurred at high redshifts and hinting at a connection with the massive galaxies seen at high-redshifts, including the passive component of the Extremely Red Object (ERO) population.  Quantitative constraints can be placed on the formation redshift from the identification of luminous early-type galaxies in clusters out to $z\\sim 1.2$, as well as the homogeneity of the colours of these bright galaxies, which imply their stars were formed at $z>2$--3 (Ellis et al.\\ 1997; Blakeslee et al.\\ 2003).  The gE's have luminosities of $\\geq 1$--10L$^\\ast_V$ (10$^{10}$--10$^{11}$L$_\\odot$) and space densities of $\\sim 10^{-4}$--$10^{-5}$\\,Mpc$^{-3}$ respectively.  The central 1-Mpc of a rich cluster can easily contain $2\\times 10^{12}$L$_\\odot$ of stars in the $\\geq$L$^\\ast$ gE population, as well as a substantial amount of metals in the intracluster medium ($7\\times 10^{10}$M$_\\odot$ of Fe, Renzini 1997).  Further insights into the mode of their assembly comes from modeling the variation in colour of the elliptical population with luminosity (the colour-magnitude relation) and the internal colour gradients within individual elliptical galaxies.  Both of these observational trends can be reproduced in models where the metallicity of the stellar populations in gE's is regulated by the on-set of a superwind in a single, massive starburst (Kodama \\& Arimoto 1998; Tamura et al.\\ 2000).  The presence of a substantial mass of metals in the intracluster medium (and little evolution in this fraction out to $z\\sim 0.8$) also suggests significant outflows of material from the cluster population at $z>1.5$ (transporting energy as well as the metals) and lends credence to the existence of intense, high-density and wind-generating starbursts during the formation of these massive galaxies.  \\begin{figure} \\centerline{ \\vbox{\\psfig{file=smail_fig1a.ps,width=1.32in,angle=0} \\psfig{file=smail_fig1b.ps,width=1.32in,angle=0}} ~~~ \\psfig{file=smail_fig1c.ps,width=3.0in,angle=0}} \\caption{\\small [left] Different views of an SMG from the SCUBA Lens Survey (from Frayer et al.\\ 2000).  The top panel shows a deep {\\it HST} F814W image, the large cross indicates the position and size of the SCUBA error-circle, which encompasses several plausible faint counterparts. The smaller cross shows the 1.4-GHz VLA position and error-box indicating that the source is close to, but not coincident with the galaxy M8.  The contour plot is a deep millimetre continuum map from the OVRO interferometer which confirms that the submm source lies in an apparently blank area of sky.  The lower panel is a contour plot of a very deep Keck K-band image of this field which detects a red K$\\sim$22.5 counterpart, M12, to the radio/mm-identified SMG.  [right] The radio fluxes (or upper limits) as a function of submm flux for SMGs from the 8-mJy survey (Ivison et al.\\ 2002) and the SCUBA Lens Survey (Smail et al.\\ 2002).  This demonstrates that  60--70\\% of SMGs brighter than $\\sim 5$mJy have radio counterparts in deep 1.4-GHz maps ($\\sigma_{1.4}\\sim 5$--8$\\mu$Jy).} \\end{figure}  What other observational tests can distinguish between the possible formation scenarios of massive galaxies? If the activity required to produce the population of massive galaxies in a rich cluster occurred over a short period of time, say $<1$\\,Gyrs (or between $z=2$--3), this would need an average star formation rate (SFR) of $>2000$\\,M$_\\odot$\\,yr$^{-1}$ in a volume of roughly 1000\\,Mpc$^{-3}$. The visibility of this activity of course depends on how it is distributed.  If the galaxies form as single entities then the bulk of this activity will occur in $\\sim 100$ massive gas-rich systems on a timescale which is probably no more than 100-Myrs -- each must therefore form stars at $\\sim $200\\,M$_\\odot$\\,yr$^{-1}$.  Locally such high levels of star-formation activity are uniquely associated with highly-obscured ultraluminous infrared galaxies (ULIRGs) such as Arp 220, which has L$_{FIR}\\sim 3\\times 10^{12}$L$_\\odot$.  Dusty, gas-rich galaxies of this luminosity would be identifiable via their strong, redshifted far-infrared emission in the submillimetre (submm) waveband. Alternatively, if the massive galaxies form through the more extended merger of many subclumps then the individual star-formation events will be much less intense, 10's M$_\\odot$\\,yr$^{-1}$.  Thus one simple route to test the formation mechanism of massive galaxies is to search for populations of strongly-clustered and gas-rich, luminous submm galaxies (SMGs) at $z>2$--3, using submm cameras such as SCUBA on the JCMT (Holland et al.\\ 1999) or MAMBO on the IRAM 30-m. ",
        "conclusions": "To summarise, we have exploited the broad overlap between bright submm sources ($>5$\\,mJy) and $\\mu$Jy radio sources ($>30\\mu$Jy) to precisely locate the sources of the submm emission.  By targeting these with LRIS-B/Keck we have measured redshifts for at least half of the bright SMG population, confirming that these are high redshift, ultraluminous infrared galaxies at $z\\sim 2.3$.  The space density of these SMGs is $\\sim 1000\\times$ higher than similar luminosity galaxies at the present-day and once the likely timescale for their visibility is taken into consideration they have a space density comparable to $\\gs $\\,L$^\\ast$ ellipticals at $z=0$.  The precise redshifts for luminous SMGs enable us to search for CO emission from these systems using millimetre interferometers.  These observations probe the gas mass and dynamics in the inner regions of the SMGs, showing that these are both massive and extremely gas-rich systems: median dynamical and gas masses of $\\sim 10^{11}$\\,M$_\\odot$. The typical bright SMG in our survey hosts a reservoir of gas which is sufficient to produce the stellar mass of an $>$L$^\\ast$ galaxy on a timescale $\\ls 100$\\,Myrs (in addition to their substantial existing stellar populations).  These characteristics are precisely those expected for the prompt formation of massive galaxies through a rapid starburst at high redshifts.  The connection between the high-redshift luminous SMG population, passive EROs at $z=1$ and massive galaxies at $z=0$ is further strengthened by the detection of strong clustering in the SMG population.  This analysis also relies on the precise redshift measurements available from our Keck survey and shows that the SMGs at $z\\sim 2$--3 may be more strongly clustered than either optical QSOs or LBGs at these epochs (although see Scannapieco \\& Thacker 2003).  If the SMG halos are really this strongly clustered then it makes it difficult to see how these systems could evolve into typical LBGs. We note that this difference in the clustering strength of bright SMGs and typical LBGs is in the same sense as the factor of $\\gs 4\\times$ mass difference implied by their respective CO and H$\\alpha$ line widths. However, claims of stronger clustering in more luminous LBGs may provide a link between these two populations, as will on-going studies of the relative distribution of SMGs and the less luminous, but more common the UV-selected population at $z\\sim 2$--3."
    },
    "0311/hep-th0311173_arXiv.txt": {
        "abstract": "We consider the behaviour of the cosmological acceleration for time-dependent hyperbolic and flux compactifications of M-theory, with an exponential potential. For flat and closed cosmologies it is seen that a positive acceleration is always transient for {\\em both} compactifications. For open cosmologies, {\\em both} compactifications can give at late times periods of positive acceleration. As a function of proper time this acceleration has a power law decay and can be either positive, negative or oscillatory. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311258_arXiv.txt": {
        "abstract": " ",
        "introduction": "During the past two years we have undertaken a systematic millimeter--wave survey of molecular gas in a sample of 150 nearby, northern dwarf galaxies with IRAS emission.  We have called this survey MIDGET (Millimeter Interferometry of Dwarf Galaxies).  MIDGET has two parts: a single--dish search for CO done at the UASO Kitt Peak 12m telescope (Leroy et al., in preparation), and an interferometric follow-up of the galaxies with strong CO emission carried out at the Berkeley--Illinois--Maryland Array (BIMA). The single--dish survey targeted the centers of northern galaxies that are nearby (V$_{\\rm LSR}\\leq1000$ km s$^{-1}$), of small mass (H{\\small I} linewidths $W_{20}\\leq200$ km s$^{-1}$), compact in size, and detected by IRAS. The single--dish observations were very successful, finding $41$ new CO emitters and more than doubling the number of previously known CO sources within the defined sample. The typical $1\\sigma$ sensitivity attained was $\\sim0.2$ K km s$^{-1}$. Carbon monoxide emission was detected in galaxies with $12 + \\log (\\mbox{[O/H]}) \\sim7.9-8.5$, a metallicity regime similar to that of the Magellanic Clouds.  We are using these data to address several questions related to molecular cloud and star formation in these objects. This paper concentrates on the issue of how CO traces molecular gas as a function of heavy element abundance $Z$. ",
        "conclusions": "Using virial mass arguments, we showed that the X$_{\\rm CO}$ factor measured in resolved extragalactic GMCs is not strongly dependent on metallicity. Using the correlation between radio continuum and molecular gas, we argued that a Galactic CO--to--H$_2$ conversion factor appears to explain the observed SFR in small galaxies as well as in large galaxies. Thus, at least for the molecular gas relevant to star formation processes, X$_{\\rm CO}$ in these metal--poor environments appears to be approximately Galactic regardless of metallicity.  A few cautionary statements are in order. First, and perhaps more importantly, this result does not mean that X$_{\\rm CO}=1$ in every conceivable environment. We are only able to obtain virial mass estimates for very special objects, the brightest extragalactic GMCs. This result thus may have little bearing on the CO--to--H$_2$ conversion factor over an entire galaxy, except perhaps that molecular gas with high X$_{\\rm CO}$ does not appear to be important for star formation. Second, the physics relating the FIR and radio continuum emission to the star formation rate are poorly understood. Therefore it is possible that other effects come into play that cancel an increase in X$_{\\rm CO}$. This cancellation, however, needs to be almost perfect. Finally, to properly make quantitative arguments such as those in Figure \\ref{rc}b it is desirable to have direct metallicity measurements on a large sample of dwarf galaxies where CO is observed. We have used a rather indirect estimator of metallicity with a large intrinsic scatter (M$_{\\rm B}$), because no other data are available."
    },
    "0311/astro-ph0311291_arXiv.txt": {
        "abstract": "The spectrum and light curve of the bright X-ray source CG X-1 in the field of the Circinus galaxy are re-examined. Previous analyses have concluded that the source is an accreting black hole of mass \\GA\\ 50~\\msun\\ although it was noted that the light curve resembles that of an AM Her system. Here we show that the short period and an assumed main sequence companion constrain the mass of the companion to $<1$~\\msun. Further a possible eclipse seen during one of the \\cha\\  observations and a subsequent \\xmm\\ observation constrains the mass of the compact object to $<60$~\\msun. If such a system lies in the Circinus galaxy, then the accreting object must either radiate anisotropically or strongly violate the Eddington limit. Even if the emission is beamed, then the companion star which intercepts this flux during eclipse will be driven out of thermal equilibrium and evaporate within $\\sim10^3$~yr. We find that the observations cannot rule out an AM Her system in the Milky Way and that such a system can account for the variations seen in the light curve. ",
        "introduction": "\\label{introduction}  \\cha\\ (see e.g. Weisskopf \\etal\\ 2002 and references therein) observations have shown that normal galaxies of the size of the Milky Way (e.g.\\ M81, Tennant \\etal\\ 2001; Swartz \\etal\\ 2003) host hundreds of discrete X-ray sources with luminosities above $10^{36}$~erg~s$^{-1}$. These sources are comprised of disparate groups of objects, including supernova remnants, X-ray binaries, and more exotic objects such as the supersoft X-ray sources. While the majority of the X-ray emitters are binaries containing a neutron star or stellar-mass black hole accreting material from a companion star, some appear to be quite luminous, with implied accretion rates well above the Eddington limit for a 1.5-\\msun\\ compact star. These peculiar bright sources are often referred to as ultra-luminous X-ray sources (ULXs).  Several explanations for ULXs have been proposed. They may be intermediate-mass black holes (IMBHs) with masses $\\sim 10^2 - 10^4$~\\msun, less than the $>10^6$~\\msun\\ inferred for active galactic nuclei (Colbert \\& Mushotzky 1999; Matsumoto \\etal\\ 2001; also Ebisuzaki \\etal\\ 2001). They may be stellar-mass black holes (SMBHs) in high-mass X-ray binaries with mildly beamed radiation during a thermal-timescale mass transfer phase (King \\etal\\ 2001), perhaps with contributions from synchrotron and inverse-Compton emission from jets in addition to the thermal X-ray emission from the accretion disks (K\\\"{o}rding, Falcke \\& Markoff 2002). They may be long-lasting outbursts in low-mass X-ray binaries, such as in microquasars (King 2002). The nature of ULXs is still under debate, in part, because the X-ray observations alone have difficulty in distinguishing amongst the various alternatives.  A ULX candidate in the Circinus galaxy ($\\alpha_{2000} = 14^{\\rm h}\\, 13^{\\rm m}\\, 12.^{\\rm s}3\\ $ and $\\delta_{2000} = -65^{\\degree}\\, 20\\arcmin\\, 13\\farcs$, hereafter CG~X-1) appears to be periodic with a period of $27\\pm 0.7$~ks (Bauer \\etal\\ 2001, Bianchi \\etal\\ 2002). Furthermore, at times there is a sharp deep dip in the X-ray light curve reminiscent of eclipses (Weisskopf 2002, and shown in Figure~\\ref{lc_365_356}). These signatures are important and are useful to determine the nature the system. Bauer \\etal\\ (2001) noted the light curve of CG~X-1 resembles that of AM Her-type cataclysmic variable (CV) system (for a review of AM Her binaries and CVs, see Cropper 1990; Warner 1995), but argued against an AM Her interpretation as follows: (1) No optical counterpart brighter than 25.3 mag is found in {\\sl HST} images, implying that if the source is an AM Her system its distance is greater than 1.2~kpc, therefore, the X-ray luminosity is too bright for a typical AM Her system. (2) The X-ray spectrum does not have the two components (blackbody and optically thin thermal plasma) typical of AM Her systems. (3) The inferred period of 27~ks (7.5~hr) is significantly larger than those of many AM Her systems, which are typically 1.5-4~hr. (4) The time variability exhibited by CG~X-1 shows flickering similar to those seen in the Galactic black-hole X-ray sources. (5) There is only a small probability ($< 0.06\\%$) of there being a foreground star in the field and an even smaller probability that such a foreground star would be an AM Her system.  Here we re-examine the spectral and timing data of CG~X-1 using data from the \\cha\\ (\\S~\\ref{chandra}), \\xmm\\ (\\S~\\ref{xmm}), and HST (\\S~\\ref{s:opt}) archives. We show that the short period and an assumed main sequence companion constrain the mass of the companion to $<1$~\\msun\\ (\\S~\\ref{parameter1}). Further a possible eclipse  constrains the compact object to $\\LA 60$~\\msun\\ (\\S~\\ref{parameter2}). It is therefore unlikely that CG~X-1 is an IMBH (\\S~\\ref{parameter2}) or a SMBH in Circinus (\\S~\\ref{parameter3}). In \\S~\\ref{cygX3} we consider the possibility of a He-burning companion. Finally, we show that the original arguments of Bauer \\etal\\ (2001) do  not conclusively rule out a CV system (\\S~\\ref{parameter4}). Although establishing that a short period ULX is present in the Circinus galaxy would be exciting, we feel that the explanation that this source is a CV in our own galaxy is also compelling and cannot be dismissed. ",
        "conclusions": "We presented spectral and timing analysis of the x-ray source CG~X-1 using archival \\cha, \\xmm, and HST data. The \\xmm\\ spectrum was  best fit by a thermal bremsstrahlung model with  $kT=7.4 \\pm 0.6$~keV. Based on the interpretation that the observed period is orbital and an assumed main sequence companion, the companion star is \\LA 1~\\msun. A possible eclipse constrains the mass of the primary compact star to less than 60~\\msun, thus ruling out an IMBH. Further, we showed that if one assumes the source is a SMBH binary in the Circinus galaxy, then the high luminosity of the X-ray source coupled with the small orbital separation and small mass of the companion would quickly drive the companion out of thermal equilibrium. The high X-ray luminosity in this case also causes the companion to evaporate within $10^3$~yrs. These arguments make it unlikely that CG X-1 is associated with the Circinus galaxy. Based on our analysis of extant HST data we have shown that optical observations do not rule out a K5 or later companion star in the Milky Way. Finally, we emphasize that the roughly square-wave light curve (Figure~\\ref{lc_374_62877}) is similar to that of AM Her systems where the accretion is funneled onto a single pole of a white dwarf and an accretion shock is formed. When the source counting rate is high throughout the orbit (as in Figure~\\ref{lc_365_356}), accretion is either occurring onto two poles of the white dwarf or a disk has formed. Such changes are naturally explained by the presence of the magnetic white dwarf (see e.g.\\ Warner 1995 and references therein). While the chance probability of finding an AM Her system in this field is small, the AM Her interpretation is consistent with all the data and cannot be ruled out. Further X-ray and optical observation are called for to firmly establish the nature of this source."
    },
    "0311/astro-ph0311544_arXiv.txt": {
        "abstract": "We present a method for ultra-fast confrontation of the WMAP cosmic microwave background observations with theoretical models, implemented as a publicly available software package called CMBfit, useful for anyone wishing to measure cosmological parameters by combining WMAP with other observations. The method takes advantage of the underlying physics by transforming into a set of parameters where the WMAP likelihood surface is accurately fit by the exponential of a quartic or sextic polynomial. Building on previous physics based approximations by Hu {\\etal}, Kosowsky {\\etal} and Chu {\\etal}, it combines their speed with precision cosmology grade accuracy. A Fortran code for computing the WMAP likelihood for a given set of parameters is provided, pre-calibrated against CMBfast, accurate to $\\Delta\\ln\\Ell\\sim 0.05$ over the entire $2\\sigma$ region of the parameter space for 6 parameter ``vanilla'' $\\Lambda CDM$ models. We also provide 7-parameter fits including spatial curvature, gravitational waves and a running spectral index. ",
        "introduction": "The impressive sensitivity of the long awaited Wilkinson Microwave Anisotropy Probe (WMAP) data allows for unprecedented constraints on cosmological models \\cite{bennett03,hinshaw03,kogut03,page03,spergel03,verde03}. The measurements have strengthened the case for the cosmological concordance model\\cite{verde03}, the inflationary $\\Lambda$CDM model, a flat Universe which is currently accelerating due to mysterious dark energy, where most matter is in the form of collisionless cold dark matter, and where the initial conditions for the fluctuations are adiabatic. Because of well known cosmic microwave background (CMB) parameter degeneracies \\cite{hu00,spergel03,bridle}, the true impressiveness of the data is most clearly demonstrated by either imposing reasonable priors, combining the data with complimentary data sets, or both. The most recent precision data-sets are the Sloan Digital Sky Survey (SDSS) galaxy power spectrum (SDSS) \\cite{SDSS} and a new Supernovae Ia compilation \\cite{tonry03}, and combining these with the WMAP constraints have further narrowed error bars \\cite{SDSS2}, giving us cosmological parameters at a precision not thought possible only few years ago.    So how are CMB data-sets used to estimate parameters? Although constraints on cosmological parameters from the CMB can in principle be extracted directly from the maps themselves, this is effectively prohibited by the huge computational power necessary to perform the required likelihood calculations (although for a novel approach, see \\cite{wandelt03}). Instead, the more efficient parameter estimation method commonly employed uses the angular power spectrum of the map as an intermediate step\\cite{tegmark97,bond87,bond98}.  Likelihoods on the basis of power spectrum constraints are much faster to calculate, the slowest step being the computation of the angular power spectrum numerically. This can be done either through integration of the full Boltzmann equation (with CMBfast \\cite{CMBfast} or the modifications CAMB\\cite{camb} or CMBEASY\\cite{cmbeasy}) or using an approximate shortcut such as DASh\\cite{dash}. Although steady improvements in both computer power and algorithm performance have made these calculations significantly faster, the CMB power spectrum likelihood calculation is still the bottleneck in any parameter estimation process by many orders of magnitude.  Another problem with the estimation process are the near-degeneracies between some cosmological parameters. Elongated, banana shaped contours on the likelihood surface make search algorithms less efficient. Several authors have advocated various transformations from cosmological parameters to ``physical\" parameters more directly linked to features in the power spectrum\\cite{hu00,kosowsky02}. Chu {\\etal} \\cite{chu02} advocated a new such set of ``normal\" parameters whose probability distributions are well fitted by Gaussian distributions.  Our key idea explored in this paper is that if the likelihood function is roughly a multivariate Gaussian in the transformed parameters, then it should be very accurately approximated by the exponential of a convenient higher-order polynomial. A Gaussian likelihood surface $\\Ell$ corresponds to the log-likelihood $\\ln\\Ell$ being quadratic, and a quadratic Taylor expansion is of course an accurate approximation of any function near its maximum. Very far from the maximum, the Gaussian approximation again becomes accurate, since both it and the true likelihood $\\Ell$ are vanishingly small. We will see that in most cases, small cubic and quartic corrections to $\\ln\\Ell$ help provide a near-perfect fit to $\\Ell$ everywhere. The WMAP team employed such quartic fits to determine reasonable step sizes for their Markov Chain Monte Carlo search algorithm \\cite{verde03}. Here we go further and show how polynomial fits can conveniently store essentially all the cosmological information inherent in the WMAP power spectra. This reduces calculation of CMB likelihoods to simply evaluating the exponential of an $n^{th}$ order polynomial. Although CMBfast still needs to be run in order to obtain a sufficient sampling of the likelihood surface, this only has to be done once for each model space and dataset. For the models explored in this paper, it has already been done and the polynomial coefficients have been obtained. The polynomial fit can thus be used to run Monte Carlo Markov Chains (MCMC) including various non-CMB data-sets  many orders of magnitude faster, dramatically reducing the need for computing power for cosmological parameter estimation purposes involving WMAP data. This is an improvement on importance sampling methods which notoriously depletes the chain of points as the added data sets shift or narrow the confidence regions. It means joint likelihood analyses between WMAP and other surveys which would previously take weeks or months of computer time on a good workstation being finished in an afternoon.  As mentioned, the last few years have seen the rise of a cosmological concordance model\\cite{xiaomin02,spergel03,SDSS2}, the $\\Lambda$CDM flat inflationary cosmology which has been confirmed and strengthened by WMAP, SDSS and new supernova observations. However as cosmological ideas and trends change, we choose to keep an open mind and allow for extensions to the 6-parameter inflationary $\\Lambda$CDM model by including models with spatial curvature, gravitational waves and a running scalar spectral index. To allow the user to include CMB polarization information separately, we also perform separate likelihood fits excluding and including WMAP polarization data for the 6 parameter scenario. ",
        "conclusions": "The big bottleneck in parameter estimation with CMB data has hitherto been the computation of theoretical power spectra by means of CMBfast, CAMB, CMBEASY, DASh or similar, for a reasonably fast computer typically requiring from $~5$ seconds to several minutes for non-flat models with massive neutrino for CMBfast, and $~1$ second for DASh. In contrast, our method requires less than a millisecond per model, since we have already precomputed the WMAP likelihoods using CMBfast and distilled the results into a convenient fitting function. After transforming into a physically motivated parameter set, the only calculation needed is the evaluation of an $n^{th}$ order polynomial with as few as 210 coefficients for the 6 parameter quartic case. As the number of coefficients may actually exceed the number of measured $C_l$s the method is clearly not a way of compressing the data. Rather we are compressing the amount of calculation necessary to convert cosmological parameters into likelihood values.  The typical error in this approach is $\\Delta\\Ell\\sim 0.01$ (with the peak likelihood normalized to unity) and $\\Delta\\ln\\Ell\\sim 0.05 - 0.5$. To place these inaccuracies in context, let us first discuss how they compare with the inaccuracies in other methods, then  the question of how good is good enough.   The latest version of CMBfast\\cite{seljak03} has an accuracy down to $\\sim 0.1\\%$ (r.m.s). It is unclear if other recent packages match this number, but they are generally good to the $1-2\\%$ level. Despite computer advances, the CMBfast code is still slow when several hundred thousand model computations are required as for grid calculations or Markov Chain Monte Carlo applications. Several shortcuts have therefore been developed in order to compute the Likelihood faster than the full CMBfast. The ``$k$-split\" method in CMBfast due to Tegmark {\\etal} \\cite{ksplit} utilizes the fact that the high-$\\l$ and low-$\\l$ ends of the power spectrum depends on different parameters, and so calculates the two parts separately. The end result is then combined into a final power spectrum. The Davis Anisotropy Shortcut (DASh)\\cite{dash} takes this method further, and creates the power spectrum by interpolation between points in a huge pre-computed grid of transfer functions. Comparing likelihoods calculated with $k$-split CMBfast and DASh with those calculated with the maximally accurate CMBfast, we found that the errors from our fitting method in both $\\Delta\\ln\\Ell$ and $\\Delta\\Ell$ are significantly smaller than the errors in these quantities from the DASh or $k$-split approximations. For instance, DASh gave r.m.s.~$\\Delta\\ln\\Ell\\approx 0.9$ for the 6-parameter case using only unpolarized WMAP information and the ksplit inaccuracies are similar, which should be compared with our value $\\Delta\\ln\\Ell\\approx 0.08$ from Table II.  How accurate is accurate enough? For most applications, the key issue is not inaccuracies in the power spectra or likelihoods, but inaccuracies in measured cosmological parameters. If an approximation shifts the measured values of all parameters by much less than their statistical uncertainties, it is for all practical purposes irrelevant.  We have also shown that the errors in the marginalized distributions are minimal, and that our method should be used by anyone who do not have excessive amounts of time and computer power on their hands. Our last six figures show that our approximation is clearly accurate enough in this sense except for the above-mentioned glitches in the tensor case. There is, however, one subtle and somewhat counterintuitive point that is worth clarifying. Since WMAP constrains the power spectrum normalization to the $10^{-3}$ level when all other parameters are held fixed, this means that even a seemingly tiny $\\sim 0.1\\%$ inaccuracy in the power spectrum can in principle cause a change of order unity in $\\chi^2$ and $\\Delta\\ln \\Ell$, \\ie, an inaccuracy larger than that of our fitting method, and one may naively expect that inaccuracies of order $\\Delta\\ln \\Ell\\sim 1$ would affect the parameter measurements at the $1\\sigma$ level. Although this would be true if only one parameter were being measured, the situation at hand is actually much better because of degeneracies. As long as the inaccuracies do not exactly mimic the effect of changing some cosmological parameter, they will alter $\\ln \\Ell$ mainly via the narrowest directions in the multidimensional degeneracy banana and hence have little effect on the final results. For instance, if a fitting inaccuracy causes a relative error in $\\Theta^E_s$ at the $10^{-3}$ level, it will have no noticeable effect on the estimates any of the vanilla cosmological parameters, whose error bars are dominated by the eigenparameters with the largest rather than the smallest uncertainties. The bottom line is therefore that although the the DASh and $k$-split $\\ln \\Ell$ inaccuracies of order unity may sound worrisome, these two approximations are nonetheless sufficient to give negligible inaccuracies in cosmological parameter measurements, and the still better accuracy of our fitting method is actually overkill.  Our method has several applications. First, it is extremely useful when combining the WMAP-data with non-CMB data such as the galaxy surveys, weak lensing data, supernovae data, {\\etc}, where likelihood calculations are fast relative to CMB power spectrum calculations. Using our method to measure parameters from the combined WMAP and SDSS data \\cite{SDSS2} enabled us to do some of the analysis with dramatically increased speed\\footnote{ The analysis in \\cite{SDSS2} used no chains from the WMAP team, but applied our fitting method to the WMAP Monte Carlo Markov Chains described in Table 8 of that paper.}. Second, even when running Markov chains for models not considered here, one can get good results by simply running the chain for long enough to acquire a sufficient (but not necessarily statistically fair) sample of the surface, then compute the polynomial fit, and go on to use this fit for the remainder of the job --- until all relevant mixing and convergence tests are fulfilled. A further obvious application is generating extremely long Markov chains, eliminating sampling errors almost completely. Finally, should the reader be intent on requiring the best accuracy that CMBfast can offer, our approach can still be helpful: It is a fundamental fact about MCMC algorithms that they do not produce completely uncorrelated points, and the correlation length of the chains can easily be several hundred points, and only gets as low as 45 even for our simply 6-parameter WMAP chains. Thus the calculations may be significantly accelerated by creating a statistically random sample by means of the polynomial fit, thinning the chain to every 200 points or so, and calculating the CMBfast power spectrum and WMAP likelihoods for these points. A final likelihood sample with the strictly correct distribution may then be obtained using importance sampling as described in \\cite{cosmomc}.  We supply Fortran routines for computing the likelihoods for all the cases given in the text at http://www.hep.upenn.edu/{$\\sim$}sandvik/CMBfit.html , and we plan to complement this work with further models and fits in the future."
    },
    "0311/astro-ph0311128_arXiv.txt": {
        "abstract": "{We study the viscosity driven bar mode (Jacobi-like) instability of rapidly rotating quark matter stars (strange stars) in general relativity. A triaxial, ``bar shaped'' compact star could be an efficient source of continuous wave gravitational radiation in the frequency range of the forthcoming interferometric detectors. We locate the secular instability point along several constant baryon mass sequences of uniformly rotating strange stars described by the MIT bag model. Contrary to neutron stars, strange stars with $T/|W|$ (the ratio of the rotational kinetic energy to the absolute value of the gravitational potential energy) much lower than the corresponding value for the mass-shed limit can be secularly unstable to  bar mode formation if shear viscosity is high enough to damp out any deviation from uniform rotation. The instability develops for a broad range of gravitational masses and rotational frequencies of strange quark stars. It imposes strong constraints on the lower limit of the frequency at the innermost stable circular orbit around rapidly rotating strange stars. The above results are robust for all linear self-bound equations of state assuming the growth time of the instability is faster than the damping timescale.  Whether the instability can grow or not  depends on many different physical quantities (e.g. value of viscosities (rather uncertain)). We discuss astrophysical scenarios where triaxial instabilities (r-mode and viscosity driven instability) could be relevant (a new born star, an old star spinning up by accretion) in strange stars described by the standard MIT bag model of normal quark matter. The spin evolution of a strange star strongly depends on the strange quark mass. Taking into account actual values of viscosities in strange quark matter and neglecting the magnetic field we show that Jacobi-like instability cannot develop in any astrophysicaly relevant temperature windows. The main result is that strange quark stars described by the MIT bag model can be accelerated to very high frequency in Low Mass X-ray binaries if the strange quark mass is consistent with values based on particle data tables. ",
        "introduction": "The first generation laser interferometric gravitational wave detectors such as LIGO, VIRGO and GEO600 should be fully operationally soon and second generation detectors like LIGO II within a few years. Triaxial instabilities of rotating compact stars (neutron stars, strange quark stars) can play an important role as emission mechanisms of gravitational waves in the frequency range of these detectors (see e.g. Stergioulas 1998, 2003 or Andersson 2003 for a review). A rapidly rotating relativistic star can spontaneously break its axial symmetry if the ratio of the rotational kinetic energy to the absolute value of the gravitational potential energy $T/|W|$ exceeds some critical value (Gondek-Rosi\\'nska \\& Gourgoulhon 2002, Di Girolamo \\& Vietri 2002, Shapiro \\& Zane 1998, Bonazzola et al. 1996, 1998).  The $m=2$, $l=2$ ``bar mode'' instability is the fastest-growing instability if a rotation rate is sufficiently high.  Neutron stars can be forced to rotate rigidly if viscosity is high enough to damp out any deviation from uniform rotation.  The viscosity-driven bar mode (Jacobi-like) instability of neutron stars has been studied in general relativity by Bonazzola, Frieben \\& Gourgoulhon (1996, 1998), who find that only the stiffest equations of state (EOS) allow symmetry breaking. Since for ``realistic'' EOS of neutron stars the ratio of $T/|W|$ has a low value (usually lower than 0.14), this instability takes place only when a star is rotating at a frequency close to the mass-shed limit and has high mass.  Strange quark stars (SS) are currently considered as a possible alternative to neutron stars as compact objects (see e.g. Weber 1999 for a review) . Claims of discovery have been recently announced (Drake et al. 2002), but the evidence is not yet decisive (Walter \\& Lattimer 2002).  It has been suggested by Gourgoulhon et al. (1999), Gondek-Rosi\\'nska et al. (2000a, b; 2001a) that a triaxial instability could develop more easily in a uniformly rotating strange star than in a neutron star. It was shown that independent of any details of the EOS of strange matter, the ratio $T/|W|$ for self-bound objects like strange stars can  be  even more than twice as high as for any models of neutron stars. No stability analysis in general relativity has been performed for strange stars untill now, because numerical constraints prevented investigators from treating stars with strongly discontinuous density profiles at the surface.  In the present paper we perform a numerical study of the secular triaxial instability of rigidly rotating strange stars in general relativity. We use the code developed recently by Gondek-Rosi\\'nska \\& Gourgoulhon (2002) for locating the secular bar mode instability points of uniformly rotating incompressible fluid bodies in general relativity. Numerous tests have been performed to assess the validity of the method and the accuracy of the numerical code (Gondek-Rosi\\'nska \\& Gourgoulhon 2002).  In this paper we shall focus on the following questions: for which gravitational masses and rotational frequencies of strange stars does the Jacobi-like instability occur? Does gravitational radiation set strong limits on spin rate of strange stars and on the frequency at the innermost stable circular orbit around rapidly rotating strange stars? What astrophysical scenarios does the viscosity driven instability allow one to develop? How do our conclusions depend on physical parameters of the MIT bag model?  The plan of the paper is as follows. The description of previous studies of the bar mode instabilities is presented in Sec. 2. Section 3 and 4 are devoted to the description of the strange matter EOS used by us and of the numerical technique we employed, respectively.  The numerical results and discussion are presented in two sections Sec. 5 and 6. In Sec. 5 we examine the necessary condition for the onset of the instability: that the frequency of a corotating mode goes through zero in a frame co-rotating with the star. This condition yields the critical angular velocity (at a given compactness) for which the instability could start to operate, provided it is not damped by other mechanisms.  The results are robust for any linear self-bound EOS.  In Sec 6. we discuss the astrophysical scenarios when triaxial instabilities could be relevant. This part includes an estimate of different timescales (uncertain).  Finally, Sec. 7 provides a summary of our work. ",
        "conclusions": "We have performed numerical investigations of the Jacobi-like bar mode secular instability of rotating strange quark stars described by the MIT bag model. In particular, we have computed the domain of the parameter space of stationary rigidly rotating configurations in which triaxial instability develops.  It has been found that, contrary to neutron stars, the triaxial instability can set in quite far towards the mass-shed limit and is allowed for small mass stars as well if the viscosity is high enough to damp out any deviation from rigid rotation.  The lower mass is the lower rotational frequency at which Jacobi-like instability can develop.  The viscosity driven instability does not set an upper limit on the spin rate (but on $T/|W|$ or oblateness) of strange stars - the configuration marginally stable against non-axisymmetric perturbation can rotate with a frequency higher than mass-shedding configurations (the maximal rotating configuration is not a Keplerian one for SS). A normal and low mass supramassive strange star gaining angular momentum always slows down before reaching the Keplerian limit. For a high-mass supramassive strange star sequence, the mass-shedding configuration is the one with the lowest rotational frequency in the sequence. The instability, if relevant, can impose a strong constrain on the lower limit of the frequency at the innermost stable circular orbit around rapidly rotating strange stars. The results are robust for all linear self-bound equations of state.  The critical rotation frequency for viscosity driven instability in the case of strange stars described by the MIT bag model is high, $0.8-1.2 {\\rm\\ kHz}$ (depending on the stellar mass and the detailed form of the SQM EOS).  We have discussed astrophysical scenarios when triaxial instabilities (r-mode and viscosity driven instability) could be relevant in strange stars described by the standard MIT bag model of normal quark matter (any kind of superconductivity being excluded). We have shown, taking into account actual values of shear viscosity, that viscosity driven instability cannot develop in any astrophysicaly relevant temperature windows.   Whether a strange star can be a good emitter of gravitational waves or not depends on CFS instability.  In our discussion we focused on the most promising CFS instability (r-mode). The spin evolution of strange stars and emission of gravitational waves strongly depend on the mass of the strange quark (caused by the strong dependence of bulk viscosity on the mass of a strange quark). If we take the strange quark mass to be $m_{\\rm s}c^2 \\sim 200\\ {\\rm MeV}$ or higher (consistent with renormalization and particle data tables) then r-mode instability is relevant only in a narrow temperature window and seems to be unimportant in Low Mass X-Ray Binaries. In this case an old star can be spun up in LMXBs to very high rotation rates and a low frequency at the marginally stable orbit can be obtained ($\\sim 1$kHz).  If on the contrary the strange quark mass is low, $m_{\\rm s}c^2 \\sim 100 {\\rm\\ MeV}$, then a star accelerated in LMXBs to a frequency of $0.4-0.6\\ $kHz can be a good emitter of persistent gravitational waves (see Andersson, Jones, Kokkotas 2002 for a discussion) that ought to be detectable with large-scale interferometers.  In this case the r-mode instability can impose strong constraints on the lower limits on rotational period and then on the maximal orbital frequency in LMXBs.  The lowest frequency at the marginally stable orbit could be obtained for slowly-rotating massive strange stars.  A young strange star would not be spun down significantly by the r-modes during its first few months of existence (because of its very rapid cooling rate) until the star has cooled to a temperature of a few times $ 10^7- 10^8$~K.   The astrophysical scenario will change if values of viscosity change or/and different physical mechanisms e.g. magnetic field will be taken into account."
    },
    "0311/astro-ph0311402_arXiv.txt": {
        "abstract": "We present an analysis of a \\Chandra\\ observation of the massive, nearby galaxy cluster Abell 2319.   A sharp surface brightness discontinuity---suggested by previous, lower angular resolution X-ray imaging---is clearly visible in the ACIS image.  This $\\sim$300~kpc  feature suggests that a major merger is taking place with a significant velocity component perpendicular to the line of sight. The cluster emission-weighted mean temperature is $11.8 \\pm 0.6 \\kev$, somewhat higher than previous temperature measurements.  The \\Chandra\\ temperature map of A2319  reveals substructure resembling that anticipated based on hydrodynamic simulations of cluster mergers, and shows an associated cool core not previously known. The map shows a separation between the intracluster medium (ICM) and galaxies of one subcluster, indicating a transient state in which the ICM has been stripped from the subcluster galaxies (and presumably the dark matter). Detailed analysis of the merger feature shows a pressure change across the surface brightness discontinuity by a factor of $\\lesssim 2.5$.  The higher density side of the front has a lower temperature, suggesting the presence of a cold front similar to those in many other merging clusters. The velocity of the front is roughly sonic.  We compare bulk properties of the ICM and galaxies in A2319 to the same properties in a large sample of clusters as a way of gauging the effects of the major merger. Interestingly, by comparing A2319 to a sample of 44 clusters studied with the \\ROSAT\\ PSPC we find that the X-ray luminosity, isophotal size, and ICM mass  are consistent with the expected values for a cluster of its temperature; in addition, the $K$-band galaxy light is consistent with the light--temperature scaling relation derived from a sample of $\\sim$100 clusters studied with 2MASS. Together, these results indicate either that the merger in A2319 has not been effective at altering the bulk properties of the cluster, or that there are large but correlated displacements in luminosity, isophotal size, ICM mass, galaxy light, and emission-weighted mean temperature in this cluster. ",
        "introduction": "\\label{sec:intro}  Galaxy cluster mergers are highly energetic events, driving shocks into the intracluster medium (ICM) of the colliding clusters.   Flattened and asymmetric X-ray morphologies are signatures of recent merging \\citep{mohr93}, and these signatures have been used to study the prevalence of merging in large samples of present-epoch clusters \\citep{mohr95,buote96}. A study of X-ray images of a flux-limited sample of 65 clusters indicates that more than half of nearby clusters show evidence of merging \\citep{mohr95}. Hydrodynamical simulations indicate that complex temperature structures should also be produced in these mergers; however, until relatively recently the required spectral and angular resolution to map this structure has not been available. \\Chandra\\ and {\\it XMM-Newton}, with their high angular resolution, are well-suited for detailed studies of merger features in galaxy clusters \\citep[e.g., ][]{markevitch00,vikhlinin01,markevitch01b, sun02,markevitch02,kempner02,maughan03}. These studies have already revealed that merger features observed in clusters may not indicate shock fronts, but rather ``cold fronts,'' wherein the cool, dense cores of clusters survive through the initial impact of the merger \\citep{markevitch00}. In fact, it now appears that many well-known merger candidates contain these cold fronts, e.g., A2142 \\citep{markevitch00}, A3667 \\citep{vikhlinin01}, A2256 \\citep{sun02}, and A85 \\citep{kempner02}.  Abell 2319 is a massive nearby cluster \\citep[$z=0.0564$;][]{abell58,struble87}.  We chose to study it with the high resolution of \\Chandra\\ because its X-ray morphology observed at lower resolution with the \\ROSAT\\ PSPC shows a strong asymmetry or ``centroid variation,'' which is a classic indicator of a recent merger.  Our goal in this study is not only to better understand the merger state of A2319, but also to determine how the ongoing merger in A2319 is affecting its bulk ICM and galaxy properties. Of particular interest is understanding how merging---which has long been known to be prevalent in the cluster population \\citep{geller82,dressler88,mohr95}---is likely to impact attempts to use cluster surveys to study cosmology \\citep[e.g., ][]{haiman01,randall02,majumdar03a,majumdar03b,hu03b}.  On the basis of galaxy spectra, \\citet{faber77} suggested that A2319 is actually composed of two clusters superimposed along the line of sight, with the smaller subcluster located $\\sim 10\\arcmin$ to the northwest of the main cluster and X-ray surface brightness peak. Additional redshift measurements led to an estimated mean velocity for the main subcluster of $\\sim 100$ members (hereafter A2319A) of 15727 km s$^{-1}$, and for the smaller subcluster of $\\sim 25$ members (hereafter A2319B) of 18636 km s$^{-1}$ \\citep{oegerle95}. This analysis suggests that there is a $\\sim 50\\%$ chance that the two subclusters are in fact not gravitationally bound and will not merge.  A2319 has been extensively studied with previous X-ray instruments, and the inferences about the cluster dynamical state have been quite varied. Emission-weighted mean temperature estimates are generally in the 9--10 keV range \\citep[e.g., ][]{david93, markevitch98, molendi99, irwin00}.   \\citet{markevitch96} produced a temperature map of A2319 using \\ASCA. These observations provided no evidence for a cold core region near the surface brightness discontinuity, though a region to the northwest of the brightness peak appeared to have a temperature  $\\sim 1.5 \\kev$ lower than the mean.  This same subcluster region was identified by \\citet{molendi99} using {\\it BeppoSAX}; it is proposed that this cool region is associated with subcluster A2319B. Using the \\ASCA\\ temperature map, \\citet{markevitch96} argued that there is no evidence of a large-scale merger in A2319. \\citet{mohr95}, however, found a value for the centroid variation of A2319 in the {\\it Einstein} IPC image that indicates an ongoing merger.  Interestingly, a combined X-ray and radio study of the cluster suggests that the two subclusters are in a premerger state \\citep{feretti97}.  This study also takes note of X-ray evidence for another merger in a late stage taking place along the northeast-southwest direction.  \\vskip20pt \\begin{inlinefigure} \\ifthenelse{\\equal{\\figtype}{EPS}}{ \\begin{center} \\epsfxsize=8.cm \\begin{minipage}{\\epsfxsize}\\epsffile{f1.eps}\\end{minipage} \\end{center}} {\\myputfigure{rosat.pdf}{1.0}{1.1}{-0}{-180}} \\figcaption{\\label{fig:rosat} \\ROSAT\\ PSPC image of A2319 with \\Chandra\\ observation footprint overlaid. North is up and east is to the left in all images. The ACIS-I footprint is roughly 17$\\arcmin$ on a side. } \\end{inlinefigure}   In this paper, we present a detailed X-ray study of A2319 based on imaging spectroscopy from \\Chandra\\ ACIS-I, providing clear evidence for an ongoing merger of  two major subclusters. In \\S\\ref{sec:obs} we present the observations. After a description of the data reduction process in \\S\\ref{sec:reduction}, we present an analysis of the overall cluster spectrum (\\S\\ref{sec:spec}) and a temperature map of the cluster (\\S\\ref{sec:temp}). In \\S\\ref{sec:merger} we analyze the merger feature in detail, including quantitative estimates of changes in the physical state of the ICM across the feature, and propose a simple dynamical model. This is followed in \\S\\ref{sec:structure} by a study of how this merger has affected the bulk X--ray properties of the cluster;  we examine how A2319---a cluster in the middle of a major merger---behaves relative to an X-ray flux-limited sample of clusters in its luminosity, isophotal size, and ICM mass. Finally, we summarize our conclusions in \\S\\ref{sec:concl}.  Throughout the paper we assume a $\\Lambda$CDM cosmology with $\\Omega_M=0.3$ and $\\Omega_\\Lambda = 0.7$, and take the Hubble parameter to be $H_0 = 70\\,h_{70}$~km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:concl}  Using \\Chandra\\ data, we have identified and studied a major merger event in A2319 that appears to be taking place along the axis connecting its two major optical subclusters.  The X-ray brightness map shows a clear discontinuity that appears similar to cold fronts found in other clusters.  Although this cold front appears to be as large as the one studied in A3667 \\citep{vikhlinin02}, it is not as sharp. This, together with previous measurements of the line of sight velocity difference between the two main optical subclusters \\citep{oegerle95}, suggests that the merger is not taking place in the plane of the sky.  We propose a merger model where the trajectory lies approximately 65$^\\circ$ out of the plane of the sky, and at this viewing angle it becomes even more challenging to make quantitative statements about the ICM properties near the cold front.  Nevertheless, we estimate that the pressure change across the front is $\\lesssim2.5$, and that the higher density ICM also has the lower temperature.  The estimated merger Mach number of $\\sim 1.1$ is likewise consistent with other merging systems such as A3667.  We propose a two body merger where A2319B merged from the southeast traveling northwest, with the A2319B galaxies and dark matter passing through the A2319A core roughly 0.4~Gyr ago.  The measured emission-weighted mean temperature of this messy, merging cluster is $T_{\\rm X} = 11.8 \\pm 0.6$ keV and the mean abundance is $Z = 0.19 \\pm 0.03$, using a hydrogen column toward the cluster of $N_{\\rm H} = 8.33\\times10^{20}$ cm$^{-2}$. The fit values deviate somewhat from previous studies of A2319. Our higher temperature is likely due in part to the small field of view of \\Chandra\\ compared with other instruments used to study this cluster. Also, we have shown that the emission-weighted mean temperature depends sensitively on the choice of energy band, which is at least partly explained by the highly nonuniform temperature structure we have revealed in this cluster.  Our temperature map shows substructure now considered typical in merging systems. The cool core of A2319A is readily visible, and the angular separation of the galaxies of A2319B and an associated cool ICM region indicates a separation of the galaxies from the ICM of this subcluster, a transient phenomenon that gives further evidence of a merger event. There is some evidence for a hot bridge of ICM between the two cores, a characteristic associated with shock heating in mergers that has been seen in simulations.  We examine how this merger affects the bulk properties of A2319. We naively apply the hydrostatic equilibrium assumption to measure a total mass within \\r2500\\ of $M_{2500} = (4.2\\pm0.5)\\times10^{14}~M_\\odot$, a factor of $1.6\\pm0.3$ lower than the mass predicted by a mass--temperature relation derived from five intermediate-redshift clusters \\citep{allen01}. This offset is the strongest indication that the structure of Abell~2319 has been significantly affected by the merger.  Our measurements for \\L2500--temperature, \\Micm--temperature, and isophotal size--temperature are compared to a sample of 44 clusters observed with \\ROSAT\\ PSPC \\citep{mohr97a,mohr99,mohr00a}. In all three cases our measured values for A2319 are within the scatter of the PSPC-derived scaling relations.  In addition, we note that the $K$-band light in the A2319 galaxy population is consistent with that expected for a cluster of this emission weighted mean temperature \\citep{lin03c}.   It is possible that changes in bulk parameters due to mergers are actually quite large but take place in a correlated way that maintain the low, observed scatter in cluster scaling relations; this has been shown for some properties in numerical simulations \\citep{ricker01,evrard02b}. However, it will require further studies to determine whether it is possible for mergers to create large, correlated displacements in five cluster parameters (i.e., emission-weighted mean temperature, luminosity, isophotal size, ICM mass, and $K$-band galaxy light) that maintain the low scatter in all four scaling relations.  Another possibility is that despite the X-ray imaging spectroscopy and optical evidence for an ongoing merger in A2319, a merger of this scale is simply not sufficient to grossly perturb the bulk properties of the cluster."
    },
    "0311/astro-ph0311152_arXiv.txt": {
        "abstract": "We have accumulated multiwavelength (X-ray, optical, radio) lightcurves for the eight black hole X-ray binaries which have been observed to enter a supposed `soft X-ray transient' outburst, but remained in the low/hard state throughout the outburst. Comparison of the lightcurve morphologies, spectral behaviour, properties of the quasi-periodic oscillations and the radio jet provides the first study of such objects as a sub-class of X-ray transients. However rather than assuming that these hard state X-ray transients are different from the `canonical' soft X-ray transient, we prefer to consider the possibility that new analysis of both soft and hard state X-ray transients in a spectral context will provide a model capable of explaining the outburst mechanisms of (almost) all black hole X-ray binaries. ",
        "introduction": "With the list of known soft X-ray transients (SXT) growing longer every year and now totalling $\\ge30$ it is inevitable that we should find a few sources which do not follow the classical behaviour that we expect. In particular, there are a number of supposed black hole SXTs which do not show a soft component in their X-ray spectra. Instead they display power-law emission across the X-ray energy spectrum, despite some sources reaching high luminosities.  Models for SXT outbursts predict the presence of a thin accretion disc to explain the soft black-body spectrum emitted by the lower energy X-rays (e.g. Esin et al. 1997). The power-law emission at higher energies is thought to be produced via reprocessing in some form of hot plasma (e.g. corona or advection dominated accretion flow), which surrounds the compact object and upscatters disc and/or stellar photons to higher energies via the inverse Compton emission mechanism. This soft spectrum has indeed been observed in the majority of soft X-ray transient outbursts.  \\subsection{The low/hard state}  Some sources, however, show power-law emission down to low X-ray energies and this has been the case for the sources considered in this paper; despite entering outburst there has been no evidence for a soft blackbody component at X-ray wavelengths. In these objects it is assumed that all X-ray emission is produced via reprocessing in the `corona' and not in the accretion disc.  The power-law X-ray energy spectrum is typical of black hole candidates in the low/hard state, most notably Cyg X-1 which spends most of the time in this state (Brocksopp et al. 1999). Corresponding X-ray power spectra for these sources show a high level (relative to the other spectral states) of low-frequency noise, a broken power-law and at least one quasi-periodic oscillation. (See e.g. van der Klis (1995) for further definition of spectral states.) The state is also characterised by the presence of a low-intensity, but extremely powerful, radio jet which emits synchrotron radiation at radio and often higher frequencies (Fender 2001; Corbel \\& Fender 2002). During the high/soft state we see the presence of the soft black-body component in the X-ray spectrum and the quenching of the radio jets, although rapid radio-emitting ejections appear to accompany changes in X-ray spectral hardness (e.g. Brocksopp et al. 2002) -- it is this behaviour that we expect in the classic black hole SXT events.  In this paper we take a closer look at eight black hole `SXTs' which, on the basis of their observed X-ray and radio properties, have been reported to have spent the whole of their outburst in the low/hard state. Although observing coverage was incomplete for the less recent outbursts, each source (other than 4U 1543$-$475) was observed at low X-ray energies for the whole outburst or at least a month after the initial rise, by which time the decay had started and no soft disc component had been discovered; while it is common for SXTs to spend some time in the low/hard state prior to softening, this period is generally only a few days (the longest initial low/hard state period observed to date is that of XTE J1650$-$500 which remained hard for $\\sim 1$ month before softening but also continued to rise in flux; Wijnands priv. comm.). Thus it is unlikely that any of the sources made an unobserved transition to a softer state. The sources are listed in Table~\\ref{sources} along with the date of the low/hard state outburst, the observations which took place and the sources from which we obtained data-points; the observing coverage at the various different frequencies can be seen in the lightcurves plotted in Fig.~\\ref{lightcurve1}. A brief summary of the properties of each source, in chronological order of low/hard state outburst, is given below.  \\subsection{V404 Cyg}  V404 Cyg is probably one of the best-known X-ray transient sources on account of its high luminosity and levels of variability (at various wavelengths) during both outburst and quiescence (e.g. Tanaka \\& Lewin 1995, Hynes et al. 2002) and many observations exist (e.g. Han \\& Hjellming 1992, Shahbaz et al. 1994). Despite the high luminosity of the outburst, the lack of soft component present in the X-ray spectrum suggests that it was indeed a low/hard state transient (e.g. Oosterbroek et al. 1997) -- we note that this has been disputed and that it may have spent a very short period in the soft state (\\.Zycki et al. 1999) but include it here for completeness.  \\subsection{A1524$-$62} A1524$-$62 has been observed in outburst twice$^{1}$, the first of which was a `normal' SXT event (Kaluzienski et al. 1975). The second outburst was observed in the hard X-ray regime by SIGMA (Barret et al. 1992) and by chance was also observed at soft X-rays during the {\\it ROSAT} All Sky Survey ten days before the SIGMA observation. The soft X-ray spectrum could be equally well fit by either a cool black body or a power-law; the presence of an ultra-soft component could not be confirmed or ruled out (Barret et al. 1995). We note, however, that the outburst was insufficient to trigger the soft X-ray all sky monitors on-board {\\it WATCH} and {\\it Ginga}, the latter providing an upper limit of 50 mCrab which was consistent with the ROSAT detection (Barret et al. 1995). Thus we include this source as a possible hard state candidate.  \\subsection{4U 1543$-$475} 4U 1543$-$475 is another recurrent transient which has been observed in outburst four times\\footnote[1]{Due to the poor positional accuracy of the initial observations of A1524$-$62, 4U 1543$-$475 and GS 1354$-$64, it is of course possible that the earlier outbursts of these systems were in fact additional sources, as discussed by e.g. Brocksopp et al. (2001) and Kitamoto et al. (1987). However, given that e.g. XTE J1550$-$564, GX 339$-$4 have indeed been observed recently to exist in a number of spectral states (Corbel et al. 2001, 2003) there is currently no reason to assume that these were necessarily different sources, particularly as they lie in relatively uncrowded regions.} and was a classical SXT during the first two of these, as well as the fourth and current outburst (Matilsky et al. 1972, Kitamoto et al. 1984, Miller \\& Remillard 2002). During the third event hard X-ray observations revealed a power-law spectrum (Harmon et al. 1992) but since there were no soft X-ray observations we cannot confirm the presence or otherwise of a soft component at lower frequencies. We include this source as a possible hard state candidate.  \\subsection{GRO J0422+32} GRO J0422+32 was first observed in 1992 when the X-ray source entered a `canonical' fast rise, exponential decay type outburst (Paciesas et al. 1992). The energy spectrum was a hard power-law (Sunyaev et al. 1994) and timing analysis revealed properties consistent with the low/hard state, although the expected anti-correlation between break frequency and power at the break was not present (van der Hooft et al. 1999). The radio counterpart and the evolution of its spectrum throughout the event are discussed in Shrader et al. (1994). The optical counterpart was discovered by Castro-Tirado et al. (1993) and the intense optical monitoring which followed resulted in the detection of two additional and unusual `mini-outbursts' (e.g. Callanan et al. 1995).  \\subsection{GRO J1719$-$24} GRO J1719$-$24 was detected simultaneously by BATSE and SIGMA and was notable for the stability of the emission at hard X-ray wavelengths (Harmon et al. 1993, Ballet et al. 1993). The low/hard state nature was confirmed through study of the X-ray energy spectra and found to be comparable with that of Cyg X-1 in the low/hard state (Revnivtsev et al. 1998); the power spectra showed the presence of a QPO which varied in frequency (van der Hooft et al. 1996). The radio and optical counterparts were discovered by Della Valle et al. (1994).  \\subsection{GRS 1737$-$31} GRS 1737$-$31 was discovered by SIGMA (Trudolyubov et al. 1999) and was also observed by $RXTE$/PCA, which confirmed the presence of the hard power-law spectrum at lower energies (Cui et al. 1997). Further X-ray observations were made by {\\it BeppoSAX} and {\\it ASCA} (Heise 1997, Ueda et al. 1997) which improved the positional accuracy and confirmed that the hard spectrum was still present (at $\\sim$ 2--20 keV) $\\sim$ 2 weeks after the initial detection; there were no observations at optical or radio wavelengths.  \\subsection{GS 1354$-$64} GS 1354$-$64 is another recurrent transient and it has been observed during four different outbursts$^{1}$ (Brocksopp et al. 2001). Assuming that all events were indeed from the same source, this was first ever SXT outburst observed (then named Cen X-2) and was also one of brightest (e.g. Tanaka \\& Lewin 1995, and references therein). The first and third events were characterised by very soft X-ray spectra; the second and fourth were hard events (Brocksopp et al. 2001 and references therein) -- on account of the observing coverage we consider just the most recent event in this paper, during which a hard power-law energy spectrum was observed throughout by {\\sl RXTE} (Revnivtsev et al. 2000).   \\subsection{XTE J1118+480} XTE J1118+480 was discovered in 2000 during an outburst (Remillard et al. 2000); the presence of the low/hard state was determined from study of the hard power-law X-ray energy spectrum and the characteristic power-law, high rms noise and QPO in the X-ray power spectrum (Revnivtsev et al. 2000a). The optical and radio counterparts were also discovered during the outburst (Uemura et al. 2000, Pooley 2001, Fender et al. 2001). ",
        "conclusions": "We have accumulated multiwavelength datasets for the eight black hole X-ray binaries which have been observed to enter an outburst but remained in the low/hard state throughout. We show that, despite very similar X-ray spectral properties, the lightcurve morphologies of these sources are very different, as are the relationships between the X-ray, optical and radio lightcurves. However, the QPO and jet properties are comparable in each system.  This comparison of the low/hard state transients suggests that some mechanism to explain the outbursts in the context of their spectral properties is necessary; the disc instability model currently lacks the flexibility to do this and needs modification in order to allow for the the observed jet. We also suggest that in light of the more recent observations of both low/hard state and soft X-ray transients (e.g. XTE J1550$-$564), the various models of sporadic mass transfer should not be ruled out, at least during the low/hard state phase of the outburst -- testing of these suggestions is work in progress."
    },
    "0311/astro-ph0311478_arXiv.txt": {
        "abstract": "Recent results of the KASCADE air shower experiment are presented in order to shed some light on the astrophysics of cosmic rays in the region of the knee in the energy spectrum. The results include investigations of high-energy interactions in the atmosphere, the analysis of the arrival directions of cosmic rays, the determination of the mean logarithmic mass, and the unfolding of energy spectra for elemental groups. \\vspace*{-1mm} ",
        "introduction": "The origin of high-energy cosmic rays is among the most interesting questions in astrophysics.  The origin of a structure in the all-particle energy spectrum around 4~PeV, the so-called knee, is generally believed to be a corner stone in the understanding of the astrophysics of high-energy cosmic rays.  The knee is proposed to be caused by the maximum energy reached in cosmic-ray accelerators or due to leakage of particles from the Galaxy. Hence, an understanding of the origin of the knee reveals hints on the acceleration and propagation of cosmic rays.  Experimental access to the understanding of the sources, acceleration and propagation mechanisms is provided by detailed investigation of the arrival directions, energy spectra, and mass composition of the ultrarelativistic particles.  While at energies below 1~PeV cosmic rays can be measured directly at the top of the atmosphere, the strongly decreasing flux as function of energy requires large acceptances and exposure times for higher energies. Presently they can be realized in ground based facilities only. There, the secondary products, generated by interactions of high-energy cosmic-ray particles in the atmosphere, the extensive air showers, are registered. It turns out that a correct description of the high-energy interactions in the atmosphere is crucial for a precise astrophysical interpretation of air shower measurements.  To investigate cosmic rays from several $10^{13}$ eV up to beyond $10^{17}$~eV the air shower experiment KASCADE (\"Karlsruhe Shower Core and Array DEtector\") \\cite{kascade} is operated since 1996. The experiment detects the three main components of air showers simultaneously.  A $200\\times 200$~m$^2$ scintillator array measures the electromagnetic and muonic components.  The 320~m$^2$ central detector system combines a large hadron calorimeter with several muon detection systems.  In addition, high-energy muons are measured by an 128~m$^2$ underground muon tracking detector. ",
        "conclusions": "The present status of the KASCADE experiment to investigate high-energy cosmic rays in the knee region has been discussed.  Continuing efforts are taken in order to improve the understanding of high-energy interactions in the atmosphere.  A search for point sources of charged particles as well as gamma rays did not reveal any significant excess. Upper limits for the large scale anisotropy have been derived.  The observations reveal an increase of the mean logarithmic mass as function of energy in the knee region.  Energy spectra for groups of elements have been derived.  To extend the investigations to higher energies up to $5\\times10^{18}$~eV data taking with an enlarged array \\cite{grande} has started in 2003.   {\\small \\footnotetext[1]{KASCADE Collaboration}"
    },
    "0311/astro-ph0311014_arXiv.txt": {
        "abstract": "I study non-radiative cooling of X-ray emitting gas via heat conduction along magnetic field lines inside magnetic flux loops in cooling flow clusters of galaxies. I find that such heat conduction can reduce the fraction of energy radiated in the X-ray band by a factor of $\\sim 1.5-2$. This non-radiative cooling joins two other proposed non-radiative cooling processes, which can be more efficient. These are mixing of cold and hot gas, and heat conduction initiated by magnetic fields reconnection between hot and cold gas. These processes when incorporated into the moderate cooling flow model lead to a general cooling flow model with the following ingredients. (1) Cooling flow does occur, but with a mass cooling rate $\\sim 10$ times lower than in old ($\\sim 10$ years ago) versions of the cooling flow model. Namely, heating occurs such that the effective age of the cooling flow is much below the cluster age, but the heating can't prevent cooling altogether. (2) The cooling flow region is in a non-steady state evolution. (3) Non-radiative cooling of X-ray emitting gas can bring the model to a much better agreement with observations. (4) The general behavior of the cooling flow gas, and in particular the role played by magnetic fields, make the intracluster medium in cooling flow clusters similar in some aspects to the active solar corona. ",
        "introduction": "\\label{sec:intro}   Recent {\\it Chandra} and {\\it XMM-Newton} observations of cooling flow (CF) clusters of galaxies have lead to a shower of papers examining reheating processes of the intracluster medium (ICM), which together with pre-{\\it Chandra} and XMM-Newton reheating papers make a long list of papers with different ideas (e.g., Binney \\& Tabor 1995; Tucker \\& David 1997; Ciotti \\& Ostriker 2001; David et al.\\ 2001; Quilis, Bower, \\& Balogh 2001; Br\\\"uggen, M. \\& Kaiser 2001; Ruszkowski \\& Begelman 2002; Nulsen et al. 2002; Fabian 2003; Fabian et al.\\ 2003; Binney 2003; Mathews et al.\\ 2003; for more references see Peterson et al.\\ 2003b). These recent observations, although in contradiction (Fabian 2003) with the cluster-CF models of 10 years ago (Fabian 1994),  are compatible with models having low mass cooling rates (see Binney 2003, for why the new results are not a surprise), such as in the moderate CF model (Soker et al.\\ 2001). The moderate CF model, which was proposed before the new results from the two X-ray telescopes, is different from many early proposed processes which aim is to prevent the CF in cluster of galaxies altogether. The main ingredient of the moderate CF model is that the effective age, i.e., the time period since the last major disturbance of the ICM inside the cooling radius $r_c \\sim 100 \\kpc$, is much shorter than the cluster age (e.g., Binney \\& Tabor 1995; see Binney 2003; Soker et al.\\ 2001). The cooling radius is defined as the place where the radiative cooling time equals the cluster age. Originally the heating in the moderate CF model was proposed to be intermittent (Soker et al.\\ 2001), but the basic idea can hold for a steady heating, or heating in short intervals (Binney 2003).  Kahn et al. (2003) and Peterson et al. (2003a), among others, find that there is less gas than that predicted by old versions of the steady-state isobaric CF model at all temperatures below the ambient gas temperature. Also, the discrepancy increases with decreasing temperature. There is not just a deficit in gas below $\\sim 1/3$ of the ambient cluster temperature, there is a deficit of gas at all temperatures below the ambient temperature (Peterson et al.\\ 2003b). Soker \\& David (2003) show that these XMM-Newton results are consistent with the expectations of the moderate CF model. Soker \\& David (2003) claim that the X-ray emitting gas in CF clusters can only be in steady-state if there exists a steady heating mechanism that scales as $H(T) \\propto T^{\\alpha}$ where $\\alpha=1-2$; this is unlikely. They show then that in the moderate CF model most of the gas within CF resides in the hottest gas, which is prevented from cooling continuously and attaining a steady-state configuration by being reheated. This results in a mass cooling rate that decreases with decreasing temperature, with a much lower mass cooling rate at the lowest temperatures, as discovered by Peterson et al. (2003a). Although there is no indication yet for a violent heating, I note the following. An intermittent reheating can work for an AGN kinetic energy of $\\sim 10^{47}~{\\rm erg}~{\\rm s}^{-1}$, and its strong activity should last $\\sim 1-2 \\times 10^7$ yr and occur every $\\sim 10^9$ yr (Soker et al.\\ 2001). A violent heating event by AGN jets was proposed by Kawano, Ohto, \\& Fukazawa (2003) to have occurred $>10^9$ yr ago in the CF clusters 2A 0335 + 096 and A 2199. Mazzotta, Edge, \\& Markevitch (2003) claim that the structure of the cluster 2A 0335+096 can be accounted for if strong AGN activity periods are followed by relative quiescence periods. O'Sullivan \\& Vrtilek (2003) propose that a heating event occured $\\sim 2 \\times 10^8 \\yr$ ago in the poor cluster MKW 4, which consequently is presently not in a steady state cooling. Such a scenario is in accord with the explanation of Soker \\& David (2003) to the observations of Peterson et al.\\ (2003a) and Kahn et al.\\ (2003). Only $\\sim 1-2 \\%$ of all CF clusters should be found during that stage. In addition, Heinz \\& En{\\ss}lin (2003) argue for bias against detecting strong shocks. Therefore, from the well studied CF clusters, only one, or at most two, are expected to be in the stage of violent heating. In Cygnus A there is a strong radio source which heats the ICM (Smith et al. 2002). Wilson, Young, \\& Smith (2003) estimate the mechanical power of the jets in Cygnus A to be $L_{\\rm jet} \\simeq 6 \\times 10^{46} ~{\\rm erg}~{\\rm s}^{-1}$. This is $\\sim 100$ times larger than the radio emission, and it is in the range required by the moderate CF model. Therefore, present observations do not rule-out a violent intermittent reheating process acting for a short time every $\\sim 10^9$ yr.  The limit on the cooling rate below a temperature $T_{\\rm min}$ inferred from X-ray observations is $< 20 \\%$ of the mass cooling rates cited in the past (Fabian 2002). That is, X-ray emission from gas below $T_{\\rm min}$ is absent in present X-ray telescopes. This minimum temperature is different for different clusters, but it is $\\sim 1/3$ of the background temperature (Kaastra et al. 2003; Peterson et al.\\ 2003b). Despite the missing X-ray emission reported by these studies, there are other indications that cooling to temperatures much below $T_{\\rm min}$ does occur, but on a much lower rate than that inferred by taking the CF age to be equal to the cluster age. Wise, McNamara, \\& Murray (2003) and McNamara, Wise, \\& Murray (2003), for example, find the cooling rate within $r \\sim 30 \\kpc$ of the CF cluster A 1068 to be about equal to the star formation rate there $\\sim 20-70 M_\\odot \\yr^{-1}$. Although the limits on the mass cooling rates cited above are within the values expected in the moderate CF model, a much better agreement with observations can be achieved if some non-radiative cooling of the X-ray gas is added to the moderate CF model (Soker, Blanton, \\& Sarazin 2003, hereafter SBS). The thermal energy released in cooling the hot gas is delivered to the cold clouds, and is radiated as optical/UV line emission at $\\sim 10^4-10^5$ K. This energy of the hot gas can be delivered to the cool gas by mixing (e.g., Oegerle et al.\\ 2001; Fabian et al.\\ 2001, 2002; Johnstone et al.\\ 2002; Bayer-Kim et al.\\ 200), or via heat conduction (SBS)\\footnote{This local heat conduction model is completely different from large scale ($\\gtrsim 100$ kpc) heat conduction models aiming at heating, rather than cooling, the X-ray emitting gas in the center of CF clusters.}  In the scenario proposed by SBS the heat conduction is regulated by reconnection between the magnetic field lines in cold ($\\sim 10^4$ K) clouds and the field lines in the ICM. A narrow conduction front is formed, which, despite the relatively low temperature, allows efficient heat conduction from the hot ICM to the cold clouds. The reconnection between the field lines in cold clouds and those in the ICM occurs only when the magnetic field in the ICM is strong enough. This occurs only in the very inner regions of CF clusters, at $r \\sim 10-30$ kpc. While in the scenario of SBS the non-radiative cooling starts with the reconnection of magnetic field lines from the pre-existing cold clouds to the X-ray emitting gas, in the present paper I examine heat conduction within pre-existing magnetic flux loops of X-ray emitting gas. The non-radiative cooling starts when small portion of the loop is radiatively cooling to $\\sim 10^4 \\K$.  Magnetic flux loops in CF clusters were studied before, e.g., in relation to global heat conduction (Norman \\& Meiksin 1996), in the amplification of magnetic fields in the center of CF clusters (Godon, Soker, \\& White 1998), and the formation of multiphase ICM (Zoabi, Soker, \\& Regev 1996, 1998; hereafter ZSR). In a recent paper Kaastra et al.\\ (2003) note the similarity of the differential emission measure in the clusters they study to that of coronal loops in the sun. The behavior of loops found by ZSR, the finding of Kaastra et al.\\ (2003), and the dynamically important magnetic field in the inner regions of cluster CFs (Eilek \\& Owen 2002; Taylor, Fabian, \\& Allen 2002), motivate me to study non-radiative cooling within magnetic flux loops. This paper is not aiming at completely suppressing CF, as there are many evidence for cooling to low temperatures, e.g., Kaastra et al.\\ (2003) find radiation from gas at temperatures of $\\lesssim 10^6  \\K$ in some clusters, Edge et al.\\ (2002; see also Edge \\& Frayer 2003; Salome \\& Combes 2003a,b; Edwards 2003) find massive cold reservoir in the form of molecular gas, probably from cooling ICM, and the star formation rate cited above (Wise et al.\\ 2003). The main goal is to study another aspect, i.e., non radiative cooling of X-ray emitting gas in magnetic flux loops, and incorporate it into the moderate cluster CF model. As such, this paper is complementary to SBS. ",
        "conclusions": "I found above that heat conduction within a magnetic flux loop can reduce the fraction of energy radiated in the X-ray band by a factor of $\\sim 1.5-2$, i.e., the X-ray luminosity is $\\sim 50- 70 \\%$ of the value expected without heat conduction. Cooling of X-ray emitting gas by conduction to cold gas was proposed to occur via mixing (e.g., Oegerle et al.\\ 2001; Fabian et al.\\ 2001, 2002; Johnstone et al.\\ 2002; Bayer-Kim et al.\\ 200) and by local heat conduction regulated by magnetic field reconnection (SBS). In the later two processes the temperature gradient is much higher than the one inside the loop studied here, and hence these mechanism are more efficient in reducing X-ray emission. However, the question is whether these processes occur, and how often. The magnetic flux loop model studied here is motivated by previous suggestions that many such loops exist in the ICM of CF clusters and that they play an active role in cooling and/or dynamics of the gas (e.g., Zoabi et al.\\ 1996; ZSB; Kaastra et al.\\ 2003).  SBS estimate that in their heat conduction scenario via magnetic field reconnection the total line emission expected for gas cooling from $\\sim 10^7$ K to $\\sim 10^6$ K might be factor of $\\sim 30-100$ lower than that expected in older versions of the CF model, where only line emission is considered for the cooling process, and where the cooling rates are much higher. This factor comes from a $\\sim 10$ times lower mass cooling rate of the hot gas in the moderate CF model (Soker et al.\\ 2001), and from another factor of $\\sim 3-10$ for cooling via heat conduction. This much lower expected radiative cooling rate at low X-ray emitting temperatures is compatible with limits inferred from observations. The observational limit, which is $\\lesssim 10-20 \\%$ of the expected cooling rate in old versions of the CF model (Fabian 2002; Peterson et al.\\ 2003a), can be met if the reduction due to heat conduction is only a factor of $\\sim 2$ as found here, with another factor of $\\sim 10$ from the moderate CF model. Possibly, the more significant role of magnetic flux loops is to stir the ICM in the center of CF clusters, by reconnection and motion caused by magnetic tension (Zoabi et al.\\ 1996; ZSR). This stirring, where some portions of loops falls fast and others buoyant outward, can lead to more reconnection between magnetic flux loops (Zoabi et al.\\ 1996), enhancing the efficiency of non-radiative cooling via local heat conduction as discussed by SBS.  If the general scenario discussed here and/or in SBS holds, then a CF model with the following ingredients emerges. \\newline {\\bf (1) CF does occur.} The evidences for CF in many clusters were summarized by Soker et al.\\ (2001) prior to the new results of the {\\it Chandra} and XMM-Newton X-ray telescopes. The new results show that the cooling rate to low temperatures is much lower than the expected values in the old versions of the CF model, but they do show that some cooling occurs (e.g., a few clusters in Kaastra et al.\\ 2003). This is compatible with the moderate CF model. On the other hand, in most CF clusters the new X-ray observations didn't reveal any violent heating events, e.g., major mergers or strong shocks, and the ICM structure is not consistent with global heat conduction from the outer hot massive clusters regions. The X-ray observations do show heating, e.g., bubbles and sound waves (e.g., Fabian et al.\\ 2003), but these can't compensate for the cooling in many cases (e.g., Kawano et al.\\ 2003). Observations using the high spatial resolution of {\\it Chandra} show most activity to occur in the very inner region, on a scale much smaller than the cooling radius (where radiative cooling equals the cluster age), compatible with much lower mass cooling rate. All these suggest that the new X-ray observations actually strengthen the CF model, although a moderate one, rather than ``kill'' it. \\newline {\\bf (2) Non-steady state evolution.} Soker \\& David (2003) show how a non-steady state situation can account for the new observations. Such a non-steady state evolution is expected in intermittent-heating models, with long enough time intervals. The gas is heated up, and then cools. Most of the heated gas does not cool to low temperatures before the next major heating event. Soker \\& David claim that no fine tuning is required. More accurate analysis may show that this non-steady state evolution can hold for heating events with shorter intervals as well. \\newline {\\bf (3) Non-radiative cooling of X-ray emitting gas.} Although not strictly a necessary ingredient, non-radiative cooling of the X-ray emitting gas after it cooled already by a factor of $\\sim 2-5$ from its initial temperature, makes the agreement of the moderate CF with observations much better (SBS; this paper). The initial radiative cooling is required, most likely, to enhance by a factor of $\\gtrsim 10$ the magnetic pressure relative to the thermal pressure (SBS). The magnetic fields play a crucial role in the scenario discussed here, and in the more efficient non-radiative cooling scenario studied by SBS. \\newline {\\bf (4) A corona-like ICM.} The outflow (buoyant bubbles and magnetic flux loops) and inflow, the strong magnetic fields, the possible importance of local heat conduction, radiative cooling in optically thin plasma, the presence of X-ray emitting gas together with optically emitting gas, all these and more, make the ICM in CF clusters similar in some aspects to the gas above the solar (and similar stars) photosphere, i.e., photosphere to corona regions. The similarity with respect to some processes relevant to magnetic field reconnection and optical line emission was mentioned by Jafelice \\& Friaca (1996) and Godon et al.\\ (1998). Sparks \\& Donahue (2003) note the morphological similarity between filaments in M 87 and the solar atmosphere. Godon et al.\\ (1998) discuss magnetic field amplification in CF clusters and the possibility of ``cluster flares'', and use hints from solar flares to project on the expected nature of violent magnetic reconnection in CF clusters. Some aspects of magnetic flux loop dynamics where studied by Zoabi et al.\\ (1996) and ZSR. Kaastra et al.\\ (2003) extend the comparison of ICM in CF cluster with stellar coronae to explain the X-ray spectra they observed in CF clusters."
    },
    "0311/astro-ph0311534_arXiv.txt": {
        "abstract": "{SECIS observations of the June 2001 total solar eclipse were taken using an Fe~{\\sc xiv} 5303 {\\AA} filter. Existing software was modified and new code was developed for the reduction and analysis of these data. The observations, data reduction, study of the atmospheric and instrumental effects, together with some preliminary results are discussed. Emphasis is given to the techniques used for the automated alignment of the 8000 images, the application of the \\`a Trous algorithm for noise filtering and the software developed for the automated detection of intensity oscillations using wavelet analysis. In line with findings from the 1999 SECIS total eclipse observations, intensity oscillations with periods in the range of 20-30~s, both inside and just outside coronal loops are also presented.} ",
        "introduction": "Coronal loops are known to be subject to different types of magnetohydrodynamic (MHD) oscillations. Review papers by Aschwanden et al. [1], Nakariakov [2], Roberts[3] and others provide an overview of the theoretical and observational progress on the detection of such oscillations in solar coronal loops. It is believed that the study of those events may provide us with more information about the physical characteristics of the corona (a subject area called coronal seismology) and help us investigate the feasibility of the coronal oscillations as a corona heating mechanism.  It has been suggested that MHD waves can be divided into two main categories. Magnetoacoustic waves, consisting of density, pressure and temperature perturbations (which in turn are divided into slow and fast modes) and incompressible Alfv\\'en waves (which are also divided into those with movements perpendicular to the magnetic field and the torsional oscillations). Roberts et al.\\ [4] provide a detailed theoretical study, using reasonable approximations for a low-$\\beta$ plasma, of the fast-mode magnetoacoustic perturbations and predict that fast magnetoacoustic (sausage-mode) oscillations with periodicity of $\\sim$1 Hz may be excited in coronal loops. The same authors also predict how the signature of a wave train, created by a pulsation and propagating to a distance {\\it h} from it source, could provide us with information of the physical parameters of the loop. More recently Nakariakov et al. [5] have confirmed the analytical work of Roberts et al.\\ [4] on the propagation of wave trains, by numerically modelling a perturbing pulse and simulating the time series as developed at a distance {\\it h} along the loop.  The coronal heating problem is addressed by a large number of authors with two main mechanisms proposed as the most likely explanation (see review article by Priest \\& Schrijver [6]). One of these is that a large number of magnetic reconnections causing current dissipation could result in micro- or nano-flare activity on a regular bases (see Parker [7] more further discussion). The other suggestion is the damping of MHD waves, caused by ion viscosity and electrical resistivity (first introduced by Hollweg [8]).  Following a long sequence of attempts to detect coronal oscillations using total solar eclipse observations (Koutchmy et al.\\ [9], Pasachoff \\& Landman [10], Singh et al.\\ [11], Cowsik et al.\\ [12], Pasachoff et al.\\ [13]), the Solar Eclipse Coronal Imaging System (SECIS) observed the 1999 total solar eclipse using an Fe~{\\sc xiv} 5303 {\\AA} filter. Phillips et al.\\ [14] developed SECIS and used the instrument to observe the total solar eclipses of 1999 and 2001. They succeeded in detecting several periodicities in the range of 4-7 s (Williams et al.\\ [15], Williams et al.\\ [16], Katsiyannis et al.\\ [17]) in the 1999 data set. A propagating wave train was detected by [16] with a phase-speed of $\\sim$2100~km~s$^{-1}$, reinforcing the identification of this perturbation as a fast-mode MHD wave. Considering the physical parameters of the loop (physical dimensions, density, etc) and making reasonable assumptions about the strength of the magnetic field, it was found that the frequencies detected are well within the range of predicted values in [4]. Work on different loops in the same active region [17] failed to detect a wave train but found intensity oscillations just outside these loops. These perturbations were in the same frequency range as those detected by [16] and with similar amplitudes.  In this paper we discuss observations taken during the June 2001 total solar eclipse, the reduction and data analysis techniques developed to process these data and some preliminary detections of oscillations in the lower solar coronal. ",
        "conclusions": ""
    },
    "0311/astro-ph0311228_arXiv.txt": {
        "abstract": "% %  Low energy X-ray emission (0.1-10 keV) from all six millisecond radio pulsars (MSPs) for which such emission has been reported support a proposed pulsar magnetic field evolution previously compared only to radiopulse data: old, very strongly spun-up neutron stars become mainly orthogonal rotators (magnetic dipole moment perpendicular to stellar spin) or aligned rotators. The neutron star properties which lead to such evolution are reviewed.  Special consideration is given to agreement between predictions and observed X-ray emission for the aligned MSP candidate PSR J0437-4715. ",
        "introduction": "\\def\\lesssim{\\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$<$}}}} \\def\\gtrsim{\\mathrel{\\hbox{\\rlap{\\hbox{\\lower4pt\\hbox{$\\sim$}}}\\hbox{$>$}}}}  Properties of the radioemission from the most rapidly spinning disk population millisecond radio pulsars (MSPs) support the neutron star model-based expectation that these pulsars should be mainly orthogonal rotators or aligned rotators. There are three kinds of observational evidence for this.  1) A large fraction of the MSPs have two radio-subpulses of comparable strength separated in time by about half a period ( P/2) , a signature property of an orthogonal rotator (Chen \\& Ruderman  1993, Jayawardhana \\& Grindlay 1996). In a sample of 8 of the fastest spinning disk-MSPs at least 4, and possibly 6, are in this category ( Chen et al 1998); among canonical radiopulsars only a few percent are ( Lyne \\& Manchester 1988).  2) The well studied original MSP, PSR 1937, is observed to have very different linear polarization properties in two narrow radio-subpulses an interval P/2 apart, just as expected from the spin-up evolution of that pulsar into an orthogonal rotator whose dipole moment is on the neutron star's (NS's) spin-axis at the NS crust-core interface (CR 93).  3) At least one pulsar in the chosen MSP octet, PSR J0437-4715, and possibly several, have only a single very broad, structured radiopulse.  Its pulse-width is so large ($\\sim$270$^{\\circ}$ =3P/4) that a plausible description of the observed radiation beam seems to require an almost aligned beam (and dipole moment) observed from a line-of-sight direction  near  that of the NS's spin-axis (Gil \\& Krawczyk 97).  Because an almost spin-aligned radiobeam from a spinning NS sweeps out a smaller solid angle on the sky then that same beam would from an orthogonal rotator, one observation of an aligned MSP out of the 8 in the sample indicates a very considerably larger aligned fraction.  In Section 2 we consider the significance of MSP low energy X-ray light-curves when these data are combined with those from the same MSPs' radiopulses. Section 3 summarizes relevant NS model predictions fro stellar magnetic field evolution as these stars were spun-up to become MSPs.  Section 4 presents model predictions for an aligned MSP's thermal X-ray emission areas and compares them to those inferred from PSR J0437 observations. In no case is there a conflict between what is expected from  special properties of NSs and what is observed. ",
        "conclusions": ""
    },
    "0311/astro-ph0311191_arXiv.txt": {
        "abstract": "The debate over whether kilometer-sized solids, or planetesimals, assemble by collision-induced chemical sticking or by gravity-driven unstable modes remains unsettled.   In light of recent work showing that gravitational growth can occur despite turbulent stirring, we critically evaluate the collisional hypothesis.  Impact speeds in protoplanetary disks reach $\\sim 50$ m/s in a laminar disk and may be larger in turbulent disks.  We consider the role of elastic and plastic deformations, restructuring of ``rubble piles,\" and sticky organic matter.  Coagulation is possible for dust grains with large surface area-to-volume ratios.  Simple energetic arguments show that bouncing, cratering, and fragmentation should dominate collisional dynamics between millimeter and kilometer sizes. ",
        "introduction": "Solid planets and gas giant cores form in a hierarchical growth process spanning some 13 orders of magnitude in radius from interstellar dust grains to spectacular giant impacts.  An important demarcation occurs at kilometer sizes, where coupling to the gas disk is weak and substantial escape speeds of $\\approx 1$ m/s allow growth by inelastic binary collisions.  Two theories aim to explain the origin of planetesimals from smaller solids with negligible surface gravity and strong gas couplings. This paper examines what we will call the Collisional-Chemical Sticking Hypothesis (hereafter CCSH), that km-sized bodies arise from a long chain of agglomerative growth, due to chemical forces during inelastic collisions. The gravitational instability (hereafter GI) hypothesis appeals the collective self-gravity of a dense midplane layer to collect small solids into larger planetesimals (Goldreich \\& Ward 1973, Youdin \\& Shu 2002, hereafter YS).  The characteristic wavelength of GI predicts planetesimal sizes $\\ga$ few km for disks with enough mass to produce a solar system like our own.  Leapfrogging intermediate sizes by direct assembly from $\\la$ cm solids avoids most problems associated with sticking and rapid radial drift (see \\S\\ref{laminar}).  The main obstacle to GI is that turbulence opposes the settling required to create a dense midplane layer.  Even an otherwise laminar disk  generates midplane turbulence, since particle settling introduces destabilizing vertical shear (Weidenschilling \\& Cuzzi 1993, hereafter WC).  Sekiya (1998) overcame this obstacle by appealing to the enhanced buoyancy of stratified, high metallicity disks. YS showed that turbulent motions can only suspend a finite surface density of solids, $\\Sigma_{\\rm p,crit}$.  GI occurs in a dense midplane layer when the actual $\\Sigma_{\\rm p}$ exceeds this saturation threshold.  However the required surface density ratio of solids to gas, $\\Sigma_{\\rm p,crit} / \\Sigma_{\\rm g}$, i.e.\\  the disk metallicity, is about an order of magnitude over solar abundances. Various enrichment mechanisms were described in YS.  A particularly robust finding is that aerodynamic radial drift has the global consequence of enhancing $\\Sigma_{\\rm p}$ as the particle disk shrinks in radius.  Youdin \\& Chiang (2003) explored these ``particle pileups\" in detail,  including inertial effects and self-consistently generated turbulent stresses in the particle layer.  Since GI can overcome the heretofore fundamental obstacle of vertical shear, we assess the problems facing the CCSH.  YS argue against the plausibility of the CCSH with evidence based on micro-gravity experiments, the size distribution of meteoritic inclusions, the limited sticking ability of icy planetary ring debris, and ``intuition.\"  Our goal here is to physically explain the ``intuitive\" notion that collisions between macroscopic solids with weak surface gravity, at speeds characteristic of freeway driving, are not cohesive.  The most fundamental obstacle to sticking is energetic: the \\emph{entire} relative kinetic energy of two bodies at large distances, $K_\\infty = m v_\\infty^2/2$, must be dissipated, where $m$ is the reduced mass and $v_\\infty$ is the relative speed at large distances.  This requires some combination of (a) efficient energy loss, and (b) binding energies, $U_{\\rm bind} < 0$, which increase kinetic energy on impact: $K_{\\rm imp} = K_\\infty + |U_{\\rm bind}|$.  The dissipative efficiency needed for sticking, $K_\\infty/K_{\\rm imp} = 1/(1+ \\theta)$, drops appreciably below unity when the focusing factor $\\theta \\equiv |U_{\\rm bind}|/K_\\infty \\ga 1$.  Gravitational focusing, which scales as $\\theta \\propto R^2/v_\\infty^2$, affects bodies with sufficiently large radius, $R$.  For small bodies, surface energies, such as van der Waals interactions, give $\\theta \\propto R^w/v^2$, where the size falloff ranges from $w = -1$, for a maximal contact area $\\sim R^2$,  to $w = -5/3$, for elastic compression (see \\S\\ref{elast}).  Over a large size range, extending from mm ---  10~m, no relevant binding energy exists.  Until an incredibly efficient, yet non-destructive, dissipation mechanism is identified, the CCSH will lack a solid physical foundation.  This proceeding is organized as follows: \\S\\ref{speed} investigates collisional speeds expected from laminar flow (\\S\\ref{laminar}) and turbulent fluctuations (\\S\\ref{turb}). Section \\ref{cm} explores collisional binding mechanisms including elastic (\\S\\ref{elast}) and plastic (\\S\\ref{plastic}) deformation as well as restructuring within ``rubble piles\" and  the role of  organic matter (\\S\\ref{agg}).  Section \\ref{concl} contains concluding remarks. ",
        "conclusions": "We studied the implications of particle velocity dispersions in protoplanetary disks  for the collisional growth of planetesimals.  Over a wide range of sizes, mm --- 10 m, sticking is strongly inhibited.  Available binding energies are negligible compared to kinetic energies, and a suitably efficient, yet non-destructive, dissipation mechanism has not been identified.  The most elegant solution to this problem is the Goldreich-Ward hypothesis, which uses collective gravity to avoid this problematic coagulation regime.  Particle concentration mechanisms, such as vortex trapping, secular instabilities, and quiescent nodes in turbulent flow (see YS for details) may play a role in planetesimal formation.  Higher densities favor gravitational instability, and could also promote collisional growth if collisional speeds are significantly reduced.  Regarding the possibility of long-lived particle-trapping vortices, only marginally coupled particles ($\\tau_{\\rm stop} \\approx 1$) are efficiently trapped.  Smaller particles remain attached to streamlines and larger bodies can plow through the vortex.  This limits the overall reduction in relative velocities.  For many years, supporters of the Goldreich-Ward hypothesis were on the defensive due to valid concerns about the stirring provided by turbulence. Proposed solutions to this problem (Sekiya 1998, YS) should reestablish the viability of particulate gravitational instabilities.  To make the debate more productive, coagulation proponents should similarly acknowledge, and explain how  to overcome, the considerable energetic obstacles to sticking."
    },
    "0311/astro-ph0311158_arXiv.txt": {
        "abstract": "We present far-infrared maps and spectroscopy for a number of deeply embedded protostellar objects (Cep\\,E, HH\\,211-MM, IC\\,1396\\,W, L\\,1157, L\\,1211, and RNO\\,15\\,FIR) from data that we acquired with the ISO instruments PHOT and LWS. Several previously undetected deeply embedded sources are found in the vicinity of our  targets. We determine temperatures and luminosities of seven objects and locate them on a L$_{bol}$-T$_{bol}$ diagram -- the equivalent to a Hertzsprung-Russell diagram for protostars. Their masses and ages, according to their location on tracks taken from our evolutionary model, are derived. L\\,1211 and Cep\\,E appear to be intermediate mass objects which will reach final masses of about 3\\,M$_\\odot$, while the other sources are in or below the solar mass range. The derived ages of 15000 to 30000\\,yr are consistent with their current Class\\,0 state. A comparison of the luminosity of the associated outflows in the 1\\,--\\,0\\,S(1) line of molecular hydrogen with the source properties (bolometric luminosity, bolometric temperature, and envelope mass) of 16 Class\\,0 sources shows no statistically significant correlations. Nevertheless, the data are consistent with a scheme in which the outflow strength and protostar evolve simultaneously. We show that the relationship is partially disguised, however, by the local properties of the surrounding material, the extinction, and short-term flux variability. ",
        "introduction": "In the earliest stages of star formation -- the so-called Class\\,0 and Class\\,1 phases -- protostars are still deeply embedded in their parental molecular cloud cores. This material absorbs almost all of the emitted radiation of the star in the optical and the near-infrared. The spectral energy distribution (SED) of Class\\,0 protostars peaks at about 100--160\\,$\\mu$m, the wavelength of the maximum of a modified blackbody at 30--80\\,K. Hence, direct observations of protostars have to be carried out in the far-infrared and in the (sub-)mm wavelength range. Sub-mm and millimeter observations of some of the sources investigated here have been obtained e.g. by Lefloch et al. \\cite{lel96}, Ladd \\& Hodapp \\cite{lh97}, Chini et al. \\cite{cwknrs01}, Gueth et al. \\cite{ggdb97}, Motte \\& Andr\\'e \\cite{ma01}, and Gueth \\& Guilloteau \\cite{gg99}.  The ISO satellite (Kessler et al. \\cite{ksa_etal96}) with its PHOT instrument had the capacity to measure the broad-band continuum in the far-infrared. Such observations, covering the peak region of the SED of protostars, help to yield some of the major properties of these objects such as their temperature, the sub-mm slope of their SED, the optical depth and the solid angle under which they emit. The latter two cannot be disentangled due to the limited spatial resolution of the ISOPHOT instrument. With higher resolution observations (e.g. SCUBA), however, we can independently determine the solid angle under which an object is seen and that way infer its optical depth. These parameters, together with the distance, enable us to calculate the total (L$_{bol}$) and sub-mm (L$_{smm}$) luminosities of each object. We may then decide whether an object really is of Class\\,0 or not by determining the L$_{smm}$/L$_{bol}$ ratio (Andr\\'e et al. \\cite{awb00}). Finally, by placing the inferred values on a temperature -- bolometric luminosity diagram -- the equivalent to a Hertzsprung-Russell diagram for protostars (Myers et al. \\cite{macs98}) -- we are able for the first time to estimate the (model dependent) ages and masses of these sources directly.  Bipolar outflows invariably accompany Class\\,0 sources: strong inflow and outflow of material are concurrent.  We thus wish to probe how the mass outflow rate  is related to the mass accretion rate onto the protostar. The outflowing material interacts with the ambient medium through  radiative shocks. Thus, the luminosity of the outflow may be correlated with some of the source properties (e.g. the bolometric source luminosity), which depend on the mass accretion rate. Therefore, we measured the luminosities of the outflows of 16 Class\\,0 sources in the 1\\,--\\,0\\,S(1) line of molecular hydrogen. This is usually the strongest  and easiest line to observe in near-infrared spectra of shocked molecular hydrogen, and due to the short cooling time of H$_2$ it is a good tracer of the present interaction of the outflow with the surrounding material. These H$_2$ luminosities are then compared with various source properties to investigate possible correlations.  Modelling of Class\\,0 protostars remains in its infancy. Schemes now exist which yield evolutionary tracks, based on relating gas accretion to the dusty envelope (Myers et al.\\cite{macs98}) and jet thrust to gas accretion (Bontemps et al. \\cite{batc96}, Saraceno et al. \\cite{saraceno96}, Smith \\cite{s98,s00}, and Andr\\'e et al. \\cite{awb00}). We combine these schemes here in order to test if the simplest assumptions, such as a spherical envelope and a single accreting object, are feasible.  In this paper, we first present our far-infrared ISO maps and spectroscopy, and then summarize the data analysis and how we derive temperatures and luminosities (Sect.\\,\\ref{datareduction}). In Sect.\\,\\ref{results}, we present our results, and comment on individual objects. A discussion of age and mass determination, and the general relationship to the outflows is contained in Sect.\\,\\ref{discussion}. A framework within which the data can be interpreted is then put forward (Sect.\\,\\ref{evolscheme}). ",
        "conclusions": "We have observed the spectral energy distributions for the six deeply embedded objects Cep\\,E, L\\,1211, IC\\,1396\\,W, L\\,1157, HH\\,211-MM, and RNO\\,15\\,FIR in the far-infrared with ISO. The inferred temperatures and L$_{smm}$/L$_{bol}$ ratios confirm the Class\\,0 nature of four of these sources, within the errors. Employing an evolutionary scheme, we are able to estimate the age, surrounding mass and the current and final mass of these sources. These estimates are, however, model dependent. Two sources,  Cep\\,E and L\\,1211 appear to develop into intermediate mass stars, while the others will become solar mass stars or lower mass objects. The comparison of the ISOPHOT and LWS observations for three of the sources reveals the existense of emission from cold dust in the immediate vicinity of the objects.  A comparison of the luminosity in the 1\\,--\\,0\\,S(1) line of H$_2$ of the related outflows for 16 Class\\,0 sources, with the source bolometric luminosity, bolometric temperature, and envelope mass was done. We found no statistically significant correlation of the outflow luminosity with each of these source parameters. This could be due to the H$_2$ luminosity mainly depending on the local properties of the surrounding gas.  The unifying scheme, however, explains the lack of correlations as due to evolutionary effects. Furthermore, the scheme which involves a redistribution of mass between envelope, disk, protostar, jets and outflow, accounts for the differences in source properties according to the Class."
    },
    "0311/astro-ph0311408_arXiv.txt": {
        "abstract": "We have made differential radial velocity measurements of the star $\\alpha$~Cen~A using two spectrographs, UVES and UCLES, both with iodine absorption cells for wavelength referencing.  Stellar oscillations are clearly visible in the time series.  After removing jumps and slow trends in the data, we show that the precision of the velocity measurements per minute of observing time is 0.42\\,\\ms{} for UVES and 1.0\\,\\ms{} for UCLES, while the noise level in the Fourier spectrum of the combined data is 1.9\\,\\cms.  As such, these observations represent the most precise velocities ever measured on any star apart from the Sun. ",
        "introduction": "The search for extra-solar planets has driven tremendous improvements in high-precision measurements of stellar differential radial velocities \\citep[e.g.,][]{M+B2000}.  In parallel, the same techniques have been used with great success to measure stellar oscillations.  Recent achievements, reviewed by \\citet{B+C2003} and \\citet{B+K2003}, include observations of oscillations in Procyon with ELODIE \\citep[][see also \\citealp{BMM99}]{MSL99}; \\bhyi{} with UCLES \\citep{BBK2001} and CORALIE \\citep{CBK2001}; and \\acena{} and~B with CORALIE \\citep{B+C2001,B+C2002,C+B2003}.  Here we report observations of \\acena{} made with UVES and UCLES which represent the most precise differential radial velocities ever measured on any star apart from the Sun. ",
        "conclusions": "We have analysed differential radial velocity measurements of \\acena{} made with UVES at the VLT and UCLES at the AAT\\@.  Stellar oscillations are clearly visible in the time series.  Slow drifts and sudden jumps of a few metres per second, presumably instrumental, were removed from each time series using a high-pass filter.  We then used the measurement uncertainties as weights in calculating the power spectrum, but we found it necessary to modify some of the weights to account for a small fraction of bad data points.  In the end, we reached a noise floor of 1.9\\,\\cms{} in the amplitude spectrum and in a future paper we will present a full analysis of the oscillation frequencies and a comparison with stellar models."
    },
    "0311/astro-ph0311314_arXiv.txt": {
        "abstract": "This article investigates the signature of the seventeen multi-connected flat spaces in cosmic microwave background (CMB) maps. For each such space it recalls a fundamental domain and a set of generating matrices, and then goes on to find an orthonormal basis for the set of eigenmodes of the Laplace operator on that space.  The basis eigenmodes are expressed as linear combinations of eigenmodes of the simply connected Euclidean space.  A preceding work, which provides a general method for implementing multi-connected topologies in standard CMB codes, is then applied to simulate CMB maps and angular power spectra for each space. Unlike in the 3-torus, the results in most multi-connected flat spaces depend on the location of the observer. This effect is discussed in detail. In particular, it is shown that the correlated circles on a CMB map are generically not back-to-back, so that negative search of back-to-back circles in the WMAP data does not exclude a vast majority of flat or nearly flat topologies. ",
        "introduction": "Among all multi-connected three-dimensional spaces, ``flat spaces''\\footnote{In this article, we follow the cosmological use and we call ``flat spaces'' the eighteen types of Euclidean spaces, and ``Euclidean space'' the simply connected universal cover $\\ETR$.} have been studied the most extensively in the cosmological context. This is due to the computational simplicity of the simplest compact flat three-manifold, the 3-torus, which has been used extensively in numerical simulations. The main goal of this article is to provide tools to compute the CMB properties and produce high resolution CMB maps for all seventeen multi-connected flat spaces~\\footnote{Test maps for these spaces are available on demand.}, following the general method introduced in our preceding work~\\cite{rulw02}.  Recent measurements show that the density parameter $\\Omega_0$ is close to unity and the observable universe is approximately flat.  CMB data obtained by the Archeops balloon experiments~\\cite{archeops} and more recently by WMAP~\\cite{wmap} place strong constraints on the curvature. In addition, WMAP~\\cite{map} and later the Planck satellite~\\cite{planck} do and will provide full sky maps of CMB anisotropies, offering an opportunity to probe the topological properties of our universe. This observational constraint on the curvature radius of the universe motivates the detailed study of flat spaces even though spherical spaces are also promising candidates~\\cite{wlu02,glluw,lwugl,luw02,prl,nat}.  At present, the status of the constraints on the topology of flat spaces is evolving rapidly driven by the release of the WMAP data. Previous analysis, based on the COBE data, mainly constrained the topology of a 3-torus (see Refs.~\\cite{sokolov93,staro93,stevens93,costa95,levin98,levin99,pf,inoue00,inoue01,inoue2,roukema} and Refs.~\\cite{lachieze95,uzan99,levin02} for reviews of different methods for searching for the topology).  The WMAP data~\\cite{wmap} possess some anomalies on large angular scales that may be explained by a topological structure. In particular, the quadrupole is abnormally low, the octopole is very planar and the alignment between the quadrupole and octopole is also anomalous~\\cite{teg1}. Besides many other potential explanations~\\cite{quad}, it was suggested that a toroidal universe with a smaller dimension on the order of half the horizon scale may explain all these anomalies~\\cite{teg1} but it was latter shown, on the basis of a finer statistical analysis, that it did not~\\cite{teg2}. Another topology was recently proposed to explain some of this anomaly in the case of a slightly positively curved space, namely the Poincar\\'e dodecahedral space~\\cite{nat}.  The first results of the search for the topology through the ``circles in the sky'' method~\\cite{cornish98} gave negative results for back-to-back or almost back-to-back circles~\\cite{teg2,cssk}. While the first applies only to back-to-back circle with no twist, the second study includes an arbitrary twist and conclude that ``it rules out the possibility that we live in a universe with topology smaller than $24 \\UUNIT{Gpc}{}$''. As will be discussed in this paper, back-to-back circles are generic only for homogeneous topologies such as e.g. 3-tori and a subclass of lens spaces. In non-homogeneous spaces the relative position of the circles depends on the position of the observer in the fundamental polyhedron.  In conclusion, as demonstrated by these preliminary results, only the toroidal spaces have been really constrained~\\cite{teg2,cssk}. Besides, a series of studies have pointed out a departure of the WMAP data from statistical isotropy. Copi {\\em et al.}~\\cite{copi} recently argued in particular that they are inconsistent with an isotropic Gaussian distribution at 98.8\\% confidence level. Previous studies pointed toward a possible North-South asymmetry of the data~\\cite{eriksen,park}. Spaces with non-trivial topology are a class of models in which global isotropy (and possibly global homogeneity) is broken. Simulated CMB maps of these spaces may help to construct estimators for quantifying the departure of the temperature distribution from isotropy, and also give a deeper understanding of recent results.  Let us emphasize that in the case where the topological scale is slightly larger than the size of the observable universe, no matching circles will be observed. This might also happen for a configuration where the circles would all lie in the direction of the galactic disk where the signal-to-noise ratio might be too low. Contrary to the simply connected case, the correlation matrix, $C_{\\ell m}^{\\ell' m'} \\equiv \\left< a_{\\ell m} a_{\\ell' m'} \\right>$, of the coefficients of the development of the temperature fluctuations on spherical harmonics, will not be proportional to $\\delta_{\\ell \\ell'} \\delta_{m m'}$. The study of this correlation matrix could offer the possibility to probe topology (slightly) beyond the horizon. Computing the correlation matrix $C_{\\ell m}^{\\ell' m'}$ for different multi-connected spaces will help design the best strategy to constrain the deviation from the simply connected case, and gives a concrete example of cosmological models in which the global homogeneity and isotropy are broken.  As described in detail in our preceding work~\\cite{rulw02}, what is needed for any CMB computation are the eigenmodes of the Laplacian \\begin{equation} \\label{Helmotz1} \\Delta \\UPSTKS{\\Gamma}{k}{} = - k^2 \\UPSTKS{\\Gamma}{k}{} , \\end{equation} with boundary conditions compatible with the given topology.  These eigenmodes can be developed on the basis ${\\cal Y}_{k \\ell m}$ of the (spherically symmetric) eigenmodes of the universal covering space as \\begin{equation} \\label{eq:1} \\UPSTKS{\\Gamma}{k}{s} = \\sum_{\\ell = 0}^\\infty \\sum_{m = -\\ell}^{\\ell} \\XITKSLM{\\Gamma}{k}{s}{\\ell}{m} {\\cal Y}_{k \\ell m} , \\end{equation} so that all the topological information is encoded in the coefficients $\\XITKSLM{\\Gamma}{k}{s}{\\ell}{m}$, where $s$ labels the various eigenmodes sharing the same eigenvalue $- k^2$.  Ref.~\\cite{rulw02} computes these coefficients for the torus and lens spaces and Refs.~\\cite{glluw,cornish99} discuss more general cases.  To summarize, this article aims at several goals. First, it will give the complete classification of flat spaces and the exact form of the eigenmodes of the Laplacian for each of them. It will also provide a set of simulated CMB maps for most of these spaces. Among other effects, it will illustrate the effect of non-compact directions and discuss the influence of the position of the observer in the case of non-homogeneous spaces, which has never been discussed before.  It also explains the structure of the observed CMB spectrum in the case of a very anisotropic (i.e., flattened or elongated in one direction) fundamental domain.  This article is organized as follows. We start by recalling the properties of the eighteen flat spaces (Section~\\ref{sec02}) as well as the eigenmodes of the simply connected three-dimensional Euclidean space $\\ETR$ (Section~\\ref{sec2}), and in particular how to convert planar waves, which suit the description of topology, to spherical waves, which are more convenient for CMB computation. Then, in Section~\\ref{SectionEigenmodesOfMulticonnected}, we explain how to extract the modes of a given multi-connected space from the modes of $\\ETR$. This method is then applied to give the eigenmodes of the ten compact flat spaces (Section~\\ref{SectionCompactSpaces}), the five multi-connected flat spaces with two compact directions (``chimney spaces'', Section~\\ref{SectionDoublyPeriodic}) and the two multi-connected flat spaces with only one compact direction (``slab spaces'', Section~\\ref{SectionSinglyPeriodic}). Applying the general formalism developed in our previous work~\\cite{rulw02}, we produce CMB maps for some of these spaces.  With three exceptions the manifolds are not homogeneous, in the sense that a given manifold does not look the same from all points. To discuss the implication of the observed CMB and the genericity of the maps, we detail in Section~\\ref{SectionLocationOfObserver} the influence of the position of the observer on the form of the eigenmodes and we study its consequences on the observed CMB maps. We show in particular that the matched circles are generically not back-to-back, but their relative position depends on the topology, the precise shape of the fundamental domain, and the position of the observer.  \\section*{Notation}  The local geometry of the universe is described by a locally Euclidean Friedmann-Lema\\^{\\i}tre metric \\begin{equation} \\label{fl_metric} \\dd s^2 = - c^2 \\dd t^2 + a^2 (t) \\left[\\dd\\CHU^2 + \\CHU^2\\dd \\omega^2 \\right] , \\end{equation} where $a(t)$ is the scale factor, $t$ the cosmic time, $\\dd \\omega^2 \\equiv \\dd \\theta^2 + \\sin^2 \\theta \\, \\dd \\varphi^2$ the infinitesimal solid angle. ",
        "conclusions": "This article has presented the tools required to compute CMB maps for all multi-connected flat spaces. We gave for each space \\begin{itemize} \\item the polyhedron and holonomy group, \\item the eigenmodes of the Laplacian. \\end{itemize} We then presented simulated maps for all of the nine compact non homogeneous spaces. On the basis of the angular power spectra we compared the effect of different topologies and different configurations for a given topology. We also implemented the effect of an arbitrary position of the observer which yields significant effects for non-homogeneous spaces. We investigated this effect both on simulated maps and angular power spectra. In particular, it shows that generically matched circles are not back-to-back and that their relative position depends on the position of the observer.  All these tools and simulations will be of great help for extending the conclusions reached on the torus and to investigate their genericity as well as for providing test maps for any method wishing to detect (an interpret) the breakdown of global isotropy."
    },
    "0311/astro-ph0311587_arXiv.txt": {
        "abstract": "{In this article we present the results from mid-infrared spectroscopy of a sample of 14 T\\,Tauri stars with silicate emission. The qualitative analysis of the spectra reveals a correlation between the strength of the silicate feature and its shape similar to the one which was found recently for the more massive Herbig Ae/Be stars by van Boekel~et~al. (2003). The comparison with theoretical spectra of amorphous olivine ([Mg,Fe]$_2$SiO$_4$) with different grain sizes suggests that this correlation is indicating grain growth in the disks of T\\,Tauri stars. Similar mechanisms of grain processing appear to be effective in both groups of young stars. ",
        "introduction": "T\\,Tauri stars are low mass young stellar objects with ages of $10^6-10^7$\\,yr. One of the characteristics of so-called classical T\\,Tauri stars is a distinctive IR excess. The origin of the IR radiation is the accretion disk which plays a crucial role in star formation by transporting material to the central star while angular momentum is transferred outwards. Different models were proposed to bring the observed facts into a general framework (\\cite{pringle74}; \\cite{adams87}; \\cite{chiang97}). Yet there remain many open questions concerning the details of configuration, chemistry and evolution of disks. A particular feature of the infrared spectra of such disks is the silicate band in the region from 8 to 12\\,$\\mu$m.  The band originates from the stretching mode of the Si-O bond of silicate minerals like olivine ([Mg,Fe]$_2$SiO$_4$), forsterite (Mg$_2$SiO$_4$), enstatite (MgSiO$_3$) and silica (SiO$_2$).  The silicate feature observed in T\\,Tauri stars appears in emission as well as in absorption (\\cite{cohen85}). Silicate emission is assumed to emerge from a warm, optically thin disk layer (disk atmosphere) which is heated by the radiation of the central star (\\cite{chiang97}, \\cite{natta00}). Honda~et~al. (2003) showed that crystalline silicates are present in T Tauri stars, indicating substantial grain processing. In this Paper we present studies on the silicate emission from a sample of 14 T\\,Tauri stars based on mid-infrared spectroscopy performed during three observation campaigns in February, June and December 2002 on the ESO 3.6\\,m telescope at La Silla. We analyse the spectra in particular with respect to a possible correlation between shape and strength of the silicate feature. Such a correlation has been recently observed for Herbig Ae/Be stars (\\cite{boekel03}) and has been interpreted as evidence of grain processing in the circumstellar disks of those stars. ",
        "conclusions": "The large difference in geometric cross section per unit mass between large and small dust grains implies that if the silicate emission band is dominated by small grains, a substantial amount of large grains could be present without being noticed.  On the other hand, the spectra with plateau-like emission are in geneal weak, likely due to a much lower number of small grains present in the upper disk regions.  In some spectra (S\\,CrA both components, RU\\,Lup) the flux at 11.3\\,$\\mu$m is even higher than the flux at 9.8\\,$\\mu$m, which results to a ratio greater than 1.  This cannot be explained by the assumption of amorphous olivine dust with the absorption spectra shown above.  It may be that crystalline forsterite (Mg$_2$SiO$_4$) is responsible for this effect, since its absorption spectrum shows a peak at 11.3\\,$\\mu$m (see \\cite{bouwman01}, \\cite{honda03}). Possible indications for crystalline forsterite are also seen in some other objects with stronger features (SU\\,Aur, GG\\,Tau).  We also note that the silicate emission emerges from warm, optically thin material, assumed here to be located in the upper layer of a flared accretion disk, which is irradiated by the central star.  The dust in these upper layers of different T\\,Tauri disks appears to vary continously between two extreme types: a) small amorphous dust grains or a mixture of small and large grains, with a spectrum very similar to the dust of the ISM and b) large amorphous dust grains with a possible addition of crystalline silicate. If we think of these differences in terms of evolution, then the reason for the difference between these two dust types could be either the removal of the small grains from the upper disk layers or the coagulation of small grains to larger ones. A removal of small dust grains by pure settling to the mid-plane of the disk is unlikely to explain the observed kind of differentiation since the settling time for larger grains is shorter than for the smaller ones. A further mechanism which could be responsible for the vanishing of small grains is by the radiation pressure of the central star.  A star of solar luminosity would be quite effective in removing tenth micron sized particles.  But in this case, a permanent diffusion of particles from the mid-plane on shorter time scales should avoid a significant decrease of the amount of small particles in the upper layer (\\cite{takeuchi03}).  The effect of coagulation is a possible explanation, if constructive collisions to larger grains dominate over destructive collisions which produce again small grains.  This should be the case in T\\,Tauri disks where the SED can be well fit by gas-rich disks in hydrostatic equilibrium.  An interesting detail of our studies is the availability of separate spectra from the components of two binary T\\,Tauri systems (AS\\,205 and S\\,CrA). We observed significant differences in strength and also differences in the shape of the silicate band of the components, with the secondary having the - relatively - stronger emission.  Since the components presumably have the same age, the difference must have other reasons than the evolutionary state of the stars."
    },
    "0311/astro-ph0311064_arXiv.txt": {
        "abstract": "I present an overview of the hierarchy of structures existing in the interstellar medium (ISM) and the possible mechanisms that cause the fragmentation of one level into the next, with the formation of stars as its last step. Within this framework, I then give an overview of the contributions to this session. Numerical work addresses, at the largest scales, the shaping and formation of structures in the ISM through turbulence driven by stellar energy injection, and the resulting star formation rate as a function of mean density. At the scales of molecular clouds, results comparing observational and numerical data on the density and velocity structure of turbulence-produced cores, as well as their mass spectra, are summarized, together with existing theories of core and star formation controlled by the turbulence. Observationally, an attempt to discriminate between the standard and turbulent models of star formation is presented, finding inconclusive results, but suggesting that both turbulence and the magnetic field are dynamically important in molecular clouds and their cores. Finally, various determinations of the magnetic field strength and geometry are also presented. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311252_arXiv.txt": {
        "abstract": "We present detailed calculations of the prompt spectrum of gamma-ray bursts (GRBs) predicted within the fireball model framework, where emission is due to internal shocks in an expanding relativistic wind. Our time dependent numerical model describes cyclo-synchrotron emission and absorption, inverse and direct Compton scattering, and $e^\\pm$ pair production and annihilation (including the evolution of high energy electro-magnetic cascades). It allows, in particular, a self-consistent calculation of the energy distribution of $e^\\pm$ pairs produced by photon annihilation, and hence a calculation of the spectra resulting when the scattering optical depth due to pairs, $\\tau_\\pm$, is high. We show that emission peaks at $\\sim1$~MeV for moderate to large $\\tau_\\pm$, reaching  $\\tau_\\pm\\sim 10^2$. In this regime of large compactness we find that (i) A large fraction of shock energy can escape as radiation even for large $\\tau_\\pm$; (ii) The spectrum depends only weakly on the magnetic field energy fraction; (iii) The spectrum is hard, $\\varepsilon^2 dN/d\\varepsilon\\propto \\varepsilon^\\alpha$ with $0.5<\\alpha<1$, between the self absorption ($\\varepsilon_{ssa}= 10^{0.5\\pm0.5}\\keV$) and peak ($\\varepsilon_{peak}= 10^{0.5\\pm0.5} \\MeV$) photon energy, (iv) and shows a sharp cutoff at $\\sim10$~MeV; (v) Thermal Comptonization leads to emission peaking at $\\varepsilon_{peak} \\gtrsim30$~MeV, and can not therefore account for observed GRB spectra. For small compactness, spectra extend to $>10$~GeV with flux detectable by GLAST, and the spectrum at low energy depends on the magnetic field energy fraction. Comparison of the flux at $\\sim1$~GeV and $\\sim100$~keV may therefore allow to determine the magnetic field strength. For both small and large compactness, the spectra depend only weakly on the spectral index of the energy distribution of accelerated electrons. ",
        "introduction": "\\label{sec:intro} \\indent  It is widely accepted that $\\gamma$-ray bursts are produced by the dissipation of kinetic energy in a highly relativistic wind, driven by gravitational collapse of a (few) solar mass object into a neutron star or a black hole (see e.g. \\citet{fireballs1,fireballs2,fireballs3} for reviews). The prompt $\\gamma$-ray emission is believed to be produced by synchrotron and inverse-Compton emission of electrons accelerated to relativistic energy by internal shocks within the expanding wind \\citep[see, however,][]{LGCR00,GLCR00}. Synchrotron emission is favored if the fireball is required to be \"radiatively efficient,\" i.e. if a significant fraction of the fireball energy is required to be converted to $\\gamma$-rays.  Over a wide range of model parameters, a large number of $e^\\pm$ pairs are produced in internal collisions, due to annihilation of high energy photons \\citep[e.g.][]{MR00, GSW01}. In fact, if the internal shocks occur at small enough radii, the plasma becomes optically thick and a second photosphere is formed \\citep{MR00,MRRZ} beyond the photosphere associated with the electrons initially present in the fireball. As we show here \\citep[\\S~\\ref{sec:compactness}, see also ][]{GSW01} requiring the emission to be dominated by $\\sim1$~MeV photons implies, within the fireball model framework, a moderate to large optical depth due to scattering by pairs. When the scattering optical depth due to pairs is high, calculation of the emergent spectrum becomes complicated. Relativistic pairs cool rapidly to mildly-relativistic energy, where their energy distribution is determined by a balance between emission and absorption of radiation. The emergent spectrum, which is affected by scattering off the pairs population, depends strongly on the pairs energy distribution, and in particular on the \"effective temperature\" which characterizes the low-end of the energy distribution. The pairs energy distribution is difficult to calculate analytically. Moreover, analytic calculation of the spectrum emerging from the electromagnetic cascades initiated by photon annihilation is also difficult. Therefore, in order to derive the emergent spectrum a numerical model is needed \\citep[see e.g.][]{Ghisellini99, Zhang03}.  Emission from steady plasma, where pair production and annihilation are taken into consideration, was numerically studied in the past in the context of active galactic nuclei (AGN's). It was found that thermal plasma, optically thin to Thomson scattering, and characterized by a comoving compactness $10 \\lesssim l' \\lesssim 10^3$, has a normalized pair temperature $\\theta \\equiv kT/m_ec^2 \\approx 10^{-2} - 10^{-1}$ \\citep{Lightman82, Svensson82, Svensson84}. The dimensionless compactness parameter $l$ is defined by $l \\equiv L \\sigma_T / R m_e c^3$, where  $L$ is the luminosity, and $R$ is a characteristic length of the object. The above result holds also in a scenario considering injection of high energy particles, which lose their energy via IC scattering of low energy photons, and thermalize before annihilating \\citep{Svensson87, LZ87}. The optical depth for scattering by pairs was found in the above analyses to be $0.1 \\lesssim \\tau_\\pm \\lesssim 15$. However, $\\tau_\\pm$ strongly depends on the comoving compactness $l'$, and sharply increases beyond these values when the compactness increases beyond $10^3$ \\citep[see][]{ LZ87, Svensson87}. When the scattering optical depth $\\tau_\\pm \\gg 1$ and $\\theta \\ll 1$, solving the Kompaneets equation gives a cutoff at normalized photon energy $h\\nu/m_ec^2 \\simeq 1/\\tau_\\pm^2$ \\citep{Sunyaev80}.  The results mentioned above are not directly applicable to GRB plasma. The GRB plasma is rapidly evolving and steady state can not be assumed. For example: The electron distribution function does not reach thermal equilibrium even at low energy (see \\S~\\ref{sec:results}); Due to expansion, the average number of scattering a photon undergoes is $\\sim\\tau_\\pm$ rather than $\\tau_\\pm^2$; The expansion (and cooling) of plasma electrons has a significant effect on the photon spectrum when $\\tau_\\pm$ is high. Moreover, significant luminosity at high energies (exceeding the pair production threshold) is expected due to synchrotron emission from energetic particles in strong magnetic fields, $B\\sim 10^5 - 10^6$~G, typical for GRB shell-shell collision phase, a phenomenon that was not considered in the analyses mentioned above. And last, in the GRB case a non-thermal high energy electron population is assumed to be produced, which leads at large compactness to the formation of a rapid electro-magnetic cascade, the evolution of which was not considered in the past.  A numeric calculation of GRB spectra that takes into consideration creation and annihilation of pairs is complicated. The evolution of electromagnetic cascades initiated by the annihilation of high energy photons occurs on a very short time scale. On the other extreme, evolution of the low-energy, mildly relativistic pairs, which is governed by synchrotron self-absorption, direct and inverse Compton emission takes much longer time. The large difference in characteristic time scales poses a challenge to numeric calculations. For this reason, the only numerical approach so far \\citep{Pilla98} was based on using a Monte-Carlo method. This model, however, did  not consider the parameter space region where pairs strongly affect the spectra. Another challenge to numerical modeling is due to the fact that at mildly relativistic energies the usual synchrotron emission and IC scattering approximations are not valid, and a precise cyclo-synchrotron emission and direct- inverse Compton scattering calculations are required.  In this work, we consider emission within the fireball model framework, resulting from internal shocks within an expanding relativistic wind. These shock waves dissipate kinetic energy, and accelerate a population of relativistic electrons. We adopt the common assumption of a power law energy distribution of the accelerated particles, and calculate the emergent spectra. We present the results of time dependent numerical calculations, considering all the relevant physical processes: cyclo-synchrotron emission, synchrotron self absorption, inverse and direct Compton scattering, and $e^\\pm$ pair production and annihilation including the evolution of high energy electro-magnetic cascades. A full description of the numerical code appears in \\citet{Pe'er03b}.  We note, that Comptonization by a thermal population of electrons (and possibly $e^\\pm$ pairs) was considered as a possible mechanism for GRB production \\citep[e.g. ][]{Ghisellini99}, following the evidence that, at least in some cases, the GRB spectra at low energy is steeper at early times than expected for synchrotron emission \\citep{Crider,Preece,Frontera}. This model is different than the common model considered in the previous paragraph in assuming that the kinetic energy dissipated in a collision between two \"shells\" within the expanding wind is continuously distributed among all shell electrons, rather than being deposited at any given time into a small fraction of shell electrons that pass through the shock wave, and in assuming that energy is equally distributed among electrons, rather than following a power-law distribution. There is no known acceleration mechanism that leads to the above energy distribution among electrons, and the spectrum predicted in this scenario does not account for the claimed steep spectra, which may be naturally explained as a contribution to the observed $\\gamma$-ray radiation of photospheric fireball emission \\citep[e.g. ][]{MR00}. It is nevertheless worthwhile to derive the spectrum that is expected from thermal Comptonization under plasma conditions typical to GRB fireballs. This will allow to determine whether this process may significantly contribute to GRB $\\gamma$-ray emission. Our numerical code enables an accurate calculation of the pair temperature in this scenario, as well as reliable calculation of the emergent spectrum.  This paper is organized as follows. In \\S\\ref{sec:parameters} we derive the plasma parameters at the internal shocks stage, and their dependence on uncertain model parameters values. In \\S~\\ref{sec:compactness} it is shown that moderate to large compactness is expected for the parameter range where emission peaks at $\\sim1$~MeV, and approximate analytic results describing the emission in this regime are given. Our numerical methods are briefly presented in \\S\\ref{sec:numerical}; a detailed description of the model is found in \\citet{Pe'er03b}. Our numerical results for the scenario of acceleration of particles in shock waves are presented in \\S\\ref{sec:results}. In \\S\\ref{sec:ghis} we present both analytical and numeric calculations of the spectra resulting from thermal Comptonization. We summarize and conclude our discussion in \\S\\ref{sec:summary}, with special emphasis on observational implications. ",
        "conclusions": "\\label{sec:summary}   Within the fireball model framework, synchrotron emission peaking at $\\sim1$~MeV may be obtained with small compactness, $l'\\sim1$, only for very short variability time, $\\Delta t\\le10^{-4.5}$~s, and large $\\Gamma$, $\\Gamma\\ge10^3$ [see Eqs.~(\\ref{eq:eps-l}) and ~(\\ref{eq:l-eps}), of \\S~\\ref{sec:compactness}]. For the longer variability time commonly assumed in modelling GRBs, $\\Delta t\\sim1$ to 10~ms \\citep[e.g.][]{fireballs1,fireballs2,fireballs3}, $l'\\sim10^2$ to $l'\\sim10^3$ is obtained for the parameter range where synchrotron emission peaks at $\\sim1$~MeV (For smaller compactness, spectra peak at lower, X-ray, energy). This result has two main consequences. First, observed GRB spectra are expected to be significantly modified by the presence of pairs. Second, peak energy $\\gg1$~MeV can not be obtained for $L\\sim10^{52}{\\rm erg/s}$, since it would imply $l'>>10^3$ [see Eqs.~(\\ref{eq:eps-l}),~(\\ref{eq:T}),~(\\ref{eq:theta_obs})], in which case most of the radiation will not escape due to large optical depth to Thomson scattering by pairs. These conclusions are consistent with the conclusions of \\citet{GSW01}, and may provide an explanation for the lack of bursts with peak energy $\\gg1$~MeV.  It should be noted at this point that GRB observations do not allow, in most cases, to identify variability on $\\sim1$~ms time scale. Rapid variability, on $\\sim1$~ms time scale, has been observed in some bursts \\citep{Bhat92,Fishman94}, and most bursts show variability on the shortest resolved time scale, $\\sim10$~ms \\citep{Woods95,Walker00}. It should be kept in mind, however, that variability on much shorter time scale would not have been possible to resolve experimentally, and can not therefore be ruled out.  We have demonstrated in \\S~\\ref{sec:compactness} that for $l'$ in the range of $10^2$ to $10^3$ the characteristics of emitted radiation are determined mainly by $l'$, with weak dependence on the values of other parameters [see, e.g., Eq.~(\\ref{eq:theta_obs})]. The peak of the specific luminosity is expected to be close to $\\sim1$~MeV, $\\sim 1 (L_{52}/\\Delta_{t,-3})^{1/5}(l'/100)^{-1}$~MeV [Eq.~(\\ref{eq:theta_obs})], and the spectrum is expected to differ significantly from an optically thin synchrotron spectrum.  These conclusions are consistent with the results of our detailed numerical calculations. We have presented numeric results of calculations of prompt GRB spectra within the fireball model framework, using a time dependent numerical code which describes cyclo-synchrotron emission and absorption, inverse and direct Compton scattering, and $e^\\pm$ pair production and annihilation. We have shown that the spectral shape depends mainly on the compactness parameter, which is most sensitive to the fireball Lorentz factor $\\Gamma$, $l'\\propto\\Gamma^{-5}$. For large compactness (small $\\Gamma$), $l'>100$, the spectra peak at $\\sim1$~MeV, show steep slopes at lower energy, $\\varepsilon^2 dN/d\\varepsilon\\propto \\varepsilon^\\alpha$ with $0.5<\\alpha<1$, and show a sharp cutoff at $\\sim10$~MeV (see fig.~\\ref{fig:results2}). The spectra depend only weakly in this regime on the power-law index $p$ of accelerated electrons and on the magnetic field energy fraction $\\epsilon_B$ (see fig.~\\ref{fig:p3}). For small to moderate compactness (large $\\Gamma$), $l'\\lesssim10$, spectra extend to $\\gtrsim10$~GeV (see figures~\\ref{fig:results1},~\\ref{fig:results3}). The spectrum at lower energy depends only weakly on $p$ (see fig.~\\ref{fig:p1}), but strongly on $\\epsilon_B$ (see fig.~\\ref{fig:xi_B}). For magnetic field close to equipartition, spectra peak at $\\sim0.1$~MeV for $l'\\sim10$ and extend to higher energy with spectral index $\\alpha=10^{0\\pm0.5}$.  Since moderate to large compactness is required for a synchrotron peak at $\\sim1$~MeV, the effects of pair-production and direct and inverse-Compton scattering by pairs are predicted to be large for observed GRBs. In this case, simple analytic approximations of synchrotron-self-Compton emission do not provide an accurate description of the emergent spectra. In particular, Compton scattering by the pairs, which are accumulated at intermediate energy $\\theta\\equiv\\varepsilon/m_e c^2 \\approx 0.1$ (see figure~\\ref{fig:elec_ad}), results in a steep slope, $0.5\\le\\alpha\\le1$, in the 1~\\keV to 1~\\MeV band. Still steeper slope, $\\alpha \\approx 3$, is obtained at lower energies, below the self-absorption frequency determined by the quasi-thermal pair distribution, $\\varepsilon_{ssa}^{ob.}\\approx 10^{0.5\\pm0.5} \\keV$ (see Eq.~\\ref{eq:eps_ssa_thermal}). These steep slopes may account for the steep slopes observed at early times in some GRBs \\citep[see][]{Frontera,Preece,Ghirlanda}.  The spectra presented in this paper (\\S~\\ref{sec:results}) describe the emission resulting from a single collision within the expanding fireball wind, for various choices of model parameters ($L$, $\\Delta t$, $\\Gamma$ etc.). Observed spectra are expected to be combinations of spectra produced by many single-shell collisions, each characterized by different parameters. This is due to the fact that at any given time a distant observer is expected to receive  radiation from many collisions taking place at various locations within the wind. Moreover, observed spectra are inferred from measurements where the signal is integrated over time intervals longer than that expected for the duration of a single collision. A detailed comparison with observations requires a detailed model describing the distribution of single-shell collision parameters within the fireball wind \\citep[see, e.g., ][]{Daigne98,Panaitescu99,GSW01}. The construction and investigation of such models are beyond the scope of this manuscript.  We have shown (\\S~\\ref{sec:ghis}, figure~\\ref{fig:ghisellini}) that the quasi-thermal Comptonization scenario \\citep[e.g. ][]{Ghisellini99}, in which kinetic energy dissipated in shocks is continuously distributed roughly equally among all electrons, leads to high peak energy, $\\gtrsim 30 \\MeV$. This scenario may not account therefore for observed GRB spectra.  Clearly, the most stringent constraints on $l'$, and hence on $\\Gamma$, will be provided by measurements of the spectra at high energy $\\gtrsim0.1$~GeV. The fluxes predicted by the model are detectable by GLAST\\footnote{http://www-glast.stanford.edu}. For large compactness, where emission is strongly suppressed above 0.1~GeV, the model predicts a steep spectrum at low energy, which is weakly dependent on other model parameters (figure~\\ref{fig:p3}). For small compactness, where strong emission is expected above 10~MeV, the low energy spectrum depends mainly on $\\epsilon_B$ (fig.~\\ref{fig:xi_B}). Comparison of the flux at $\\sim1$~GeV and $\\sim100$~keV (available, e.g.  with SWIFT\\footnote{http://www.swift.psu.edu}) may therefore allow to determine the fireball magnetic field strength. Stringent constraints on $p$, the spectral index of the energy distribution of accelerated electrons, will be difficult to obtain, due to the weak dependence of spectra on this parameter."
    },
    "0311/astro-ph0311246_arXiv.txt": {
        "abstract": "{ We have made self-consistent models of the density and temperature profiles of the gas and dust surrounding embedded luminous objects using a detailed radiative transfer model together with observations of the spectral energy distribution of hot molecular cores. Using these profiles we have investigated the hot core chemistry which results when grain mantles are evaporated, taking into account the different binding energies of the mantle molecules, as well a model in which we assume that all molecules are embedded in water ice and have a common binding energy. We find that most of the resulting column densities are consistent with those observed toward the hot core G34.3+0.15 at a time around 10$^4$ years after central luminous star formation. We have also investigated the dependence of the chemical structure on the density profile which suggests an observational possibility of constraining density profiles from determination of the source sizes of line emission from desorbed molecules. ",
        "introduction": "Hot cores are compact (diameters $\\leq 0.1$ pc), dense (n$\\geq 10^7$ cm$^{-3}$), hot (T$\\geq 100$ K), and dark ($A_v\\geq 100$ mag) molecular cloud cores. At least some of them are believed to be precursors to ultracompact HII regions and in the early stage of forming massive stars because luminous infrared sources, massive bipolar outflows and/or inflow motion are observed associated with them (e.g., Henning et al. 2000; Kurtz et al. 2000; Churchwell 2002). They are observationally characterized by anomalously large abundances of some molecular species including hydrogenated, saturated, small molecules such as H$_2$O and H$_2$S, and large complex molecules such as HCOOCH$_3$ and (CH$_3$)$_2$O, in contrast to colder dark clouds. It is thought that these molecules are produced under conditions specific to star formation: the freeze-out of gas-phase species onto dust grains and the grain surface reactions in the cold collapsing pre-protostellar phase, and the evaporation of the grain mantle material and the subsequent gas-phase reactions in the hot protostellar phase (e.g., Millar 1993, 1997). Therefore, the observed large column densities of some molecules in hot cores are also expected to be an indicator of luminous high mass star formation which yields a large mass of high temperature gas. Previously, chemical models of hot cores have been studied together with the pre-stellar phase or with some initial conditions of evaporated grain mantle species, and have been successful in reproducing their observational properties (e.g., Brown et al. 1988; Millar et al. 1991; Charnley et al. 1992; Caselli et al. 1993; Charnley 1997; Millar et al. 1997; Viti \\& Williams 1999; Rodgers \\& Charnley 2001; Doty et al. 2002).  Physical properties of hot cores have been investigated on the basis of observational studies of molecular line and dust continuum emission. Some physical models together with radiative transfer calculation and/or chemical calculation have succeeded in reproducing the observations of spectral energy distributions and/or molecular column densities (e.g., Millar et al. 1997; Kaufman et al. 1998; Osorio et al. 1999; Hatchell et al. 2000; van der Tak et al. 2000a; Doty et al. 2002). On the other hand, the dynamical processes and the formation mechanisms of massive stars are not yet established because of the difficulty of treating in a self-consistent manner young stellar evolution, the accretion flow and radiation pressure. Non-spherical accretion to protostellar cores, for example, through accretion disks, and/or coalescence of lower mass stars may be involved in massive star formation, but there are, as yet, no accepted theories (e.g., Bonnell et al. 1998; Stahler et al. 2000; Yorke \\& Sonnhalter 2002). The nature and timescale of the dissipation of the surrounding envelopes of young massive stars, for example, by powerful outflows and ionizing strong winds which stop the accretion, are also uncertain but important in determining the mass of forming stars in the case of the isolated star formation (e.g., Shu et al. 1987; Nakano et al. 1995).  In this paper, we have modeled the density and temperature profiles of hot cores, making use of radiative transfer calculations, and then investigated the hot core chemistry, taking into account the temperature dependent grain mantle evaporation. In the following section, we present the physical model of the density profiles constrained by observational results and the temperature profiles calculated by solving the radiative transfer equation. In Sect. 3 we introduce the chemical model along with the temperature dependent grain mantle evaporation process. These models are applied to the hot core G34.3+0.15, and the results of chemical calculation are described and compared with the observations in Sect. 4. Also, we investigate the dependence of the chemical structure of hot cores on (1) the trapping of mantle molecules in water ice and (2) the density profile generally to discuss the chemical and physical condition of massive-star-forming clumps in Sect. 5. Finally, the results are summarized in Sect. 6. ",
        "conclusions": "\\subsection{Trapping of mantle molecules in water ice}   \\begin{table} \\caption[]{Initial abundances of mantle molecules for Model B}\\label{T3} $$ \\begin{array}{p{0.11\\linewidth}c|p{0.12\\linewidth}c|p{0.12\\linewidth}c} \\hline Species & {\\rm Abundance} & Species & {\\rm Abundance} & Species & {\\rm Abundance}  \\\\ \\hline C$_2$H$_2$ & 5.0\\ 10^{-7} & CO$_2$ & 3.0\\ 10^{-6} & O$_2$ & 1.0\\ 10^{-6} \\\\ CH$_4$ & 2.0\\ 10^{-7} & H$_2$CO & 2.0\\ 10^{-6} & H$_2$O & 2.8\\ 10^{-4} \\\\ C$_2$H$_4$ & 1.0\\ 10^{-8} & CH$_3$OH & 4.0\\ 10^{-7} & NH$_3$ & 8.0\\ 10^{-7} \\\\ C$_2$H$_6$ & 1.0\\ 10^{-8} & C$_2$H$_5$OH & 1.0\\ 10^{-8} & H$_2$S & 3.0\\ 10^{-7} \\\\ & & & & OCS & 1.0\\ 10^{-7} \\\\ \\hline \\end{array} $$ \\end{table} \\begin{table} \\caption[]{Molecular column densities of Model B and observations.}\\label{T4} $$ \\begin{array}{p{0.13\\linewidth}|cr|p{0.17\\linewidth}|cr} \\hline & \\multicolumn{2}{|c|}{N_{\\rm mol}\\ [{\\rm cm}^{-2}]} & & \\multicolumn{2}{|c}{N_{\\rm mol}\\ [{\\rm cm}^{-2}]} \\\\ \\cline{2-3} \\cline{5-6} Species & {\\rm Model\\ B^{\\mathrm{a}}} & \\multicolumn{1}{c|}{\\rm Obs.} & Species & {\\rm Model\\ B^{\\mathrm{a}}} & \\multicolumn{1}{c}{\\rm Obs.} \\\\ \\hline CO & 2.0\\ 10^{20} & >2.9\\ 10^{19} & H$_2$S & 2.4\\ 10^{16} & >1.1\\ 10^{16} \\\\ C$_2$H & 3.2\\ 10^{15} & >8.2\\ 10^{15} & SO & 9.7\\ 10^{15} & >2.4\\ 10^{15} \\\\ CH$_3$C$_2$H & 1.3\\ 10^{12} & 1.8\\ 10^{16} & SO$_2$ & 2.3\\ 10^{16} & 9.3\\ 10^{15} \\\\ H$_2$O & 3.5\\ 10^{19} & 1.3\\ 10^{19} & CS & 1.5\\ 10^{15} & 1.0\\ 10^{16} \\\\ H$_2$CO & 2.3\\ 10^{17} & >2.3\\ 10^{16} & OCS & 1.9\\ 10^{16} & >1.8\\ 10^{16} \\\\ HCO$^+$ & 5.5\\ 10^{14} & >1.4\\ 10^{15} & H$_2$CS & 1.8\\ 10^{14} & >3.8\\ 10^{16} \\\\ HCO & 6.4\\ 10^{13} & >5.3\\ 10^{13} & HCS$^+$ & 5.6\\ 10^{12} & >3.0\\ 10^{13} \\\\ NO & 1.6\\ 10^{16} & >5.8\\ 10^{15} & NS & 5.0\\ 10^{13} & >4.0\\ 10^{13} \\\\ NH$_3$ $^{\\mathrm{b}}$ & 2.6\\ 10^{18} & >2.7\\ 10^{18} & CH$_3$OH & 7.4\\ 10^{16} & 2.6\\ 10^{16} \\\\ HNC & 1.3\\ 10^{15} & <3.3\\ 10^{14} & C$_2$H$_5$OH & 1.8\\ 10^{15} & 1.7\\ 10^{15} \\\\ HCN & 5.0\\ 10^{15} & 1.2\\ 10^{15} & CH$_2$CO & 1.3\\ 10^{14} & 6.7\\ 10^{14} \\\\ CN & 1.3\\ 10^{15} & >5.9\\ 10^{14} & CH$_3$CHO & 1.3\\ 10^{14} & 2.4\\ 10^{14} \\\\ HC$_3$N & 1.1\\ 10^{15} & >2.5\\ 10^{13} & HCOOCH$_3$ & 1.0\\ 10^{15} & 1.4\\ 10^{15} \\\\ CH$_3$CN & 3.4\\ 10^{12} & >2.4\\ 10^{14} & CH$_3$OCH$_3$ & 5.5\\ 10^{13} & 3.7\\ 10^{15} \\\\ \\hline \\end{array} $$ \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] Calculated results at $10^4$ yr. \\item[$^{\\mathrm{b}}$] In the case of $\\Theta_{\\rm beam}=0.025$ pc. \\end{list} \\end{table}  It is known that water is one of the most abundant ice component in dense clouds. Some observations of ice absorption features in infrared spectra towards young stellar objects have suggested that a certain amount of CO is frozen in H$_2$O-rich matrices (e.g., review by Allamandola et al. 1999), and laboratory experiments have shown that solid CO is able to diffuse into the porous structure of H$_2$O and be trapped in it (Collings et al. 2003). These studies suggest that mantle molecules other than CO also could be trapped in water ice. Thus, in this subsection we examine the effects of this trapping on the chemical structure of the hot core G34.3+0.15. In order to see the extreme case, we put the injection radius of each mantle molecule as $r_{{\\rm inj},i}\\equiv r_{{\\rm inj, H}_2{\\rm O}}$, and assume that it is completely trapped in H$_2$O matrices outside $r_{{\\rm inj, H}_2{\\rm O}}$, except CO, for which we assume that $x({\\rm CO})=0.3x({\\rm H}_2{\\rm O})$ is trapped outside $r_{{\\rm inj, H}_2{\\rm O}}$ based on observations (e.g., Chiar et al. 1998; cf. Doty et al. 2002). As a result of the time-dependent chemical calculation described in Sect. 3, we get the molecular column densities of the hot core consistent with most of the observations at around $10^4$ yr (Table 4). The initial condition listed in Table \\ref{T3} is used in this calculation, which is slightly different from that in Table 1, in order to fit the results to the observations. Although there are few differences in the integrated molecular column densities, we find that very different structure appear in the radial profiles of molecular abundances at around $10^4$ yr between the results in Sect. 4 (Fig.4: Model A) and those obtained from this calculation (Fig.5: Model B). Especially the differences are remarkable in the following molecules: (1) the binding energies of the molecules or their parent species are low, that is, the injection radii listed in Table \\ref{T1} are large, such as H$_2$CO, HCO, H$_2$S, and SO, or (2) daughter molecules whose production are suppressed where H$_2$O is evaporated from grain mantle, such as CH$_3$C$_2$H and CH$_3$CN. For comparison, the resulting radial abundance profiles of SO and CH$_3$CN from Model A (dotted lines) and B (solid lines) are plotted in the linear radial scales in Fig.6. Detailed observations of radial profiles of molecular abundances may make it possible to examine what fraction of molecules is trapped in water ice in cooler region of hot cores.  \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{H4682f6a.eps}} \\resizebox{\\hsize}{!}{\\includegraphics{H4682f6b.eps}} \\caption{The comparison of the radial abundance profiles of (a) SO and (b) CH$_3$CN in the linear radial scales between the models without (dotted lines: Model A) and with (solid lines: Model B) the trapping of mantle molecules in water ice. We can find the difference in the location of jump in (a) and that the abundance is low through the clump in Model B in (b).} \\end{figure}   \\subsection{Dependence of chemical structure on density profile -- General model}  The comparison between the results of chemical calculation and the observations of various kinds of molecular lines will be useful to constrain the physical condition of massive star forming regions (e.g., Viti \\& Williams 1999). The density profile of hot molecular cores is one of the possible tracers of the evolution process just before and after massive star formation. In this section we will investigate the dependences of the chemical structure of hot cores on the density profile to discuss the physical condition of massive-star-forming clumps.  \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{H4682f7a.eps}} \\resizebox{\\hsize}{!}{\\includegraphics{H4682f7b.eps}} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{H4682f7c.eps}} \\caption{(a) Density and (b) temperature profiles in the logarithmic scales for the models with $L_*=10^5L_{\\odot}$ and $r_{\\rm c}=$ 0.01 (dotted lines), 0.05 (solid lines), and 0.1 pc (dashed lines). The temperature profiles directly reflect the difference in the density profiles. (c) Temperature profiles in the linear scales represent clearly the dependence on the density models around the temperatures of grain mantle evaporation. Profiles of the model with $L_*=10^3L_{\\odot}$ and $r_{\\rm c}=$ 0.05 pc (dot-dashed lines) are also plotted for comparison.} \\end{figure}  The chemical structures of three different types of molecular clupms are investigated here. Their density profiles are set by Eq. (\\ref{eq1}) with core radii of $r_{\\rm c}=$ 0.01, 0.05, and 0.1 pc, and plotted in Fig.7(a) with dotted, solid, and dashed lines, respectively. Making use of these density profiles, we calculate the temperature profiles by solving the radiative transfer equation as described in Sect. 2.2. The resulting temperature distributions are shown in Fig.7(b). We can see from the figure that the temperature profiles directly reflect the difference in the density profiles. This is because the temperature profiles of hot cores are determined by the balance between the absorption of radiation from the central star and the reemission of the radiation by dust grains in the core as described by Eq. (\\ref{eq2}).  The inner regions of the cores are optically thick enough to the radiation from the central stars so that the photons from the stars diffuse via interactions with the surrounding dusty material. If we simply assume that the dust opacity depends on the frequency of radiation as $\\kappa_{\\nu}\\propto \\nu^n$ and the density profile of a clump is set as $\\rho\\propto r^{-p}$, for example, the temperature profile of the clump illuminated by a central star is derived as $T\\propto r^{-(p+1)/(4-n)}$ in the optically thick limit from the equation of local radiative equilibrium and the radiative transfer equation (e.g., Adams \\& Shu 1985). Meanwhile, since the evaporation process of grain mantle species is very sensitive to the temperature as shown in Sect. 3, it is to be expected that the density profiles affect the radial abundance profiles of molecules in hot cores via the temperature profiles. In order to see the dependence on the density models clearly, the temperature profiles in the region of grain mantle evaporation are plotted with linear scales in Fig.7(c).  The luminosity of the central star is another factor that influences the temperature profile of the clump. As we can see from Fig.7(b) and (c), the temperature distribution of the clump whose core radius is $r_{\\rm c}=0.05$ pc and whose central star luminosity is $L_*=10^3L_{\\odot}$ (dot-dashed lines) is similar to that of the model with $r_{\\rm c}=0.01$ pc and $L_*=10^5L_{\\odot}$. The difference in luminosity will be, however, distinguishable by observing the spectral energy distribution. The hydrogen column density will not have very much effect on the temperature structure in the optically thick region because basically $T(r)$ depends not on the absolute density but on the steepness of the density profile as we have shown in the above example. The optically thick region spreads (shrinks) if the hydrogen column density is high (low) so that it will affect the temperature profile at the outer radii.  \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{H4682f8a.eps}} \\resizebox{\\hsize}{!}{\\includegraphics{H4682f8b.eps}} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{H4682f8c.eps}} \\resizebox{\\hsize}{!}{\\includegraphics{H4682f8d.eps}} \\resizebox{\\hsize}{!}{\\includegraphics{H4682f8e.eps}} \\caption{The radial abundance profiles of some parent and daughter molecules of the models in the linear radial scales with $L_*=10^5L_{\\odot}$ and $r_{\\rm c}=$ 0.01 (dotted lines), 0.05 (solid lines), and 0.1 pc (dashed lines). The abundance jumps appear at the inner radii in the models with the smaller core radii, that is, having steeper density and temperature profiles.} \\end{figure}  We calculate the gas-phase chemical reactions in the clumps with the above density and temperature profiles and $L_*=10^5L_{\\odot}$ by means of the chemical model in Sect. 3. Fig.8 gives the resulting radial abundance profiles of some molecules at $10^4$ yr in the linear radial scales (while those in Fig.4 are scaled in the logarithmic manner). They are parent or daughter species and have dramatic changes at certain radii related to grain mantle evaporation. H$_2$O, CO$_2$, H$_2$CO, H$_2$S, C$_2$H$_2$, and CH$_3$OH are possible parent molecules, and have abundance jumps at their injection radii. The daughter species HC$_3$N, CH$_3$CN, and HCOOCH$_3$ are produced mainly from C$_2$H$_2$, HCN, and CH$_3$OH, respectively, in this model, and their abundance jumps appear at the same radii as those of the parent species. Comparison between reasonably high resolution observations of abundance jumps and precise radiative transfer and physical-chemical models of hot cores may make it possible to test the formation processes of these molecules.  From Fig.8 we can see that the abundance jump of each molecule occurs at a smaller radius in the clump with a smaller core radius $r_{\\rm c}$, that is, with a steeper density and therefore a steeper temperature profile. The molecules in Fig.8 can keep high fractional abundances only in the inner region of the clump ($< 0.1$pc) if $r_{\\rm c}=0.01$pc. This result suggests that the chemically rich hot cores, whose observational beam-averaged molecular column densities are large, have flat density profiles rather than steep ones. On the basis of the low and high mass star formation theories, the central density profiles of clumps are thought to be steep just after the star formation, at least if accretion of the surrounding material to the central protostellar cores is important (e.g., Shu 1977; Yorke \\& Bodenheimer 1999; Masunaga \\& Inutsuka 2000; Yorke \\& Sonnhalter 2002). Subsequently, dissipation by outflows will let the clumps gravitationally unbound and their central density profiles will be flattened (e.g., Shu et al. 1987; Nakano et al. 1995). In the case of massive star formation, powerful outflows with high mass loss rates as well as radiation pressure and strong winds from the stars may accelerate the dissipation process. If massive stars are formed as a result of coalescence of lower mass stars, the envelopes will have different density profiles from those of isolated star-forming clumps (e.g., Bonnell et al. 1998). Statistical observations of various kinds of molecular lines towards hot cores could constrain their density profiles and the massive star formation theories.  The effects of inflow motion of hot cores have not been considered in this paper. If the inflow has a velocity of a few km s$^{-1}$ as is estimated from the observed atomic and molecular lines (e.g., Keto 2002), it will not affect the structure on the scale of about 0.25 pc, which corresponds to the observational beam size for most of the molecular lines we have used in this paper, within the timescale of $10^4$ yr. Comparison between  model calculations and the forthcoming high spatial resolution and high sensitivity observations could reveal the dependence of the chemical structure of hot cores on the inflow motion. In addition, the temporal evolution of young massive stars in hot cores is also a factor that may possibly influence the fractional abundances of some species (cf. Viti \\& Williams 1999)."
    },
    "0311/astro-ph0311070_arXiv.txt": {
        "abstract": "We model the infrared (IR) emission from the double-ring disk of $\\hda$, using a porous dust model that was previously shown successful in reproducing the spectral energy distributions (SEDs) of the disks around $\\beta$ Pictoris and $\\hra$. The dust consists of either unaltered or highly processed interstellar materials and vacuum with a volume fraction of $\\sim 90\\%$. Together with a population of polycyclic aromatic hydrocarbon (PAH) molecules of which a large fraction is charged, the porous dust model made of coagulated but otherwise unaltered protostellar interstellar grains provides an excellent fit to the entire SED from the mid-IR to millimeter wavelengths including the PAH mid-IR emission features. The nondetection of the 21$\\mum$ crystalline silicate feature predicted from the porous dust model composed of highly-processed grains (mainly crystalline silicate dust) imposes an upper limit of $\\sim 10\\%$ on the mass fraction of crystalline silicates. Spectroscopic and broadband photometric predictions are made for SIRTF observations which will provide further tests of the applicability of the porous dust model to circumstellar disks. ",
        "introduction": "} Since dust disks around young stars, depending on their age, are the source material or the remnants of newly-formed planets, the physical, chemical, and dynamical properties of circumstellar disks and their constituent grains are crucial in understanding the formation process of planetary systems.  Over the past 2 decades, impressive evidence has been assembled for the commonality of circumstellar dust disks around pre-main-sequence stars (T Tauri stars and Herbig Ae/Be stars), main sequence (MS) stars, post-MS stars (red giants), and a white dwarf (see Backman \\& Paresce 1993, Habing et al.\\ 2001, Zuckerman 2001, Lagrange, Backman, \\& Artymowicz 2001 for reviews). To date, at least 15\\% of the A--K type stars are found to have dust disks around them (see Lagrange, Backman, \\& Artymowicz 2001). The majority of these disks are spatially extensive, with their dust quantities and intensities of scattered-starlight and of dust thermal emission gradually decreasing outward from the central stars. However, so far there exist four exceptions: (1) $\\epsilon$ Eridani -- a $\\sim 800\\Myr$ old (Henry et al.\\ 1996) K2V MS star -- has a nearly face-on dust ring with a mass and radius ($r \\sim 60\\AU$) similar to the solar system Kuiper Belt and is thought to be an analog to the young solar system (Greaves et al.\\ 1998); (2) Formahault ($\\alpha$ PsA) -- a $\\sim 200\\Myr$ old (Barrado y Navascu\\'{e}s et al.\\ 1997) A3V MS star -- exhibits a ``doughnut''-like dust annulus with a radial distance of $\\sim 150\\AU$ from the star and a width of $\\sim 60\\AU$ (Holland et al.\\ 1998, 2003); (3) $\\hra$ -- a $\\sim 8\\pm3\\Myr$ old A0V star (see Jura et al.\\ 1998) -- displays a ring-like disk peaking at $\\sim 70\\AU$ from the central star and abruptly truncated both inward and outward with a width of $\\simlt 17\\AU$ (Schneider et al.\\ 1999; Koerner et al.\\ 1998; Jayawardhana et al.\\ 1998; Telesco et al.\\ 2000); (4) $\\hda$ -- a nearby (distance to the Earth $d\\approx 99\\pc$; van den Ancker, de Winter, \\& Tjin A Djie 1998) $\\sim 5\\pm3\\Myr$ old (Weinberger et al.\\ 2000) pre-MS Herbig Ae/Be star with a spectral type of B9.5V, an effective temperature of $\\Teff\\approx 10000\\K$ and a surface gravity of $\\lg g\\approx 4.0$ (Dunkin, Barlow, \\& Ryan 1997).  The presence of dust around $\\hda$ was first reported by Jaschek, Jaschek, \\& Egret (1986), de Grijp, Miley, \\& Lub (1987), and Walker \\& Wolstencroft (1988) based on the IRAS ({\\it Infrared Astronomical Satellite}) detection of infrared (IR) excess over expected from its photosphere. It has recently been imaged in scattered light and in thermal IR emission. With the discovery of its large and morphologically complicated debris disk (see below), $\\hda$ has recently aroused considerable interest.  The disk around $\\hda$ appears to have two ring-like structures: an outer ring at $\\sim 325\\AU$ from the star with a width of $\\sim 150\\AU$; and an inner ring at $\\sim 185\\AU$. The outer ring was reported independently by Augereau et al.\\ (1999a) and by Weinberger et al.\\ (1999), respectively based on the 1.6$\\mum$ and 1.1$\\mum$ imaging of the scattered-light. They both used the {\\it Near-IR Camera and Multiobject Spectrometer} (NICMOS) on the {\\it Hubble Space Telescope} (HST). The inner ring was discovered by Weinberger et al.\\ (1999) at 1.1$\\mum$ using the same instrument. Weinberger et al.\\ (1999) also reported the detection of a gap or ``dip'' (a region of depleted material) at $\\sim 250\\AU$ from the star with a width of $\\sim 60\\AU$. They interpreted the gap as a sign of one or more planets orbiting within the disk (the planets do not have to be in the gap), but other explanations such as dust migration also exist (e.g. see Takeuchi \\& Artymowicz 2001).  This ring-gap-ring structure was later confirmed by optical coronagraphic observations made by Mouillet et al.\\ (2001) at an effective wavelength of 5850$\\Angstrom$ using the HST {\\it Space Telescope Imaging Spectrograph} (STIS). With a higher spatial resolution and a higher signal-to-noise ratio than the NICMOS/HST data, Mouillet et al.\\ (2001) found that the inner ring, with a FWHM (full width of half maximum) of $\\sim 50\\AU$, actually peaks at $\\sim 200\\pm 5\\AU$. Strong brightness asymmetry was also reported by Mouillet et al.\\ (2001). This asymmetry, interpreted as a result of the gravitational perturbation caused by putative massive planets within the disk, by the two stellar companions $\\hdb$ and $\\hdc$, or both types of objects (Mouillet et al.\\ 2001), was confirmed by the ground-based 2.2$\\mum$ near-IR imaging observation at the Palomar 200 inch telescope using the adaptive optics system (Boccaletti et al.\\ 2003).  In addition to the two annuli, the existence of a third dust component -- a population of warm, mid-IR emitting dust within $\\sim 100\\AU$ from the star -- was recently reported by Fisher et al.\\ (2000). It was shown that the 10.8 and 18.2$\\mum$ emission observed with the University of Florida {\\it Observatory Spectrometer and Camera for the Infrared} (OSCIR) on the Keck II telescope extend out to $\\sim 100\\AU$ from the star (Fisher et al.\\ 2000). The existence of such an extended, warm source close to the star, originally suggested by Augereau et al.\\ (1998) to account for the disk IR emission at wavelengths $\\lambda \\simlt 10\\mum$, was later confirmed by the 12.5, 17.9, and  20.8$\\mum$ imaging using the JPL {\\it Mid-InfraRed Large-well Imager} (MIRLIN) on the Keck II telescope (Marsh et al.\\ 2002).  More recently, the HRC ({\\it High Resolution Channel}) coronagraphic B, V, and I images of the $\\hda$ disk obtained with the HST {\\it Advanced Camera for Surveys} (ACS) revealed a tightly-wound spiral structure with two arms, one of which stretches outward to reach the nearby companion stars (Clampin et al.\\ 2003).  In a previous paper (Li \\& Lunine 2003), we have shown that the IR emission from the $\\hra$ disk can be well fit by a model invoking highly porous cometary-type dust. The porous cometary dust model was originally proposed for the $\\beta$ Pictoris disk and was shown successful in reproducing its spectral energy distribution from the near-IR to millimeter wavelengths including the 10$\\mum$ amorphous and the 11.3$\\mum$ crystalline silicate features (Li \\& Greenberg 1998). In this paper we will model the observed IR and millimeter photometric and mid-IR spectroscopic signatures of $\\hda$. The objectives of this paper are three-fold: (1) We wish to infer the physical and chemical properties of the dust in the $\\hra$ disk and its relationship to the formation of planets. (2) We wish to know how widely applicable is the porous dust model; if it can be shown that the porous dust model is also valid for many other disks at different evolutionary stages and with different geometrical structures, it will be a valuable guide for interpreting future data sets. This is of particular importance given that the incoming {\\it Space Infrared Telescope Facilities} (SIRTF) will accumulate a rich set of IR photometric and spectroscopic data for a large number of dust disks; we would hope that the porous dust model presented in this and related papers could serve as a starting point in analyzing these data. (3) We wish to make near- and mid-IR spectral and broadband photometry predictions for the $\\hda$ disk; these predictions can be compared to future SIRTF observations in order to test further the validity of the porous dust model.  This paper is organized as follows: we first summarize in \\S\\ref{sec:fluxobs} the photometric and spectroscopic data available for the $\\hda$ disk that are relevant in modeling the IR emission of the disk. We then give an overview discussion in \\S\\ref{sec:model} of the proposed dust model with emphasis on dust composition and morphology from the point of view of protoplanetary disk evolution, followed by detailed discussions of the PAH component in \\S\\ref{sec:pah} and the porous dust component in \\S\\ref{sec:pd}. The IR emission modeling technique is presented in \\S\\ref{sec:iremmethod}. In \\S\\ref{sec:ionrec} the PAH photoelectric emission rates and the electron recombination rates are calculated, using the methods described in \\S\\ref{sec:photophys} and \\S\\ref{sec:ne}. The key model parameters are discussed and specified in \\S\\ref{sec:para}. Model results are presented in \\S\\ref{sec:results}. In \\S\\ref{sec:discussion} we discuss (1) previous studies of the SED modeling of the $\\hda$ disk, (2) comparison of the $\\hda$ disk with the $\\hra$ disk, (3) exploration of parameter space (\\S\\ref{sec:robust}), (4) PAH photodestruction (\\S\\ref{sec:pahdes}), and (5) dust loss through radiation expulsion and Poynting-Robertson drag (\\S\\ref{sec:rppr}). In \\S\\ref{sec:sirtf} we calculate the dust IR flux densities integrated over the SIRTF/IRAC and MIPS bands predicted for our best-fitting models. Spectroscopic predictions are also made for SIRTF/IRS observations. The major conclusions are summarized in \\S\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} We have modelled the mid-infrared to submillimeter spectral energy distribution of the double-ring disk around $\\hda$ in terms of the porous dust model previously shown successful in reproducing the SEDs observed for the disks around $\\beta$ Pictoris (Li \\& Greenberg 1998) and $\\hra$ (Li \\& Lunine 2003). The dust is modelled as porous aggregates of either unaltered or heavily processed interstellar materials. While the former (``cold-coagulation'' dust) is made of amorphous silicate dust, carbonaceous dust, H$_2$O-dominated ices, and vacuum, the latter (``hot-nebula'' dust) consists of a mixture of crystalline silicates, ices, and vacuum (\\S\\ref{sec:model}). For both dust types, a vacuum volume fraction of 90\\% is assumed (\\S\\ref{sec:para}). The dust spatial distribution is taken to be that derived from imaging observations of scattered light and dust thermal emission (Eqs.[\\ref{eq:dndr1}-\\ref{eq:dndr3}]; \\S\\ref{sec:iremmethod}). We take a simple power-law for the dust size distribution ($dn/da\\propto a^{-\\alpha}$) in the size range of $1\\mum \\simlt a \\simlt 1\\cm$. In addition, a population of PAH molecules with a log-normal size distribution characterized by a peak-parameter $a_0$ and a width-parameter of $\\sigma$ is invoked to account for the observed 7.7$\\mum$ and 11.3$\\mum$ ``unidentified infrared'' emission features that are generally attributed to PAHs. Our principal results are: \\begin{enumerate} \\item The ``cold-coagulation'' dust model with a porosity of $P=0.90$, a size-distribution power index of $\\alpha\\approx 3.3$, and a total dust mass of $\\md \\approx 3.56\\mearth$, together with a population of charged PAHs with $a_0\\approx 2.5\\Angstrom$, $\\sigma\\approx 0.3$ and a total mass of $\\mpah\\approx 7.10\\times 10^{-6}\\mearth$ provides an excellent fit to the entire SED, including the PAH mid-IR emission features (\\S\\ref{sec:results}; Figure \\ref{fig:cold90}). \\item Although the $P=0.90$ ``hot-nebula'' dust model provides an overall good fit to the observed SED, it predicts a strong crystalline silicate emission feature at $21\\mum$ which is not seen in the currently available observational spectrum (\\S\\ref{sec:results}; Figure \\ref{fig:hot90}). This places an upper limit of $\\sim 10\\%$ on the fraction of ``hot-nebula'' dust (i.e. crystalline silicates; \\S\\ref{sec:results}; Figure \\ref{fig:coldhot90}). \\item The PAH molecules in the $\\hda$ disk are expected to be negatively charged (Figure \\ref{fig:ionrec}; \\S\\ref{sec:ionrec}; \\S\\ref{sec:photophys}), assuming that the cosmic ray ionization of H$_2$ is the dominant source of electrons (\\S\\ref{sec:ne}). But it has also been shown that models consisting of a mixture of both neutral and charged PAHs are also capable of reproducing the observed PAH IR emission features, provided that the fraction of neutral PAHs is smaller than 50\\% (Figure \\ref{fig:ionneupah}; \\S\\ref{sec:robust}). \\item Approximating the photodestruction rate of a PAH molecule by the photoejection rate of an acetylene molecule, PAHs smaller than $\\sim 4.6\\Angstrom$ will be photodestroyed in a time scale shorter than the age of the $\\hda$ system (Figure \\ref{fig:pahdes}; \\S\\ref{sec:pahdes}; \\S\\ref{sec:photophys}). In order to have a stable distribution of small PAHs ($3.5<a<4.6\\Angstrom$) in the $\\hda$ disk, we require a continuous replenishment of this component at a rate of $\\sim 1.11\\times 10^{-3}\\mearth/{\\rm yr}$. Such a high replenishment rate may not be essential, since (1) the photoejection of a C$_2$H$_2$ molecule does not necessarily lead to the complete destruction of a PAH molecule, and (2) the repair of a PAH molecule with a broken C-C bond through reacting with gas-phase C atoms and/ions is ignored in deriving the PAH destruction time scale. Furthermore, we have shown in \\S\\ref{sec:pahdes} (Figure \\ref{fig:cold90paha46}) that a model exclusively consisting of PAHs larger than $\\sim 4.6\\Angstrom$ also achieves a good fit to the observed SED. \\item Grains smaller than $\\sim 10-20\\mum$ will be radiatively expelled from the disk; grains at $r=55\\AU$ in the size range of $\\sim 10-20\\mum \\simlt a \\simlt 100-200\\mum$ will also be removed from the disk due to the Poynting-Robertson inward spiralling drag (at a {\\it closer} [{\\it further}] distance from the star, {\\it larger} [{\\it smaller}] grains will be removed). Collisions of planetesimals must continuously replenish the dust in the disk at a rate of $\\approx 7.88\\times 10^{-6}\\,m_\\oplus/{\\rm yr}$ (\\S\\ref{sec:rppr}). PAHs will also be rapidly removed from the disk by radiation pressure and, to a less degree, by Poynting-Robertson drag. A replenishment rate of $\\approx 3.28\\times 10^{-8}\\,m_\\oplus/{\\rm yr}$ is needed to maintain the PAH disk (\\S\\ref{sec:rppr}). \\item Spectroscopic and broadband photometric predictions are made for SIRTF observations (\\S\\ref{sec:sirtf}). \\end{enumerate}"
    },
    "0311/astro-ph0311593_arXiv.txt": {
        "abstract": "We present new sub-mm observations of Vega excess stars, and consistent modelling for all known Vega excess stars with sub-mm data. Our analysis uses dust grain models with realistic optical properties, with the aim of determining physical parameters of the unresolved disks from just their SEDs. For the resolved targets, we find that different objects require very different dust grain properties in order to simultaneously fit the image data and SED. Fomalhaut and Vega require solid dust grains, whilst HR4796 and HD141569 can only be fitted using porous grains. The older stars tend to have less porous grains than younger stars, which may indicate that collisions in the disks have reprocessed the initially fluffy grains into a more solid form. \\epseri appears to be deficient in small dust grains compared to our best fitting model. This may show that it is important to include all the factors which cause the size distribution to depart from a simple power law for grains close to the radiation pressure blowout limit. Alternatively, this discrepancy may be due to some external influence on the disk (e.g.\\ a planet).  When the model is applied to the unresolved targets, an estimate of the disk size can be made. However, the large diversity in dust composition for the resolved disks means that we cannot make a reliable assumption as to the composition of the grains in an unresolved disk, and there is corresponding uncertainty in the disk size. In addition, the poor fit for \\epseri shows that the model cannot always account for the SED even if the disk size is known. These two factors mean that it may not be possible to determine a disk's size without actually resolving it. ",
        "introduction": "In 1983, the \\iras satellite detected a large infrared excess from Vega ($\\alpha$ Lyr) during a routine calibration observation \\nocite{1984ApJ...278L..23A}({Aumann} {et~al.} 1984).  Soon, a similar excess was detected from three more stars (\\betapic, Fomalhaut and \\epseri). The excess could only be explained by a shell or ring of dust, heated to between 50 and 125~K by the central star. Quite by accident, \\iras had uncovered the first examples of solid material orbiting a normal main sequence star other that our own Sun.  The four stars mentioned above have become the prototypes for the `Vega excess' phenomenon \\nocite{1993prpl.conf.1253B,2000prpl.conf..639G}({Backman} \\& {Paresce} 1993; {Lagrange}, {Backman} \\&  {Artymowicz} 2000).  A systematic study of the \\iras all-sky survey shows that around 15~per cent of all main-sequence stars have an IR excess due to circumstellar dust \\nocite{1997Ap&SS.255..103D,1999ApJ...520..215F}({Dominik} \\& {The Hjhvega Consortium} 1997; {Fajardo-Acosta} {et~al.} 1999).  Coronographic and sub-mm imaging has shown the dust is typically in a ring or disk geometry \\nocite{1984Sci...226.1421S,1998Natur.392..788H}({Smith} \\& {Terrile} 1984; {Holland} {et~al.} 1998) and that these systems have much in common with our own solar system, i.e.\\ solid material in stable coplanar orbits distributed within a region roughly the size of the Kuiper belt.  Though the Vega phenomenon is common, the poor resolution of current far-IR and sub-mm instruments means that few Vega excess stars are close enough to be resolved. For the cases where the disks are resolved, detailed analysis using realistic models for the optical properties and size distribution of the dust can be performed. This type of work has been attempted for \\betapic \\nocite{1998A&A...331..291L}({Li} \\& {Greenberg} 1998a), HR4796 \\nocite{1999A&A...348..557A, 2003ApJ...590..368L}({Augereau} {et~al.} 1999b; {Li} \\& {Lunine} 2003a) and Fomalhaut \\nocite{2002MNRAS.334..589W}({Wyatt} \\& {Dent} 2002). These models use the observed spatial distribution of dust, and fit the SED by varying the dust composition and size distribution. However, as these objects have been analysed on a case-by-case basis, different models and assumptions are used, which makes it somewhat difficult to make direct comparisons. In addition, there are so few resolved targets that it is impossible to look for statistical trends between disk size, age, stellar mass, etc.  However, there are many unresolved targets with good SED information, and new instruments such as SIRTF and SOFIA will produce many more such objects. A method which could retrieve the physical parameters of a disk from just its SED would be invaluable. To accomplish this, the detailed properties of the dust grains in Vega excess disks must be fully understood.  We have tried to tackle these problems in two ways. Firstly, we have made sub-mm photometric observations of Vega excess stars, to help constrain their SEDs. Secondly, we have attempted to model all known Vega excess stars with sub-mm data, using a consistent model with realistic dust grains.  This allows us to directly compare the results of modelling for resolved targets, and use the knowledge of dust grain size and composition gained to estimate the disk size for the unresolved targets. ",
        "conclusions": "We have presented new observations and modelling of Vega excess disks with sub-mm data. Our observations expand the sample of Vega excess stars detected in the sub-mm from 5 to~10, with a further 4~objects which were not detected but have useful upper limits. We have fitted the observed SEDs with models based on realistic dust grain composition and size distribution. We find that dust grain composition varies significantly between different objects, with younger disks having more porous grains. For \\epseri, our model fails to fit the data unless there are less small grains in the system than expected from a collisional cascade. This discrepancy may be due to an inadequacy in our model of the size distribution, or the result of an external influence such as a planet. For the unresolved targets, disk size is estimated assuming solid dust grains, but given the observed diversity in dust composition and the problems fitting \\epseri, these size estimates may not be reliable.  \\vspace{7mm}  The James Clerk Maxwell Telescope is operated by the Joint Astronomy Centre on behalf of the UK Particle Physics and Astronomy Research Council. We wish to thank the staff at the JCMT for enabling these observations, and an anonymous referee for their helpful and positive comments."
    },
    "0311/astro-ph0311300_arXiv.txt": {
        "abstract": "We report the findings of a new search for RR Lyraes in an M31 halo field located 40 arcminutes from the nucleus of the galaxy along the minor axis.  We detected 37 variable stars, of which 24 are classified as RR Lyraes and the others are ambiguous.  Estimating a completeness fraction of $\\sim 24\\%$, we calculate that there are approximately 100 RR Lyraes in the field, which is consistent with what is expected from deep HST color-magnitude diagrams.  We calculate a mean magnitude of $\\mean{g} = 25.15 \\pm 0.03$, which we interpret to mean that the mean metallicity of RR Lyraes is significantly lower than that of the M31 halo as a whole.  The presence of ancient, metal-poor stars opens the possibility that, initially, the M31 halo appeared much like the Milky Way halo. ",
        "introduction": "The Andromeda galaxy (M31) provides a unique opportunity to study the structure and evolution of massive spirals, as it is the closest such system that can be studied from outside (i.e., line of sight depth is not a major problem).  As such, a comparison to our Galaxy allows us to address the question of variety in the evolutionary histories of massive spirals.  For example, van den Bergh (2000, 2003) has suggested that the chemical evolution of M31 was much more rapid than that of the Galaxy, and this might point to the origin of M31 as the early merger of two (or more) relatively massive metal-rich ancestral objects.  Variable stars provide a unique and powerful tool to inspect the stellar populations and and star formation histories of nearby galaxies (Mateo 1998; 2000; Saha 1999), and here we employ observations of variable stars in the M31 halo to provide new constraints on the evolution of M31.  \\citet{mou86} made the first study of the stellar populations in the M31 halo.  An important result of that study was that the M31 halo red giant branch is much redder than those of Galactic globular clusters.  This was interpreted as a sign that the mean metallicity was $\\mean{\\feh} = -0.6$, a result that has been confirmed by deeper imaging \\citep{dur01}.  Shortly after this work, a second milestone study of the M31 halo was carried out by Pritchet \\& van den Bergh (1987; hereafter PvdB), who imaged a field 40 arcminutes from the nucleus along the SE minor axis to search for variable stars.  They reported the discovery of 30 RR Lyraes, and assuming a completeness of 25\\% estimated that the field contained 120 such objects.  Compared with the number of red giants, this implied that the frequency of RR Lyraes in the M31 halo is comparable to that found in RR Lyrae-rich globular clusters.  As noted in their work, this frequency was surprisingly high given that the \\citet{mou86} metallicity was significantly higher than that measured in RR Lyrae-rich Galactic globular clusters.  The number density of RR Lyraes estimated by PvdB has been called into question by deep HST color-magnitude diagrams of the M31 halo, which show few stars in the region where one would expect RR Lyraes.  A recent ultra-deep ACS CMD \\citep{bro03} is the best such example, finding roughly an order of magnitude fewer stars in that part of the CMD than would be expected from the PdvB RR Lyrae frequency.  Several possibilities exist for the discrepancy in the RR Lyrae frequency.   The pointings used by PvdB and \\citet{bro03} were different, and findings of streams in the M31 halo \\citep{iba01} raise the possibility of position-dependent populations.  A second possibility is that accurate photometry of stars at $B = 25.7$ was at or beyond the capabilities of CCDs in the mid 1980s, and perhaps the earlier RR Lyrae experiment detected specious objects in addition to bona fide RR Lyraes.  To address this discrepancy with a straightforward approach, we repeated the PvdB RR Lyrae experiment and present our findings in this paper.  Section \\ref{sec_obs} describes our data and reduction procedures, and the following sections describe the number of RR Lyraes and their properties. ",
        "conclusions": "We have presented the results of a search for RR Lyraes in the M31 halo using observations taken over four nights on the WIYN 3.5m telescope.  Our $9.6 \\times 9.6$ arcminute field was selected to include the smaller field used by a previous study by \\citet{pri87}, which was 40 arcminutes southeast of M31's nucleus along the minor axis.  Applying a combination of PSF-fitting and image difference photometry techniques, we were able to detect 37 variable stars, including 24 RR Lyraes.  We estimate an completeness fraction of 24\\%, which implies that roughly 100 RR Lyraes are present in this field -- slightly more than 1 RR Lyrae per square arcminute.  This number is much less than that calculated by \\citet{pri87} ($\\sim 17$ per square arc minute), and is consistent with the frequency of RR Lyraes expected from deep HST color-magnitude diagrams \\citep{bro03}.  The mean magnitude of the RR Lyraes is $\\mean{g} = 25.15 \\pm 0.03$, which is consistent with the accepted M31 distance modulus of $(m-M) = 24.44$ only if RR Lyraes have a significantly lower metallicity than the mean value of $\\logz = -0.8$ \\citep{dur01}.  We interpret this as an indication that RR Lyraes are predominantly created by an old, low-metallicity population within the M31 halo (i.e. very similar to the population that appears to dominate the Milky Way halo) that accounts for at least 10\\% of the red giants."
    },
    "0311/astro-ph0311136_arXiv.txt": {
        "abstract": "It is a well-established fact that massive cosmological objects exhibit a ``geometrical bias'' that boosts their spatial correlations with respect to the underlying mass distribution.  Although this geometrical bias is a simple function of mass, this is only half of the story. We show using numerical simulations that objects that are in the midst of accreting material also exhibit a ``temporal bias,'' which further boosts their clustering far above geometrical bias levels. These results may help to resolve a discrepancy between spectroscopic and clustering mass estimates of Lyman Break Galaxies, a population of high-redshift galaxies that are caught in the act of forming large numbers of new stars. ",
        "introduction": "Large-scale structure in the Universe is believed to have originated from a primordial Gaussian random field of matter fluctuations, the product of quantum fluctuations that were shifted to larger spatial scales during cosmological inflation [{\\em e.g.\\ } 1].  Cosmic Microwave Background observations show that when the Universe was 100,000 years old, the gaseous component was extremely smooth, with temperature variations $\\sim 10^{-5}$.  These tiny inhomogeneities, in concert with inhomogeneities in the unseen massive dark-matter, were amplified through gravitational instability, eventually forming the galaxies, clusters, and other cosmological objects we see today.  While the precise details of structure formation are highly complex, the simplicity of the initial random field allows us to easily compute the overall distributions of cosmological objects.  Galaxies and galaxy clusters represent the peaks in the initial density distribution which accrete matter at the expense of the diffuse regions between them, and thus their number densities can be simply related to the number densities of the peaks of the initial random field.  This technique has been applied most cleanly to galaxy clusters, whose densities and evolution provide strong constraints on the overall matter density \\cite{ek96}.  Similarly, because the peaks in a random field are more clustered than the overall distribution, the spatial clustering of cosmological objects is stronger than the underlying mass distribution. Furthermore, this ``geometrical bias'' is a systematic function of the mass of these structures, an effect that has been well-studied analytically and numerically [{\\em e.g.\\ } 3, 4, 5].  Yet, this is only half of the story.  Here we conduct a detailed numerical simulation that shows that the spatial correlation function of objects that are in the midst of accreting substantial amounts of material is significantly enhanced over that of the general population.  This temporal bias causes them to mimic the properties of higher-mass structures, with important astrophysical implications as discussed below.  The structure of this work is as follows: In \\S 2 we describe our numerical simulation, discuss our group-finding algorithms, and develop a robust definition of accreting groups.  In \\S 3 we present our results for the spatial correlation functions of these samples, and in \\S 4 we discuss the astrophysical implications of our results. Further details of this study are given in \\cite{sc03}. ",
        "conclusions": ""
    },
    "0311/hep-ph0311365_arXiv.txt": {
        "abstract": "\\PRE{\\vspace*{.1in}} Black hole production in elementary particle collisions is among the most promising probes of large extra spacetime dimensions.  Studies of black holes at particle colliders have assumed that all of the incoming energy is captured in the resulting black hole.  We incorporate the inelasticity inherent in such processes and determine the prospects for discovering black holes in colliders and cosmic ray experiments, employing a dynamical model of Hawking evolution.  At the Large Hadron Collider, inelasticity reduces rates by factors of $10^3$ to $10^6$ in the accessible parameter space, moderating, but not eliminating, hopes for black hole discovery.  At the Pierre Auger Observatory, rates are suppressed by a factor of 10. We evaluate the impact of cosmic ray observations on collider prospects. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311466_arXiv.txt": {
        "abstract": "Cosmic rays with energies above $10^{18}$~eV are currently of considerable interest in astrophysics and are to be further studied in a number of projects which are either currently under construction or the subject of well-developed proposals.  This paper aims to discuss some of the physics of such particles in terms of current knowledge and information from particle astrophysics at other energies. ",
        "introduction": "Cosmic rays (CR) are the non-photon particles with which we learn about astrophysics.  Their composition ranges over the known nuclei (and antiprotons) to electrons (and positrons) and neutrinos of all flavours, and (perhaps) to exotic particles as yet unobserved in accelerator physics. Cosmic ray studies are complementary to photon astrophysics since many astrophysical photons are produced in processes, such as synchrotron emission, which involve charged cosmic ray particles.  In this review we concentrate on the highest energy cosmic rays, which are unaffected by heliospheric modulation, but are strongly affected by propagation effects through our galaxy. In the case of an extragalacic origin for the highest energy cosmic rays, their propagation through intergalactic photon and magnetic fields will have a profound influence on what we observe.  Radio astronomical studies probe cosmic rays at their sources and we are led to associate cosmic ray acceleration with energetic radio objects, pulsars, supernovae, AGNs etc.  The next generation of large radio projects (LOFAR, SKA etc.) will certainly add considerably to our understanding of cosmic ray sources, and may even offer opportunities to develop new detection techniques through the radio-frequency fields of the air showers of cosmic rays arriving at Earth (e.g.\\ Falcke and Gorham 2003).  The radio emission is dominated by cosmic ray electrons which rapidly lose energy in the emission processes.  The more massive nuclei do not have a clear radio signature but propagate with little energy loss and are readily observed at Earth.  They are believed to fill (at least) the volume of our galaxy. However, being charged particles, their paths are dictated by astrophysical magnetic fields and only at the highest energies do we expect magnetic deflections to be sufficiently small for directional astrophysics to be possible.  At modest energies (up to about $10^{14}$eV), cosmic rays are sufficiently plentiful that balloon experiments can have enough collecting area to directly study properties of the beam. Indeed, even satellite experiments extend up to 2~TeV/amu (for the CRN experiment on Spacelab 2, M\\\"uller et al 1991) which for iron is about $10^{14}$eV per nucleus.  At higher energies, we depend on techniques which investigate the cascades of particles (extensive air showers) resulting from the impact of cosmic ray particles on our atmospheric gas.  Since a single cosmic ray particle can produce millions (or billions) of secondary particles in this way, and those particles scatter in the atmosphere to hundreds of metres from the original cosmic ray trajectory, such studies offer uniquely efficient opportunities for studying fluxes down to levels of primary CR particles per square kilometre per century.  The cosmic ray cascades may be directly detected with groups of spaced large-area radiation detectors on the ground (ground arrays) or indirectly through collecting (with large telescopes) the photons which they produce, fluorescence light from atmospheric nitrogen or Cerenkov emission from the bulk atmospheric gas.  In the future, such light collectors may well be satellite-based.  The Ultra High Energy cosmic rays (UHE CR) which are the focus of this review are currently of particular interest.  Spectral data at energies above 1 EeV ($10^{18}$eV) and directional results, notably from the AGASA project, are very suggestive of fascinating, unexpected physics (Hayashida et al.\\ 1999).  Furthermore, this field of research is experimentally challenging and there is controversy in discrepancies between experimental data from experiments currently operational or recently discontinued.  A new era in the field is beginning with the commissioning of the Pierre Auger Observatory (Auger Collaboration 2001) which comprises a pair of 3000~km$^2$ arrays (one under construction in Mendoza Province, Argentina, and one planned for Utah, U.S.A) employing both particle and optical detectors, followed by large-scale Japanese projects, and possibly the space-based Extreme Universe Space Observatory project (EUSO is a Europe/Japan/US collaboration currently under Phase A study as an ESA mission with the goal of a three-year mission on the ISS starting in 2009, see http://www.euso-mission.org/).  At the present time, our ideas concerning cosmic rays at these highest energies are predominantly based on the idea that such particles are likely to be of extragalactic origin.  Their sources are presumed to be found in some extreme astrophysical environment such as the most energetic radio or gamma-ray sources.  We have been forced to these ideas by our failure to identify any models for galactic objects capable of accelerating particles to within one thousandth of the required energies, and also by the failure of our galactic source models to reproduce the directional isotropy of the observed beam.  Nonetheless, extragalactic scenarios have their own problems.  The intergalactic magnetic field has largely unknown properties.  It could be strong enough, with a structure which makes it impenetrable to particles from nearby clusters of galaxies in realistic periods of time.  Also, intracluster magnetic fields may limit the ability of particles even to leave galactic clusters.  This paper begins by presenting a brief overview of the observed properties (energy spectrum, arrival directions, and composition) of the cosmic ray beam.  There follows a physical discussion of the cosmic ray acceleration process and resulting limits to the maximum achievable energies.  Energy losses in the following propagation through astrophysical fields are fundamental in determining the observed beam properties and these are then discussed.  It is possible that the origin of the highest energy cosmic rays lies in exotic particles and that field of research is introduced.  Whatever the sources of these particles may be, their observed directional properties will be closely linked with the properties of the intervening magnetic fields.  A discussion of this key factor completes our review.  For an earlier review see Nagano and Watson (2000).     \\subsection{The Detection of Ultra High Energy Cosmic Rays}  Cosmic ray observatories record the arrival directions and distribution of energies of incident particles.  They also attempt to measure a remaining parameter, which is the mass composition of the primary cosmic ray particle.  It is in the latter respect that cosmic ray studies differ from others in astrophysics.  The mass composition is a key parameter in our astrophysical understanding since the charge on the particle (closely related to mass, for a nucleus), with the momentum (or energy) determines the propagation path of the particle.  At the highest energies, the cosmic ray flux is extremely low and the particle recording must be through processes which enable detection to be achieved at large distances from the path of the primary cosmic ray.  In practice, this involves making use of the particle cascade which results from the interaction of the primary cosmic ray particle with our atmosphere.  That cascade is initiated by a primary cosmic ray particle when its first atmospheric interaction occurs.  A cascade of secondary particles is then fed by degrading energy from the primary particle as it repeatedly interacts in its atmospheric passage. Those secondary particles cause energy to be deposited into the atmospheric gas, with some remaining energy reaching the ground. The cosmic ray detection process then consists in either a measurement of the passage of the energy as it is carried by particles through the atmosphere (through any emitted Cerenkov light or by any induced nitrogen fluorescence light) or in direct detections in radiation detectors of remaining particles which reach the ground.  The arrival direction of the primary particle is deduced from the direction of the path of the atmospheric cascade and characteristically has a resolution of a fraction of a degree. The primary particle energy is deduced by attempting to make sufficient measurements on the cascade such that a calorimetric accounting (to a few tens of percent) can be made of the various energy sinks.  This is most directly done through nitrogen fluorescence studies, but the signal is weak and the technique works best at the highest energies.  Mass composition measurements are the most difficult, and cannot be made on an individual cascade basis.  They depend on statistical studies of the cascade developments as a function of energy.  Crudely speaking, massive nuclei have a large cross section and interact early.  Their cascades also degrade in energy relatively rapidly. As a result, early developing cascades are signatures of 'heavy' primaries and late developing cascades are indicative of 'light' (probably proton) primaries.  Massive amounts of cascade modelling puts flesh on these arguments.  Nitrogen fluorescence studies were pioneered by the Fly's Eye experiment (Baltrusaitis et al.\\ 1985) and its successor, the High Resolution Fly's Eye (HiRes).  Many ground detector arrays have been used.  The largest one in fully developed use is the AGASA (Hayashida et al.\\ 1999) experiment in Japan.  Both of the HiRes and AGASA experiments are sufficiently large to probe energies in excess of $10^{20}$eV.  The next of these huge experiments will use both techniques.  The Pierre Auger Observatory (Auger Collaboration 2001) is designed to operate above $10^{19}$eV.  It will have a 3000 square kilometre collecting area instrumented with 1600 large-area particle detectors and will also have twenty-four 4~m diameter Schmidt optical telescopes for fluorescence detection.   \\subsection{The Cosmic Ray Energy Spectrum}  \\begin{figure} \\epsfig{file=fig_1.eps,width=120mm} \\caption{ The cosmic ray energy spectrum (dotted) as measured from the Earth (after Bhattacharjee and Sigl 2000).} \\label{fig:spectrum} \\end{figure}  The cosmic ray energy spectrum is shown in Fig.~\\ref{fig:spectrum}.  It is remarkable both in its range of energies and in its range of fluxes.  It covers over ten decades of energy and thirty decades of flux in a form close to a power law with an index of about -2.7.  Deviations from that power law are relatively small but are generally regarded as physically significant.  There is a steepening at about $10^{16}$eV, known as the knee, and a flattening at about $10^{18}$eV known as the ankle (e.g.\\ Abu-Zayyad et al.\\ 2001).  The knee is often argued to be associated with an energy limit of acceleration from supernova remnant sources, although it may well be related to a loss of ability for our galactic magnetic fields to retain (and build up internally) the cosmic ray flux (e.g.\\ Clay 2002).  The ankle is usually associated with the onset of a dominant, flatter, extra-galactic cosmic ray spectrum. It is important to note that, in this model, our galaxy produces particles with energies up to those of the ankle of the spectrum.  That is already above 1~EeV$\\equiv$$10^{18}$eV, and is well above energies which are easily accessible for any present galactic acceleration models.  A key region of the spectrum is its very highest energies.  The flux here is so low that the low event statistics in our observations to date leave us uncertain of the spectral structure above the key energy of $6\\times 10^{19}$eV.  Here there is a predicted spectral downturn, the Greisen-Zatsepin-Kuzmin (GZK) cut-off (Greisen 1966, Zatsepin and Kuzmin 1966) for particles which have travelled more than a few tens of Mpc, due to interactions with the 2.7K cosmic microwave background radiation (CMBR).  However, several experiments have reported CR events with energies above $10^{20}$~eV (Takeda et al.\\ 2003) with the highest energy event having 300~EeV (Bird et al.\\ 1995).  Very recent data from the two largest aperture high energy cosmic ray detectors are contradictory: AGASA (Takeda et al.\\ 2003) observes no GZK cut-off while HiRes (Abu-Zayyad et al. 2002) observes a cut-off consistent with the expected GZK cut-off.  A systematic over-estimation of energy of about 25\\% by AGASA or under-estimation of energy of about 25\\% by HiRes could account the discrepancy (Abu-Zayyad et al. 2002), but the continuation of the UHECR spectrum to energies well above $10^{20}$ eV is now far from certain.  Future measurements with Auger (Auger Collaboration 2001) should resolve this question.  If the spectrum does extend well beyond $10^{20}$ eV, determining the origin of these particles could have important implications for astrophysics, cosmology and particle physics.   \\subsection{Arrival Directions}   A key observation in cosmic ray astrophysics is the directional distribution of the particles.  That distribution will depend on any galactic magnetic fields and hence will be energy (rigidity) dependent.  However, with very limited exceptions, which are not individually statistically significant, there is no observed deviation from isotropy above the knee of the energy spectrum, and any anisotropies at lower energies are themselves very small (Smith and Clay 1997, Clay, McDonough and Smith 1997).  Recently, the AGASA experiment (Hayashida et al.\\ 1999) found a non-uniform distribution of arrival directions, suggestive of a source direction, in the energy range $10^{18.0}$eV to $10^{18.4}$eV.  That observation is potentially very important, particularly as there is supporting evidence in data from the SUGAR array (Bellido et al.\\ 2001).  However, neither of those observations on their own is clearly statistically significant. Still, those data are regarded by many as the possible beginning of a new era in cosmic ray astrophysics in which we can begin directional cosmic ray astronomy.  The possibility of having a source to observe may indeed open up new frontiers.   \\subsection{Composition of the UHECR}  The highest energy cosmic rays show no major differences in their air shower characteristics to cosmic rays at lower energies.  One would therefore expect the highest energy cosmic rays to be protons particularly if, as is most likely, they are extragalactic in origin.  However, it is still possible that they are not single nucleons.  Obvious candidates are heavier nuclei (e.g. Fe), $\\gamma$-rays and neutrinos.  Surprisingly, it is even more difficult to propagate nuclei than protons, because of the additional photonuclear disintegration (Tkaczyk et al. 1975, Puget et al. 1976, Karakula and Tkaczyk 1993, Elbert and Sommers 1995, Anchordoqui et al.\\ 1997, Stecker and Salamon 1999). The possibility that the 300 EeV Fly's Eye event is a $\\gamma$-ray has been discussed (Halzen et al. 1995) and, although not completely ruled out, the air shower development profile seems inconsistent with a $\\gamma$-ray primary.  Weakly interacting particles such as neutrinos will have no difficulty in propagating over extragalactic distances, of course.  This possibility has been considered, and generally discounted (Halzen et al. 1995, Elbert and Sommers 1995), mainly because of the relative unlikelihood of a neutrino interacting in the atmosphere, and the resulting great increase in the luminosity required of cosmic sources.   \\subsection{Cosmic Ray Sources}   Figure~\\ref{fig:hillas_plot} is a well known diagram first produced by Hillas (1984).  It reminds us that acceleration, associated with magnetic structures, requires the field and its dimensions to be sufficient to contain the accelerating particle through the acceleration process.  The lines simply give the magnetic field vs.\\ gyroradii $r_g$ for protons (solid) and iron nuclei (dashed), and this gives the minimum size for scattering centres moving at at speeds $\\sim c$.  More realistically, for scattering centres moving at speeds $v_{\\rm scat} < c$ the size would need to be a factor $c/v_{\\rm scat}$ larger.  This puts a limit on the product of the source field and its physical dimensions.  Strong fields with large-scale structure are attractive for acceleration to the highest energies.  The acceleration is thought to be likely to be associated with astrophysical shocks.  \\begin{figure} \\epsfig{file=fig2.eps,width=120mm} \\caption{ Proposed sites for cosmic ray acceleration related to their likely dimensions and magnetic field strength for scattering centres moving with $v_{\\rm scat}\\sim c$ (after Hillas 1984).  The lines represent plausible limits for $10^{20}$eV cosmic ray containment in the sources (solid line -- protons, dashed line -- iron nuclei) for scattering centres moving at at speeds $\\sim c$.  For scattering centres moving at speeds $v_{\\rm scat} < c$ the size would need to be a factor $c/v_{\\rm scat}$ larger.} \\label{fig:hillas_plot} \\end{figure}  One of the very few plausible acceleration sites may be associated with the radio lobes of powerful radio galaxies, either in the hot spots (Rachen and Biermann 1993) or possibly the cocoon or jet (Norman et al. 1995).  One-shot processes such as magnetic reconnection (e.g. in jets or accretion disks) comprise another possible class of sources (Haswell et al. 1992, Sorrell 1987).  Acceleration at the termination shock of the galactic wind from our Galaxy has also been suggested by Jokipii and Morfill (1985), but due to the lack of any statistically significant anisotropy associated with the Galaxy it is unlikely to be the explanation.  However, a re-evaluation of the world data set of cosmic rays has shown that there is a correlation of the arrival directions of cosmic rays above 40 EeV with the supergalactic plane (Stanev et al. 1995), lending support to an extragalactic origin above this energy, and in particular to models where ``local'' sources ($<100$ Mpc) would appear to cluster near the supergalactic plane.  Such a correlation would also be consistent with a Gamma Ray Burst (GRB) origin as two type Ic hypernovae (supernovae with broad absorption features indicating high velocity ejected material and a rather large explosion energy) have now been identified with GRB (SN 2003dh/GRB 030329 Kawabata et al.\\ 2003; SN 1998bw/GRB980425 Galama et al.\\ 1998).  The expanding fireball may have ultrarelativistic components (e.g.\\ $\\Gamma \\sim 300$) and this may lead to production of UHECR through relativistic shock acceleration (Vietri 1995) or some other process (see e.g.\\ Dermer 2002 for a discussion and references to earlier work).   Because of the resulting flat spectrum of particles (including gamma-rays and protons) extending up to GUT (grand unified theory) scale energies, topological defect models have been invoked to try to explain the UHE CR.  Propagation of the spectra of all particle species over cosmological distances is necessary to predict the cosmic ray and gamma-ray spectra expected at Earth. Propagation over cosmological distances to Earth (as would be the case in some topological defect origin models) results in potentially observable gamma-ray fluxes at GeV energies in addition to cosmic rays.  Massive relic particles on the other hand, would cluster in galaxy halos, including that of our Galaxy, and may give rise to anisotropic cosmic ray signals at ultra high energies.  The suggestion that neutron stars might accelerate cosmic rays followed soon after the discovery of pulsars (see, e.g.\\ Gold 1975 and references therein).  Voltages up to $\\sim 10^{12}$--$10^{15}$V (depending on pulsar period and magnetic field) are expected in a pulsar's magnetosphere.  These could accelerate nuclei with a resulting flat spectrum extending up to $\\sim$10$^{16}$~eV, and could possibly explain the knee in the cosmic ray spectrum (e.g.\\ Bednarek \\& Protheroe 2002).  The pulsar wind shock has been proposed as an acceleration site (e.g., Berezhko 1994, Bell and Lucek 1996) and might in principle accelerate particles to $10^{15}$--$10^{19}$eV.  Blasi et al.\\ (2000) suggest that the UHE CR events are due to iron nuclei accelerated from young strongly magnetized neutron stars through relativistic MHD winds of neutron stars whose initial spin periods are shorter than $\\sim 10$ ms.  More recently, the acceleration to ultra-high energies has been discussed in the context of fast re-connection in millisecond pulsars, but it appears that to reach $\\sim 10^{20}$~eV magnetic fields $\\sim 10^{12}$--$10^{15}$~G, and special geometries are required (de Gouveia Dal Pino and Lazarian, 2001).  In our opinion, while neutron stars may contribute to the observed cosmic ray spectrum it seems unlikely that they are responsible for the UHE CR, with the possible exception of transient shocks in pulsar winds of neutron stars formed during a hypernova explosion (i.e. a GRB as already discussed).  We shall delay a more detailed discussion of radio galaxies, active galactic nuclei and topological defects as sources of the UHE CR, and discuss next diffusive shock acceleration and cosmic ray propagation. ",
        "conclusions": "Cosmic ray astrophysics at the highest energies is entering a new era with major new facilities being proposed and developed.  The tiny flux of particles in this range of energies challenges our understanding of the astrophysics which we apply to many other studies.  We are led to examine processes in extreme regions such as the various components of AGN, to ask how particle acceleration may take place there, and to ask for detailed information on photon and magnetic fields throughout our local universe.  This challenging work now truly complements studies throughout photon astrophysics."
    },
    "0311/astro-ph0311185_arXiv.txt": {
        "abstract": "Using our semi-linear inversion method, we measure the mass profile of the lens galaxy in the Einstein ring system 0047$-$2808. The lens is modelled as a baryonic component following the observed light, embedded in a dark matter generalised NFW halo. The semi-linear method makes full use of the information content in the ring image. We determine an unevolved B-band mass to light ratio for the baryons of $3.05^{+0.45}_{-0.86} \\,h_{65}\\,M_{\\odot}/L_{B\\odot}$ (95\\% CL), accounting for 65\\% of the total projected mass within the radius of $1.16''$ traced by the ring. This result is obtained without need of dynamical measurements. The inner logarithmic slope of the halo is found to be $0.87^{+0.69}_{-0.61}$ (95\\% CL). We find that the halo is fairly well aligned with the light but has only half the ellipticity. ",
        "introduction": "Gravitational lensing provides a simple and well understood tool for the determination of galaxy mass profiles. It has the advantage that the deflection angle of a photon passing a massive object is independent of the dynamical state of the deflecting mass. Dynamical methods for measuring galaxy mass profiles rely on a number of assumptions that can be difficult to test, such as dynamical equilibrium, circular orbits of gaseous disks, and, frequently, spherical symmetry of the mass distribution.  In contrast to weak lensing, where a statistical detection of galaxy-galaxy lensing can place limits on mass profiles (eg. Hoekstra, Yee, \\& Gladders 2003), strong lensing involves the analysis of individual systems and allows smaller scales to be probed. In strong lens systems, multiple images of a background source are observed.  Once the redshift of the source and lens galaxy are known, mass models for the lens can be fit by comparing predicted image positions with those measured. Strong lens systems where the source has extended structure give rise to multiple arc-shaped images. These arcs trace out the Einstein ring which, as Kochanek, Keeton, \\& McLeod (2001) show, allows stronger constraints to be placed on the lens mass profile.  Unlike strong lens systems where the source is point-like, the analysis of Einstein rings requires modelling of the source surface brightness distribution. In this way, the properties of both the source and the lens are minimised to give the best fit to the observed ring.  In the analysis of Einstein rings, the way in which the source is modelled has a significant effect. Assuming a simple analytical surface brightness profile for the source can bias the lens model solution. Real sources typically have complicated structure and an over-simplified source forces the fitted lens model to compensate.  A non-parametric form, such as one in which the source surface brightness is pixelised, almost entirely removes this problem but creates a new one: The surface brightness in each source pixel must be adjusted in addition to the lens model parameters to find the minimum thus hugely slowing the procedure.  Wallington, Kochanek, \\& Narayan (1996) describe a method based on maximum entropy that uses a pixelised source surface brightness distribution. This has been applied by Wayth et al. (2003) to the Einstein ring 0047$-$2808. At the heart of the method are two nested minimisation cycles. For a given lens model stipulated by the outer cycle, the surface brightnesses in the source pixels are adjusted in the inner cycle until an image is obtained which gives the best fit to the observed ring image. The inner cycle is a slow non-linear search for the minimum requiring many iteration to reach convergence.  In a previous paper (Warren \\& Dye 2003), we developed a new technique called the semi-linear method. The semi-linear method uses a pixelised source plane but replaces the inner cycle of the maximum entropy method with a single-step linear inversion. This has the advantage that it is very fast and, more importantly, that the correct source minimum is guaranteed for a given lens model. The possibility of minimising to an incorrect source, thereby biasing the minimised lens model, is thus completely removed.  In the work presented here, we apply the semi-linear method to WFPC2 observations of the Einstein ring 0047$-$2808. We model the lens as a baryonic component embedded in a dark halo and show how the contribution from each component can be determined to allow measurement of the baryonic M/L. ",
        "conclusions": "We have applied our semi-linear method to HST observations of the Einstein ring system 0047$-$2808 to determine the lens galaxy mass profile. Modelling the lens as a baryonic Sersic profile + point source nested in a dark matter generalised NFW halo, we find a M/L of $3.05^{+0.45}_{-0.86} \\,h_{65}\\,M_{\\odot}/L_{B\\odot}$ and an inner slope $\\gamma=0.87^{+0.69}_{-0.61}$ (95\\% CL). This M/L was obtained without any dynamical measurements. The reconstructed source surface brightness distribution shows two distinct source objects, highlighting the need for non-parametric sources to obtain unbiased lens mass profiles.  One of the main predictions of the Cold Dark Matter (CDM) model is that halo mass profiles are cuspy. Recent CDM simulations show that the inner slope should lie somewhere in the range $1 \\simless \\gamma \\simless 1.5$ (Navarro, Frenk, \\& White 1996; Moore et al. 1999). Our measurement of $\\gamma$ for 0047$-$2808 is not inconsistent with these.  Nevertheless there are two important factors which must be considered. The first is that the simulations are only just starting to reliably probe scales of $\\sim$ few kpc. The most recent simulations seem to indicate that $\\gamma$ continues to decrease at smaller radii with no sign of reaching a convergent value (Power et al. 2003; see also Navarro et al. in these proceedings). For complete reassurance in comparing these predictions with observational constraints, higher resolution simulations are needed.  The second complication is that the majority of simulations model only pure dark matter distributions.  The effect of a baryonic potential well alters the shape of the dark matter halo in a non-trivial way. Clearly, these shortcomings must be addressed before CDM can be tested with confidence."
    },
    "0311/astro-ph0311051_arXiv.txt": {
        "abstract": "{ We analyse the $L_x-\\sigma_v$ relation for the new Mulchaey et al. group Atlas. We find that once we take into account the possible statistical bias introduced by the cutoff in luminosity, we recover a relation which is consistent with that of clusters, ie., $L_x \\propto \\sigma^{4.1}$. The larger scatter of this relation for groups of galaxies could be attributed to an orientation effect, due to which the radial velocity dispersion of groups oriented close to orthogonal to the line of sight, would be underestimated. This effect could also contribute in the direction of flattening the slope of the group $L_x-\\sigma_v$ relation. ",
        "introduction": "Most galaxies in the universe occur in small groups (cf. Geller \\& Huchra 1983; Tully 1987; Nolthenius \\& White 1987), which in many respects could be considered as poor clusters of galaxies. The detection of X-ray emission from some groups has increased considerably the interest to study them since they proved to be real entities and not projection effects (for an extensive review see Mulchaey 2000). Solinger \\& Tucker (1972) showed that if the source of the X-ray emission is hot gas bound in clusters, then the X-ray luminosity, $L_{x}$, should be correlated with the optical radial velocity dispersion, $\\sigma_v$. Simple theoretical arguments show that in a virialized, isothermal aggregation of gas, which emits thermal bremsstrahlung emission ($L_{x} \\propto \\int \\rho_{gas}^2 T(r)^{1/2} {\\rm d}V$), we can obtain that $L_x \\propto M^{4/3}$ (where $M$ is the total of the system) and using virial arguments ($M \\propto \\sigma^3$) we then have that $L_x$ should be roughly proportional to the fourth power of $\\sigma$: $\\log L_{x} \\propto \\log \\sigma_v^{4}$ (cf. Navaro, Frenk \\& White 1995). Quintana \\& Melnick (1982) first showed that the X-ray luminosity of clusters of galaxies obey the expected correlation.  Numerical simulations have shown that the relationship between $L_x$ and $\\sigma_v$ for groups should be similar to that of clusters (cf. Navarro, Frenk, \\& White 1997), or even steeper if one takes into account radiative cooling which significantly reduces the amount of the hot gas fraction at low-$\\sigma_v$'s (Dav\\'e, Katz \\& Weinberg 2002). Such a steep relation for systems with $L_x<10^{43}$ ergs s$^{-1}$, has been advocated by Mahdavi \\& Geller (2001), although they notice an erratic behavior of the poor groups of galaxies.  In some observational studies a consistency has been found between the $L_x-\\sigma_v$ relation of groups and clusters with a slope $\\sim 4$ (cf. Ponman et al 1996; Mulchaey \\& Zabludoff 1998; Helsdon \\& Ponman 2000), while in others a shallower slope has been found, i.e., groups appear to have a relatively enhanced X-ray emission to what predicted from the $L_x - \\sigma_v$ relation deduced from clusters of galaxies (cf. Dell`Antonio et al. 1994; Mahdavi et al. 1997, 2000; Xue \\& Wu 2000). Attempts to explain such deviations of the group from the cluster $L_x-\\sigma_v$ behaviour have invoked a possible contribution of individual galaxy halos to the group X-ray luminosity (Mahdavi et al. 2000), poorly determined $\\sigma_v$'s and/or $L_x$'s (Zimer, Mulchaey \\& Zabludoff 2001), a large scatter due to an non-equilibrium galaxy velocity distribution (Mahdavi \\& Geller 2001).  An alternative explanation, based on a possible orientation effect, was proposed by Tovmassian, Tiersch, \\& Yam (2002) [see also Tovmassian, Martinez, \\& Tiersch 1999; Tovmassian 2002]. They considered the relatively small RASSCALS and HCG X-ray group samples (Mahdavi et al. 2000; Ponman et al. 1996, respectively) and showed that the flattening of groups with relatively small velocity dispersions is, on average, larger than those of groups with higher velocity dispersions. They suggested that the shallow shape of $\\log L_{x}-\\log \\sigma_v$ of groups of galaxies could be partly the result of an underestimation of the velocity dispersion of elongated groups when seen roughly orthogonally to the line of sight, while when seen edge-on, they will have higher and probably nearer to their true $\\sigma_v$ values. This correlation could be explained if member galaxies in groups move preferentially along the group elongation, since groups have been found to have a prolate-like shape (Hickson et al. 1984; Malykh \\& Orlov 1986; Oleak et al. 1998) as in the case of clusters (cf. Plionis, Barrow \\& Frenk 1991).  Recently Mulchaey et al. (2003) published an X-ray Atlas of groups of galaxies, which is the largest sample of groups studied to date having X-ray ROSAT PSPC pointed observations. In this paper we address two questions: is there sufficient evidence to support claims that groups have enhanced, with respect to clusters, X-ray emission (ie., that the power law index of the $\\log L_x - \\log\\sigma^\\alpha$ relation has $\\alpha <<4$) and what is the origin of the larger scatter of the $L_x - \\sigma$ relation for groups.  \\section {Data and results} In order to address the issue of the possibly enhanced group X-ray emission we have selected from the Atlas of Mulchaey et al (2003) those with detected X-ray emission that contain less than 20 members (since a larger membership should rather define poor clusters). Note that the $L_{x}$ values were determined after the removal of point sources. We exclude three groups: NGC 2484 and NGC 6251 because they consist of only two members and the NGC 6329 group due to its very large projected length $a$ ($\\approx~11$Mpc)\\footnote{Lengths are determined assuming $H_{o}$ = 72 km s$^{-1}$Mpc$^{-1}$.}, which is suspected that it could be a superposition of two or even more groups and hence, its $\\sigma_v$ value may be unreliable. We are left with 43 groups in total out of which we have three triplets. Although, Focardi \\& Kelm (2002) argue that triplets are a distinct class of low-$\\sigma_v$ objects, in our case they span the whole $\\sigma_v$ range.  In Fig. 1 we plot $\\log L_{x}$ versus $\\log \\sigma_v$ for the considered groups (open symbols) and for groups with members $N>20$ (squares), which could be considered as poor clusters. The latter are closely located along the line $L_{x} \\propto \\sigma_v^{4}$ (solid line).  \\noindent We perform a direct least-square regression to the $N<20$ groups of the type \\begin{equation}\\label{eq:1} \\log L_x = \\alpha \\log\\sigma + C_1 \\end{equation} and find a strong correlation with a correlation coefficient $R=0.65$ and a probability of random occurrence $3\\times 10^{-6}$. The fitted parameters are: $$ \\alpha=1.72 \\pm 0.31 \\;\\;\\;\\;\\;\\;\\;\\ C_1 = 37.6 \\pm 0.75 \\;.$$ The determined slope is much shallower than the expected value for clusters $\\alpha\\simeq 4$, while the value found from the $N>20$ groups is $\\alpha \\simeq 3.2$. As discussed in the introduction, the shallower slope of the group $L_x-\\sigma_v$ relation already been noticed in other studies as well and indeed it has stimulated speculations on the reason for the apparently enhanced X-ray emission of groups. Hence, we could have concluded from our results that this sample supports the claims for a relatively enhanced X-ray emission for groups. \\begin{figure}[t] \\includegraphics[width=9cm]{plio_fig1.ps} \\caption{$L_x - \\sigma_v$ relation for groups: those with $N\\le20$ are represented as circles, while those with $N>20$, supposed to be poor clusters, are marked by filled squares. The solid line is the $\\log L_{x} \\propto \\log \\sigma_v^{4}$ relation of clusters of galaxies. The dashed line is the best direct regression fit for the groups ($\\log L_{x} \\propto \\log \\sigma_v^{1.7}$) and the dot-dashed line is is the best inverse regression fit for the groups ($\\log \\sigma \\propto (1/4.1) \\log L_{x}$).} \\end{figure}  \\subsection{Statistical bias ?} However, what we are witnessing is the result of a statistical bias, resembling the Malmquist bias, which appears because of the low-$L_x$ limit (either due to the lower mass limit necessary for the ICM to light up, or due to the X-ray flux limit in the construction of the sample). This bias is of the same nature that enters in scaling relation (eg. Tully-Fisher, Faber-Jackson, etc) where the magnitude or flux limit imposes a bias such that the slope of the derived relation is shallower than the nominal one. The larger the scatter of the relation, the larger the bias imposed. This bias may or may not appear in the corresponding cluster relation depending on the amplitude of the scatter around the nominal relation. To understand this bias let us imagine that the cluster relation $L_x \\propto \\sigma^4$ is obeyed by groups as well and let us consider a subset of our sample of groups that have $\\sigma_v \\simeq 600$ km s$^{-1}$ so that they are typically quite brighter than the lower $L_x$ limit. It is clear that their $\\log L_x$ value will be distributed around the mean $\\langle \\log L_x \\rangle$ value with some dispersion $\\sigma$. However, if we take a subsample with $\\sigma_v \\simeq 100$ km s$^{-1}$ so that they are on average quite faint with $\\langle \\log L_x \\rangle \\simeq \\log (L_x)_{limit}$, then the only groups that will appear in the sample are those with $\\log L_x > \\langle \\log L_x \\rangle$ and none with $\\log L_x < \\langle \\log L_x \\rangle$. This will induce the above mentioned bias.  To see this more clearly, we have performed Monte-Carlo simulations in which we assume a relation $L_x \\propto \\sigma_v^{4}$, and a Gaussian scatter of $\\delta \\log L_x = 0.8$. We then perform a direct (ie., eq.\\ref{eq:1}) and an inverse least-square regression fit to the resulting data (ie., we fit $\\log\\sigma_v = (1/\\alpha) \\log L_x + C_2$). In Fig. 2a we show such a simulation with 5000 ``groups'', where the solid line is the input $L_x \\propto \\sigma^4$ relation and the dashed line is the recovered from the direct regression. The inverse regression recovers exactly the input slope.  Furthermore, to study more accurately the effect of this statistical bias on our sample we have performed a series of simulations in which we use the observed $L_x$ values of the groups and we derive the values $\\sigma_v$, assuming a Gaussian scatter of the $\\log L_x - \\log\\sigma_v$ relation of varying magnitude. We have performed 1000 such simulations for each different value of the Gaussian scatter. In Fig. 2b we plot the derived values of the slope $\\alpha$ for both regression methods as a function input scatter. The vertical line indicates the scatter of the relation derived from the $N\\le 20$ sample of groups. The direct regression method (star-like symbols) underestimates severely the input slope, with the underestimation increasing with increasing scatter, while the inverse regression method (open symbols) recovers it accurately. Note also that we plot (squares) the results of the inverse regression in case that there is a sharp cutoff at $\\sigma_v=1000$ km s$^{-1}$ , in which case a similar, although of smaller magnitude, bias is introduced in the opposite direction.  It is evident that: \\begin{itemize} \\item The inverse regression recovers correctly the input slope of the $\\log L_x - \\log \\sigma_v$ relation for all values of the scatter. \\item for the observed amount of scatter ($\\delta C_1 \\simeq 0.8$) the expected slope of the relation once we use the direct regression is around $\\sim 2$, as indeed observed. \\end{itemize}  \\begin{figure*}[t] \\includegraphics[width=9cm]{plio_fig2a.ps} \\hfill \\includegraphics[width=9cm]{plio_fig2b.ps} \\caption{(a) Manifestation of the statistical bias for 5000 ``groups'' that have been drawn from a $L_x \\propto \\sigma_v^4$ relation (solid line), which have a scatter of 0.8 in $\\log L_x$. The recovered direct regression line is the dashed one. (b) Recovered values of the $L_x - \\sigma_v$ slope ($\\alpha$) for different amounts of scatter of the relation for the direct regression ($\\log L_x \\propto \\alpha \\log \\sigma_v$; star symbols), for the inverse regression ($\\log \\sigma_v \\propto (1/\\alpha) \\log L_x$; open symbols) and for the inverse regression once we have imposed an upper limit in $\\sigma_v=1000$ km s$^{-1}$.} \\end{figure*}  As a further test of this bias, we have re-analysed the group data of Xue \\& Wu (2000). Their sample consists of 60 groups that have velocity dispersion and X-ray luminosity data. Performing the direct regression fit we recover their results (listed in their table 3 - OLS method), ie. $\\alpha\\simeq 1$ and $C_1 \\simeq 40$. However, if we perform an inverse regression fit to their data we recover a very different slope, ie., $\\alpha \\simeq 6.7$ and $C_1 \\simeq 25.7$. In Fig.3 we plot the $L_x - \\sigma$ correlation for this sample and the fitted lines for both regression methods. Using our Monte-Carlo procedure we have seen that such a dichotomy between the direct and inverse regression methods can be approximately accommodated if there is a scatter of $\\sim 1$ in $\\log L_x$ and a cutoff at $\\sigma\\simeq 600$ km s$^{-1}$. \\begin{figure}[t] \\includegraphics[width=9cm]{plio_fig3.ps} \\caption{Results of the Xue \\& Wu (2000) sample of groups.} \\end{figure}  Guided by our Monte-Carlo analysis we performed an inverse regression to our sample of $N\\le 20$ groups, and we have indeed recovered a slope $$\\alpha \\simeq 4.1 \\pm 0.6 \\;.$$ We therefore conclude that this sample of X-ray groups is absolutely consistent with the relation found from clusters of galaxies.  \\subsection{The scatter in the $L_x - \\sigma$ relation} Now we move to address the issue of the larger scatter of the $L_x - \\sigma_v$ relation of groups than of clusters. This is an important issue because, as we have seen previously, the larger scatter will induce an artificial flattening of the $L_x - \\sigma_v$ relation.  We will show below that the larger scatter is not only statistical, due to possible measurement uncertainties especially in the low-$L_x$ and $\\sigma_v$ regime, but has a probably large intrinsic component which is due to an orientation effect. To guide the reader through our arguments we need first to define the axis ratio $b/a$ of the groups studied \\footnote{$a$ is the angular distance between the most widely separated galaxies in the group, and $b$ is the angular distances $b$ of the third galaxy of the group consisting of three galaxies from the line $a$ joining the most separated two galaxies, or is the sum of the angular distances $b_{1}$ and $b_{2}$ of the most distant galaxies on either side of the line $a$ joining the most separated galaxies (Rood 1979).}. The ratio $b/a$ was determined by using the positions of member galaxies which are mentioned in the corresponding references of Table 1 in Mulchaey et al. (2003). For those groups that Mulchaey et al. (2003) used the NASA Extragalactic Database (NED), we selected the members of the corresponding groups by their redshifts taking into account the membership number, mentioned in Mulchaey et al. (2003). Note that that member galaxies are not drawn from a well-defined and uniform magnitude limited sample. This can introduce differences in the depth coverage within groups, but such an effect would probably introduce a random error in the determination of the group elongation. In Table 1 we present $b/a$ values for those groups with number of members $N<20$. This is the sample we will investigate in detail.  \\begin{table*}[] \\hfill \\caption[]{The $b/a$ ratios and the number of galaxy members for groups with $2<N<20$.} \\tabcolsep 8pt \\begin{tabular}{|lrr|lrr|lrr|lrr|} \\hline Group & N & $b/a$ & Group & N & $b/a$ & Group & N & $b/a$ & Group & N & $b/a$ \\\\ \\hline NGC 315  & 4 &0.53 & NGC 1407 & 8&0.36  & NGC 4065&7&0.38 & HCG 68 & 5& 0.53 \\\\ NGC 326  & 9 &0.56 & NGC 1587 & 4&0.03  & NGC 4104&8&0.66 & ARP 330 &8& 0.44 \\\\ NGC 524  & 8 &0.54 & NGC 2300 &13&0.58  & NGC 4125&3&0.44 & NGC 6338&11& 0.46 \\\\ HCG 12   & 5 &0.63 & HCG 37 &5 &0.39    & NGC 4261&8&0.26 & HCG 90 &16& 0.67 \\\\ NGC 720  & 4 &0.79 & HCG 48 & 3&0.08    & SHK 202&5&0.75 & UGC 12064&9& 0.79 \\\\ HCG 15   & 6 &0.49 & CGCG 154-041&4&0.13& NGC 4291&11&0.42 & HCG 92&4& 0.14 \\\\ HCG 16   & 9 &0.34 & NGC 3607 &7  &0.33 & NGC 4636&12&0.79 & IC 1459&5& 0.22 \\\\ UGC 1651 & 3 &0.02 & NGC 3647 & 6&0.24  & NGC 5044&9&0.63 & NGC 7619&7& 0.51 \\\\ NGC 1044 &13 &0.50 & NGC 3665 & 4 &0.19 & NGC 5171&15&0.26 & HCG 97 &14& 0.62 \\\\ IC 1880  & 7 &0.48 & HCG 57 & 7 &0.34   & HCG 67 &14 &0.73 & NGC 7777&4 & 0.63 \\\\ UGC 2775 & 5&0.39 & NGC 3923 & 5 &0.06  & NGC 5322& 8&0.36  &          &     \\\\ \\hline \\end{tabular} \\end{table*}  We have noticed that there is a fairly significant correlation ($R=0.37$ and ${\\cal P}=0.025$) between flattening and velocity dispersion, with flatter systems having lower velocity dispersions (see Fig. 4). \\begin{figure}[t] \\includegraphics[width=9cm]{plio_fig4.ps} \\caption{Group axial ratio verus group $\\log \\sigma_v$.} \\end{figure} Taking the 10 groups with the lowest and the highest values of $\\sigma_v$ ($\\log\\sigma_v \\le 2.07$ and $\\ge 2.6$, respectively) we find that the median and 68\\% quantiles of their $b/a$ ratio is $0.22^{+0.04}_{-0.16}$ and $0.49^{+0.07}_{-0.10}$, respectively.  This difference in the group flattening is also accompanied by a significant difference in the number of group members, with mean values of $\\sim 5.2$ and 9.2 respectively. This could be expected either (a) due to the number of galaxies - mass correlation (higher number of galaxies implies higher mass which then implies a higher velocity dispersion) or (b) due to the fact that sparser groups are more likely, merely by chance, to show higher elongations than more dense systems. The latter possibility has been ruled out after performing a large set of Monte-Carlo simulations in which we have generated the same number of groups and group members, as in the two observed subsamples, by randomly placing group members in a sphere and then projecting them to a plane. The resulting median axial ratio for the low and high-$\\sigma$ subsamples was found to be: $\\sim 0.62 \\pm 0.18$ and $\\sim 0.7 \\pm 0.15$, respectively, significantly larger than the observed values. However, neither the former possibility can explain the significant difference in the flattening of the two extreme subsamples of groups. Why should poorer groups be flatter than richer ones ? A possible explanation could be the different level of group virialization. If groups accrete material anisotropically along one dimensional structures, like filaments, then one may expect that flatter systems are dynamically younger (thus they have lower values of $\\sigma_v$ and $L_x$) while when virialization takes place, after the groups have accreted enough material, it will drive groups to more spherical configurations, with higher velocity dispersions.  This straight-forward explanation cannot account however for all the observables. For example, selecting only the groups with $N\\le 5$ members (in total 10 groups), which as expected from the above discussion should be very flat (indeed their median $b/a$ is $0.14^{+0.05}_{-0.06}$), we find that although in this subsample there is no correlation of $\\sigma_v$ with the number of group members, there is a significant correlation between their flattening and $\\sigma_v$ (correlation coefficient R=0.8 and ${\\cal P}_{random} = 0.006$), with low-$\\sigma_v$ groups being flatter. In the virialization paradigm discussed above such a correlation cannot be explained. In order to explain such correlations we suggest that an {\\em orientation effect} is at work, i.e., if galaxy members move along their group major axis (for example, infalling in their common center of mass or rotating in elongated orbits around their gravitational center; cf. Tovmassian 2001, 2002), then flat groups oriented close to orthogonally to the line of sight (small $b/a$) will have small values of $\\sigma_v$, while the opposite is true for flat groups seen edge-on.  According to this view the length of the apparent major axis, $a$, of the flattest groups (small $b/a$ ratio) should also depend on orientation. Groups oriented close to the line of sight will have on average small $a$ and high $\\sigma_v$, and groups oriented close to the orthogonal to the line of sight will have large $a$ and small $\\sigma_v$. Tovmassian (2002) showed that such anti-correlation between the $a$ and $\\sigma_v$ is observed in the case of HCGs groups. In order to investigate this possibility in our sample, we select the 11 groups with $b/a<0.35$. In Fig. 5 we plot their apparent value of $a$ versus $\\sigma_v$. The expected trend is indeed apparent, with $a$ increasing as $\\sigma_v$ decreases, and has a probability of occurring randomly of ${\\cal P}=0.07$. Performing a bootstrap resampling technique in order to estimate the uncertainty of the above significance level, we find $\\delta {\\cal P}=0.02$.  The existence of the orientation effect does not imply that the virialization arguments, discussed previously, are incorrect. Most probably both are at work since the orientation effect is not apparent in groups with large number of members, exactly because these groups are richer and most probably in a more advanced dynamical state.  \\begin{figure}[t] \\includegraphics[width=9cm]{plio_fig5.ps} \\caption{The flat ($b/a<0.35$) group $a$ - log $\\sigma_v$ correlation.} \\end{figure} ",
        "conclusions": "We have investigated the X-ray luminosity - velocity dispersion relation in the new group Atlas of Mulchaey et al. (2003). A direct regression shows that there is a strong correlation with a slope ($\\alpha \\simeq 1.7$) significantly shallower than that found in clusters of galaxies. However, we attribute this to a statistical bias, resembling the Malmquist bias, that enters in the direct regression approach and which is due to a limit in $L_x$. We have quantified this using Monte-Carlo simulations and once we use the more accurate inverse regression method we obtain a slope of $\\sim 4.1$, absolutely consistent with the cluster relation. We have also investigated the apparently larger scatter of the $L_x - \\sigma$ relation of groups with respect to clusters. We find that at least part of this scatter is intrinsic in nature and due to an orientation effect by which flat groups seen edge-on have their $\\sigma_v$ values underestimated. This can be understood by noting that groups of galaxies have a roughly prolate spheroidal shape and thus if member galaxies move along their major axis, being accreted for example to their common centre of mass, then the correlation between $\\sigma_v$ and their major axis, could be due to the different group orientations with respect to the line of sight. Flat groups that are oriented roughly orthogonally to the line of sight will have low values of $\\sigma_v$ , while when oriented close to the line of sight or at intermediate angles, they will have higher values of $\\sigma_v$."
    },
    "0311/astro-ph0311098_arXiv.txt": {
        "abstract": "{The environmental influence of cluster media on its member galaxies, known as Butcher--Oemler effect, has recently been subject to revision due to numerous observations of strong morphological transformations occurring outside the cluster virial radii, caused by some unidentified gas removal processes. In this context we present new XMM-Newton observations of M86 group.  The unique combination of high spatial and spectral resolution and large field of view of XMM-Newton allows an in-depth investigation of the processes involved in the spectacular disruption of this object.  We identify a possible shock with Mach number of $\\sim 1.4$ in the process of crushing the galaxy in the North-East direction. The latter is ascribed to the presence of a dense X-ray emitting filament, previously revealed in the RASS data. The shock is not associated with other previously identified features of M86 X-ray emission, such as the plume, the north-eastern arm and the southern extension, which are found to have low entropy, similar to the inner 2 kpc of M86.  Finally, mere existence of the large scale gas halo around the M86 group, suggests that the disruptions of M86's X-ray halo may be caused by small-scale types of interactions such as galaxy-galaxy collisions. ",
        "introduction": "X-ray studies of galaxies in nearby clusters provide important clues on environmental effects on the properties of galaxies. X-ray luminous early-type galaxies appear often surrounded by group-size mass concentrations, whose interaction with the cluster environment is not yet completely understood from both observational and modeling sides. In fact, on one hand, recent observations find quite a bit of substructure in merging systems, demonstrating the survival of the core of the accreting objects (e.g. Vikhlinin et al. 2001).  On the other hand, detailed studies of cluster galaxies show that strong interactions take place among the infalling galaxies at the virial radius (Poggianti et al. 1999; Goto et al. 2003). In addition, recent studies on the evolution of galaxies prior to their infall onto clusters reveal strong effects of local environment on the appearance of galaxies (Kodama et al. 2001). Therefore, the morphology of galaxies appear to be disturbed in the so called quiet accretion mode of cluster formation.  M86 is a prototype for this type of investigation. It is a giant elliptical galaxy falling at high speed toward M87 (relative velocity $\\sim 1500$ km s$^{-1}$). Furthermore it sits in the core of a larger structure, M86 groups of galaxies, which exhibits a diffuse X-ray emission on a few hundreds kpc scale (B\\\"ohringer et al. 1994). An X-ray study of M86 with Einstein (Forman et al. 1979) revealed a peak of emission centered on M86 and a plume extending northwest of the galaxy. Ram-pressure stripping of the hot gas from M86 as the galaxy traverses the Virgo cluster was suggested to explain the observed structure (Forman et al. 1979; Fabian, Schwartz, \\& Forman 1980; Takeda, Nulsen, \\& Fabian 1984).  Spatial analysis of the X-ray brightness of M86 from the ROSAT PSPC and HRI data was presented in Rangarajan et al. (1995). In addition to the plume, they identified a southern extension (extending from the center by about $3^\\prime$ to the south), a void (north to the galaxy), and a northeastern arm (extending from the center by more than $5^\\prime$ to the northeast). Detailed spectral analysis of M86 emission using ROSAT PSPC and ASCA SIS results was presented in Finoguenov \\& Jones (2000). The temperature variations were found to lie in the 0.6--1.1 keV range. Cool regions include the galaxy center and extend eastward with slightly increasing temperature. In the plume, spectral variations were detected on the 0.1 keV level.   \\begin{figure*} \\includegraphics[width=5.9cm]{f1a.ps}\\includegraphics[width=5.9cm]{f1b.ps}\\includegraphics[width=5.9cm]{f1c.ps}  \\figcaption{{\\it Left panel.} Identification of the previously reported features of M86 on the smallest scale of the wavelet decomposition of the image in the 0.5--2 keV band. {\\it Middle panel.} Wavelet-decomposed image of M86 with contours indicating the selection of regions for the spectral analysis. Black contours denote the regions excluded from the analysis due to either low statistics or a presence of the point source. Numbers on the figure indicate the approximate location of the larger zones, reported in Tab.\\ref{t:mo}. {\\it Right panel.} Hardness ratio map with surface brightness contours overlaid. Dark grey color corresponds to temperature of 1.1--1.2 keV, grey to 0.9--1 keV, light grey to 0.7--0.9 keV, white to 0.5--0.7 keV. \\label{imh} } \\end{figure*}  As it became evident from the temperature and detailed surface brightness maps, M86 does not fit into the simple picture suggested by early observations.  For this reason Rangarajan et al. (1995) proposed that M86 owes its appearance to the presence of a conical shock seen in projection and due to the galaxy interaction with Virgo intracluster medium (ICM).  In this paper we present new XMM-Newton observations of M86 galaxy. By taking advantage of the unique combination of high spatial and spectral resolution of XMM-Newton, we study the thermodynamic properties of a total of 120 regions comprising the environment in and around M86.  Based on these, we were able to reveal the remarkable presence a of a weak shock in the process of crushing the galaxy interstellar medium (ISM).  In view of a $2.4\\pm 1.4$ Mpc distance of M86 to the core of Virgo (Neilsen \\& Tsvetanov 2000), it is, however, quite unlikely for such shock to be driven by the impact of the galaxy on Virgo ICM, even considering the uncertainties in the distance determination. A conclusion that M86 group is located outside the M87 cloud is further supported by observation of the X-ray halo extending over a degree in diameter (B\\\"ohringer et al. 1994).  We have therefore carried out a closer inspection of the large scale gas distribution surrounding M86 which led us to associate the counteracting structure with a filament in the neighborhood of Virgo cluster, already discovered by B\\\"ohringer et al. (1994).  The picture of M86 system emerging from these observations argues for a series of interactions at different stages of completeness. Although the weak shock is probably driven by an interaction with the surrounding large scale medium, the large-scale emission of M86 group is mostly unchanged. On the contrary, the core of M86 galaxy has been already strongly perturbed. This suggests that additional significant dynamic interactions were responsible for disturbing the X-ray morphology of M86 galaxy and that they took place on small scales, of order 10--50 kpc. Galaxy-galaxy encounters may offer a possible explanation.  The paper is organized as follows: in \\S\\ref{s:data} we present a detailed description of the data, including the reconstruction of maps of thermodynamic quantities in \\S\\ref{s:maps} that lead to the identification of a weak shock in \\S\\ref{s:shock}. The results are discussed in \\S\\ref{s:disc} and summarized in \\S\\ref{s:conc}.  In the following we will assume a 19.5 Mpc distance to M86 (Neilsen \\& Tsvetanov 2000), for which $1^\\prime=5.67$ kpc. ",
        "conclusions": "\\label{s:conc}  An in-depth understanding of the processes in the hot interstellar medium of M86, provided by observations of M86 has suggested a number of past and on-going interactions to take place. Comparison with the position of M86 in the Virgo cluster, promotes the environment outside the virial radius of M87 cloud to bear the prime responsibility for the apparent morphological transformation of the X-ray appearance of M86. Some of the past interactions in M86 are characterized by 10--50 kpc scale, which is only suitable for galaxy-galaxy interactions on one hand, setting limits on a degree of interaction, on the other, thus providing us with understanding of the processes leading to establishment of the morphology-density relations in clusters of galaxies.   \\smallskip  {\\it Acknowledgments.} This paper is based on observations obtained with XMM-Newton, an ESA science mission with instruments and contributions directly funded by ESA Member States and the USA (NASA). The XMM-Newton project is supported by the Bundesministerium f\\\"{u}r Bildung und Forschung/Deutsches Zentrum f\\\"{u}r Luft- und Raumfahrt (BMFT/DLR), the Max-Planck Society and the Heidenhain-Stiftung, and also by PPARC, CEA, CNES, and ASI. AF acknowledges receiving the Max-Plank-Gesellschaft Fellowship.  FM was partially supported by the European Community through the Research and Training Network - `The Physics of the Intergalactic Medium' - under contract HPRN-CT2000-00126 RG29185. AF acknowledges useful discussions with Hans B\\\"ohringer, Bernd Vollmer, Jacqueline van Gorkom and Cristina Popescu. The authors thank the anonymous referee for useful criticism.  \\vspace*{-0.5cm}"
    },
    "0311/astro-ph0311267_arXiv.txt": {
        "abstract": "{ A multipole spin-down equation based on a monopolar term is derived from the general expression $\\dot{\\nu} = -f(\\nu,t)$ and used to study pulsar evolution. We show that the time-independent version of such equation cannot reproduce the observed properties of pulsars and conclude that there is no equation of the form $\\dot{\\nu} = -f(\\nu)$ consistent with the \\pp diagram and braking index measurements. We explore the time-dependent model under the hypothesis of decaying magnetic fields, showing that an inverse linear decay gives reasonable evolutionary trajectories. This model distinguishes the evolution of Vela from that of the other three young pulsars considered. We discuss the origin of the monopolar term, which cannot be attributed to radiative processes, pointing to the importance of particle acceleration and/or mass loss processes in the dynamical evolution of pulsars. ",
        "introduction": "Pulsars are accepted to be rapidly rotating and highly magnetized neutron stars powered by the loss of rotational kinetic energy. Yet, the mechanisms through which pulsars lose their rotational energy are not very well understood. The discovery of pulsars was preluded by the rotating magnetic dipole model of Pacini (\\cite{pacini67}), which predicted a spin-down relation $\\dot{\\nu}\\propto \\nu^{3}$, where $\\nu$ is the rotational frequency. Pulsar spin-down has then been studied through the power-law differential equation, \\begin{equation} \\dot{\\nu}=-k\\nu^n \\, , \\end{equation} where $k$ is a constant and $n$ is called the ``braking index'', determined by the physical mechanisms spinning-down the star. For a pure magnetic dipole $n$ is equal to 3 (Gold \\cite{gold}; Pacini \\cite{pacini68}), while a pure gravitational or electromagnetic quadrupole one has $n=5$ (Ostriker \\& Gunn \\cite{ostriker}; Ferrari \\& Ruffini \\cite{ferrari}). The expansion of the Larmor equation into electromagnetic moments higher than dipolar gives $n > 3$. Departures from these basic models have been studied. For example, the deformation of magnetic field lines by corotation with the magnetosphere might produce $1\\leq n\\leq 3$ (Manchester \\& Taylor \\cite{manchester77}), pulsar winds $n<3$ (Blandford \\& Romani \\cite{blandford}; Manchester et al. \\cite{manchester85}), while magnetic field decay (Chanmugam \\& Sang \\cite{chanmugam89}) and the alignment between rotation and magnetic axis tend to give $n>3$ (Goldreich \\cite{goldreich}). {Physical processes such as phase transitions in the interior of the fast rotating neutron stars can produce changes in the moment of inertia $I$ and led also to deviations from $n=3$ even if the pulsar is spinning-down by pure dipolar radiation (Chubarian et al. \\cite{chubarian2000}; Glendenning et al. \\cite{glendenning97}). However, these phase transitions are significant only for very fast pulsars, with periods $P<1.5$ ms. For pulsars with periods $P\\ga 3 $ ms, these changes in the moment of inertia can be neglected.} But while physical mechanisms to produce $n<3$ have been proposed, {the reconciliation of pulsar evolution with $n<3$ remains to be studied}. This {question, addressed to the case of young classical pulsars,} is one of the main objectives of this work.  The braking index $n$ is frequently defined in terms of observational quantities, with a definition generalized to a second braking index, $m$: \\begin{equation} n \\equiv {\\ddot{\\nu}\\nu \\over \\dot{\\nu}^2} \\, , \\qquad m\\equiv {\\stackrel{...}{\\nu}\\nu^2\\over\\dot{\\nu}^3} \\, . \\label{index} \\end{equation} Accurate measurement of both braking indices requires long-term timing measurements, a difficult task due to the existence of glitches and timing noise in young pulsars, affecting the measurement of $\\ddot{\\nu}$ and $\\stackrel{\\ldots}{\\nu}$. Out of more than 1300 detected pulsars, only a handfull have reliable braking index measurements (Table~\\ref{tbl}). Still, measurements continue to be made and objects like PSR~J1119--6127 and PSR~J1846--0258 might probably be recent additions to the list (Camilo et al. \\cite{camilo00}; Mereghetti et al. \\cite{mereghetti}). The second braking index $m$ is even more difficult to measure and has only been estimated for PSR~B1509--58~(Kaspi et al. \\cite{kaspi}) and the Crab pulsar~ (Lyne et al. \\cite{lyne88}) (Table~\\ref{tbl}). It is well-known that measured values of $n$ and $m$ are below, but close, to those of a magnetic dipole in vacuum ($n=3$ and $m=15$).  \\begin{table} \\caption[]{Timing parameters and stationary multipole model fitting parameters for the four young pulsars selected.} \\label{tbl} $$ \\begin{array}{lcccc} \\hline \\noalign{\\smallskip} Pulsar & \\mathrm{Crab}^{(1)} & 1509-58^{(2)} & 0540-69^{(3)} & \\mathrm{Vela}^{(4)} \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} P ($ms$) & 33.5  & 150.9 & 50.3 & 89.3 \\\\ t_{dyn} ($yr$)  &  1258 & 1553 & 1664 & 11300 \\\\ n & 2.509 & 2.837 & 2.01 & 1.4 \\\\ & \\pm 0.005 & \\pm 0.001 & \\pm 0.02 & \\pm 0.2 \\\\ m & 10.23 & 14.5 & \\ldots & \\ldots \\\\ & \\pm 0.03 & \\pm 3.6 & \\ldots & \\ldots \\\\ \\hline \\noalign{\\smallskip} g ($Hz$^{-3}) & 3.5\\cdot 10^{-19} & 6.2\\cdot 10^{-16} & 6.6\\cdot 10^{-18} & 3.8\\cdot 10^{-18} \\\\ r ($Hz$^{-1}) & 1.0\\cdot 10^{-14} & 1.6\\cdot 10^{-13} & 7.1\\cdot 10^{-15} & 1.3\\cdot 10^{-15} \\\\ s ($Hz$) & 3.4\\cdot 10^{-12}  & 2.0\\cdot 10^{-12} & 5.7\\cdot 10^{-12} & 1.2\\cdot 10^{-12} \\\\ \\noalign{\\smallskip} \\hline P_{birth} & 9.9 & 14.2^{b} & 17.6 & 50 \\\\ $(ms)$^{a} &  &  &  & \\\\ \\hline \\end{array} $$ \\begin{list}{}{} \\item[$^{\\mathrm{a}}$]$P_{birth}$ defined as the period one dynamical time ago. \\item[$^{\\mathrm{b}}$]Integration is stopped at $t_{dyn}$ = 1~year. (1) Lyne et al. (\\cite{lyne88}); (2)~Kaspi et al. (\\cite{kaspi}); (3)~Manchester \\& Peterson (\\cite{manchester89}); (4)~Lyne et al. (\\cite{lyne96}). \\end{list} \\end{table}  Pulsar evolution is mostly studied through the \\pp diagram, which shows a sparse distribution of pulsars clustered roughly at $\\log P\\sim -0.3$ and $\\log\\dot{P} \\sim -15$. In principle it should be possible to link the bulk of the pulsar population in the \\pp diagram with the younger pulsars through evolutionary tracks consistent with a general spin-down equation. However, neither the magnetic dipole model or a constant braking index equation can link these groups and it has been argued that pulsars might be born with small period derivatives, in a region of the \\pp diagram where no pulsar detection has been reported (Camilo \\cite{camilo96}). Alternatively, magnetic dipole moments might decay during the pulsar life (Camilo \\cite{camilo96}; Tauris \\& Konar \\cite{Tauris}). On these lines, {Colpi et al.~(\\cite{colpi2000}) studied the distribution and evolution in the \\pp diagram of anomalous X-ray pulsars (AXPs), considering dipolar spin-down of magnetars. They conclude that magnetic field decay is required to account for the observed distribution and X-ray luminosities.}  We propose here a spin-down model derived from the very general spin-down law $\\dot{\\nu}=-f(\\nu,t)$. Although the derivation is empirical, the model accomodates the usual spin-down mechanisms, namely magnetic dipole radiation and gravitational radiation. In addition, a monopolar term is introduced, aiming to reconcile braking indices with trajectories in the \\pp diagram. We calculate evolutionary trajectories for PSR~B1509--58, PSR~B0540--69, the Vela pulsar (PSR~B0833--45) and the Crab pulsar (PSR~B0531+21). The model is presented in Section~2, while in Section~3 we apply the model to the four pulsars just mentioned, assuming time-independent spin-down mechanisms. In Section~4 we study the model considering time-dependent spin-down mechanisms. In Section~5, a discussion of possible physical mechanisms related to the spin-down terms is given, to lead to the concluding Section~6. ",
        "conclusions": "A multipole equation for pulsar evolution can be derived naturally as the Taylor expansion of the general equation $\\dot{\\nu} = -f(\\nu,t)$. While it might give insight into short term evolution of pulsars, the time independent version of the multipole equation is inconsistent with observations, indicating that no equation of the form $\\dot{\\nu} = -f(\\nu)$ is able to fully model pulsar spin-down evolution. A time dependent multipole equation considering exponentially decaying magnetic-fields cannot reproduce the data without shortening the active life of pulsars below $10^{5}$~years. Introducing an inverse linear magnetic field decay with timescales of the order of $10^4$ years gives good agreement with observations and allows pulsars to remain active about $10^7$ years, in agreement with pulsar and supernova birth rates. The model is consistent with observations of the Vela pulsar only if its braking index is $\\ga 1.5$. The different trajectories of this pulsar relative to those of the Crab, PSR~B1509--58 and PSR~B0540--69 suggest the existence two types of pulsars: those with relatively short decay timescales which evolve to periods $P_{\\infty} \\la 1~\\rm s$ and Vela-like pulsars, with more persistent magnetic fields evolving to the larger period pulsars observed in the \\pp diagram. We believe the need to introduce a monopolar term in the spin-down equation is an indirect observational evidence that particle acceleration and/or mass-loss processes must follow $\\dot{E} \\propto \\Omega^{2}$."
    },
    "0311/astro-ph0311233.txt": {
        "abstract": "We present an extensive set of stellar population synthesis models, which  self-consistently include the (optical--far-UV) continuum emission from stars as well as the resulting emission-line spectrum from photoionized gas surrounding massive stars during their main- sequence-life time. The models are presented as  a compiled library, \\ade, available electronically.  \\ade\\  contains the equivalent widths and the intensities of the lines \\oii\\ $\\lambda\\lambda$3727, \\hb, \\oiii\\ $\\lambda$5007, \\ha, and \\nii\\ $\\lambda$6584, which were calculated assuming a metallicity of 0.2, 0.4, 1.0, and 1.5 \\zsun\\ and a wide range of ionization parameters. We investigate the regime of continuous star formation, assuming a Salpeter initial mass function, whose upper mass limit is an input parameter.  The calculated equivalent width models, which are function of the burst age, are compared with the Jansen et al. atlas of integrated spectra for nearby galaxies.  We reproduce the observed properties of galaxies along the full Hubble sequence and suggest how the metallicity and age of such galaxies might be roughly estimated. ",
        "introduction": "The presence of emission lines is a conspicuous feature of optical spectra of any galaxy that is currently forming stars. The lines result from the reprocessing of the UV ionizing radiation from the most massive of the young stars.  Therefore, emission lines carry imprints of both phases, the newly formed stars versus the `excited' interstellar medium (ISM), in a single observable quantity, the integrated optical spectrum. It is for this reason that the modeling of the emission-line spectra occurring in star-forming galaxies must include the interaction between both phases (stars $+$ gas) in a self-consistent manner. Partly for this reason, this paper will emphasize the study of equivalent widths\\footnote{Equivalent widths present the advantage of relating any line flux with an observable quantity at the {\\it same} wavelength, namely the underlying continuum. There are quite a few papers in the literature that are concerned with the calibration of the \\oii\\ equivalent width and whether it can be used as a reliable tracer of star formation. Such diagnostics of \\wl(\\oii) in the UV would be particularly valuable at redshifts where \\hb\\ or \\ha\\ become inaccessible.} (EWs or \\wl). The optical stellar continuum underlying the emission lines consists of much more than just the stars that ionize the gas. Actually, in any evolved stellar population, the optical and red continua are dominated by red giants and low-mass main sequence stars. In order to model the observed continuum properly, one must therefore consider the past stellar formation history of the whole stellar population, an essential step if the emission-line equivalent widths are to be calculated reliably.   Working toward these goals, Garc\\ia a Vargas, Bressan \\& D\\ia az (1995), Stasi\\'nska \\& Leitherer 1996, Stasi\\'nska, Schaerer \\& Leitherer (2001) built models in which the spectral energy distributions (SEDs) generated by stellar evolutionary synthesis codes were used as input to photoionization calculations of emission-line strengths.  These models aimed mainly at constraining the ages and metallicities of isolated extragalactic \\hii\\ regions that are characterized by the absence or weakness of an old stellar population, suggesting a present-day strongly enhanced star formation rate.  In short, these earlier models followed the evolution of spectral properties of an instantaneous burst of star formation. Recently, Kewley \\etal (2001), Zackrisson \\etal (2001) and Moy, Rocca-Volmerange \\& Fioc (2001) introduced models in which the star formation rate   is constant for a few megayears, without interruption. These were used to derive the properties of starburst galaxies.  In an attempt to include the forgotten dust phase, Magris (1993) combined the spectral evolution synthesis code of Bruzual \\& Charlot (1993) with the photoionization code \\mapI\\ (Binette, Dopita \\& Tuohy 1985). The particularity of these models is that they included the effects of dust on the line- and continuum-emission in a self-consistent manner.  On the other hand, only the solar metallicity case was studied, since detailed non-solar metallicity tracks and stellar spectral libraries were not complete at that time.  Barbaro \\& Poggianti (1997) have presented a model that is used to calculate the emission lines of \\oii\\ $\\lambda\\lambda$ 3727 and of the Balmer series along with photometric quantities. The purpose was to estimate the present day star formation rate in late type galaxies. In their model, the star formation rate and the chemical evolution of the stellar phase for each Hubble type have been derived from their own chemical evolution model, assuming, on the other hand, a fixed solar composition for the ionized gas.  Charlot \\& Longhetti (2001) computed the line and continuum emission from star-forming galaxies and presented a family of estimators of the star formation rate, the nebular oxygen abundance, and the optical depth of the dust, in galaxies, as a function of the available spectral information.  Their model is generally successful in reproducing the observed properties\\footnote{See Charlot \\etal (2002) for a more recent application of this model to the determination of the star formation rate.} of template galaxies of different types. Their analysis, however, did not include the modeling of equivalent widths that are particularly suitable for analyzing the interrelated properties of the stellar and nebular phases.  We have developed a model along similar lines, which calculates, for an arbitrary star formation rate, both the nebular emission-line spectrum of the \\hii\\ regions associated to the newly born massive stars and the SED of the underlying stellar population.  However, rather than discussing line ratios and stellar energy distributions separately, our work will emphasize the relationship between the two phenomena, by presenting our model results as a library, \\ade, of emission-line equivalent widths. As supplementary information, the tables also include the line luminosities of the most prominent species observed in normal star-forming nearby galaxies, i.e. \\oii\\ $\\lambda\\lambda$3727, \\hb, \\oiii\\ $\\lambda$5007, \\ha, and \\nii\\ $\\lambda$6584. The models are characterized by a metallicity range of 0.2, 0.4, 1, and 1.5 times solar (gas $+$ star) and consider rates of star formation that are constrained by both the line and the stellar continuum emission, of galaxies of all Hubble types.  The method used to derive our evolutionary models is presented in \\S 2.  A description of the results from \\ade\\ follows in \\S 3 accompanied by a comparison of the models with the Jansen et al. (2000a,b; hereafter \\joo) atlas of nearby galaxies. The conclusions are summarized in \\S 4. A more extensive version than the adjoining tables of  \\ade\\ models is available on the Web. ",
        "conclusions": "We have presented a new library of models, \\ade, which combine the stellar population history with the emission line-strengths.  These models can be used to interpret the properties of objects that are characterized by either a very recent and strong burst or by a secular and lower star formation rate.  Starting from simple assumptions about the distribution of \\hii\\ regions in disk galaxies, namely, a uniform distribution of dust-free gas and stars, in which young stars are distributed in mass in a similar fashion as the older stars, we have shown how it is possible to compute several properties of star forming galaxies. A photoionization code is used to compute the intensity of the emission lines, assuming no temporal evolution of $U$, while an evolutionary synthesis code provides the number luminosity and hardness of the MS ionizing photons. These, in conjunction with the photoionization code and the integrated stellar distribution, allow us to compute the temporal evolution of the emission line equivalent widths.  Diagnostic diagrams built by comparing pairs of equivalent widths of various lines, show that our model predictions fall in figures in the same region as the observed galaxies of \\joo.  In particular, we conclude the following. \\begin{enumerate} \\item A factor of 10 in $U$ and a range in metallicity between 0.2 and 1.5 \\zsun\\ seems appropriate to describe the observed properties of a representative sample of nearby galaxies. \\item The hardness of the ionizing continuum is constant after $10^7$ yr in  case of a continuous SFR. \\item Our favored {\\it continuous} SFR model with $\\msup = 125 \\msun$ presents a hardness $\\eta$ very similar to that produced by an instantaneous burst with an IMF having only $\\msup = 60 \\msun$. \\item Galaxies with $\\wl(\\oii)\\gtrsim 30$ \\AA\\ and \\wl$(\\ha) \\lesssim 70$ \\AA\\, appear to have subsolar abundances. \\item Galaxies with $\\wl(\\ha) \\lesssim 100$ \\AA\\ have a dominant underlying stellar population older than at least 1 Gyr. Only galaxies with a recent burst of star formation ($t \\lesssim 1$ Gyr) are able to show \\ha\\ EW in excess of 100 \\AA. \\item The information provided by the \\oiii\\ line is of limited use, owing to the strong \\msup--$Z$ degeneracy, which complicates the interpretation of this high-excitation indicator, particularly in the case of low \\oiii\\ EWs. \\end{enumerate} %An enlarged version of the \\ade\\ model predictions  is available at\\\\ %\\url{http://www.cida.ve/$^\\sim$magris/ademis}"
    },
    "0311/astro-ph0311117_arXiv.txt": {
        "abstract": "Low Surface Brightness (LSB) galaxies are dominated by dark matter. High-resolution rotation curves suggest that their total mass-density distributions are dominated by constant density cores rather than the steep and cuspy distributions found in Cold Dark Matter (CDM) simulations.  The data are best described by a model with a soft core with an inner power-law mass-density slope $\\alpha = 0.2 \\pm 0.2$. However no single universal halo profile provides an adequate description of the data. The observed mass profiles appear to be inconsistent with $\\Lambda$CDM. ",
        "introduction": "Cosmological numerical Cold Dark Matter (CDM) simulations predict a specific and universal shape for the dark matter density distributions of galaxies. This was investigated in detail by Navarro, Frenk \\& White (1996, 1997) [NFW] who found that the mass-density distribution in the inner parts of simulated dark-matter-only CDM halos could best be described by a $r^{-1}$ power law. This is however not what is typically observed: the mass-density distributions of, for example, dwarfs and late-type galaxies (which do include baryons) are best described by an almost constant density central core $\\rho(r) \\sim r^{-0.2 \\pm 0.2}$ (de Blok et al.\\ 2001b).  This disagreement has become known as the ``cusp/core problem''.  There has been much discussion of the validity of both observations and simulations. My intention here is not to give an exhaustive review of this debate. Rather I want to highlight some of the concerns and criticisms of both observers and simulators, and discuss them from an observer's point of view (and in the spirit of the Sydney panel discussion I will do this in a slightly provocative manner). ",
        "conclusions": "\\begin{itemize}  \\item The available data overwhelmingly prefer soft cores.  \\item There is a large range in inner slopes. Observationally there is no evidence for a ``universal profile''.  \\item Systematic effects play no significant r\\^ole.  \\item Observations probe a regime where simulations have no predictive power (yet).  \\item Crucial (baryon) physics are still missing from the simulations.  \\end{itemize}"
    },
    "0311/astro-ph0311321_arXiv.txt": {
        "abstract": "The cosmological gamma-ray burst (GRB) phenomenon is reviewed.  The broad observational facts and empirical phenomenological relations of the GRB prompt emission and afterglow are outlined.  A well-tested, successful fireball shock model is introduced in a pedagogical manner. Several important uncertainties in the current understanding of the phenomenon are reviewed, and prospects of how future experiments and extensive observational and theoretical efforts may address these problems are discussed. ",
        "introduction": "The study of astrophysical objects is overwhelmingly done by astronomers using the temporal and spectral information contained in the electromagnetic signals that these objects emit. Gamma-ray bursts (GRBs) are, by definition, electromagnetic signals in the gamma-ray band (in the spectral domain) with short durations (in the temporal domain). They are, however, unusual in having most of their electromagnetic output in gamma-rays, typically at sub-MeV energies, and having most of it concentrated into a brief episode, typically lasting tens of seconds.  The road leading to an understanding the nature of these objects has been bumpy, mainly due to the limited information contained in these abrupt gamma-ray episodes. For comparison, almost all the other astrophysical objects are observed in a broader spectral band during a much longer observation time. Historically, the lack (until recently) of observational breakthroughs on GRBs had left ample freedom for modelers to play around, so that the cumulative list of models championed to interpret GRBs has been more numerous than for any other astrophysical phenomenon.  As in other fields in science, experimental (or observational) progress eventually constrains and eliminates most predictive models. At a certain stage of the development of a field, it is useful to ask oneself questions such as what is it that we know, what is it that we still do not know, and how could we find the answers for the unknowns. This is the main emphasis of this review. There exist already a number of comprehensive reviews on GRB observations\\cite{fm95,vkw00,kul00}, on theories of GRB\\cite{p99,m02} and on some general or selected topics\\cite{m01,cl01,wzh02,hsd02,dermer02,zmw02,ghis03,d03,w03,mkrz03}. Our purpose here is to provide a review which differs from others in several aspects. For example, we summarize the current observational progress in an organized, itemized form, making a distinction between solid facts and empirical laws. The introduction of the standard theoretical model is pedagogical, the content of which overlaps other reviews. However, we also dedicate a lot of space to discuss the uncertainties in the current models and highlight some active debates that are going on in the community.  We also attempt to reflect in an unbiased manner the work from various groups, rather than focusing on our own work, although some ``selection effects'' must still exist due to our incomplete survey of the literature.  The structure of the review is as follows.  We start with a brief summary of the GRB research history as well as the unique role of GRB study in general astrophysics. We then compile the observational facts and empirical laws (\\S\\ref{sec:prog1}), and introduce the standard theoretical framework, i.e. the fireball shock model (\\S\\ref{sec:prog2}). In \\S\\ref{sec:prob} we outline some outstanding issues in the GRB field, which comprises most of the hot topics (and sometimes hot debates) in the current GRB scene. These issues include whether the fireball is dominated by the conventional (hadronic) kinetic energy content or by the less-studied magnetic energy; whether the GRB prompt emission occurs before the fireball deceleration (internal) or at the deceleration radius (external); whether the GRB jets are uniform or with a quasi-universal angular structure; whether the GRB progenitors give rise to a one-step or to a (long delay) two-step collapse, etc. After outlining these issues, we attempt to foresee in \\S\\ref{sec:pros} developments in the upcoming new era for GRB study, led by future missions such as {\\em Swift} and {\\em GLAST}, complemented by other space- and ground-based detectors for both electromagnetic and non-electromagnetic signals.  We review some of the theoretical predictions related to these observational capabilities and discuss how future observational campaigns could constrain and possibly settle some of the outstanding issues listed in \\S\\ref{sec:prob}.   \\subsection{Major Landmarks in the Study of GRB}  As a framework for the following discussions, we list below some of the most important events in the development of the understanding of GRB.  There may be many other important events which are not included in this list, and we apologize for any omission. Most of the papers selected below are marked by their high citation counts in the ADS archives, an indication that the works are significant and widely accepted. For the theoretical papers, we only include those that survived the later observational data and correctly contributed to the current knowledge of the GRB phenomenon, or those which, although not fully tested yet, represent ``popular\" views and draw broad attention from the GRB community.  The papers extend up to October 2003, when this review is completed.  They include both observational and theoretical developments, which have mutually influenced, and have benefited from, each other.  A list of key advances on the observational side is: \\footnote{Notice that all the afterglow-related breakthroughs are for the so-called ``long'' GRBs (see \\S\\ref{sec:prog1} for definition).  No afterglow from a ``short'' GRB has so far been firmly identified.}.  \\begin{itemlist} \\item GRBs were discovered serendipitously in the late 1960s, and the data were reported several years later\\cite{kso73,m74}; \\item The {\\em Compton Gamma-ray Observatory} ({\\em CGRO}) was launched in 1991, with the {\\em Burst and Transient Experiment} ({\\em BATSE}) on board; {\\em BATSE} provided evidence for an isotropic spatial distribution of GRBs, giving significant support to a cosmological origin interpretation\\cite{m92}; \\item Two categories of GRBs, ``long'' ($t_\\gamma\\simg 2$ s) and ``short'', ($t_\\gamma \\siml 2$ s) were identified\\cite{k93}; \\item Systematic analyses of GRB spectral data indicate that a so-called ``GRB function'' (or ``Band function'') fits reasonably well most of the GRB spectra (for both classes of bursts)\\cite{b93}; \\item Another detector on board of {\\em CGRO}, the Energetic Gamma Ray Experiment Telescope ({\\em EGRET}), detected some bursts in the hard gamma-ray band; one burst, GRB 940217, was detected to have long-lived GeV emission extending for 1.5 hrs\\cite{h94}; \\item A signature consistent with cosmological time dilation was detected\\cite{norris94}; \\item In 1997, the Italian-Dutch satellite {\\em Beppo-SAX} pinpointed the first GRB low energy afterglow (for GRB 970228) in the X-ray band\\cite{c97}, facilitating optical\\cite{van97} and radio\\cite{f97} detections; \\item The first measurments of the GRB redshifts were obtained (for GRB 970508\\cite{m97} and GRB 971214\\cite{k98}, giving a solid proof that GRBs are at cosmological distances; \\item A likely GRB-supernova association (GRB 980425 vs. SN 1998bw) was discovered\\cite{g98,k98b}; \\item A bright and prompt optical flash and a radio flare were discovered to accompany the energetic burst GRB 990123\\cite{a99,k99a}; \\item An achromatic steepening break in the afterglow lightcurves was found in several bursts, hinting that at least some GRB fireballs are likely to be collimated\\cite{k99,h99}; \\item X-ray spectral features with moderate significance were discovered in several GRB X-ray afterglows\\cite{p00,r02}; \\item Assuming the lightcurve steepening is due to the jet break, a collection of the burst jet data revealed that GRBs have a quasi-standard energy reservoir\\cite{f01}; \\item A class of so-called ``X-ray flashes'' (XRFs) were identified\\cite{heise,kippen}, which appear to be closely related to GRB; \\item Prompt optical afterglows from two more GRBs, GRB 021004 and GRB 021211, were detected by ground-based robotic telescopes\\cite{f03a,f03b,li03}; \\item The prompt gamma-ray emission was reported to be strongly polarized in GRB 021206\\cite{cb03} (cf Refs.\\refcite{rf03,bc03}); \\item The first clear GRB-SN association was identified: an unambiguous supernova light-curve bump and spectrum in the GRB 030329's optical afterglow lightcurve was discovered\\cite{sta03,h03}; \\item A distinct high energy spectral component which appears to evolve separately from the usual sub-MeV component was identified in GRB 941017\\cite{g03}; \\item Broadband observations of GRB 030329 suggest the there might exist a multi-component jet structure in some bursts, the total energy output being standard\\cite{b03,sheth03}. \\end{itemlist}   On the theoretical side, some of the major steps in the development towards the contemporary ``standard'' model are listed below.  This list does not include many neutron-star-based Galactic GRB models (see Ref.~\\refcite{h91} for a review), many of which were proposed before the proofs of a cosmological origin of GRBs became solid.  \\begin{itemlist} \\item A generic argument about the ``compactness problem'' was raised\\cite{r75} in 1975; \\item A self-similar solution of the relativistic blast waves was found in 1976\\cite{bm76}. Although this study was not aimed at GRBs, it laid a mathematical foundation for relativistic GRB blast wave models, and has been extensively applied since the discovery of GRB afterglows; \\item Some generic features of a gamma-ray fireball were discussed in 1978\\cite{cr78}; \\item A cosmological model of GRBs was formally proposed in 1986\\cite{p86,g86,p90}, involving neutron star - neutron star (NS-NS) mergers as the power source. A pure photon-pair fireball was found to expand relativistically, leaving a quasi-thermal photospheric emission; \\item It was found that by including a small amount of baryon contamination, the fireball thermal energy is essentially converted into kinetic energy of the expanding gas\\cite{sp90}; \\item It was pointed out that the kinetic energy is re-converted to heat and radiation via shocks\\cite{rm92,mr93}. This suggestion laid the foundation for the current fireball-shock models. The specific scenario proposed in Refs.~\\refcite{rm92,mr93} involves the shock forming when the fireball is decelerated by the ambient interstellar medium (ISM), which is known as the ``external shock scenario''. Synchrotron radiation was proposed as the mechanism of GRB emission, noting that this would soften and fade as a power law in time. \\item Around the same time, several major GRB progenitor models were further developed or proposed: the NS-NS merger model was further studied in some detail\\cite{e89,npp92}; a ``failed'' supernova model was proposed\\cite{w93}, which is now known as the ``collapsar'' model; and a millisecond magnetar (neutron stars with surface magnetic field of order $10^{15}$ G) model was proposed\\cite{u92}; \\item Another fireball shock scenario involving collisions among the individual shells in the fireball outflow, known as the ``internal shock scenario'' was proposed\\cite{rm94,px94}, as a way to resolve the very short time variability problem. \\item The hydrodynamical evolution of a fireball was extensively modeled both numerically and analytically\\cite{mlr93,mr93b,psn93,katz94,sp95,kps99,mr00,mrrz02}; \\item The detailed reverse and forward shock synchrotron and inverse Compton spectral components of external shocks were explored\\cite{mrp94,mr94}, indicating the presence of a reverse shock optical synchrotron component (now identified as responsible for prompt optical flashes) and a GeV component due to inverse Compton. \\item Based on the synchrotron radiation assumption, GRB radio and optical afterglows were suggested\\cite{pr93,k94}. A detailed self-consistent afterglow calculation\\cite{mr97} stressing also the X-ray afterglow was performed shortly before the observational discovery of afterglows.  The predictions include the power-law decay of the optical lightcurves as well as prompt optical flashes arising from the reverse external shock during the early stage of the afterglows. All these were verified by the later observations\\cite{van97,a99}; \\item Prompted by the discovery of the first GRB afterglows\\cite{c97,van97,f97} the fireball model was studied and tested extensively, and found to be successful in interpreting the observations\\cite{wrm97,v97,w97a,w97b,t97}; \\item It was suggested\\cite{p98} that GRBs are likely originated from ``hypernovae'', and are associated with star forming regions; \\item The simplest standard afterglow model was set up\\cite{mrw98,spn98}, and was further developed by incorporating additional effects\\cite{mrw98,rm98,dl98,dl99,cl99}; \\item Stimulated by the increasing evidence that GRBs are associated with supernovae, the GRB progenitor models involving collapses of massive stars were further developed. The ``collapsar model'' was extensively studied numerically\\cite{mw99,a00,zwm03} and analytically\\cite{mr01,wm03}. A variant of the massive-star-collapse model, invoking a delayed GRB with respect to the SN explosion, was proposed\\cite{vs98}, known as the ``supranova'' model; \\item The scenario involving GRB prompt gamma-rays from internal shocks and an afterglow from the external shock as a unified sequence requiring a joint fit was analyzed\\cite{sp97a,sp97b}. In parallel to this, extensive efforts were made to model the GRB prompt emission both within the internal shock models\\cite{kps97,dm98,psm99,spm00,gsw01} and the external shock models\\cite{pm98,cd99,dm99,dbc99}; \\item The consequences of GRBs originating from collimated jets was proposed and studied\\cite{r97,r99}, leading to an extensive campaign of theoretical jet modeling\\cite{sph99,pm99,msb00,hgdl00}; \\item Prompted by the detection of the optical flashes associated with GRB 990123\\cite{a99}, GRB 021004\\cite{f03a} and GRB 021211\\cite{f03b,li03}, the reverse shock emission was studied in greater detail\\cite{sp99a,sp99b,mr99,ks00,k00,kz03a,zkm03,wdhl03,kz03b,kmz03}. The recipes to use early afterglow data to systematically infer fireball parameters were recently proposed\\cite{zkm03,kz03b,kmz03}; \\item The possibility of producing Fe X-ray lines in supranova and in collapsar stellar funnel models was investigated\\cite{lcg99,rm00,weth00,mr01,vietri01,lrr02,kmr03}. \\item Broadband modeling of GRB afterglows was carried out, leading to constraints on some unknown model parameters\\cite{wg99,pk01,pk02,harrison01,yost03}; \\item Prompted by the finding of a standard energy reservoir in long GRBs\\cite{f01}, it was proposed that long GRBs may have a quasi-universal structured jet configuration\\cite{rlr02,zm02b,zdlm03}. \\end{itemlist}  \\subsection{The multi-disciplinary nature of the GRB field}  The GRB field is almost unique in astrophysics in its multi-disciplinary nature. Involving stellar-scale events located at cosmological distances, GRBs straddle the traditional distance scales, and are a high energy phenomenon which emits a broad-band electro-magnetic spectrum extending over at least fifteen decades, as well some possible non-electromagnetic signals (such as cosmic rays, neutrinos and gravitational waves). This makes the GRB field an intersection of many branches in astrophysics.  GRBs were discovered at about the same time when quasars (or broadly speaking active galactic nuclei, AGNs) and pulsars, two other major developments in the 1960's, were discovered. The nature of those two classes of objects was agreed upon soon after their discoveries. AGNs are extra-galactic sources believed to be powered by gigantic black holes, while pulsars are compact neutron stars located in our Galaxy. Since the lack of GRB observational breakthroughs hampered the progress in this latter field, scientists from both the AGN and the pulsar community brought to bear their collective wisdom from their own fields to tackle the GRB problem. The GRB neutron star models were developed to fairly sophisticated levels\\cite{h91}, especially motivated by the reported detection of absorption and emission features in some GRB spectra\\cite{m81,m88}. It turned out that the ``classical\" GRBs are of cosmological origin, so that the wisdom borrowed from the AGN community (especially about blazars, the most energetic type of AGNs) finally bore the most rewarding fruit. Nonetheless, neutron star models eventually turned out to be partially correct, i.e., one sub-class of such models that invokes ultra-strong magnetic fields\\cite{dt92} turned out to be useful in interpreting the class of soft, repeating GRBs\\cite{} (which appear to observationally distinct from the ``classical\" cosmological GRBs). These sources, known as ``soft gamma-ray repeaters'' (SGRs), together with another type of objects called ``anomalous X-ray pulsars'' (AXPs), are now believed to be ``magnetars'', neutron stars with super-strong ($\\sim 10^{15}$ G at surface) magnetic fields\\cite{td95,td96}. Even within the current paradigm, pulsar models can still provide helpful insights for the understanding of cosmological GRBs. For example, the recent claim of strong polarization of gamma-ray prompt emission in GRB 021206\\cite{cb03} (cf. Refs. \\refcite{rf03,bc03}) suggests a strongly magnetized central engine. A similar conclusion of a strong but less extreme magnetization was reached through a combined reverse-forward shock emission analysis for GRB 990123\\cite{zkm03} and GRB 021211\\cite{zkm03,kp03}. This hints that magnetic fields and/or Poynting flux may play an essential role in GRBs, and that some insights from pulsar wind nebula theories may turn out to be important to unveil the GRB mystery. We discuss this in more detail in \\S\\ref{sec:mag}.  Besides its close intimacy with the AGN and the pulsar fields, the GRB field is also broadly related to several other branches of astrophysics.  First, in the stellar context, since long GRBs are related to the deaths of massive stars, as supported by many observational facts, GRB study is closely related to the fields of stellar structure and evolution\\cite{heger03}, supernovae\\cite{wzh02b,i98}, and supernova remnants (e.g. the external shock invoked in GRB models is the relativistic version of the supernova remnant shock\\cite{a01,ed02,lp03}). Progenitor studies have stimulated stellar population studies\\cite{fwh99}. Central engine studies, on the other hand, have greatly promoted studies about mechanisms for extracting energy from accretion disks or spinning black holes\\cite{bz77,mr97b,pwf99,lwb00,npk01,van01}.  Second, within the galactic context, GRB afterglow lightcurves\\cite{cl99,wl00,dl02} and spectral features\\cite{p00,r02,s03} probe the properties (e.g. density profile and chemical abundance) of the ambient interstellar medium (ISM) or the prestellar wind. Studying the properties of the GRB host galaxies as well as the GRB locations within the host galaxies brings valuable information about the global star forming history of the universe and the nature of GRB progenitors\\cite{bloom98,f99,bloom02}. Third, GRBs are also important objects within the cosmological context. Not only can they play a role similar to AGNs in probing the observed low-redshift universe, but they may also form and can be detectable at much higher redshifts\\cite{lr00,cl00,bl02,gou03}. Thus, detecting high-$z$ GRBs would allow us to see into deeper and earlier epochs of the universe, and to probe how the universe ends its dark age and gets re-ionized\\cite{m-e03}. Fourth, in high energy astrophysics, the origin of ultra high energy cosmic rays (UHECRs) has remained a mystery for decades. Among many other models, a GRB origin\\cite{w95,v95} is a leading model for the so-called ``bottom-up models''. In this scenario, cosmic ray protons are believed to be accelerated in GRB shocks. The same processes can also generate high energy neutrinos\\cite{px94,wb97,bm00,mw01}, so that GRBs are currently regarded by many as the top potential high energy neutrino sources for large area neutrino telescopes being built. Finally, several GRB progenitor scenarios are believed to generate gravitational wave signals\\cite{kp93,km03}, and GRBs are also one of the major targets for large gravitational wave detectors.  On balance, it would appear that hardly any other field has such a broad range of interactions with other branches of astrophysics. ",
        "conclusions": "Our understanding of GRBs has been greatly advanced since their discovery about 30 years ago. Extensive observational efforts have revealed a rich phenomenology of both prompt emission and afterglows (\\S\\ref{sec:prog1}), and a successful theoretical framework has been set up, which is able to interpret most of the observational data so far (\\S\\ref{sec:prog2}). However, there remain many questions at the present stage of GRB studies, which greatly stimulate further observational and theoretical efforts (\\S\\ref{sec:prob}). Some of these questions will undoubtedly be addressed in the coming years with the advent of a number of upcoming space- and ground-based experiments, both in the electromagnetic and non-electromagnetic channels (\\S\\ref{sec:pros}). As suggested by past experience, new challenges and surprises are bound to emerge, which will stimulate further extensive observational campaigns and theoretical efforts. The GRB field is likely to remain one of the most active fields in contemporary astrophysics in the next decades.  This research was partly supported through NASA ATP NAG5-13286, NSF AST0098416 and the Monell Foundation. We thank S. Kobayashi, T. Lu, D. A. Frail, J. P. Norris, E. D. Feigelson, Z. G. Dai, Y. F. Huang, and S. Razzaque for helpful comments."
    },
    "0311/hep-ph0311172_arXiv.txt": {
        "abstract": "The Earth matter effects on supernova (SN) neutrinos can be identified at a single detector through peaks in the Fourier transform of their ``inverse energy'' spectrum.  The positions of these peaks are independent of the SN models and therefore the peaks can be used as a robust signature of the Earth matter effects, which in turn can distinguish between different neutrino mixing scenarios.  Whereas only one genuine peak is observable when the neutrinos traverse only the Earth mantle, traversing also the core gives rise to multiple peaks. We calculate the strengths and positions of these peaks analytically and explore their features at a large scintillation detector as well as at a megaton water Cherenkov detector through Monte Carlo simulations.  We propose a simple algorithm to identify the peaks in the actual data and quantify the chances of a peak identification as a function of the location of the SN in the sky. ",
        "introduction": "\\label{intro}  The neutrino spectra that arrive at the Earth from a core collapse supernova (SN) have information about the neutrino masses and mixings encoded in them. The $\\sim 20$ neutrinos observed from SN~1987A were extensively used trying to constrain the solar neutrino parameters as well as $\\theta_{13}$ and the neutrino mass hierarchy~\\cite{SN87a,SN87b,SN87c,SN87d}.  The determination of neutrino parameters depends crucially on our understanding of the primary neutrino fluxes produced inside the SN.  In spite of large uncertainties on these primary fluxes, some of their robust features may be exploited to identify the type of neutrino mass hierarchy and put bounds on the mixing of $\\nue$ in the ``third'' neutrino mass eigenstate~\\cite{ds,cicilia2,observ}.  When neutrinos pass through the Earth before arriving at the detector, their spectra may get modified due to the Earth matter effects. The presence or absence of these effects can distinguish between different neutrino mixing scenarios~\\cite{cairo}.  The comparison of neutrino spectra at two different detectors can clearly give signatures of the matter effects, which can be used not only for the determination of the neutrino parameters~\\cite{cicilia1,takahashi}, but also to extract information about the density structure of the Earth core~\\cite{earthtomography}.  The measurements of the Cherenkov glow at IceCube may also be combined with the signal at a water Cherenkov detector like Super- or Hyper-Kamiokande to identify the Earth effects~\\cite{icecube}.  It is also possible to ascertain the presence of these matter effects using the signal at a single detector.  It has recently been pointed out \\cite{wiggles} that the Earth matter effects on supernova neutrinos traversing the Earth mantle give rise to a specific frequency in the ``inverse energy'' spectrum of these neutrinos.  This frequency, which may be identified through the Fourier transform of the inverse energy spectrum, is independent of the initial neutrino fluxes and spectral shapes.  Therefore, its identification serves as a model independent signature of the Earth matter effects on SN neutrinos, which in turn can distinguish between different scenarios of neutrino masses and mixings.  If the SN neutrinos reach the detector ``from below,'' they have to travel through the Earth matter. If the nadir angle of the SN direction at the detector is less than 33$^\\circ$, the neutrino path crosses the Earth core. An investigation of the effect of the core on the observed neutrino spectra is therefore necessary for a complete understanding of the Earth effects. As we shall see later in this paper, the passage through the core increases the chances of the identification of the Earth effects. The core density is almost twice the mantle density and this sudden density jump gives rise to new features in the spectra.  When the SN neutrinos traverse the Earth core in addition to the mantle, one does not get a single specific frequency as in the mantle-only case, indeed as many as seven distinct frequencies are present in the inverse energy spectrum.  We study the strengths of these frequency components analytically and show that three of these frequencies dominate.  These three frequencies can also be clearly observed in the Fourier transform of the inverse energy spectrum when averaged over many simulated SN neutrino signals.  Although it is difficult to isolate these frequencies individually from the background fluctuations from a single SN burst, we suggest a procedure that can identify the presence of these frequency components in a sizeable fraction of cases.  Certain characteristics of the distribution of the frequency components in the background fluctuations are identified and used to reject the null hypothesis of the absence of Earth effects.  We quantify the efficiency of this algorithm by simulating the SN neutrino signal at a large scintillation detector like LENA \\cite{lena} and at a megaton water Cherenkov detector like Hyper-Kamiokande.  Whereas the scintillation detector has the advantage of a much better energy resolution, this is compensated in part by the larger number of events in a megaton water Cherenkov detector.  This paper is organized as follows. In Sec.~\\ref{analytic}, we discuss the positions and the strengths of the frequencies that characterize the ``inverse-energy'' spectra of the neutrinos crossing the Earth mantle as well as the core.  In Sec.~\\ref{peaks}, we simulate the SN neutrino spectra at the detectors and study the features of the peaks with the background fluctuations averaged out.  In Sec.~\\ref{bkgnd}, we introduce a method to identify the peaks in the presence of background fluctuations and make a quantitative estimation of the probability of peak identification as a function of the location of the SN in the sky.  In Sec.~\\ref{conclusions}, we summarize the results. ",
        "conclusions": "\\label{conclusions}  When neutrinos coming from a core-collapse supernova pass through the Earth before arriving at the detector, the spectra may get modified due to the Earth matter effects.  The presence or absence of these effects can distinguish between different neutrino mixing scenarios. We have seen that these Earth matter effects on supernova neutrinos can be identified at a single detector through peaks in the Fourier transform of their ``inverse energy'' spectrum. The position of these peaks are independent of the initial neutrino fluxes and spectral shapes.   We have performed an analytical study of the positions and the strengths of the frequencies that characterize the inverse-energy spectrum of the neutrinos for different neutrino trajectories through the Earth. In the case that the SN neutrinos only traverse the mantle a single peak shows up in the power spectrum. In contrast we have observed that if both the mantle and the core are crossed before the neutrinos reach the detector as many as seven distinct frequencies are present in the inverse energy spectrum. However only three peaks are dominant in the  power spectrum. This increase in the number of expected peaks leads to an easier identification of the Earth matter effects.   In order to illustrate the qualitative features of the present analysis we have considered the power spectrum resulting from averaging 1000 SN simulations for different SN models and different detector capabilities. In particular we have assumed a 32 kton scintillator detector and a megaton water Cherenkov detector.  We have shown how the energy resolution turns out to be crucial in detecting the modulation introduced in the neutrino spectra by the Earth matter effects. First, the better resolution of the scintillator detector compensates for the larger water Cherenkov detector size.  On the other hand, the worse energy resolution in water Cherenkov detectors does not only imply the need of a larger volume but also suppresses significantly the peaks at higher frequencies, in contrast to the case of scintillator detectors.   We have considered two different SN models as an illustration of the current uncertainties in the initial fluxes. We have observed that the strength of the peaks is larger in those SN models with bigger differences between $\\bar\\nu_e$ and $\\bar\\nu_\\mu$ spectra. However, we have found that the position of the peaks is model independent.  Therefore their identification serves as a clear signature of the Earth matter effects on SN neutrinos, which in turn can help to discard the neutrino mass scheme with inverted mass hierarchy and $\\sin^2\\Theta_{13}\\gsim 10^{-3}$.  We have introduced a simple algorithm to identify the peaks in the presence of background fluctuations. This method is based on the integration of the area around the expected position of the peak.  By comparing the area distribution without and with the spectral modulations induced by the Earth matter effects we have analyzed the statistical significance of the result. As expected the presence of the core as well as a larger difference in the initial spectra enhance the probability of identifying the Earth effects. We have also presented a variation of the algorithm which improves its efficiency significantly in the case of a water Cherenkov detector for neutrino trajectories passing through the core. Therefore we believe that more efficient algorithms could be developed. The results we presented should be considered therefore as conservative lower limits."
    },
    "0311/astro-ph0311447_arXiv.txt": {
        "abstract": "We discuss the different physical processes contributing to the infrared continuum of AGN, assuming that both photoionization from the active center and shocks ionize and heat the gas and dust contained in an ensemble of clouds surrounding the nucleus. In our model, radiation transfer of primary and secondary radiation throughout a cloud is calculated consistently with collisional processes due to the shock. We consider that the observed continuum corresponds to reprocessed radiation  from both dust and gas  in the clouds. Collisional processes are important in the presence of shocks.  The grains are sputtered crossing the shock front.  The models are constrained by sputtering as well as by the far-infrared data. The model is applied to the continuum of Seyfert galaxies from which best estimate of the nuclear, stellar subtracted, emission is available. The results show that radiation-dominated high velocity clouds are more numerous in Seyfert 1-1.5 whereas shock-dominated low velocity clouds are dominant in Seyfert type 2. This result is in full agreement with the unified model for AGN, by which high velocity clouds, placed deeper into the central region and therefore reached by a more intense radiation, should play a more significant role in the spectra of broad line objects. We could therefore conclude that in type 2 objects, radiation is partly suppressed by a central dusty medium with a high dust-to-gas ratio. Once the model approach is tested, a grid of models is used to provide a phenomenological analysis of the observed infrared spectral energy distribution. This empirical method is a useful tool to rapidly access the physical conditions of the AGN emitting clouds. For this, analytical forms are derived for the two processes contributing to the infrared emission: dust emission and thermal bremsstrahlung produced by the NLR clouds. Their relative contribution provides a measurement of the dust-to-gas ratio. ",
        "introduction": "The presence of a dusty torus around the central engine was proposed by Antonucci \\& Miller (1985)  to explain the broad emission-lines observed in the polarized light of Seyfert 2 galaxies. The torus absorbs a large fraction of the nuclear emission and reradiates it in the infrared. This is the basis of the so-called AGN unification scheme largely analysed  in the literature and that provides a geometrical explanation for the observed differences between broad line and narrow line objects.  After the pioneering work by Edelson  \\& Malkan (1986), the characteristics of the continuum due to dust heated by the optical-UV nuclear radiation has been analyzed by several authors (Barvainis 1992, Pier \\& Krolik 1992; Laor \\& Draine 1993; Granato \\& Danese 1994; Granato, Danese \\& Franceschini 1997,  Nenkova, Ivezi\\'c \\& Elitzur 2002).  The infrared emission  is generally associated to a dusty torus surrounding the central source of the AGN. Although proposed by Pier and Krolik (1992,1993), the calculation of the radiative transfer radiation throughout a clumpy   torus has only been  recently performed by Nenkova et al. (2002). The AGN obscuring region is modelled as a toroidal distribution  of dusty clouds placed at different radii from the center, each with  same optical depth. The optically thick clumps are heated by  the nuclear radiation as well as from the radiation originated in other clumps. The dust temperature is much higher in their illuminated side. The authors claim  that heating by diffuse radiation is highly inefficient and can be neglected. Following their results, a clumpy torus accounts for   both, the often observed broad IR bump extending up to 100 \\mum ~and  the absence of the 10 \\mum ~silicate feature in type 1 - on axis view- objects.  In general,  two type of models are used in the literature: a compact torus (radial dimension smaller than a few pc) with a large optical depth, and  a more extended (up to hundreds pc) moderately thick disk. A common characteristic  to all  is that dust survival in the vicinity of the AGN is set by its evaporation limit. Therefore, dust can reach temperatures up to 1000-1500 K. Overall, these models provide a reasonable fit to the observed spectral energy distribution (SED) The torus models usually assume a high equatorial opacity leading to  flat spectra for type 1 and steep spectra for type 2 galaxies. However, as pointed out by Alonso-Herrero et al. (2003), the large number of spectral indices observed in both types of galaxies, as well as the large number of intermediate and type 1 objects  with  steep SEDs, can not be explained by the torus model.  Other simpler descriptions of the IR SED include e.g., that  by Blain, Barnard, \\& Chapman (2003)  who describe the far-IR to submillimeter  SED of dusty galaxies  with different forms of  modified black-body models.  Although most available models provide in a way or other a fair representation of the IR SED, we believe that a more sounded physical approach is needed. An important problem affecting all  these models - those mentioned  above, as well as in the case of Blain's et al. approach, is that  the results are independent of the dust-to gas ratio.  For several years, we are pursuing the development of self-consistent models aimed at explaining both the emission-line and SED  of AGN. These models have proved to be successful in a number of representative  Seyfert galaxies (Contini, Prieto, Viegas 1998a,b, Contini, Viegas, \\& Prieto 2002). In our modelling approach, the infrared-optical continuum is due to dust emission and free-free emission from the clouds powered  by the nuclear radiation, diffuse radiation  and by shocks. Shocks are due to the cloud motion through a dilute gas (Contini \\& Viegas-Aldrovandi 1990, Viegas \\& Contini 1994, Contini \\& Viegas 2000). An important aspect considered in that approach  are the conditions of dust survival when effects as sputtering and grain-gas collisions are taken on board. Thus, these models offer an alternative to the torus model, particularly for the objects showing infrared emission as extended as the narrow-line region.  Overall,  we consider a model to be  well  constrained when both the complete SED and the  emission-line spectrum are simultaneously modeled. When possible, that has been our approach, e.g. in the references quoted above. However, in this paper, we  focus on the optical-IR continuum  only, since in this wavelength range the   best estimate of the nuclear flux emission  and of shock velocities for a reasonable number of Seyfert galaxies of different types (the sample by Alonso-Herrero et al. 2003) is available.  The  paper is structured as follows. A  short review of the processes involving dust is given in Sect. 2. Numerical simulations accounting for the presence of both shocks and photoionization in the NLR are presented  and applied to the  SEDs of Seyfert  galaxies, described in Alonso-Herrero et al. (2003),  in Sect. 3. Concluding remarks appear in Sect. 4. An empirical method  derived from the models,  aimed at  providing a first assessment of the nature of the observed IR  continuum, is  presented in the Appendix. ",
        "conclusions": "In the last few years we have shown that composite models that account for the simultaneous effect of photoionization by a central source and by shocks produce a fair representation of  the SED  of quite a large number of different  AGN types, including  Seyfert 1.5-2 galaxies, NLS1, galaxies where the AGN coexists with starbursts, active galaxies with very extended regions and  luminous infrared galaxies.  In comparison with other models  in the literature, our approach does not resort to any particular  dust/gas geometry but explore the possibilities of  two  mechanisms expected to be present in any AGN, bremsstrhalung and dust reprocessing, to  reproduce the  IR-optical SED.  We discuss the different contributions to the IR continuum of AGN considering that both photoionization from the active center and shocks ionize and heat gas and dust within the surrounding AGN clouds. Collisional processes start to be important in the presence of shocks with velocities of  about 100 \\kms.  Therefore dust grains are heated by collision with gas to relatively high temperatures, depending on \\Vs. Moreover, the grains are sputtered crossing the shock front to different radius depending on the grain initial radius, the plasma density, and the slab width covered by the grains before being decelerated and stopped by collisional drag.  The geometrical thickness of the slabs is determined by the assumption that the physical conditions in the slab are constant, and are constrained by the choice of \\Vs, \\n0, and a$_{gr}$ suitable to the NLR conditions.  Observations indicate that there is a gradient of densities and velocities in the NLR, both increasing towards the active center. In the internal region of the NLR the radiation should be more intense.  We found, however, that in most Seyfert type 2 objects the high velocity-density clouds are shock-dominated.  If these clouds are closer to the centre, the effect of radiation must have been partly cancelled by the effect of an intervening dusty medium. Such dusty medium could be clouds with  a very high d/g, which will emit in the far-infrared.   On the other hand, the  clouds contributing to intermediate and type 1 Seyfert galaxies   are mainly radiation-dominated. These clouds have a direct view of the central radiation as expected from the unified model.  In this paper, the model approach is applied to a sample of galaxies of different Seyfert types for which the best estimates of the nuclear- stellar subtracted- optical-IR SED are available.  With a limited range of parameter values  - shock velocity, pre-shock density, ionizing flux, grain size and dust-to-gas ratio - which are dictated from our previous detailed  modelling of representative Seyfert galaxies - the model results are encouraging as they produce a fair representation of  the SED for all Seyfert type, in particular, they easily produce steep IR SED for broad-line type objects.  Further, a first estimate of the dust-to-gas ratio for some of the clouds is provided. Still, far-IR data  are needed to constrain better those values. Fine tunning of the modelling can be achieved by cross cheking it with the respective nebular spectrum and the SED at other wavelengths.  Nevertheless, the analysis of the optical-IR continuum of a  sample of galaxies provides a general view of the mechanisms producing the emission, as well as on how they differ in the two types of Seyfert galaxies.  A comparison of Seyfert 1.8-2 with Seyfert 1.5-1 SED models confirms that continuum fluxes are generally higher for Seyfert 1.5-1 and that the flux decrease at higher frequencies can be reproduced for most Seyfert 1.8-2 galaxies.  Recall that in multi-cloud models the single-cloud models are summed up adopting relative weights.  We explain, therefore, the flux decrease at higher frequencies by lower relative weights of clouds with higher velocities and/or reached by a strong flux from the active center.  For all type of galaxies, the continuum between 1 and 10 \\mum is found to be  produced by thermal bremsstrahlung emitted by gas with relatively high density, \\n0 ~(1000\\cm3), heated by a  low  velocity shock, \\Vs ~(100 \\kms). These clouds produce  weak lines and therefore their presence can  mostly be accessed from the SED analysis.  Because of their low velocity, we expect that dust  re-emits in the far-infrared.   To provide a first analysis of an AGN SED, an analytical method is presented in the Appendix. In general, galaxy continua are explained by multi-cloud models, leading to extended IR bumps, which must be decomposed into single components. Indeed, the SED profile in the IR of a galaxy depends on the ratio between the dust reradiation peak and the bremsstrahlung emission in the IR, i.e. the dust-to-gas ratios in the single-cloud models, but depends also on the relative weights adopted to sum up the different contributions, which  are constrained by the line spectra.  A linear fit of the bremsstrahlung emission (normalized to the value at 12 \\mum) is adequate for log($\\nu$) $<$ 13.50, while the emission by dust can be represented by a black body (see Appendix)."
    },
    "0311/astro-ph0311392_arXiv.txt": {
        "abstract": "% Previous analytical and numerical work shows that the {\\it generalized interacting stellar winds} model can explain the observed bipolar shapes of planetary nebulae very well. However, many circumstellar nebulae have a multipolar or point-symmetric shape. With two-dimensional calculations, Icke showed that these seemingly enigmatic forms can be easily reproduced by a two-wind model in which the confining disk is warped, as is expected to occur in irradiated disks. In this contribution we present the extension to fully three-dimensional adaptive mesh refinement simulations of such an interaction. ",
        "introduction": "The interaction between a slow disk-shaped inner AGB nebula and a fast stellar wind as a mechanism to produce bipolar bubbles is by now well established. In this {\\it generalized interacting stellar winds} model (see Balick \\& Frank 2002 for a review) the shape of the dense gas around the star is {\\it assumed} to be a disk or a toroid. Analytical (Icke 1988; Icke et al. 1989) and numerical (e.g. Soker \\& Livio 1989; Icke et al. 1992; Mellema et al. 1991) work shows that this mechanism works very well.  However, many circumstellar nebulae have a multipolar or point-symmetric shape (Schwarz 1993; Sahai \\& Trauger 1998). Icke (2003, this volume) proposed that these nebulae are formed in a wind-disk interaction where {\\it the disk confining the fast wind is warped}. One can produce such a warp around a single star, from the combined effects of irradiation and cooling (e.g. Pringle 1996; Maloney et al. 1996).  Whereas Icke's computations were restricted to a two-dimensional proof-of-principle, we now present a first series of fully three-dimensional hydrodynamic computations of such a wind-disk interaction. ",
        "conclusions": ""
    },
    "0311/astro-ph0311501_arXiv.txt": {
        "abstract": "We present \\ha-derived star-formation rates (SFRs) for the galaxy cluster CL~J0023+0423B at $z = 0.845$.  Our $3\\sigma$ flux limits corresponds to a star-formation rate of 0.24~\\smy, and our minimum reliable \\ha~+ [N~II] equivalent width is $> 10$~\\AA, demonstrating that near-infrared narrow-band imaging can sample the star-forming galaxy population in distant clusters. Comparison with spectroscopy shows that the number of false detections is low ($9 \\pm 6$\\%) and that our \\ha \\ equivalent widths are correlated with spectroscopically determined [O II] equivalent widths. A magnitude-limited spectroscopic survey conducted over the same area missed 70\\% of the star-forming galaxies and 65\\% of the integrated star formation. Using Hubble Space Telescope Wide Field Planetary Camera 2 Archive images, we fit Sersic profiles to all galaxies with significant narrow-band equivalent widths and find that equivalent width decreases as the steepness of galaxy profile increases. We find no significant population of early type galaxies with ongoing star formation. The integrated SFR per cluster mass of \\cj \\ is a factor of ten higher than that of the three $z \\sim 0.2$ clusters in the literature with available \\ha\\ observations. A larger sample of $z \\sim 0.8$ clusters spanning a range of cluster masses is needed to determine whether this variation is due to a difference in cluster mass or redshift. ",
        "introduction": "The observed variation of galaxy properties as a function of environment \\citep{hubble31} is evidently an important clue to how galaxies form and evolve. The two classic observations that drive this field of study are that nearby galaxy clusters contain a higher fraction of E/S0 galaxies than observed in the field \\citep[e.g.][]{dressler80,whitmore93} and that the fraction of blue galaxies in a cluster increases as a function of redshift \\citep{butcher84}. The subsequent interpretation of these results has been varied, often conflicting, and not yet distilled. The difficulty, in part, lies in the lack of direct observations that trace star formation in large samples of cluster and field galaxies.  Broad-band colors are most often used to trace star-formation rates \\citep[for example,][]{butcher84}, but these offer a crude measure and are affected by dust and metallicity. Comparisons of emission-line measurements of star formation are complicated by the use of different lines at different redshifts, by the difficulty in getting large samples at high ($> 0.5$) redshifts, and by the dust and metallicity sensitivity of [O~II], the most popular indicator at high redshift. We present our initial attempt to compile a sample of \\ha\\ measurements of star formation in a sample of high-redshift ($\\sim 0.7 - 0.8$) galaxy clusters by presenting the results for our first completed cluster, \\cj.  The advantages of these \\ha\\ data are that they are directly comparable to $z < 0.3$ studies, they directly measure the ionizing flux from young stars, and they are less sensitive to extinction than the common [O~II] measurements.  Even at low redshift, \\ha\\ observations of cluster populations are limited. \\citet{kenn83b} and \\citet{kenn84} use \\ha \\ spectroscopy to study the star-formation properties of four nearby clusters and find that cluster and field spirals of a given morphology have similar star-formation rates (SFRs) in three of four clusters. In contrast, an objective prism study of nearby Abell clusters, \\citet{moss93,moss00} find that the difference in star-formation properties between field and cluster spirals depends on morphology. Cluster \\ha \\ studies have now been pushed out to redshifts of $0.1 - 0.3$ using imaging \\citep{balogh00} and spectroscopy \\citep{couch01,balogh02}.  At intermediate and high redshift the results have depended on the less robust [O II] measure of star formation. The CNOC \\citep{balogh97,balogh98} and MORPHS \\citep{smail97,dressler99, poggianti99} surveys agree that cluster galaxies of all Hubble types have lower SFRs than the same type field galaxies. However, photometric modeling of the CNOC clusters favors a slow decline in star formation \\citep{ellingson91, kodama01b} as one approaches the cluster, while the MORPHS spectroscopy reveal a large population of starburst and post-starburst cluster galaxies, which reflect sudden and dramatic changes in star formation rates. In a cluster at $z = 0.83$, \\citet{vandokkum99} find that although the observed merger rate is significantly higher than the field, there is no sign of excess star formation. \\citet{postman98,postman01} study four $z \\sim 0.9$ clusters and find that cluster galaxies have systematically lower SFRs than field galaxies at similar redshifts. The only result that appears to be consistent throughout all of these studies and across redshift is that cluster galaxies have lower rates of star formation than field galaxies.  The origin of this empirical result is still debated.  Given this situation, one might naturally conclude that significant progress could be made at any redshift, and that all things being equal one should focus at low redshift. However, there are two developments at low redshift that will vastly improve the current situation. SDSS and 2dF have already begun to provide \\ha\\ measurements of star formation in local galaxies as a function of environment \\citep{gomez03, lewis02}, and an ongoing deep \\ha\\ selected survey of nearby clusters will complete the picture \\citep{sakai02}. For high-redshift, one wants to push as high as possible to maximize the evolutionary effects. The redshift regime of $0.7 - 0.8$ is the highest for which there exists a significant number of clusters with complementary data, such as redshifts and morphologies;  however, \\ha\\ moves beyond the optical window into the $J$ band.  To observe the \\ha\\ line, we are undertaking a near-infrared, narrow-band \\ha \\ imaging survey of ten $z \\sim 0.8$ clusters. Here we present results for the first cluster in our sample, \\cj. The goals of this initial paper are to demonstrate the accuracy and sensitivity of the technique and to provide the first comparison between \\ha \\ imaging results for a cluster at $z \\sim 0.8$ and similar surveys of lower-redshift clusters.  We choose CLJ0023+0423B as part of our cluster sample because its redshifted \\ha \\ emission lies in a standard near-infrared narrow-band filter (Figure \\ref{jcont}) and because it has Hubble Space Telescope ($HST$) Wide Field Planetary Camera 2 (WFPC2) imaging and Keck spectroscopy \\citep{postman98,lubin98a}. These additional data provide a necessary complement to our imaging. For example, \\citet{postman98} present a mass estimate for the cluster, $\\rm \\sim 1-5\\times 10^{14} ~{\\it h}^{-1}_{100}~ M_\\odot$, and  \\citet{lubin98a} use WFPC2 imaging to measure an early-type fraction of $\\le 33\\%$, a value comparable to that of the field.  \\begin{figure} \\plotone{f1.eps} \\caption{Trace of narrow-band filter (solid line) plotted with shaded histogram of spectroscopic redshifts from \\cite{postman98}.} \\label{jcont} \\end{figure}  The organization of this paper is as follows.  In \\S\\ref{clj0023_reduction} we describe the observations and data reduction procedure. In \\S\\ref{clj0023_fluxcal} we explain the calibration of the narrow-band images, and in \\S\\ref{clj0023_contsub} we describe the continuum-subtraction process and model the associated uncertainties.  Results from the narrow-band imaging are presented in \\S\\ref{clj0023_results}, and we discuss these results in the context of existing observations in \\S\\ref{clj0023_discussion}.  The paper is summarized in \\S\\ref{clj0023_summary}. We assume $\\Omega_0 = 0.3$ and $\\Omega_\\Lambda = 0.7$ throughout and express results in terms of $h_{100}$, where $\\rm H_0 = \\it h_{100} \\ \\rm {100~km~s^{-1}~Mpc^{-1}}$. ",
        "conclusions": "\\label{clj0023_discussion} In this section we compare the star-formation properties of \\cj \\ with comparable observations for other clusters at lower redshifts. Comparing spectroscopic and imaging surveys is difficult, and comparing SFR measurements obtained with \\ha \\  vs. [O~II] is even more problematic. Because of these difficulties, we limit our comparison to three $z \\sim 0.2$ clusters from the literature for which \\ha \\ imaging or spectroscopy is available.  The three surveys are summarized in Table \\ref{cljlitsfrs}.  Analogous to the Butcher-Oemler blue fraction measurement is the fraction of emission-line galaxies. To provide a fair comparison to the low-redshift \\ha \\ studies, we have to apply four criteria to all clusters.  First, we must sample the same radial fraction of each cluster relative to $R_{200}$.  Then, we must apply the same EW, SFR, and absolute magnitude cuts to all data.  The low-redshift spectroscopic samples are sensitive in terms of SFR and EW but have limited radial coverage; the spectroscopic data for Abell~1689 extends to only $0.27 \\times R_{200}$ \\citep{balogh02}. The imaging survey of \\citet{balogh00} has good areal coverage relative to $R_{200}$ but is sensitive to only the most actively star-forming galaxies with $EW > 50$~\\AA. Thus, a fair comparison among the clusters in Table \\ref{cljlitsfrs} requires that we sample only within $0.27 \\times R_{200}$ and include only galaxies with $EW > 50$~\\AA \\ as emission line. In addition, we must correct the non-emission line galaxy counts because a large faction \\citep[$\\sim$50\\%;][]{maihara01} are not cluster members, and we do not have a large enough field area to do statistical subtraction. The fraction of emission-line galaxies is therefore complicated to calculate and uses only a small fraction of the available data.    \\begin{figure} \\epsscale{1} \\plotone{f12.eps} \\caption{Integrated SFR per cluster mass versus cluster redshift for \\cj \\ and three $z \\sim 0.2$ clusters from the literature.  All $>3\\sigma$ sources within $0.5 \\times R_{200}$ are included. } \\label{clj0023_sfrz} \\end{figure}    \\begin{figure} \\plotone{f13.eps} \\caption{Integrated cluster SFR per cluster mass versus cluster velocity dispersion for CL~J0023+0423B and clusters from literature. All $>3\\sigma$ sources within $0.5\\times R_{200}$ are included. } \\label{clj0023_sfrsigma} \\end{figure}   Perhaps a simpler way to quantify the evolution of cluster SFRs is in terms of the integrated SFR per cluster mass. No correction for background/foreground galaxies is required, and this measure is less sensitive to incompleteness because the integrated SFR is dominated by galaxies whose SFR is easiest to detect. To calculate, we apply the same radial cut to all four clusters, including only those galaxies that lie within a projected radial distance of $0.5 \\times R_{200}$. We choose a maximum radial extent of $0.5\\times R_{200}$ to approximate the radial survey size of AC~114 \\citep[see Table \\ref{cljlitsfrs};][]{couch01}. The radial coverage of Abell 1689 does not extend to $0.5 \\times R_{200}$, so we multiply the integrated SFR by 1.35 to correct for incomplete sampling within $0.5 \\times R_{200}$, where we assume that the galaxy distribution follows the dark matter profile. This correction will still underestimate the integrated SFR if there is a strong increasing radial gradient in SFR. The integrated SFR for Abell 1689 listed in Table \\ref{cljlitsfrs} is corrected for volume. The integrated SFR per cluster mass might be expected to be a function of both redshift and cluster mass (even beyond the normalization by cluster mass one might expect that more massive clusters are more evolved), and in Figure \\ref{clj0023_sfrz} we plot SFR versus redshift for the four clusters and versus cluster velocity dispersion in  Figure \\ref{clj0023_sfrsigma}. These plots exemplify a key difficulty in assessing the evolution of cluster star formation, which is present in all studies of the Butcher-Oemler effect. A larger sample of $z \\sim 0.8$ clusters is needed to break this degeneracy between mass and redshift.   We present \\ha-derived star-formation rates for the galaxy cluster CL~J0023+0423B ($z = 0.845$) and demonstrate that near-infrared narrow-band imaging is an effective and efficient method to sample the star-forming galaxy population in distant clusters. Comparison with spectroscopy shows that the number of false detections is low ($9\\pm 6$\\%), and that our \\ha \\ equivalent widths are correlated with the spectroscopically determined [O~II] equivalent widths. We find that the magnitude-limited spectroscopic survey of this cluster by \\citet{postman98} misses $\\sim$70\\% of the star-forming galaxies and $\\sim$65\\% of the integrated star formation within our image area. We fit Sersic profiles to $HST$ WFPC2 images of all galaxies with significant EW$_R$ and find that EW$_R$ decreases as the light profile steepens.  No significant population of early-type galaxies with ongoing star formation is detected. We compare \\cj \\ to three $z \\sim 0.2$ clusters from the literature for which \\ha \\ data are available and find that the integrated, mass-normalized SFR of \\cj \\ is a factor of ten higher than that of the lower-redshift clusters. Interpreting this difference is complicated by the strong correlation between integrated SFR and cluster mass that we find. Imaging for a larger sample of $z \\approx 0.8$ clusters drawn from the EDisCS  sample\\footnote{http://www.mpa-garching.mpg.de/\\~{}ediscs/project.html}, which spans a range in cluster mass, is ongoing."
    },
    "0311/astro-ph0311209_arXiv.txt": {
        "abstract": "We study the effect of galactic outflows on the statistical properties of  the \\op\\ forest and its correlation with galaxies. The winds are modelled as fully ionised spherical bubbles centered around the haloes in an  N-body simulation of a $\\Lambda$CDM model. The observed flux probability distribution and flux  power spectrum limit the volume filling factor of  bubbles to be less than  $10\\%$. We have compared the mean flux as a function of distance from haloes with the Adelberger \\etal (ASSP) measurement. For a model of bubbles of constant size surrounding the most massive haloes, bubble radii of $\\bsim 1.5\\hmpc$ are necessary to match the high  transmissivity at separations $\\leq 0.5\\hmpc$ but the increase of the transmissivity at small scales is more gradual than observed. The cosmic variance error due  to the finite number of galaxies in the sample increases rapidly with  decreasing separation.  At separations $\\leq 0.5 \\hmpc$  our estimate  of the cosmic variance error is  $\\Delta{\\bar F}\\sim 0.3$, 30\\% higher than  that of ASSP.  The difficulty in matching the rise in the transmissivity at  separations smaller than the size of the fully ionised bubbles surrounding the  haloes is caused by residual absorption of neutral hydrogen lying physically  outside the bubbles but having a redshift position similar to the haloes. The  flux level is thus sensitive to the amplitude  of the coherent velocity shear near halos and to a smaller extent to the amplitude of thermal motions. We find that the velocity shear increases with halo mass in the simulation. A model where  LBGs are starbursts in small mass haloes matches the observations with smaller bubble radii than a model where massive haloes host the LBGs.  If we account for the uncertainty in the redshift position of haloes, a starburst model  with a bubble  radius of $1\\hmpc$ and a volume filling  factor of 2\\% is consistent with the ASSP measurements at the 1-1.5$\\sigma$ level.  If this model is correct the sharp rise of the transmissivity at separations  $\\leq 0.5 \\hmpc$ in the ASSP sample is due to cosmic variance and is expected to become more moderate for a larger sample. ",
        "introduction": "\\label{lyman}   The \\op forest is now widely believed to originate from  an undulating warm ($~10^{4}$K) photoionized IGM.  The \\op optical depth in redshift space due to resonant scattering can be expressed as a convolution of the real space \\hi\\ density along the line of sight with a Voigt profile ${\\cal H}$ (Gunn \\& Peterson 1965, Bahcall \\& Salpeter 1965), \\begin{eqnarray} \\tau(w) &=& \\frac{c \\sigma_0}{H(z)}\\nonumber \\\\ & & \\times \\int_{-\\infty}^{+\\infty}\\dhi(x){\\cal H} \\left[w-x-\\vhi(x),b(x) \\right] dx\\; . \\label{depth} \\end{eqnarray} where $\\sigma_0=4.45\\times 10^{-18}\\cmm$ is the effective cross-section for resonant line scattering and H$(z)$ is the Hubble constant at redshift $z$, $x$ is the real space coordinate, $v_p(x)$ is the  line of sight component of the \\hi\\ peculiar velocity field, ${\\cal H}$ is the Voigt profile, and $b(x)$ is the Doppler parameter due to  thermal/turbulent broadening. For thermal broadening $b(\\vx)=13\\left({{T(\\vx)}\\over 10^4\\kel}\\right)^{1/2}\\kms$, where T is the temperature of the gas.  Both $x$ and $w$ are measured in $\\kms$.  The Voigt profile is normalized such that $\\int{\\cal H}=1$ and can be  approximated by a Gaussian for moderate optical depths, ${\\cal H}=1/\\left(\\pi^{1/2} b\\right)\\times\\exp\\left\\{-[\\omega-x-v_p(x)]^2/b^2\\right\\}$.  The normalized flux obtained by ``continuum fitting'' of the observed spectrum is  related to the optical depth along the line of sight as, \\begin{equation} F(w)\\equiv I_{\\rm obs}(w)/I_{\\rm cont}={\\rm e}^{-\\tau(w)}\\; , \\end{equation} where $\\tau$ is the optical depth, $w$ is the redshift space coordinate along the line of sight, $I_{\\rm obs}$ is the observed flux and $I_{\\rm cont}$ is the flux emitted from the source (quasar) that would be observed in the absence of any intervening material.  On scales larger than a filtering scale which is related to the  Jeans scale , $x_{\\rm J}$, the IGM traces the DM distribution very well. For moderate overdensities the balance between  adiabatic cooling and photo-heating of the expanding IGM  establishes  power-law relations between  temperature, neutral hydrogen density and total gas density (Katz \\etal 1996, Hui \\& Gnedin 1997, Theuns \\etal 1998), \\begin{equation} T_{\\rm g}\\propto \\rho_{\\rm g}^\\beta \\quad {\\rm and}\\quad \\dhi\\propto \\rho_{\\rm g}^\\alpha \\; , \\label{state} \\end{equation} where the parameter $\\beta$ is in the range $0-0.62$, $\\alpha=2-0.7\\beta$, and $\\rho_{\\rm g}$ and $T_{\\rm g}$ are the density and temperature of the gas, respectively.  On scales smaller than the Jeans scale pressure dominates over gravity.  The gas pressure smoothes  the gas distribution compared to the distribution of collisionless dark matter ({\\it e.g.}  Theuns, Schaye \\& Haehnelt 2000). We thus  assume that $\\rho_{\\rm g}(\\vx)\\propto \\rhof(\\vx)$,  where $\\rhof(\\vx)$ is the dark matter density smoothed in a way  that mimics the pressure effects.  We adopt the smoothing used  in SPH simulations (e.g. Springel, Yoshida \\& White 2001)  to estimate the gas density from the dark matter particle distribution.  The \\hi\\ density $\\dhi$ and the smoothed dark matter field $\\rhof$ are then related as \\begin{equation} \\dhi(\\vx)= \\dhihat \\left[\\frac{\\rhof(\\vx)}{{\\bar \\rho}_{\\rm dm}^{\\rm s}}\\right]^{\\alpha}\\; , \\label{neutral} \\end{equation} where ${\\bar \\rho}_{\\rm dm}^{\\rm s}$ is the (volume) average of $\\rhof$ , and $\\dhihat$ is the \\hi\\ density at $\\rhof={\\bar \\rho}_{\\rm dm}^{\\rm s} $. ",
        "conclusions": "\\label{discussion}  We used mock spectra obtained from numerical simulations to investigate the effect of fully ionised bubbles around the galactic haloes hosting LBGS on the \\hi\\ distribution in the IGM as probed by the \\op\\ forest in QSO spectra. By matching the probability distribution and the power spectrum of the flux of mock spectra obtained from a N-body simulation of a $\\Lambda$CDM model with those of observed spectra we derive an upper limit of $10\\%$ on the volume filling factor of bubbles. We further considered the mean transmitted flux at a given distance from galactic haloes.  To model the observational data we computed the mean transmitted flux $\\bar F$ from an annular average of the two-dimensional conditional flux ${\\bar F}(\\pi,\\sigma)$, in the same way as done by Adelberger et al. (2003). As previous investigations we find that the observed decrease of the transmission at separation $1\\leq s\\leq 5\\hmpc$ to a halo is well reproduced by the simulated spectra. It is due to the increased matter density around haloes. We further find that an increase of the transmissivity at the smallest separations $s\\leq 0.5\\hmpc$ requires bubble radii of $\\rw=1-1.5 \\hmpc$. This is two to three times larger than the separation at which the increase of the transmissivity is observed. As discussed by Kollmeier \\etal 2003 this is because in redshift space the velocity shear of the material infalling onto the halo fills in the cavity which is caused by the fully ionised wind bubble in real space. For a model where the haloes are surrounded with bubbles of size $\\rw=1-1.5 \\hmpc$ there is no effect on either the flux PDF or the flux power spectrum.  The flux PDF and the flux power spectrum of our simulated spectra are fully consistent with those of observed spectra. This is because the volume filling factor of the ionised bubbles is still small. We have investigated a variety of scalings of bubble size with halo mass. We also compared a model where LBGs are long-lived and are hosted by the most massive haloes with a model where LBGs are short-lived starburts and are hosted by more numerous less massive haloes.  Unlike Kolatt \\etal (1999) and Wechsler \\etal (2001), we do not identify starbursts as ``colliding haloes'' but choose a fraction of less massive haloes at random. Matching the observed transmission at scales $1\\leq s\\leq 5\\hmpc$ is not difficult and there is certainly more than one way to do so. Introducing redshift errors at the expected level improves the agreement with the observations.  The increase in transmission at small separation in our model spectra spectra is, however, generally more gradual than that of the observed measurements. As other authors we found it challenging to match the observations at the smallest separation.  Numerical simulations including galactic winds and simple analytical estimates as the one in Appendix~\\ref{cooltimescale} agree that the required bubble sizes are difficult to achieve (Croft et al. 2002, Bruscoli et al. 2002). We have thus particularly looked into models which are consistent with the observed measurements and are most ``economical'' in terms of bubble size. The velocity shear around haloes increases with increasing mass. Haloes of lower mass require thus smaller bubble radii to explain the observed increase of the transmission at $s\\leq 0.5\\hmpc$. This may favour a model where LBGs are starburst galaxies hosted by more numerous less massive haloes. We have taken special care in estimating the errors on the flux near galaxies in a similar way as done in the observations.  For small separations we found that the expected error due to cosmic variance is about 30\\% larger than estimated by ASSP. Our starburst model is consistent with all measurements for a bubble radius of $\\rw=1 \\hmpc$ and a volume filling factor of 2\\%.  A bubble size of $\\rw=1 \\hmpc$ appears, however, still difficult to achieve for galactic winds.  If for some reason the velocity shear around haloes is smaller than suggested by our (DM only) simulation this would further reduce the required bubble size and filling factor.  The correlation length for our ``starburst'' model is $2.5\\hmpc$, slightly lower than the observations. Identifying the LBGs with satellites infalling on high mass haloes, as in Kolatt \\etal (1999), would further improve the agreement with the observed clustering properties of LBGs. However, this would be at the expense of matching the observed small-scale flux $\\bar F$, since the velocity shear increases with halo mass.  Therefore we advocate the model in which starbursts occur in small mass haloes. In a $\\Lambda$CDM cosmology of normalisation $\\sigma_8=0.9$, this model yields a reasonable match to the observed clustering of LBGs and transmitted flux $\\bar F$.  The mean transmitted flux of our simulated spectra at the smallest separation is 30 percent smaller than that of the ASSP sample. We would thus expect the observed mean transmission to decrease substantially for larger samples if the model is correct.  A sample with 12 LBGs at separation $s\\leq 0.5\\hmpc$ should reduce the error to $\\Delta{\\bar F}\\sim 0.1$. This should help to discriminate against alternative models like the suggestion of Croft \\etal (2002) that LBGs are located preferentially in low-density environments."
    },
    "0311/gr-qc0311078_arXiv.txt": {
        "abstract": "A method to construct exact general relativistic thick disks that is a simple generalization of the ``displace, cut and reflect'' method commonly used in Newtonian, as well as, in Einstein theory of gravitation is presented. This generalization consists  in  the addition of a new step in the above mentioned method. The new method can be  pictured as a ``displace, cut, {\\it fill} and reflect'' method. In the Newtonian case, the method is  illustrated in some detail  with the Kuzmin-Toomre  disk.  We obtain a thick disk with acceptable physical properties. In the relativistic case two solutions of the  Weyl equations, the Weyl gamma metric (also known as Zipoy-Voorhees metric) and the Chazy-Curzon metric are used to construct thick disks. Also  the Schwarzschild metric in isotropic coordinates is employed to construct another family of thick disks. In all the considered cases  we have non trivial ranges of the involved parameter  that yield thick disks in which all  the energy conditions are satisfied. ",
        "introduction": "Exact solutions of the Einstein equations are associated  to highly idealized physical  systems that have some exceptional geometrical properties. In some cases with a simple exact solution  one can capture a significant part of the physical properties of non trivial systems. Also, in nonlinear theories like general relativity  and fluid dynamics the exact solutions play an important role in numerical analysis. These solutions can be used to test numerical codes and its  outcomes.  Also they can be employed  as initial conditions to describe  more realistic situations, e.g., a static solution can be used as part of the initial conditions  for a full dynamical simulation.  Since the natural shape of an isolated  self-gravitating fluid is axially symmetric, the solutions of Einstein's field equations with this symmetry play a particularly important role in the  astrophysical  applications of general relativity. In particular disk like configurations of matter are of great interest, since they can be used as models of galaxies or accretion disks. Also these disks can be used as starting point to represent  more realistic models in which the bulge and halo of the galaxy are considered.  Solutions for static thin disks without radial pressure were first studied by Bonnor and Sackfield \\cite{BS}, and Morgan and Morgan \\cite{MM1}, and with radial pressure by Morgan and Morgan \\cite{MM2}.  The first solution represent disks made of pressureless dust whereas the second disks with azimuthal pressure but without radial pressure. And the third a disk made of an anisotropic fluid with nonzero radial pressure. The Bonnor-Sackfield disk has a singular rim. These disks are finite.  Several classes of exact solutions of the Einstein field equations corresponding to static \\cite{LP,LO,LEM,LL1,BLK,BLP,LL2,LL3} and stationary \\cite{BL,LZB,GL2} thin disks have been obtained by different authors, with or without radial pressure. Thin disks with radial tension \\cite{GL1}, magnetic fields \\cite{LET1} and magnetic and electric fields \\cite{KBL} have been also studied. The non linear superposition of a disk  and a black hole was first considered by Lemos and Letelier \\cite{LL1}. This solution and its generalizations has been studied in some detail in \\cite{SZZ1,SZZ2,SZ1,ZS1,SZ2,SEM1,SEM2,SEM3}.  Recently the stability of circular orbits of particles moving around black holes surrounded by axially symmetric structures have been considered in \\cite{LET2}. For a recent survey on relativistic gravitating thin disks, see \\cite{SEM4}.  Except for the pressureless disks all the other disks have as source matter with azimuthal pressure (tension)  different from the radial pressure (tension).  However, in some cases these disks can be interpreted as the superposition of two counter-rotating perfect fluids.  A detailed study of the counter-rotating model for the case of static thin disks  is  presented in \\cite{GAG}. Recently, more realistic models of thin disks and thin disks with halos  made of perfect fluids were considered in \\cite{DANLET}.  In all  the disks mentioned above an inverse style method was used to solve the Einstein equations. The metric representing the disk is guessed and then it is used to  compute the source  (energy-momentum tensor). This method was  named by Synge as the g-method \\cite{SYN} in contrast to the  t-method or direct method in which  the source is given and the Einstein equations are solved. The  t-method,  has  been used to generate disks by the Jena group \\cite{KLE1,NM,KR,KLE2,FK,KLE3,KLE4}. Essentially, they  are obtained by solving a Riemann-Hilbert problem. These solutions are highly non trivial, but they deserve special attention because of their clear physical meaning.  In the solutions obtained by the g-method  the well known ``displace, cut and reflect'' method is used. The idea of the method is simple. Given a solution of the vacuum Einstein equations, a cut is make above all singularities or sources. The identification of this solution with its mirror image yields relativistic models of disks. In general, these disks are of infinite extension and   finite  mass.  The aim of this paper is to consider disks beyond the thin disk limit to add a new degree of reality to these geometric  models  of galaxies. Even though in first  approximation the galactic disks can be considered to be  very thin , e.g., in our Galaxy the radius of the disk is  10$kpc$ and its thickness is 1$kpc$. In a more realistic model the thickness of the disk need to be considered. Also it is well known in fluid mechanics  that the addition of a new dimension can make dramatic  changes in the  dynamics of the fluid. In principle, this new dimension    will also change the dynamical properties  of the  the disk source, e.g., its  stability.  In  this paper we  generalize the ``displace, cut and reflect'' method in order to obtain thick disks models from vacuum solutions of Einstein equations.  We shall  replace the surface of discontinuity of the metric derivatives  with a ``thick'' shell in such a way that the matter content of the disk will be described by continuous functions  with continuous first derivatives. This generalization can be named ``displace, cut, {\\it fill} and reflect'' method. The disks obtained with this method, in general, will be of infinite extension and finite mass. Also as in the case of thin  disks, the matter that form the thick disks will not obey   simple equations of state and  in some regions of the disk  the pressure can change  sign given rise to  tensions. Although the energy condition will be  fulfilled. The models of thick disks presented can be considered aS generalizations of models of thin disks studied in references  \\cite{KBL} and \\cite{DANLET}.  The article is divided as follows. In Sec. II we present, in some detail, the main idea of the ``displace, cut, {\\it fill} and reflect'' method in Newtonian gravity. The method is then applied, in Sec. III, to construct relativistic thick disks in Weyl coordinates. We also study  the general expression for the energy-momentum tensor of the disks. The method is illustrated by taking two simple Weyl solutions that lead to thick disks in agreement with all the energy-conditions. In Sec. IV we apply the method to the Schwarzschild solution in isotropic cylindrical coordinates. The disk obtained also satisfy all the energy conditions, this disk have equal azimuthal and radial pressures and different vertical pressure. Finally, in Sec. V, we summarize our main results. ",
        "conclusions": "We presented a method to obtain exact general relativistic thick disks as a generalization of the ``displace, cut and reflect'' method commonly used to obtain Newtonian and relativistic thin disks. The generalization was be done by means of the transformation $z \\to h (z) + b$, where $h (z)$ is an even function of $z$ and $b$ is a positive constant. The function $h(z)$ must be selected in such a way that the metric tensor and and its first derivatives will be continuous across the plane $z = 0$.  All the cases considered leads to thick disks with similar behavior of the energy and Newtonian effective densities: a maximum at the center of the central plane of the disks, the $z = 0$ plane, and then monotonously decreasing for increasing $r$ and $z$.  We found that, when the method is applied to vacuum Weyl spacetimes, the thick disks presents radial tension and vertical pressure. The azimuthal stress is negative at the central region of the disks, then have a positive maximum and finally decrease monotonously for large values of $r$ and $z \\to \\pm a$, where $2a$ is the thickness of the disk. The disks obtained are in full agreement with all the energy conditions.  On the other hand, when the method is applied to the Schwarzschild isotropic metric all the stresses are pressures and have a behavior like the densities. The disks obtained are also in full agreement with all the energy conditions.   We plan  to extend the models of thick disk presented along this lines by considering more elaborate functions $h(z)$ and by the  incorporation of new properties like rotation, either electric or magnetic fields or both. Also we believe that the study of stability in this disks can produce some non trivial results."
    },
    "0311/astro-ph0311515_arXiv.txt": {
        "abstract": "The galaxy power spectrum is now a well-known tool of precision cosmology. In addition to the overall shape, baryon oscillations and the small-scale suppression of power by massive neutrinos capture complimentary information on cosmological parameters when compared to the angular power spectrum of cosmic microwave background anisotropies. We study both the real space and redshift space galaxy power spectra in the context of non-linear effects and model them  based on the halo approach to large scale structure clustering. We consider potential systematic in the cosmological parameter determination when non-linear effects are ignored and the galaxy power spectrum is described with the linear power spectrum scaled by a constant bias factor. We suggest that significant improvements can be made when non-linear effects are taken into account as a power-law contribution  with two additional parameters to be determined from the data. In addition to cosmological parameters through galaxy clustering, such an approach allow a determination of useful information related to astrophysics on how galaxies occupy dark matter halos. ",
        "introduction": "Through a statistical measurement of information related to inhomogeneities in the universe, the galaxy power spectrum provides a strong tool in the present era of precision cosmology. In addition to information related to the primordial power spectrum, characterized by a slope and a normalization, through a turnover at $k_{\\rm eq} \\equiv \\sqrt{2\\Omega_mH_0^2(1+z_{\\rm eq})}$, the overall shape of the galaxy power spectrum captures cosmological information related to the horizon at the matter-radiation equality. In addition to the overall shape, additional features in the galaxy power spectrum allow improved measurements of several cosmological parameters. Similar to the oscillatory features in the CMB anisotropy power spectrum, the matter power spectrum is expected to exhibit the presence of baryons through oscillations. These features are associated with the sound horizon at the end of the Compton-drag epoch  and capture information related to $\\Omega_mh^2$ and $\\Omega_bh^2$ \\cite{EisHu98}.  Unlike well-known oscillatory features in the angular power spectrum of cosmic microwave background (CMB) anisotropies, baryon related oscillations in the matter power spectrum of the large-scale structure are highly suppressed due to the low baryon-to-dark matter ratio. For currently favored $\\Lambda$CDM cosmologies, baryon oscillations in the matter power spectrum have amplitude variations at the level of $\\sim$ 8\\% with effective widths in Fourier space of order $\\Delta k \\sim 0.03$ h Mpc$^{-1}$. While the ideal way to detect baryon oscillations is to measure the three-dimensional dark matter power spectrum directly, unfortunately, there are no useful probes of this quantity. The effects such as the weak gravitational lensing,  which trace matter fluctuations, only provide information related to the projected power spectrum over a  broad window function in the redshift space (e.g., Bartlemann \\& Schneider 2001).  This leads to  an  averaging of  small features such as baryon oscillations such that  they remain undetectable even considering most favorable scenarios.  Galaxy redshift surveys, on the other hand, allow a measurement of the three dimensional power spectrum of the galaxy distribution. Since galaxies are expected to trace matter fluctuations, at least in large and linear scales, the galaxy power spectrum has been pursued for an  observational detection of baryon oscillations. When detected, baryon oscillations are expected to allow significant improvements in parameter determination by breaking various degeneracies associated with the interpretation of CMB data  \\cite{Eisetal98,BlaGla03,Eis02,Coo02a}. While expectations for precision parameter measurements are generally high, these have been tested to some extent with measured power spectra from redshift surveys such as the 2dF Galaxy Redshift Survey (2dFGRS; Colless et al. 2001) when combined with recent results from CMB anisotropy experiments (e.g., Percival et al. 2002). The same surveys have allowed first attempts to detect baryon oscillations, though, there is still no significant evidence for them (e.g., Percival et al. 2001; also, see, Miller Nichol \\& Batuski 2001). This can be understood by the fact that three-dimensional redshift surveys, so far, lack the required volume, with sizes of order $L \\sim 2\\pi/\\Delta k \\sim 300$ h$^{-1}$ Mpc in each dimension, to reliably resolve them.  In addition to galaxy redshift surveys, a number of additional observational efforts are either under way or planned to image the large-scale structure out to a redshift of a few.  These wide-field surveys typically cover tens to thousands of square degrees on the sky and include weak gravitational lensing shear observations with instruments such as the Supernova Acceleration Probe (SNAP; Massey et al. 2003) and the Large Aperture Synoptic Survey Telescope (LSST; Tyson, Wittmann \\& Angel 2000), and observations of the Sunyaev-Zel'dovich effect \\cite{SunZel80} with dedicated small angular scale CMB telescopes such as the South Pole Telescope (SPT; Stark et al. 1998). These surveys are expected to produce catalogs of dark matter halos, which in the case of lensing and SZ surveys are expected to be essentially mass selected \\cite{Witetal01,Holetal00}.  While the halo number count as a function of redshift is a well-known cosmological test \\cite{Haietal01},  one can also consider the additional information supplied by the clustering of halos. In particular, similar to the approach in galaxy redshift surveys, one can measure the power spectrum of clustering associated with these halos, as a whole, and use that power spectrum for cosmological studies. In addition to the shape, cosmological information comes from the redshift evolution of {\\it rulers} that can be calibrated through CMB data \\cite{Cooetal01,BlaGla03}.  The cosmological studies based on the galaxy power spectrum and related halo clustering information, unfortunately, is affected by non-linearities. While sources are expected to trace inhomogeneities in mass, due to the non-linear evolution of gravitational perturbations at late times, the galaxy power spectrum at small scales departs significantly from the linear description corrected for source bias. The enhancement of power due to non-linearities erases oscillatory features \\cite{Meietal99}, while effects such as redshift-space distortions further contribute to their disappearance. Due to the cosmological significance of baryon oscillations, suggestions have  already been made to detect them at redshifts of a few since the non-linear scale at high redshift is expected to move to smaller scales than today \\cite{Eis02,BlaGla03}.  In the case of current measurements, the cosmological interpretation is restricted only for large scales where clustering is expected to be linear. The standard approach to describe the galaxy power spectrum involves the linear power spectrum scaled by a constant bias factor (e.g., Percival et al. 2001). While  slight modifications to the bias description have been considered \\cite{ElgLah03}, there is still no guidance as to how non-linearities affect cosmological parameter measurements and what approaches can be made to improve such estimates. Additionally, by restricting cosmological interpretation to specific range of wave numbers, it is likely that we have only extracted limited information from current data. For example, in the case of the 2dFGRS, analyzes have only been considered out to a wave number of 0.15 h Mpc$^{-1}$, while measurements span to considerably smaller scales.  It is clear that further studies are needed on galaxy clustering well in to the non-linear regime and to understand what information can be extracted from the non-linear regime of clustering, or at least, from transition regime where non-linearities become important.  The purpose of this paper is to understand the onset of non-linearities in the galaxy power spectrum. While detailed numerical simulations are useful for this purpose, here, we make use of the halo approach to galaxy clustering \\cite{Sel00,PeaSmi00,Scoetal01,SheDia01,BerWei02}. The halo approach provides a useful tool to understand certain aspects of clustering and is physically consistent model without any prior assumptions on the biasing nature of galaxies with respect to the dark matter distribution (for a review, see, Cooray \\& Sheth 2002). We study how cosmological parameter measurements can be improved by extracting information from mildly non-linear scales that are  currently ignored. In particular, we pay attention to systematics that are introduced to cosmological parameter measurements when the galaxy power spectrum is modeled simply with a scaled version of the linear power spectrum. We show that parameter estimates are  biased away from the true values and that these departures can be significantly reduced when the non-linear contribution, at least near the regime of transition from linear to non-linear clustering, is modeled as a power-law.Since parameters related to this description is determined from data simultaneously with cosmological parameter estimates, this approach leads to a slight degradation of errors associated with cosmological parameters alone. The main advantage is that these estimates, however, remain either unbiased or biased at a level that is negligible. The suggested approach can be easily implemented with current studies and when interpreting results related to completed 2dFGRS and SDSS surveys \\cite{Yoretal00}, among others. The suggested approach also avoids an arbitrary distinction as to what scales correspond to  the linear regime of clustering while also allowing the maximum extraction of cosmological information from data.  The paper is organized as following. In the next Section, we outline main considerations related to the halo approach to galaxy clustering. We refer the reader to review by Cooray \\& Sheth (2003) for full details of this approach. We make use of the Fisher matrix approach to quantify biases and expected errors on parameter measurements in Section~3. In Section~4, we conclude with a summary.  \\begin{figure*} \\centerline{\\psfig{file=Fig1.eps,width=110mm,angle=-90}} \\centerline{\\psfig{file=Fig2.eps,width=110mm,angle=-90}} \\caption{The galaxy power spectrum in real space (solid lines) and a comparison to the clustering in dark matter (dotted lines). We show the 1- and 2-halo contributions to the galaxy power spectrum and the 1-halo contribution to the dark matter power spectrum. For reference, the dot-dashed line shows the linear power spectrum. The measured data points represent the power spectrum of the IRAS Point Source Catalog Redshift Survey (PCSz; Saunders et al. 2000) as measured by Hamilton \\& Tegmark (2002). In the bottom panel, we consider the transition regime between linear to non-linear clustering. As we discuss later, the transition to non-linear clustering can be described via a combination of the linear power spectrum scaled with a constant, scale-free, bias and an additional contribution which is a power-law.} \\label{fig:power} \\end{figure*} ",
        "conclusions": "The galaxy power spectrum is now a well known tool of precision cosmology. Ongoing surveys such as 2dF and SDSS have already provided a wealth of information related to cosmological parameters based on clustering of galaxies in the presumed linear regime, though measurements span down to much smaller scales than used for cosmological purposes. We have shown here that there is considerable amount on information in the regime where clustering transits from linear to non-linear. The extraction of this information, however, cannot be done alone with the usual approach involving the linear power spectrum and a constant scale-free bias. We have investigated the non-linear behavior of the galaxy power spectrum in both real and redshift space based on the halo approach and have suggested that the on set of non-linearities can be modeled as a simple power-law. The normalization and the power-law slope provide certain information related to halo occupation description in the halo models of galaxy clustering. We have shown that significant improvements can be made when non-linear effects are taken into account through this power-law contribution  with two additional parameters to be determined from the data simultaneously with cosmological parameters. The suggested approach can easily be implemented in current cosmological studies of galaxy clustering and to extract most information from galaxy clustering data to date."
    },
    "0311/astro-ph0311386_arXiv.txt": {
        "abstract": "This paper considers some effects of foreground  Galactic gas on radiation received from extragalactic objects, with an  emphasis on the use of the 21cm line to determine the total $N_{HI}$. In general, the opacity of the 21cm line makes it impossible to derive an accurate value of $N_{HI}$ by simply applying a formula to the observed emission, except in directions where there is very little interstellar matter. The 21cm line can be used to estimate the likelihood that there is significant $H_2$ in a particular direction, but carries little or no information on the amount of ionized gas, which can be a major source of foreground effects. Considerable discussion is devoted to the importance of small-scale angular structure in HI, with the conclusion that it will rarely contribute significantly to the total error compared to other factors (such as the  effects of ionized gas) for extragalactic sight lines  at high Galactic latitude. The direction of the Hubble/Chandra Deep Field North is used as an example of the complexities that might occur even in the absence of opacity or molecular gas. ",
        "introduction": "The Interstellar Medium (ISM) regulates the evolution of the Galaxy. It is the source of material for new stars and the repository of the products of stellar evolution.    But it is a damned nuisance to astronomers seeking to peer beyond the local gas. In this article I treat the ISM as if it were simply an impediment to knowledge, and suggest ways that one might estimate its effects. This topic has taken on increasing importance in recent years as more and more experiments are requiring correction for the ``Galactic foreground'' (e.g., Hauser 2001).  Here the emphasis will be on the use of the 21cm line to determine a total $N_{HI}$, for the 21cm line is our most general tool, and $N_{HI}$ is an important quantity which can be used to estimate  $N_{He}$, E(B--V)  and $S_{100\\mu}$, as well as the  likelihood that there is molecular hydrogen along the line of sight. Some of the points treated here are discussed in more detail in reviews by Kulkarni \\& Heiles (1987), Dickey \\& Lockman (1990; hereafter DL90), Dickey (2002), and Lockman (2002). ",
        "conclusions": "The 21cm line can be a powerful tool for estimating the influence of the Galactic ISM on extragalactic observations, but it must be used with some thought, and it does not give the complete picture.  Unresolved angular structure in the 21cm data is unlikely to dominate the error budget in most directions. For $N_{HI} > 6 \\times 10^{20}$ cm$^{-2}$  it is probable that there is some  molecular hydrogen  along the sight line and the total $N_H$ will be difficult to determine.  But in my opinion, it is the poorly-understood ionized component of the ISM which introduces the most serious uncertainties for directions with little molecular gas, and it affects observations in all directions. We are entering an era when highly accurate 21cm data will be available over the entire sky, and then, (though likely even now) the limits on understanding the influence of the ISM on extragalactic observations will lie not in the uncertainties in $N_{HI}$, but in $N_{H_2}$ and $N_{He}$, and in the relationship between the dust and the gas.  Every observation of an extragalactic object is an opportunity to learn something about the ISM.  We should not let such opportunities go to waste."
    },
    "0311/astro-ph0311179_arXiv.txt": {
        "abstract": "We performed H$\\alpha$ imaging observations of 22 luminous infrared galaxies to investigate how the distribution of star-forming regions in these galaxies is related to galaxy interactions. Based on correlation diagrams between H$\\alpha$ flux and continuum emission for individual galaxies, a sequence for the distribution of star-forming regions was found: very compact ($\\sim$ 100~pc) nuclear starbursts with almost no star-forming activity in the outer regions (type~1), dominant nuclear starbursts $\\lesssim$ 1~kpc in size and a negligible contribution from the outer regions (type~2), nuclear starbursts $\\gtrsim$ 1~kpc in size and a significant contribution from the outer regions (type~3), and extended starbursts with relatively faint nuclei (type~4). These classes of star-forming region were found to be strongly related to global star-forming properties such as star-formation efficiency, far-infrared color, and dust extinction. There was a clear tendency for the objects with more compact distributions of star-forming regions to show a higher star-formation efficiency and hotter far-infrared color. An appreciable fraction of the sample objects were dominated by extended starbursts (type~4), which is unexpected in the standard scenario of interaction-induced starburst galaxies. We also found that the distribution of star-forming regions was weakly but clearly related to galaxy morphology: severely disturbed objects had a more concentrated distribution of star-forming regions. This suggests that the properties of galaxy interactions, such as dynamical phase and orbital parameters, play a more important role than the internal properties of progenitor galaxies, such as dynamical structure or gas mass fraction. We also discuss the evolution of the distribution of star-forming regions in interacting galaxies. ",
        "introduction": "Many statistical studies of interacting galaxies have shown that galaxy interactions can enhance star-forming activity \\citep[e.g.,][]{larson78,kennicutt87,hummel87,bushouse88,xu91}. Luminous infrared galaxies (LIRGs; $L_{\\rm IR}$[8--1000 \\micron] $\\geq$ $10^{11}$ $L_\\odot$) are galaxies which emit the bulk of their energy in the far-infrared (FIR) and often show signs of interaction, such as tidal tails, multiple nuclei, or disturbed outer envelopes \\citep{sanders96}. LIRGs are considered to be extreme objects, where strong starbursts are induced by galaxy interactions, because in many such objects star formation can account for the infrared emission.  Theoretical work supports the idea that interactions play an important role in stimulating starbursts in galaxies. Numerical simulations of merging gas-rich spiral galaxies show that during the merging process, gas clouds lose their angular momentum and flow into the circumnuclear region of the host galaxies \\citep[e.g.,][]{barnes91,mihos96}. The resulting high concentration of molecular gas can fuel star-formation activity in the circumnuclear region. This is consistent with compact nuclear starbursts and gas condensation in ultra-luminous infrared galaxies (ULIRGs; $L_{\\rm IR}$ $\\geq$ $10^{12}$ $L_\\odot$), as observed in mid-infrared \\citep[MIR,][]{soifer00}, radio continuum \\citep[hereafter C91]{condon91}, and CO emission \\citep[e.g.,][]{downes98,bryant99} data. Therefore, a nuclear starburst triggered by gas inflow has generally been assumed to be the mechanism for producing enhanced star formation in interacting galaxies.  However, the diversity of star-forming and morphological properties in LIRGs, and the details of the physical processes that may enhance star-forming activity, are not fully understood. For example, several paired systems of LIRGs show intense star formation in the region of overlap of the two galaxies \\citep{xu00,mirabel98,haas00,soifer01,lefloch02}, which is not predicted by numerical simulations. Moreover, some observational studies have suggested that the distribution of star-forming regions in LIRGs depends on pair separation \\citep[hereafter H99]{hwang99}, gas content \\citep{mihos98}, or degree of morphological disturbance \\citep{xu00}. Although these observations yield intriguing clues, the samples are small (4--8 objects) and the mechanism that determines the distribution of star-forming regions in LIRGs remains unclear. As for the morphological properties, observational studies of $IRAS$ galaxies have revealed a weak dependence of global star-forming properties, such as infrared luminosity and FIR color, on the interaction strength or morphological class \\citep[e.g., ][]{xu91,lutz92}. These studies suggest that in order to better understand the physical processes in interaction-induced star-forming activity it is necessary to determine how the star-forming properties of interacting galaxies depend on the dynamical or physical properties of the system. On the other hand, even if objects with similar $L_{\\rm IR}$ or FIR color are considered, there is a significant variation in morphological appearance. For example, 20--30\\% of LIRGs with 10$^{11}$ $L_\\odot$ $<$ $L_{\\rm IR}$ $<$ 10$^{12}$ $L_\\odot$ are apparently single galaxies, while the remaining objects are closely interacting or merging galaxies \\citep{sanders96}. It is unclear why these galaxies are actively forming stars at the same level as strongly interacting galaxies, and whether the single-galaxy LIRGs have nuclear starbursts.  The distribution of star-forming regions is expected to be connected to the process by which star-formation activity is enhanced \\citep[e.g.,][]{mihos98}, and therefore it is important to investigate the distribution of star-forming regions in LIRGs and its dependence on morphological, physical, and dynamical properties. For this purpose, we have undertaken an H$\\alpha$ imaging survey of LIRGs. In contrast to previous H$\\alpha$ imaging surveys of interacting galaxies \\citep{kennicutt87,bushouse87,combes94,domingue03} and infrared-selected galaxies \\citep[hereafter AHM; \\citealt{lehnert96,xu00,dopita02}]{armus90}, all of our sample objects are LIRGs, and they are not selected based on morphological properties or FIR color. This should enable us to focus on activated star formation, and may be useful for examining the relationships between the distribution of star-forming regions and other properties. As LIRGs are dust-rich systems, the H$\\alpha$ emission, especially the total H$\\alpha$ flux \\citep{dopita02}, can be significantly affected by dust extinction. However, \\citet{lehnert96} presented evidence that the half-light radius of H$\\alpha$ emission is a useful indicator of the size of the star-forming region. Therefore, H$\\alpha$ emission may be used to trace the spatial distribution of star-forming regions within individual galaxies. An advantage of H$\\alpha$ emission is its large range of spatial scales, i.e., it is sensitive to small ($\\sim$ 1\\arcsec) and large ($\\sim$ 1\\arcmin) scale emission. Although observations in MIR or radio continuum emission can provide the highest spatial resolutions ($\\sim$ 0\\farcs1 or smaller), such observations are insensitive to diffuse emission \\citep[C91;][]{soifer00,soifer01}, which can extend well beyond several arcseconds. Continuum images also offer important information, since red images taken through a tunable filter are almost free from emission lines and can therefore be used to trace the distribution of old stars. This enables a direct comparison of the distribution of old and newly-formed stars.  In \\S~2, we present the sample selection and the observational details. In \\S~3, we describe the results of the observations for the individual objects and analyses of the resulting dataset. The observational results are discussed in \\S~4, and the conclusions and summary are given in the final section. ",
        "conclusions": "\\subsection{Distribution and Global Properties of Star Formation} \\label{section_distribution_property}  In \\S~\\ref{section_ha_con_diagram}, it is shown that the size of the nuclear star-forming region increases from type~1, through type~2, to type~3 objects. In the type~3 objects, there are substantial contributions from the outer regions to the total H$\\alpha$ flux. The star-forming activity in the type~4 objects is dominated by extended starbursts. If the distribution of molecular gas clouds is similar to that of star-forming regions, the density of gas clouds may also vary from type to type, being highest in the type~1 objects and lowest in type~4 objects. The higher density of gas clouds in galaxies with compact star formation will lead to a higher SFE, because of the non-linear dependence of the star-formation rate on gas density \\citep{kennicutt98}. The high efficiency of star formation will then produce more photons to heat the dust and, consequently, raise the dust temperature \\citep{young86}. In Figure~\\ref{global_differ}a, galaxies with a strong concentration (type~1 and type~2) show the largest values of $f_{60}/f_{100}$, i.e., the highest dust temperature, whereas galaxies that are dominated by extended starbursts (type~4) show the smallest values, with type~3 lying between the two extremes. Except for type~1, the $f_{25}/f_{60}$ ratio shows a similar tendency to $f_{60}/f_{100}$, although the difference between the types is less clear. The small values of $f_{25}/f_{60}$ in the type~1 objects are also expected to be caused by the compact distribution of nuclear star-forming regions, as discussed in \\S~\\ref{section_type12}. Therefore, the difference in the distribution of star-forming regions between the types can explain their segregation in Figure~\\ref{global_differ}a.  In addition to the compactness of the nuclear star-forming regions ($\\sim$ 100 pc), the type~1 objects are remarkable in the small amount of molecular gas present and the nearly complete lack of star-forming activity in the outer regions. These properties suggest that they are in a late stage in the star-formation history of LIRGs \\citep{gao99}, and that most of the remaining gas has fallen to the nuclear region. In addition, all of the type~1 objects are either OH megamaser sources or OH absorbers. This suggests that there is a strong radio source in the type~1 objects, located behind an extended molecular-gas screen. It is notable that, although LIRGs with $L_{\\rm IR}/L_{\\rm H\\alpha}$ as large as type~1 are rare, Arp~220, the prototypical ULIRG, shows comparably faint H$\\alpha$ emission and large $L_{\\rm IR}/L_{\\rm H\\alpha}$ ratio (AHM). The FIR color and the presence of an OH megamaser source in Arp~220 are also similar to type~1 objects. Thus, it is probable that, as in Arp~220, the energy source of type~1 objects is a compact starburst deeply embedded in dusty molecular clouds \\citep[see also \\citealt{baan87}]{downes98,sakamoto99}. This would cause the large $L_{\\rm IR}/L_{\\rm H\\alpha}$ and the small $f_{25}/f_{60}$ seen in type~1 objects.   From the above discussion, we conclude that differences in the distribution of star-forming regions causes the differences in star-forming properties among the types of object. Therefore, the H$\\alpha$--continuum diagram combined with the H$\\alpha$-flux histogram (Figure~\\ref{ha_cont}) is a useful tool for investigating the properties of star-forming activity in LIRGs.   \\subsection{Morphological Consideration} \\label{section_morph_type}  In this subsection, we investigate the relationship between star-forming properties and the morphological class of the sample objects. As the object types defined in \\S~\\ref{section_ha_con_diagram} were found to be strongly correlated with the distribution and global properties of star formation, it may be useful to compare our type classification with the morphological class of the galaxies.  In Figure~\\ref{morph_type_hist}, histograms of the types are shown for each morphological class. The type~1 and type~2 objects are relatively small in number and are both dominated by nuclear starbursts, so the objects of the two types are combined and labeled as \"Type 1/2\" in Figure~\\ref{morph_type_hist}. The morphological classes are defined in \\S~\\ref{section_hacon_image}. Dynamical disturbance due to galaxy interactions is expected to be the strongest in \"mergers\" and to gradually become weaker in the order \"close pairs\", \"pairs\", and \"singles\". Therefore, Figure~\\ref{morph_type_hist} implies that the more severely disturbed objects tend to have a more compact distribution of star-forming regions. Although there are several exceptions to this rule, we suggest that the degree of dynamical disturbance plays an important role in determining the distribution of star-forming regions.  From MIR imaging of seven LIRGs with morphological properties indicating the early/intermediate stages of merging, H99 found that the peak-to-total flux ratio, which is a measure of the concentration of star-forming regions (see \\S~\\ref{section_ha_con_diagram}), increases as the projected separation of the interacting pairs becomes smaller. Morphologically, their entire sample corresponds to the class of \"close pairs\" in our classification scheme, and their FIR colors are similar to our type~2 and type~3 objects. In addition, the objects with larger projected separations have FIR colors similar to type~3 objects, and the objects with smaller projected separations have FIR colors similar to type~2 objects. Therefore, H99's and our results are complementary: whereas H99 revealed the detailed dependence of the distribution of star-forming regions on the interaction stage in a morphologically selected sample, our results show that the dependence can be extended to a sample with a broader range of morphological properties. H99 interpreted their results as a manifestation of the evolutionary sequence: as the merging process advances, star formation becomes more active and develops a more concentrated distribution. A similar interpretation may also be applicable to our sample, with the exception of the single galaxies. This issue will be discussed in the next subsection.  \\citet{mihos98} and \\citet{xu00} also found connections between dynamical states and the distribution of star-forming regions. From imaging Fabry--Perot observations of four ULIRGs, \\citet{mihos98} found that the H$\\alpha$ emission is more spatially concentrated in the later stages of a galaxy interaction. \\citet{xu00} found that relatively undisturbed systems show evenly distributed MIR and H$\\alpha$ emission in the disks of component galaxies, while severely disturbed objects show dominant nuclear components and enhanced emission in the overlap regions of the pairs. As the star-forming activity of their sample objects differs significantly from ours, we do not make detailed comparisons with these studies. Nevertheless, it is notable that these studies obtained qualitatively similar results to ours. Therefore, it is almost certain that the strength of dynamical disturbance plays a critical role in determining the distribution of star-forming regions. Our sample size, the restricted range of star-formation rates, and the broad range of morphological properties of our sample objects explain why such a clear relation emerges in Figure~\\ref{morph_type_hist}.  With respect to global star-forming properties, several statistical studies of $IRAS$ galaxies, covering a wide range of morphological properties, have shown that they are weakly connected to the interaction strength \\citep[e.g., SSS;][]{xu91,lutz92}. Our results suggest that this is because the interaction strength controls the distribution of star-forming regions: the stronger the dynamical disturbance, the more compact the region to which the star-forming activity is confined, resulting in the more extreme properties of higher SFE and higher dust temperature.   \\subsection{Evolution of the Distribution of Star-forming Regions} \\label{section_evolution}  In this subsection, we consider the relation between the type and the star-formation history of LIRGs.   There are remarkable differences in the size of nuclear star-forming regions and in the relative contribution of the outer regions to the total H$\\alpha$ flux between the different types of galaxy. These properties, combined with the connection between type and morphological class, suggest a sequence in the distribution of star-forming regions in interacting galaxies. At an early stage of interaction, active star formation occurs in both the nuclear and outer regions (type~3). As the process of interaction proceeds, the infall of gas clouds activates nuclear star formation, making it the dominant energy source of the system (type~2). By the final phase, most of the gas clouds have fallen into the nuclear region, and the star-forming activity is confined to the nucleus (type~1). In this process, the size of the nuclear star-forming region also becomes gradually smaller.   The evolutionary sequence above is similar to that described by \\citet{scoville01}, which invokes star formation induced by cloud-cloud collisions \\citep{scoville86}. In their sequence, the cloud collision rate is increased by tidally enhanced velocity dispersion in the individual disks in the initial phase, later by the passage of the disks through each other, and lastly by concentration of the gas at smaller radii in the merged nucleus. Therefore, star formation in the outer regions is expected to be important at an earlier phase of galaxy interaction, which is consistent with the observed sequence described above. As for the variation in size of the nuclear star-forming regions, a probable cause of the slightly larger ($\\gtrsim$ 1 kpc) size in type~3 objects is a nuclear resonance ring, such as that observed in NGC~7771 (see \\S~\\ref{section_type3}). The circumnuclear star-forming ring might shrink with time as a result of the combined effect of a non-axisymmetric potential and the self-gravity of gas \\citep{wada92} or due to the friction exerted by background stars on gas clouds \\citep{combes92}. The timescale of these processes is of the order of $10^8$ years. As this is shorter than the merger timescale ($\\sim$ $10^9$ years), and in some cases strong gaseous inflows can occur at early stages of galaxy interactions \\citep{mihos96}, interacting galaxies in the type~3 phase can evolve into the type~2 or type~1 phase before the final merger. Therefore, the absence of a one-to-one correspondence between the morphological class and the type (Figure~\\ref{morph_type_hist}) does not conflict with this evolutionary sequence.  Since $L_{\\rm IR}$ may vary significantly throughout the star-formation history of LIRGs \\citep[e.g.,][]{xu00,murphy01}, type~3 objects may not evolve simply into type~2 and then into type~1. Our suggestion is that the distribution of star-forming regions in interacting galaxies follows the sequence above and, during some phases, the galaxies are recognized as LIRGs. Although this evolutionary model seems to succeed in explaining the processes that occur in the evolution of the distribution of star-forming regions, detailed studies of individual objects are needed to examine the validity of these conclusions.   As many of the type~4 objects are singles, we consider them separately. With the exception of NGC~828, the star-forming activity in type~4 objects is characterized by either prominent spiral arms (NGC~877, NGC~958, NGC~3110) or the brightest H$\\alpha$ source being located in the outer regions (NGC~695, NGC~834, NGC~5653), with both of these groups showing a knotty appearance in H$\\alpha$. In addition, their morphological properties are related to the distribution of star-forming regions. All of the objects with prominent spiral arms have a small companion at a projected separation of 20--30~kpc, assuming that the faint galaxy $\\sim$ 70\\arcsec \\ southeast of NGC~958 has the same redshift as NGC~958. The other objects have peculiar morphologies, with no galaxy-like object within a projected radius of 100~kpc. Since strong spiral arms can be induced by the presence of a small companion \\citep[e.g.,][]{toomre81}, this result suggests that the star-forming activity in the first group is enhanced by the presence of such tidally excited spiral arms. However, it is unclear whether such a process is efficient enough to induce active star formation and produce a LIRG. As for the latter group, in spite of the fact that they appear to be isolated galaxies, the structural peculiarities and the asymmetric distribution of star-forming regions suggest that their star-forming activity is also induced by galaxy interactions. This result could be explained if they are remnants of either minor mergers or high speed tidal encounters. The lack of detailed observational information about the type~4 objects makes it difficult to examine these possibilities. Two-dimensional kinematic observations are needed to reveal the nature of these objects, and may be useful in further developing our understanding of star-forming activity in LIRGs."
    },
    "0311/astro-ph0311453_arXiv.txt": {
        "abstract": "The aperiodic X-ray variability in neutron star and black hole X-ray binaries (XRBs), and active galactic nuclei (AGN) shows a characteristic linear relationship between rms amplitude and flux, implying a multiplying-together or `coupling' of variability on different time-scales.  Such a coupling may result from avalanches of flares, due to magnetic reconnection in an X-ray emitting corona.  Alternatively this coupling may arise directly from the coupling of perturbations in the accretion flow, which propagate to the inner emitting regions and so modulate the X-ray emission.  Here, we demonstrate explicitly that the component of aperiodic variability which carries the rms-flux relation in the accreting millisecond pulsar SAX~J1808.4-3658 is also coupled to the 401~Hz pulsation in this source.  This result implies that the rms-flux relation in SAX~J1808.4-3658 is produced in the accretion flow on to the magnetic caps of the neutron star, and not in a corona.  By extension we infer that propagating perturbations in the accretion flow, and not coronal flares, are the source of the rms-flux relations and hence the aperiodic variability in other XRBs and AGN. ",
        "introduction": "Aperiodic X-ray variability, occuring over a broad range of time-scales in the form of `flicker noise' or `band-limited noise' is a ubiquitous characteristic of X-ray binary systems \\citep{vdk95} and active galaxies (\\citealt{mch88,utt02,mar03}).  The origin of this variability is not well understood.  Recently, a new fundamental property of the aperiodic variability in XRBs and AGN was discovered: the rms-flux relation \\citep{utt01}.   Specifically, there exists a strong linear correlation between the amplitude of X-ray variability on short time-scales and the X-ray flux as measured on longer time-scales (\\citealt{utt01,ede02,vau03}).  Since the flux variations on all the time-scales probed in these studies are dominated by the aperiodic variability, this linear `rms-flux relation' implies that aperiodic variations on longer time-scales modulate the aperiodic variations on shorter time-scales, i.e. the aperiodic variations on different time-scales are coupled together.  Importantly, the observed linear rms-flux relations, which occur over a broad range of time-scales, cannot be explained if the variability is produced by simple shot-noise models where the light curve is produced by the summation of randomly occuring shots or flares which are independent of one another (see \\citealt{utt01} for a detailed discussion). Instead the shots must be coupled together. For example, the shots may correspond to fractal structures, e.g. flares that break into smaller structures, which in turn break up and so on, producing a coupling between larger slower variations and smaller faster ones. Alternatively, the required coupling may be produced if flares trigger avalanches of smaller flares (e.g. similar to the model of \\citealt{ste96}). A natural interpretation of such models is that they correspond to flares due to magnetic reconnection in the corona \\citep{pou99}, so we refer to this class of explanations for the variability as the `coronal-flare' or CF model.  Alternatively, a linear rms-flux relation can be naturally produced by the model of \\citet{lyu97}, where the variability is caused by variations in accretion rate occuring at different radii (with slower variations occuring at larger radii), which propagate inwards and modulate the X-ray emitting region and hence the X-ray light curve.  In this model, the amplitudes of accretion rate variations produced at a given radius scale with the local accretion rate at that radius (which is driven by the slower variations that originated at larger radii) and hence variations on different time-scales are coupled together.  In fact, any model of propagating perturbations in the accretion flow (e.g. density waves, \\citealt{mis00}) could produce a linear rms-flux relation provided the perturbations can couple linearly together.  We will call this class of explanations for the variability the `propagating perturbation' or PP model.  In this Letter, we point out that the CF and PP models can be distinguished in the case of accreting neutron star XRBs, which show a linear rms-flux relation in their aperiodic X-ray variability, provided that it can be shown that the aperiodic X-ray variability containing the rms-flux relation originates on or close to the neutron star surface, and hence cannot be associated with coronal flares. We demonstrate explicitly that the accreting millisecond X-ray pulsar SAX~J1808.4-3658 fulfils these criteria, so that the CF model can be ruled out in this case and the PP model is the most likely explanation for the aperiodic variability and the rms-flux relation in this system. We then infer that the PP model is also the correct explanation for the aperiodic variability which shows a linear rms-flux relation in other XRBs, and AGN, before discussing the further implications of this result. ",
        "conclusions": "We have demonstrated that the aperiodic variability which carries the linear rms-flux relation in the accreting millisecond pulsar SAX~J1808.4-3658, is coupled to the 401~Hz pulsation in this source, and hence the component of X-ray emission which contains the rms-flux relation is produced at the magnetic caps of the neutron star.  We conclude that the aperiodic variability is produced by inward-propagating perturbations in the accretion flow on to the neutron star, with the rms-flux relation produced by a coupling of perturbations produced on different time-scales.  By extension, this result suggests that the same mechanism produces the rms-flux relations observed in black hole XRBs and AGN, so that the aperiodic X-ray variability in all these diverse systems is caused by perturbations in the accretion flow, and not by flares due to magnetic reconnection in the corona.  \\subsection*"
    },
    "0311/astro-ph0311103_arXiv.txt": {
        "abstract": "For crucial tests between theory and experiment, ab initio close coupling calculations are carried out for photoionization of several oxgyen ions: O II, O III, O IV, O V. The relativistic fine structure and resonance effects are studied using the R-matrix and its relativistic variant the Breit Pauli R-matrix (BPRM) approximations. Extremely detailed comparison is made with high resolution experimental measurements carried out in three different set-ups: Advanced Light Source at Berkeley, synchrotron radiation experiment with a new modulator at University of Aarh\\\"{u}s, and synchrotron radiation experiment at University of Paris-Sud. The comparisons illustrate physical effects in photoionization such as (i) fine structure, (ii) resolution, and (iii) metastable components. Photoionization cross sections $\\sigma_{PI}$ of the ground and a few low lying excited states of these ions obtained in the experimental spectrum include combined features of these states. The general features of the measured $\\sigma_{PI}$ were predicted in the earlier calculations in LS coupling. However, due to higher resolution achievement in the recent experiments, the theoretically calculated resonances need to be resolved with much finer energy mesh for precise comparison. In addition, prominent resonant features are observed in the measured spectra that can form from transitions allowed with relativistic fine structure, but not in LS coupling. Present results include calculations from both LS coupling and the relativistic BPRM method to decipher the features in photoionization spectra. $\\sigma_{PI}$ are obtained for ground and metastable (i) $2s^22p^3(^4S^o,^2D^o, ^2P^o)$ states of O II, (ii) $2s^22p^2(^3P,^1D,^1S)$ and $2s2p^3(^5S^o)$ states of O III, (iii) $2s^22p(^2P^o_J)$ and $2s2p^2(^4P_J)$ levels of O IV, and (iv) $2s^2(^1S)$ and $2s2p(^3P^o,^1P^o)$ states of O V. Most of the computed theoretical resonance structures are seen experimentally, although there are some discrepancies. One of the main conclusions of the present study is that resonances in photoionization cross sections of the ground and metastable states can be a diagnostic of experimental beam composition, with potential applications to astrophysical and laboratory plasma environments. ",
        "introduction": "\\label{sec:level1}   Photoionization cross sections of many atoms and ions were studied in detail with resonance structures for ground and many excited states under the Opacity Project (OP, \\cite{op}). The close coupling (CC) R-matrix method, as employed under the OP, accurately considers the numerous autoionizing resonances in cross sections along interacting (overlapping) Rydberg series. The study under the OP resulted in new features in the cross sections, such as prominent PEC (photo-excitation of core) resonances. The radiative work under the OP has been extened to collisional and radiative processes under the Iron Project (IP, \\cite{ip}) with inclusion of relativistic effects using the BPRM method. However, while the methods used are state-of-the-art, possible deficiencies are: (i) the resonance structures in the OP/IP cross sections are not always fully resolved, and (ii) the important near-threshold resonances may not be accurate owing to inadequate eigenfunction expansions. A prime need therefore is to benchmark the vast amount of theoretical OP/IP data that is now being utilized for global applications such as calculation of stellar opacities, photoionization rates in astrophysical models.  On the experimental side, accurate photoionization cross sections are being measured at several places, such as the merged ion-photon beam facility at the University of Aarh\\\"{u}s (e.g. for C II \\cite{kjl}), the Electron Cyclotron Resonant Ion Source (ECR) at the University of Paris-Sud (e.g. for Xe V - Xe VIII \\cite{fw}), and the Advanced Light Source (ALS) at Berekely (e.g. for O II \\cite{als}). Each of these three sophisticated experimental set-ups has different advantages in measurements, and the results often complement each other. Cross sections for O II have been measured at ALS for near threshold features in $\\sigma_{PI}$ of metastable states \\cite{als}, including a wider photon energy range \\cite{o2a}, and at the University of Aarh\\\"{u}s \\cite{o2b}. Detailed photoionization cross sections have recently been measured for the ground and low excited states of multiply charged oxygen ions for the first time for O III, O IV, and O V, at the synchrotron radiation source at University of Paris-Sud \\cite{oions}. Study of oxygen ions is of great importance because of their abundances in terrestrial and planetary atmospheres and in astrophysical objects.  Comparison of all these results is extremely crucial to evaluate the accuracy of both the theoretical and experimental results, and to study the resonance phenomena. Theoretical work can also provide spectroscopic identification of resonances, and a quantitative measure of the mixture of states in the experimental beam. The resonant structures and the background shapes of the recently measured photoionization cross sections of oxygen ions are found to be in close agreement overall with the some ealier theoretical cross sections \\cite{snn1}. However, certain features in the measured cross sections are not seen prominently or are missing in the calculated cross sections.  One of the aims of this study is to study relativistic effects and fine structure. We distinguish these in terms of relativistic mixing and fine structure splitting with reference to pure LS coupling results. While relativistic effects per se might not be of great importance for these light ions, fine structure splitting manifests itself in many experimentally observed features. In addition, relativistic mixing via E1 transitions (e.g. singlet-triplet, doublet-quartet) results in additional resonances not allowed in LS coupling. The present study investigates the LS coupling and fine structure effects together with resonance phenonmena in detail. Thus to begin with it is necessary to do LS coupling calculations at high resolution. Comparison with experiment then reveals whether fine structure is also to be considered in order to analyze experimental data. In such cases the relativistic BPRM method is employed to compute cross sections.  Detailed comparison of experimental and theoretical resonance positions and shapes also highlights the importance of correlation effects in theoretical calculations. Resonances may be shifted with respect to their respective thresholds of convergence, indicating missing correlation configurations in the wavefunction expansion. This again is crucial since the inherent accuracy of the close coupling calculations depends on the completeness of the eigenfunction expansions, in particular short-range correlation at low energies.  Another motivation for these studies is that the R-matrix method, as developed for the OP/IP work, has been extended to a theoretically self-consistent treatment of photoionization and unified electron-ion recombination including radiative and dielectronic recombination in an ab initio manner\\cite{np,znp,petal}. The total recombination cross sections have been benchmarked for accuracy by comparing with several experimental cross sections \\cite{znp,petal}. But while recombination cross sections require total contributions of photorecombination cross sections from all bound levels, obtained from the correnpoiding $\\sigma_{PI}$ using detailed balance, comparison of photoionization cross sections with experiment involves individual features of $\\sigma_{PI}$ of single states. For instance, photoionization cross sections including relativistic effects in the close coupling approximation, as developed under the IP, was first benchmarked by comparing with the measured $\\sigma_{PI}$ of C II \\cite{snn2}.  We outline the theory and calculations in Sections II and III respectively, followed by a discussion of results for the four ions O~II~-~O~V in Section IV, and the conclusions in Section V. ",
        "conclusions": "Photoionization cross section of oxygen in several stages of ionization have been obtained for the ground and low lying excited states. These cross sections are compared with measured cross sections at three different experimental set-ups.  While the experimental measurement includes combined features of resonances belonging to different Rydberg series of states, and of a mixture of different states, theoretical work can evaluate them for individaul states.  Certain features in the measured cross sections are not seen prominently or are missing in previous calculations. Present report shows that the missing features are due to (i) lower resolution of previous calculations or (ii) due to relativistic fine structure effects, not included in the previous calculations. For most cases of this relatively low Z element, $LS$ coupling calculations are often adequate. However, for some metastable states, fine structure is crucial for the low energy resonances which allowed only the couplings of fine structure levels. It is the fine structure components, rather than the relativistic effects, that is important for the formation of these high peak resonances with nearly zero backgound. Inclusion of relativistic effects will certainly be more exact and accurate way to carry out the calculations."
    },
    "0311/astro-ph0311429_arXiv.txt": {
        "abstract": "In this brief overview of the stellar populations of galaxies in clusters I highlight some of the most recent results, including the downsizing of the k+a population and the star formation--local density relation.  I conclude discussing some open questions regarding the interpretation of the observational results, and speculating on the current meaning of the words ``primordial'' and ``environmental''. ",
        "introduction": "Studying the evolutionary histories of galaxies in clusters has always been an endeavour to understand how galaxies (all galaxies) came to be as we observe them. In recent years, this quest can be better described as an effort to comprehend how galaxies formed and evolved in a ``changing'' environment. In fact, the common framework for these studies is now a hierarchical universe, in which galaxies can and often change environment during their evolution. When a galaxy infalls into a cluster, it can do so as a single isolated galaxy, as a pair, as part of a group or as member of another merging cluster.  A galaxy can experience very different environmental conditions throughout its evolution. A major challenge for today's research is to identify the effects produced by the various environments on the observable properties of galaxies. It is especially hard to discriminate between minor and major influences. If all environments are potentially expected to have {\\it some} impact on their galaxies, we would like to discern secondary effects (``cosmetics'') from the primary causes that determine how galaxies are.  In this search we can rely on some objective and measurable properties, such as mass in stars and gas, their chemical content, their motion within the galaxy and the structures they design.  The star formation activity is one of the most outstanding characteristics.  Though a precise measurement of the current and past star formation activity is no easy task for any galaxy, star formation affects the accessible observables, such as colors and spectra, in such a manifest way that is often the most immediate evidence for evolution.  In this contribution I focus on the build-up of the stellar populations in cluster galaxies and on the evolution with redshift of their star formation activity.  The subject has been very rich in results and papers published during the past 25 years, and it is impossible to give an exhaustive summary in a short text. I will therefore only touch upon some of the issues, and highlight just a couple of the most recent results. Overviews on the evolution of the stellar populations in cluster galaxies can be found in Ellingson 2003, Poggianti 2003a and several reviews in the third volume of the Carnegie Observatories Astrophysics Series: Clusters of Galaxies: Probes of Cosmological Structure and Galaxy Evolution (eds. J.S. Mulchaey, A. Dressler, and A. Oemler (Pasadena: Carnegie Observatories, http://www.ociw.edu/ociw/symposia/series/symposium3/proceedings.html). ",
        "conclusions": ""
    },
    "0311/astro-ph0311273_arXiv.txt": {
        "abstract": "{ Three-dimensional numerical simulations with \\mbox{\\textsf{CO$^\\mathsf{5}$BOLD}}, a new radiation hydrodynamics code, result in a dynamic, thermally bifurcated model of the non-magnetic chromosphere of the quiet Sun. The \\mbox{3-D}~model includes the middle and low chromosphere, the photosphere, and the top of the convection zone, where acoustic waves are excited by convective motions.  While the waves propagate upwards, they steepen into shocks, dissipate, and deposit their mechanical energy as heat in the chromosphere. Our numerical simulations show for the first time a complex \\mbox{3-D} structure of the chromospheric layers, formed by the interaction of shock waves. Horizontal temperature cross-sections of the model chromosphere exhibit a network of hot filaments and enclosed cool regions. The horizontal pattern evolves on short time-scales of the order of typically \\mbox{20 - 25} seconds, and has spatial scales comparable to those of the underlying granulation. The resulting thermal bifurcation, i.e., the co-existence of cold and hot regions, provides temperatures high enough to produce the observed chromospheric UV emission and -- at the same time -- temperatures cold enough to allow the formation of molecules (e.g., carbon monoxide). Our \\mbox{3-D}~model corroborates the finding by \\citet{carlsson94} that the chromospheric temperature rise of semi-empirical models does not necessarily imply an increase in the average gas temperature but can be explained by the presence of substantial spatial and temporal temperature inhomogeneities. ",
        "introduction": "\\label{sec:intro}  Three-dimensional, time-dependent radiation hydrodynamics simulations of solar and stellar surface convection have now reached a level of sophistication which goes far beyond that of idealised numerical experiments, and allows a direct confrontation of such models with real stars \\citep[e.g.,][]{stein98,asplund00,cobold,ludwig02}. Extending this kind of simulation to include the low chromosphere, it is possible to study \\-- in a single model and based on first principles \\-- the generation of waves by the convective flow as well as the wave propagation and dissipation in the higher layers. Extended simulations of this type may then be utilised to explore the hitherto poorly understood \\mbox{3-D} thermal structure and dynamics of the non-magnetic chromospheric internetwork regions, and to obtain an independent theoretical estimate of the amount of chromospheric heating due to acoustic waves.  A strong motivation for three-dimensional time-dependent modelling arises from the need to reconcile apparently contradictory solar observations: Carbon monoxide absorption lines imply gas temperatures as low as $\\approx 3700$~K in the chromosphere of the quiet Sun \\citep[see ][ and references therein]{ noyes72a,ayres81b,solanki94,uitenbroek94,uitenbroek00a,ayres02}, whereas chromospheric UV emission features require much higher temperatures at the same heights \\citep[e.g.,][]{ayres76,carlsson97b}.  Semi-empirical models which have been constructed based on UV and microwave observations (\\citealp[e.g.,][ hereafter VAL]{val81}; \\citealp[][]{maltby86}; \\citealp[][ hereafter FAL]{fal93}) commonly feature a temperature minimum of $T_{\\mathrm{min}} \\approx 4200\\,-\\,4400$~K at a height of $z \\approx\\,500$~km above optical depth unity and an outwardly increasing temperature above. On the other hand, models based on CO observations \\citep[e.g.,][]{wiedemann94} show a monotonic decrease of temperature with height.  These conflicting observations and the inferred representative models have led to a controversy about the nature of the chromosphere of the non-magnetic quiet Sun which is going on for many years now \\citep[see, e.g.,][]{kalkofen01}: Is the chromosphere of the average quiet Sun a time-dependent phenomenon with a mostly cool background and large temperature fluctuations due to upward propagating shock waves? Or is it persistent and always hot with only small temperature fluctuations? In short: Is the non-magnetic solar chromosphere hot or cool?  A large number of observations show that the chromosphere of the quiet Sun is indeed a very dynamic phenomenon \\citep[e.g.,][]{carlsson97b,muglach01, krijger01, wunnenberg02}. Obviously, static one-dimensional models can only describe selected time-averaged properties. More realistic modelling should therefore be time-dependent.  Starting in the late 1960s, the pioneering work on \\mbox{1-D} time-dependent numerical models of chromospheric heating by acoustic and magneto-hydrodynamic waves is due to Ulmschneider and collaborators.  In a long series of papers \\citep[e.g.,][]{ulmschneider71,ulmschneider77,ulmschneider78,muchmore85, ulmschneider87,ulmschneider89,cuntz94}, they studied in detail the chromospheric energy balance between dissipation of prescribed short-period (mostly monochromatic) acoustic waves and radiative emission. In their models, the acoustic energy flux is supplied by a piston acting as a lower boundary condition. Assuming that the generation of acoustic waves by the ``turbulent'' flows in the upper convection zone can be described by the Lighthill-Stein theory \\citep[][]{lighthill52,stein67,stein68,musielak94,ulmschneider96, ulmschneider99}, they compute dynamic chromospheric models not only for the Sun but also for a sample of main-sequence stars and giants. Based on these models, they conclude that the observed ``basal flux'' from the chromospheres of late-type stars \\citep[][]{schrijver87,rutten91} is fully attributable to the dissipation of acoustic wave energy \\citep[][]{buchholz98}, and that the observed variation of chromospheric emission can be explained by the additional heating of magnetohydrodynamic shock waves \\citep{ulmschneider01}.  The detailed radiation hydrodynamics simulations by \\citet[][ hereafter CS]{carlsson94,carlsson95,carlsson97a} are another prominent example of sophisticated \\mbox{1-D} time-dependent modelling. These authors successfully explained the \\ion{Ca}{ii}~\\mbox{H$_{\\rm 2v}$} bright points as a result of propagating shock waves. In their model, the waves are excited by a piston which is driven by a velocity variation derived from observed oscillations at the photospheric level.  Instead of a temperature minimum and a monotonic temperature increase above, as characteristic of the VAL and FAL models, CS find a chromosphere with a mostly cool background and large temperature fluctuations due to upward propagating shocks. Even more remarkable is the fact that they are able to reproduce the rise of the radiation temperature without an increase of the mean gas temperature. Basic reasons are the nonlinear temperature dependence of the Planck function in the UV and the extreme temperature peaks associated with the shock waves.  This led CS to the conclusion that the chromosphere of the quiet Sun is not persistent but a spatially and temporally intermittent phenomenon which -- if averaged over space and time -- is mostly cool and not hot.  Although the one-dimensional models of the non-magnetic solar chromosphere mentioned above are highly elaborate, including a fully time-dependent \\element{H} ionisation and detailed NLTE radiative transfer, they suffer from the need for an external prescription of the wave excitation, and of course they cannot account for horizontal inhomogeneities and the associated effect of dynamic cooling on the atmospheric energy balance.  In this regard, the  three-dimensional self-consistent modelling by \\citet{skartlien00c} can be considered as a major improvement. The idea of Skartlien and co-workers was to extend the standard radiation hydrodynamics simulations of the solar granulation \\citep{stein98} into the chromosphere, where local thermodynamic equilibrium (LTE) is known to be a poor approximation. In order to adapt it to chromospheric conditions, \\citet[][]{skartlien00b} upgraded the radiative transfer part of the Nordlund-Stein code by implementing an iterative method to treat coherent isotropic scattering in \\mbox{3-D}.  The simulations enabled \\citeauthor{skartlien00c} to analyse the generation, propagation, and dissipation of acoustic waves in three dimensions.  The main emphasis of their study was on the excitation of transient wave emission resulting from the collapse of small granules, and the dynamic response of the chromospheric layers to such acoustic events.  In the present paper, we present similar time-dependent \\mbox{3-D}~models which extend from the upper convection zone to the middle chromosphere. The radiation hydrodynamics simulations are performed with \\mbox{\\textsf{CO$^\\mathsf{5}$BOLD}}, a new radiation hydrodynamics code developed by B.~Freytag and M.~Steffen \\citep{cobold}.  In this exploratory simulation, we treat the radiative transport in LTE with grey opacities (see Sect.~\\ref{sec:rad} and discussion in Sect.~\\ref{sec:discuss}). This simplification allows us to work at a significantly higher spatial resolution ($140 \\times 140 \\times 200$ cells) than \\citeauthor{skartlien00c} ($32 \\times 32 \\times 100$ grid). We find that the \\mbox{3-D} structure of the non-magnetic chromospheric layers is characterised by a complex pattern of interacting shocks, forming a network of hot filaments and enclosed cool ``bubbles''. This chromospheric pattern and its implications are chosen as major subject of this paper since the topology and the dynamics of the pattern are likely not to be too sensitive to the LTE simplification. We conclude that the low chromosphere exhibits a prominent thermal bifurcation: hot and cool regions exist side by side.  Surprisingly, this small-scale (non-magnetic) network was not mentioned by \\citeauthor{skartlien00c}; presumably, it was not noticed due to the poor (horizontal) spatial resolution of their numerical model.  In Sect.~\\ref{sec:cobold} we will give a short overview of the numerical details of \\mbox{\\textsf{CO$^\\mathsf{5}$BOLD}}.  The \\mbox{3-D}~model is described in Sect.~\\ref{sec:model}, followed by the results in Sect.~\\ref{sec:results}. Finally, a discussion and conclusions are presented in Sect.~\\ref{sec:discuss} and Sect.~\\ref{sec:conclusion}, respectively. ",
        "conclusions": "\\label{sec:conclusion}  Based on a detailed \\mbox{3-D} simulation of the solar granulation and the overlying atmosphere, we have studied the generation of waves by the time-dependent convective flow, and the wave propagation and dissipation in the higher layers. The most important improvements compared to previous numerical simulations are (i) self-consistent dynamics without a need for a driving piston like done by CS and (ii) a high spatial resolution which is obviously necessary for modelling the small-scale structure of the solar chromosphere. On the other hand, the LTE treatment of the thermodynamics and the radiative transfer is certainly unrealistic in the chromospheric layers. We have presented evidence that some of the basic features seen in our model are nevertheless representative of the (non-magnetic) internetwork regions of the solar chromosphere.  The main result of the present investigation is the discovery of a complex network of hot filaments pervading the otherwise cool chromospheric layers. Caused by interaction of standing and propagating hydrodynamic waves of large amplitude, the model chromosphere is a highly dynamical, spatially and temporally intermittent phenomenon. Its temperature structure is characterised by a thermal bifurcation: hot and cool regions co-exist side by side. Temperatures in the hot filaments are high enough to produce chromospheric emission lines, and the cool ``bubbles'' are cold enough to form molecular features. Thus, the chromosphere is {\\it hot and cold} at the same time.  This picture of the \\mbox{3-D} structure of the solar chromosphere has the potential to explain the apparently contradictory observational diagnostics which cannot be understood in the framework of \\mbox{one-dimensional} theoretical or semi-empirical models.  The presence of strong spatial and temporal temperature fluctuations has a remarkable consequence: the temperature minimum and the outward directed temperature rise inferred from semi-empirical models might be artifacts in the sense that they do not necessarily imply an increase of the average gas temperature with height.  Our model suggests that the radiative emission can be sustained by the hot propagating shock waves even though the main fraction of the chromospheric layers is cool and the mean gas temperature profile shows an almost monotonic decrease \\-- a conclusion already reached  by \\citet{carlsson94,carlsson95,carlsson97a} on the basis of \\mbox{one-dimensional} hydrodynamical simulations.  We conclude that improved \\mbox{3-D} radiation hydrodynamic simulations of the kind presented in this work are likely to lead the way towards a consistent physical model of the thermal structure and dynamics of the non-magnetic solar chromosphere which eventually can explain the various observational diagnostics."
    },
    "0311/astro-ph0311045_arXiv.txt": {
        "abstract": "The aim of the Gemini Deep Deep Survey is to push spectroscopic studies of complete galaxy samples (both red and blue objects) significantly beyond $z=1$; this is the redshift  where the current Hubble sequence of ellipticals and spirals is already extant. In the Universe at  $z=2$ the only currently spectroscopically confirmed galaxies are blue, star-forming and of fragmented morphology. Exploring this transition means filling the `redshift desert' $1<z<2$ where there is a dearth of spectroscopic measurements. To do this we need to secure redshifts of the oldest, reddest galaxies (candidate ellipticals) beyond $z>1$ which has led us to carry out the longest exposure redshift survey ever done: 100 ksec spectroscopic MOS exposures with GMOS on Gemini North. We have developed an implementation of the CCD ``nod \\& shuffle'' technique to ensure precise sky-subtraction in these ultra-deep exposures. At the halfway mark the GDDS now has $\\sim 36$ galaxies in the redshift desert $1.2<z<2$ extending up to $z=1.97$ and $I<24.5$ with secure redshifts based on weak rest-frame UV absorption features complete for both red, old objects and young, blue objects. The peak epoch of galaxy assembly is now being probed by direct spectroscopic investigation for the first time. On behalf of the GDDS team  I present our first results on the properties of galaxies in the `redshift desert'. ",
        "introduction": "The `redshift desert' is the well known gap between $z=1$ and $z=2$ where there is a dearth of spectroscopic redshifts. This is illustrated in Figure~1 which shows a color-redshift plot of some  of the deepest Keck spectroscopic samples now available. Generally galaxies fall in too two groups: pure magnitude selected samples which now reach as faint as $R_{AB}<24$ and redshifts $z<1.2$ and at higher redshifts the `Lyman Break' galaxies. These are color-selected via their Lyman break and are easily found for $z>3$ (Steidel et al. 2003) and are being extended down to $z=2$ (Erb et al. 2003).  \\begin{figure}[t] \\begin{center} \\plotfiddle{fig1.ps}{1.5in}{-90}{30}{30}{-120}{170} \\caption{Current deep spectroscopic surveys. The  $z<1.2$ points are from Cohen et al. (2001) and the $z>2$ points are Lyman break galaxies from Steidel et al. (1995). The tracks represent the colors expected of non-evolving galaxy SEDs: E$/$S0, Sbc, Scd (red -- blue)} \\label{fig1} \\end{center} \\end{figure}  These samples are selected optically and as a result the galaxies selected at high-redshift have turned out to be those with the strongest ultraviolet (UV) emission because for $z>1$ the observed optical probes the rest-frame UV. Of note are the Lyman break galaxies which all have spectra characteristic of starburst galaxies: blue UV continua, ISM absorption lines and evidence of outflows.  Other high-redshift populations, which emit less strongly in the UV, have been found. Examples include the sub-mm selected SCUBA sources (Chapman et al. 2003) and the class of objects known as `Extremely Red Objects' selected via their strong near-infrared (NIR) emission and red optical-NIR colors (e.g. Smail et al. 2002). While some spectra of such objects have been achieved there has been a lack of success in securing large samples of both red and blue galaxies above $z=1$.  The reason for this is primarily that the K-correction above $z=1$ has such a strong effect. Most galaxies apart from strong starbursts are significantly less luminous in the near-UV than at optical wavelengths, this is especially true of galaxies with SEDs characteristic of old stellar populations. Furthermore there is a lack of strong emission lines, even in starforming galaxies, between [OII] (3727\\AA) and Lyman-$\\alpha$ (1216\\AA). Thus spectroscopy must reach sufficient signal:noise to robustly detect the continuum and hence achieve absorption line redshifts. ",
        "conclusions": "The Gemini Deep Deep Survey has succeeded in obtaining spectroscopic redshifts of galaxies in the former redshift desert.  We have closed the $1<z<2$ gap, but more importantly we are highly complete for both red and blue galaxies. A catalog paper (Abraham et al. 2004) is in preparation. Our results to date are:  \\begin{enumerate} \\item An abundance of massive galaxies at $z\\ga 1.5$. (Glazebrook et al. 2004, in prep.) These may be problematic for $\\Lambda$CDM models of heircharchical growth of massive halos. \\item Many of the massive red galaxies at $z>1$ have spectra of pure old stellar populations (McCarthy et al. 2003, in prep.) \\item Massive starburst galaxies at $1.3<z<2$ have high metal columns and are close to solar in abundances (Savaglio et al. 2004 in press). \\end{enumerate}  At the time of this meeting we had just received the Gemini data for the final two of our four fields. The complete sample will be more than double the size of the one presented here. To come we have been awarded Cycle 12 time to image our fields with the Advanced Camera for Surveys on the Hubble Space Telescope. Thus we hope to connect our results to physical galaxy morphology. It is clear we are now living in an age where the formation of galaxies and the Hubble Sequence can be observed directly."
    },
    "0311/astro-ph0311363_arXiv.txt": {
        "abstract": "It has been known that at least some of the observed gamma-ray bursts (GRBs) are produced at cosmological distances and the GRB production rate may follow the star formation rate. We model the BATSE detected intensity distribution of long GRBs in order to determine their space density distribution and opening angle distribution. Our main results are: the lower and upper distance limits to the GRB production are $z\\approx 0.24$ and $z>10$, respectively; the GRB opening angle follows exponential distribution and the mean opening angle is about 0.03 radian; the peak luminosity appears a better standard candle than the total energy of a GRB. ",
        "introduction": "Recent research results indicate that radiation from gamma-ray bursts (GRBs) may be constrained mostly in narrow beams \\citep{wa98,fr99,wo01}, which suggests that GRBs may be standard candles at cosmological distances \\citep{fra01}. Previous studies of the GRB space density distribution suggested that GRB production rate may follow the observed star formation rate for small values of redshift \\citep{wij98}, but may increase monotonically for very high values of redshift  \\citep{sch,no02}; however neither the lower limit nor the upper limit of redshift for GRB production could be determined reliably previously because these studies are either limited by statistics or have to rely on some empirical relationships between some statistical properties of selected samples of GRBs. Because both the space density distribution and beaming effect play important roles for the observed GRB intensity distribution, here we make use of the BATSE/CGRO GRB intensity distribution, which is the largest sample of GRBs collected so far, to determine directly the GRB space density and opening angle distribution. ",
        "conclusions": "First we compare our study with previous studies of GRB cosmological models with BATSE intensity distribution \\citep{fen93,fen95}\\\\ \\citep{wic93,wij98}. All these studies concluded that the dimmest GRBs are at redshift around unity, probaly following the measured star formation rate which peaks also at redshift around unity \\citep{st99}. The main differences between our approach and previous studies are the following:  \\begin{itemize} \\item We take the GRB opening angle distribution as part of the cosmological model. This is very important, since the observed GRB flux depends strongly on the opening angle, as well as the redshift.  \\item We separate the long and short GRBs as two different classes in our study. Our conclusions are only applicable to long GRBs. For short GRBs, no conclusion can be drawn due to the limited statistical quality of data.  \\item One simplification of our study is that we did not make the energy spectral correction to the fluence or peak flux as a function of $z$ (K-correction); we just take the fluence or peak flux values as listed in the catalog. Technically this correction is not straightforward, because we can no longer uniquely convert the fluence or peak flux directly to redshift, because we believe that GRB radiation is beamed, and the opening angle follows a certain distribution. Therefore the observed fluence or peak flux is due to both the redshift and the opening angle of the GRB. The good fit of the assumed simple cosmological model with only several free-parameters indicates that this correction may not impact our conclusions significantly. In fact, since the energy spectra of GRBs in the BATSE band follows approximately a power-law with an index of $-$2, the K-correction is not important.  \\end{itemize}  Frail {\\it et al.} (2001) suggested that the total energy of GRBs may be a constant, based on the opening angles of several GRBs determined from their afterglow light curve break. Our result is not in major conflict with their suggestion, but favors the peak luminosity as the standard candle. Perhaps the direction of the GRB beam is not completely stable, as indicated from observations of other relativistic jets and outflows from stellar mass and supermassive black holes; only when the beam is pointed directly at us can we observe the maximum flux from the GRB. However it has been argued that for some GRBs the lower limit of the total energy released is $10^{52}$ {\\it ergs} as determined from the redshifted X-ray emission lines \\citep{Ghi02}. However the reality of the detection of the emission lines is still under debate, and an alternative interpretation of the detected emission lines also exists \\citep{cer}.  \\clearpage  \\begin{figure}\\label{fig:fig2} \\centerline{ \\epsscale{0.7} \\plotone{f2.eps} } \\caption {The data for the obsered GRB opening angles are taken from Bloom  {\\it et al.} 2003; arrows indicate lower and upper limits of GRB opening angles. The power-law distribution fits the data for large opening angles, but fails to predict the absence of GRBs detected with opening angles smaller than 0.04 radian. The Gaussian distribution agrees with the peak reasonably well, but cannot predict enough GRBs with opening angles greater than 0.2 radian. The exponential distribution is consistent with the data.} \\end{figure}  \\clearpage  The opening angles of some GRBs with afterglow detections have been determined previously from the GRB afterglow data \\citep{fra01,bloom2003}; we call this Frail-distribution. In $\\emph{fig.2}$ we compare the Frail-distribution with the opening angle distributions determined with the observed GRB intensity distribution data, just as a sanity check. The power-law distribution agrees with the Frail-distribution for large opening angles, but fails to predict the absence of GRBs detected with opening angles smaller than 0.04 radian. No GRBs with opening angles narrower than 0.04 radian was directly observed, but the recent report on the polarization of GRB021206 suggests that it is possible for some GRBs to have opening angles narrower than 0.01 radian \\citep{cob03,wax03}. The Gaussian distribution agrees with the peak of the Frail-distribution distribution reasonably well; however this model cannot predict enough GRBs with opening angles greater than 0.2 radian. The exponential distribution is consistent with the Frail-distribution distribution for both standard candle models, but favors slightly the peak luminosity standard candle model. It was suggested that for the model in which the GRBs follow the star formation rate strictly, the most possible opening angle should be around 0.1 radian \\cite{bloom2003}. However, if the majority of faint GRBs originate at higher redshift than the peak of star formation rate as our model predicts, then a narrower average opening angle is required.  We therefore conclude that GRB opening angles follow an exponential distribution with a mean opening angle of $1/\\lambda=0.03$ radian, and the GRB space density follows a power-law with an index of $-$3.5 between $z_0=0.24\\pm0.07$ and $z_{break}>10$. The general trend of the GRB space density distribution determined here as a function of $z$ is consistent with previous studies \\citep{sch,no02}; however our result is statistically much more robust, and in particular $z_0$ is determined for the first time and the lower limit to $z_{break}$ is constrained more strictly than before. Our result therefore demonstrates that GRBs are no longer produced frequently in the nearby and present universe at $z<0.24\\pm0.07$, and GRBs are predominantly produced in the early universe at $z>10$. Indeed most GRBs with direct redshift measurements have redshifts $z>0.3$ \\citep{djo01}, with the exceptions of GRB 980425 ($z$=0.0085) and GRB 030329 ($z$=0.1687) which may be associated with two peculiar but similar supernovae SN1998bw \\citep{sn1998}, and SN2003dh \\citep{sn03}. It remains to be seen if these GRBs belong to a subset of GRBs which are different from the majority of GRBs which seem to be produced at much larger redshift.   It has been suggested that GRBs follow star formation rate, because it is commonly believed that GRBs are produced as the final gravitational collapse of massive stars \\citep{djo01,pac98}. The GRB space density determined here follows the measured star formation rate between $z\\approx 0.24$ to $z\\approx 1$. The determination of $z_0=0.24\\pm0.07$ suggests that the massive stars capable of producing the majority of long GRBs are no longer formed frequently today. Future GRB experiments capable of measuring the redshift of many more GRBs and thus providing direct distance measurements of GRBs, such as the {\\it Swift} mission to be launched in late 2003 \\citep{swift}, will test this prediction firmly.  The monotonic increase of the GRB space density up to at least $z=10$ suggests that the formation rate of massive stars capable of producing GRBs should peak at significantly higher $z$ than the measured star formation rate. This implies that the ``dark ages\" of the universe ended much earlier than previously believed at $z\\sim 4-5$ based on the measured space density distribution of high-redshift QSOs \\citep{hook}. There are also some previous independent works showing that an essential amount of long GRBs should be at $z > 4$ \\cite{cho02,bag03} or even up to $z \\cong 20$\\cite{mes96}. A recently detected GRB031203 by INTEGRAL was reported to be located at $z > 9$ (formally $z \\geq 11$) \\cite{grb03}, consistent with our model prediction. Therefore GRBs may be a probe of the epoch of reionization of the early universe \\citep{lamb03}. Currently the GRB intensity distribution data cannot constrain the exact upper limit of $z_{break}$, because not enough extremely faint GRBs have been observed. Future GRB experiments, such as the proposed ``Next Generation GRB Mission\" \\citep{exist} would detect many more GRBs originated in the extremely early universe, which may determine the exact ending time of the ``dark ages\" of the universe in order to test the prediction of some numerical simulations \\citep{first-star} which suggests that the first stars may have been formed at $z\\sim 20$. GRBs may prove to be the most sensitive and ultimate probes into the early epoch of the universe when the first structures were formed.  \\noindent {\\bf Acknowledgement: } We thank Drs. Zigao Dai, Tan Lu, Yongfeng Huang and Daming Wei for valuable comments to the manuscript. This study is supported in part by the Special Funds for Major State Basic Research Projects (10233010) and by the National Natural Science Foundation of China. SNZ also acknowledges supports by NASA's Marshall Space Flight Center and through NASA's Long Term Space Astrophysics Program."
    },
    "0311/astro-ph0311155_arXiv.txt": {
        "abstract": "An estimate of the convergence radius of a simulated CDM halo is obtained under the assumption that the peak phase-space density in the system is set by discreteness effects that operate prior to relaxation. The predicted convergence radii are approximately a  factor 2 larger than those estimated for numerical convergence studies.  A toy model is used to study the formation of sheets of the cosmic web, from which DM haloes form later. This model demonstrates the interplay between phase mixing and violent relaxation that must also be characteristic of spherical collapse. In the limit that sheets contain arbitrarily many particles, it seems that power-law profiles are established in both distance and energy. When only a finite number of particles is employed, relaxation is prematurely terminated and the power laws are broken. In a given simulation, the sheets with the highest peak phase-space densities are those that form from the longest waves. Hence simulations with little small-scale power are expected to form the cuspiest haloes. ",
        "introduction": "Cold Dark Matter (CDM) is thought to dominate structure formation through gravitational clustering. As its name implies, this material starts out very cold, so that to an excellent approximation particles are confined to a three-dimensional surface $[\\b r,\\b p(\\b r)]$ in six-dimensional phase space. The distribution function $f$ is infinite on this surface and zero elsewhere.  The fine-grained distribution function is constant because it obeys the collisionless Boltzmann equation. Hence, very high phase-space densities are in principle possible right up to the present epoch.  Studies of structure formation rely heavily on $N$-body simulations to follow the clustering of dark matter (DM). Galaxies are believed to form from baryons that have fallen in to the centres of DM haloes, so the implications of various theories of DM for the observable properties of galaxies depend sensitively on the small-scale structure of dark haloes. Hence, it is important to understand how accurately $N$-body simulations represent the inner parts of dark haloes, and this problem has attracted considerable attention (Jing et al.\\ 1995; Klypin et al.\\ 2001; Taylor \\& Navarro 2001; Navarro 2003; Power et al.\\ 2003; Fukushige, Kawai \\& Makino 2003; Hayashi et al.\\ 2003).  On sufficiently small scales it is clear that the simulations will become untrustworthy because they employ only a finite number of particles (Kuhlman et al., 1996; Melott et al., 1997; Splinter et al., 1998; Binney \\& Knebe, 2002). Generally this scale is determined by resimulating a given halo with improved resolution, in terms of the number of particles employed, the number of waves used to define the initial conditions, the smallness of the gravitational softening parameter $\\epsilon$, and the smallness of the timesteps. Based on such simulation series, Power et al.\\ (2003) define a radius $r_{\\rm conv}$ outside which a given simulated halo should be reliable, by calculating the smallest radius at which three criteria are all met: (i) the local orbital time is much longer than the numerical timestep; (ii) the centripetal acceleration of a circular orbit is smaller than $0.5V_{200}^2/\\epsilon$, where $V_{200}$ is the halo's characteristic velocity; and (iii) the local two-body relaxation time is longer than the Hubble time.  While it is clear that satisfaction of all three criteria is a necessary conditions for a N-body simulation to be credible, it is not clear that it is a sufficient condition: errors in the integration of the equations of motion at any cosmic epoch can invalidate the final halo. We argue here that a particularly important epoch to get right is the one that ends when overdensities first collapse to form sheets that are virialized in only one direction. We present evidence that discreteness effects at this primary stage, rather than, for example, two-body relaxation in virialized haloes, are what limit the central densities of simulated haloes.  We use a toy model and analytic calculations to clarify what happens when virialized CDM structures form, and to understand the extent to which N-body simulations of structure formation are compromised by artifacts due to discreteness. The model explains why cuspy profiles form even when power at high spatial frequencies is removed from the initial fluctuation spectrum. It also suggests that in the absence of discreteness effects, virialized haloes would have power-law cusps. Finally analytic arguments are used to estimate the radius in a simulated halo within which the central profile is flattened by discreteness effects that operate prior to virialization, and we show that these estimates are in good agreement with the latest numerical convergence studies.  \\begin{figure*} \\centerline{\\psfig{file=out0.ps,width=.46\\hsize}\\quad \\psfig{file=out4.ps,width=.46\\hsize}} \\centerline{\\psfig{file=out16.ps,width=.46\\hsize}\\quad \\psfig{file=out30.ps,width=.46\\hsize}} \\centerline{\\psfig{file=out38.ps,width=.46\\hsize}\\quad \\psfig{file=out96.ps,width=.46\\hsize}} \\caption{From top left to bottom right, six stages in the violent relaxation of 301 particles. The number in the top left corner of each panel is the time since the initial conditions of equations (\\ref{Xinit}) were imposed. The units of time are the turnaround time $\\tau$ of the underlying homogeneous distribution of particles.\\label{Nbdy1}} \\end{figure*} ",
        "conclusions": "Semi-analytic work on the formation of cosmic structure usually focuses on the collapse of spherical over-densities (Gunn \\& Gott 1972; Fillmore \\& Goldreich 1984; Bertschinger 1985; Subramanian, Cen \\& Ostriker, 2000). Actually one direction collapses before the other two, so the crucial first phase of the virialization process is one-dimensional (Zel'dovich, 1970; Shandarin et al., 1995).  Fillmore \\& Goldreich (1984) obtained a self-similar solution of the collapse problem in this planar symmetry. Unfortunately, their results are not directly comparable to ours because (i) they assumed that the transverse expansion continued unperturbed, and (ii) their initial overdensity was constrained to diverge at the centre of the sheet.  Consequently, for all $t>0$ their solution has a collapsed core, and the phase-space spiral has a time-invariant form: its scale changes with time, but it does not wind up in the manner of Fig.~\\ref{Nbdy1}.  Our toy model of the planar collapse of CDM reveals a strong interplay between phase mixing and violent relaxation, in which an initially small phase lag of the oscillations between the centre and periphery of an overdensity causes energy to flow outwards, and this outward energy flow hastens the growth of the phase difference, which in turn causes the exchange of energy to repeat on ever smaller scales. The numerical evidence suggests that in the limit of arbitrarily many particles, power-law profiles in space and energy would be established as a natural consequence of the large dynamic range between the smallest scale of energy interchange and the overall size of the system. Power laws are also consistent with the original infinite phase-space density persisting at the centre of the relaxed system, while admixture of `air' ensures that the phase-space density is finite and outwards-decreasing elsewhere.  When the system is represented by only a finite number of particles, random motions develop prior to virialization. Consequently, as collapse begins, the phase-space density is finite and the interplay between phase mixing and violent relaxation is prematurely terminated. The central phase-space density of the relaxed system is comparable to the phase-space density just before collapse, while further out it differs little from the value that would hold in the limit of an arbitrary number of particles. Consequently, the real-space density profile must turn over near the centre from the slope it would have in that limit.  An interesting result from the one-dimensional model is that, in a given simulation, the peak phase-space density in simulated sheets of the cosmic web scales with particle number as $f\\sim n^{3/2}$. Hence, the sheets with the highest surface density achieve the largest phase-space density. The surface density of the sheets that form is increased if we suppress small-scale power in the initial fluctuation spectrum, so there is a suggestion that simulations that have little small-scale power will produce the cuspiest haloes. This finding is consistent with the counter-intuitive results of Knebe et al.\\ (2002).  The extension to the formation of DM haloes of these results for the formation of the sheets of the cosmic web is not straightforward. Some insight can be gained by imagining that the original process repeats as each of the other directions experiences collapse. This idealization is problematic because the collapses of these directions are rarely well separated in time. However, it is probably legitimate to argue that in the limit of infinite particle number, the final density profile should be a power law\\footnote{Figure 4 of Hayashi et al.\\ (2003) argues against this conclusion, however.} because the collapsing sheet brings no relevant scale into the relaxation process, and the latter has no intrinsic scale. Also it seems likely that violent relaxation will be less efficient in later collapses than it was in the first, because the phase spirals associated with these directions will wind up very quickly. This reasoning allows us to set an upper limit ${3\\over2}$ on the power-law index of the final density cusp, which is clearly greater than ${1\\over2}$.  If we assume that in the limit of infinite particle number the slope would be $\\alpha=1$, as in the NFW profile, then we can estimate the radius $r_{\\rm min}$ at which the maximum phase-space density sampled equals the phase-space density that developed prior to virialization. This radius forms a reasonable estimate of the radius at which the simulated density profile turns away from a line of unit slope. We find that $r_{\\rm min}$ is larger by a factor $\\sim2$ than the convergence radius $r_{\\rm conv}$ of Hayashi et al.\\ (2003). This finding strongly suggests that processes that take place prior to virialization determine the smallest radius at which a simulated halo can be trusted.  The importance of knowing which structures in cosmological simulations are reliable, is clearly very great. Brute-force convergence studies such as those of Power et al.\\ (2003) do not constitute a straightforward tool for validating N-body simulations because the limit in which reality is approached involves sending three quantities to zero: the mass of an individual particle, the softening length, and the wavelength of the highest-frequency mode included in the initial conditions. Before an entirely convincing convergence study can be made, one needs to know along which path in the three-dimensional space of these parameters one should approach the origin (e.g., Splinter et al.\\ 1998). Hitherto there has been a tendendency for N-body simulators to focus on the impact that discreteness has on the dynamics of relaxed systems. This study suggests that dicreteness may have its dominant impact prior to virialization."
    },
    "0311/astro-ph0311225_arXiv.txt": {
        "abstract": "We study the accretion on to the neutron star in Be/X-ray binaries, using a 3D SPH code and the data imported from a simulation by \\citet{oka2} for a coplanar system with a short period ($P_{\\rm orb}=24.3\\ \\rm{d}$) and a moderate eccentricity $(e=0.34)$, which targeted the Be/X-ray binary 4U\\,0115+63. For simplicity, we adopt the polytropic equation of state. We find that a time-dependent accretion disc is formed around the neutron star regardless of the simulation parameters. In the long term, the disc evolves via a two-stage process, which consists of the initial developing stage and the later developed stage. The developed disc is nearly Keplerian. In the short term, the disc structure modulates with the orbital phase. The disc shrinks at the periastron passage of the Be star and restores its radius afterwards. The accretion rate on to the neutron star is also phase dependent, but its peak is broader and much lower than that of the mass-transfer rate from the Be disc, unless the polytropic exponent is as large as 5/3. Our simulations show that the truncated Be disk model for Be/X-ray binaries is consistent with the observed X-ray behaviour of 4U\\,0115+63. ",
        "introduction": "\\label{sec:intro}  The Be/X-ray binaries represent the largest subclass of high-mass X-ray binaries. About two-thirds of the identified systems fall into this category. These systems consist of, generally, a neutron star and a Be star with a cool ($\\sim 10^{4}K$) equatorial disc, which is geometrically thin and nearly Keplerian. The orbit is wide $(16\\,{\\rm d} \\la P_{\\rm{orb}} \\la 243\\,\\rmn{d})$ and usually eccentric ($e\\ga 0.3$) (\\rm{e.g.,} \\citealt{zi}).  Most of the Be/X-ray binaries show only transient activity in the X-ray emission and are termed Be/X-ray transients. Be/X-ray transients show periodical (Type I) outbursts, which are separated by the orbital period and have the lumiocity of $L_{\\rmn{X}}=10^{36-37}\\rmn{erg\\,s}^{-1}$, and giant (Type II) outbursts of $L_{\\rmn{X}} \\ga 10^{38} \\rmn{erg\\,s}^{-1}$ with no orbital modulation. The transient nature in the X-ray activity of Be/X-ray binaries is considered to result from the interaction between the accreted material from the circumsteller disc of the Be star and the rotating magnetized neutron star \\citep*{st}. If the accreted material is dense enough to make the magnetospheric radius smaller than the corotation radius, the accretion on to the neutron star causes a bright X-ray emission (the direct accretion regime). Otherwise, the magnetospheric radius is larger than the corotation radius, and the accretion is prevented by the centrifugal inhibition. The system is then in quiescence (the propeller regime).  During the outbursts, the neutron stars in most Be/X-ray binaries show spin-up (e.g., 4U\\,0115+63, GS\\,1843-02, EXO\\,2030+375 and XTE\\,J1946+26) (\\citealt{bi}; \\citealt{fi2}; \\citealt{wi1}; \\citealt{wi2}), and some even show quasi-periodic oscillations (QPOs) (e.g., A\\,0535+26 and EXO\\,2030+375) (\\citealt{fi1}; \\citealt{wi1}). These observational features strongly suggest the presence of an accretion disc during the outbursts.  The accretion disc is also indicated to exist between the outbursts in several systems (e.g., A\\,0535+26) (\\citealt{fi1}; see also \\citealt{zi}). On the other hand, \\citet{cl} searched the optical/infrared contribution from the accretion disc in A\\,0535+26 during the X-ray quiescent phase, and found no signature of the accretion disc in propeller regime. Obviously, more work is needed to determine whether the accretion disc also exists in the quiescent state.  Recently, \\citet{oka2} studied the interaction between the Be-star disc and the neutron star, using a three-dimensional (3D) smoothed particle hydrodynamics (SPH) code (\\citealt{be1}; \\citealt*{ba}). They also qualitatively discussed the relation between the X-ray luminosity and the accretion rate using the simulated mass-capture rate by the neutron star. In their simulations, however, the neutron star was modelled by a sink particle \\citep{ba} with the size of the Roche lobe. Hence, no information was available how the captured material accretes on to the neutron star. Moreover, no direct comparison with the observed X-ray data was possible.  In this paper, we study the accretion on to the neutron star in Be/X-ray binaries, using a 3D SPH code. The simulation should ideally take account of all the processes at work in a Be/X-ray binary, such as the disc evolution around the Be star, the mass transfer from the Be-star disc to the neutron star, and the accretion on to the neutron star. Such a simulation, however, would require very intense computations.  Therefore, we confine ourselves to simulate only the accretion flow around the neutron star, with an outer boundary condition constructed by the data of the particles captured by the neutron star from a high-resolution simulation by \\citet{oka2}. This situation is illustrated in Fig.~1, where the right panel taken from their simulation shows the mass transfer from the Be-star disc to the neutron star at periastron, while the left panel shows the accretion disc formed around the neutron star in one of our simulations.   \\begin{figure} \\includegraphics*[width=8.5cm]{khfig1.ps} \\caption{Illustrative diagram of a Be/X-ray binary. At the periastron passage, the gas in an outer part of the Be-star disc is transferred to the neutron star and forms an  accretion disc. The right panel is taken from a simulation by \\citet{oka2}, while the left panel is from one of our simulations.} \\label{figure1} \\end{figure} ",
        "conclusions": "\\label{sec:summary}  For the first time, we have performed numerical simulations of accretion on to the neutron star in Be/X-ray binaries, taking into account the phase-dependent mass transfer from the Be disc. We have adopted the mass-transfer rate from a high-resolution simulation by \\citet{oka2} for a coplanar system with a short period ($P_{\\rmn{orb}}=24.3\\,\\rmn{d}$) and moderate eccentricity $(e=0.34)$, which targeted 4U\\,0115+63, one of the best studied Be/X-ray binaries. We have found that a nearly Keplerian, non-steady accretion disc is formed around the neutron star regardless of the parameters adopted. The formation of the accretion disc involves a two-stage process, which consists of the initial developing stage and the later developed stage. The disc structure shows a strong dependence on the orbital phase, because both the tidal potential and the mass-transfer rate change periodically. The disc is also significantly eccentric, because the gas particles from the Be disc originally have elliptical orbits.  In 4U\\,0115+63, Type~I X-ray outbursts have occurred only as a short series after occasional Type~II X-ray outbursts \\citep{ne1}. Except for such occasions, the system is in the quiescent state ($10^{33} - 10^{35}\\ \\rmn{erg\\,s}^{-1}$) \\citep{ca}. The peak X-ray luminosity estimated from our simulations corresponds to the transition regime between the direct accretion regime and the propeller regime  and is consistent with the observed X-ray behaviour. Note, however, that the discussion based on the X-ray luminosity estimated from the simulations only gives the necessary condition for the quiescent state.  In order to see whether our model satisfies a more stringent condition, we also have estimated the magnetospheric radius by using a general formula with the mass-accretion rate from our simulations and compared with the corotation radius for 4U\\,0115+63. The result was that our models with $\\Gamma \\le 1.2$ are consistent with the observed X-ray behaviour.  Other models with $\\Gamma = 5/3$ and $k \\la 0.6$ were ruled out, because in these models the system would show an X-ray outburst after every periastron passage, which has never been observed. We have to admit, however, that the radius of our inner simulation boundary is about 60-100 times as large as the corotation radius of the neutron star in 4U\\,0115+63. To make the constraint on the system parameters more reliable, we need numerical simulations with a much smaller inner boundary.  It is important to note that our model is not only consistent with the quiescent state of 4U\\,0115+63. It can also give a natural explanation for a short series of Type~I X-ray outbursts after occasional Type~II X-ray outbursts. The mass-transfer rate from the Be disc in the X-ray outbursts is likely much higher than in the quiescent state. Suppose that the mass-transfer rate from the Be disc in 4U\\,0115+63 is temporarily increased by an order of magnitude. Then, the estimated X-ray luminosity in our model increases by an order of magnitude to a level of several $\\times 10^{36}\\,\\rmn{erg\\,s}^{-1}$, a typical X-ray luminosity of Type~I X-ray outbursts. At the same time, the magnetosphric radius $R_{\\rmn{M}}$ decreases by a factor of about two, for which $R_{\\rmn{M}}$ is almost always smaller than the corotation radius, $R_{\\Omega}$, if $k \\sim 0.5$ (see Fig.~\\ref{fig:mdot_acc}(b)). This means that the system is in the direct accretion regime. Thus, our model explains the Type~I X-ray outbursts in 4U\\,0115+63 as a temporal phenomena with an enhanced mass-transfer from the Be disc.  Be/X-ray binaries are distributed over a wide range of orbital periods $(16\\,{\\rmn{d}} \\la P_{\\rm{orb}} \\la 243\\,{\\rmn{d}})$ and eccentricities $(e \\la 0.9)$. It is interesting to consider here the effects of orbital parameters. In this paper, we have seen that the main ingredients which determines the X-ray behaviour of a system are the viscous time-scale relative to the orbital period and the mass-transfer rate from the Be disc. Since the former is insensitive to the orbital period itself (see equation~(\\ref{eq:tratio})), the latter is most important if the viscosity parameter and the equation of state are similar in accretion discs of Be/X-ray binaries.  In systems with $\\alpha_{\\rmn{SS}} \\la 0.1$, the viscous time-scale in an outer part of the accretion disc is much longer than the orbital period unless the cooling is ineffective and the gas obeys an adiabatic equation of state. In these systems, persistent accretion discs are expected to form. Then, the mass-transfer rate from the Be disc determines whether the system shows Type~I X-ray outbursts. For example, for a system with a moderate orbital eccentricity of $e = 0.34$ and a typical density of Be disc, the peak mass-transfer rate is $\\sim 2 \\times 10^{-10} M_{\\sun}\\,{\\rmn{yr}}^{-1}$, for which the system stays in quiescence. It is obvious, however, that the situation should be different in systems with much higher rate of mass-transfer. A high mass-transfer rate is expected to result from a high orbital eccentricity and/or an intrinsically dense Be disc. Therefore, it is not strange that many systems which have regularly shown Type~I X-ray outbursts have orbital eccentricities higher than that of 4U\\,0115+63 studied in this paper. In a subsequent paper, we will perform numerical simulations of accretion in systems with high orbital eccentricities.  Above we have assumed that the viscous time-scale is much longer than the orbital period in all Be/X-ray binaries. However, it is possible that an accretion disc has a relatively short viscous time-scale by some reason. One possibility is that the disc has a low cooling efficiency and becomes hot, particularly in an inner region. Our model~4 with $\\Gamma = 5/3$ gives an extreme example of such discs (see also Fig.~\\ref{fig:chardisc}). As shown in Fig.~\\ref{fig:mdot_234}, in a system with a short viscous time-scale, the mass-accretion rate has a high and sharp peak shortly after every periastron passage. The resemblance between the accretion rate profile in model~4 and the X-ray light curves of regular Type~I X-ray outbursts from systems such as EXO\\,2030+375 \\citep{wi1} implies that these are systems with relatively short viscous time-scales.  Be/X-ray binaries are an ideal group of objects for studying the physics of accretion under the circumstances which have been studied little. Unlike close binaries, many Be/X-ray binaries have eccentric orbits. They can also be highly inclined systems. Such systems provide us a valuable oppotunity to study the effects of the periodically-changing tidal potential and mass-transfer rate and the inclination angle on the structure and evolution of the accretion flow. Another important characteristic of Be/X-ray binaries is that they are systems with double circumstellar discs (the Be disc and the accretion disc). The interaction in such systems is threefold: the interactions between the neutron star and the Be disc, the Be star and the accretion disc and the Be disc and the accretion disc (mainly via the mass transfer). Much more work is desirable both theoretically and observationally in order to understand this interesting and important group of objects."
    },
    "0311/astro-ph0311013_arXiv.txt": {
        "abstract": "We present preliminary results from our spectroscopic search for planets within 1 AU of metal-poor field dwarfs using NASA time with HIRES on Keck I. The core accretion model of gas giant planet formation is sensitive to the metallicity of the raw material, while the disk instability model is not. By observing metal-poor stars in the field we eliminate the role of dynamical interactions in dense stellar environments, such as a globular cluster. The results of our survey should allow us to distinguish the relative roles of the two competing giant planet formation scenarios. ",
        "introduction": "The two proposed models for the formation of gas giant planets make different, and testable, predictions. For example, core accretion requires several Myr to grow a solid core massive enough to accrete a gaseous envelope (Lissauer~1993). In contrast, disk instability leads to Jupiter-mass clumps (e.g., Boss~1997), which can survive and give rise to actual protoplanets (Mayer et al.~2002), within $\\sim 10^3$ yr. The time-scale for core accretion to proceed depends strongly on the initial surface density of solids (Pollack et al. 1996), so that formation by this mechanism may be enhanced in metal-rich stars. Instead, disk instability is remarkably insensitive to the primordial metallicity of the protoplanetary disk (Boss 2002). Several observational tests can be performed to probe giant planet formation models (Boss 2003). In order to unambiguously establish whether high primordial metallicity is a requirement, and thus determine which is the dominant formation mechanism for gas giant planets, we are conducting a spectroscopic search for planets orbiting a sample of metal-poor stars in the field, thus eliminating possible sources of interference with planet formation, or with migration to close-in orbits, or with planet survival, such as dynamical interactions in dense stellar environments. ",
        "conclusions": ""
    },
    "0311/astro-ph0311539_arXiv.txt": {
        "abstract": "Results are presented from a wide area, high resolution HI synthesis survey of the outer regions of the nearby M81 group, where internal (galactic) and external (group-related) evolution processes can be studied simultaneously in great detail. The survey encompasses the star forming dwarf galaxies M81dwA, UGC4483, and HoII, where evidence of ram pressure stripping was recently discovered. The data do not reveal any intergalactic HI, but the outer parts of HoII are reminiscent of tidal tails. We argue however that those structures are equally consistent with the latest ram pressure models including cooling. The case for a hot intergalactic medium in this poor, spiral-only group is thus still open. The survey also puts tight constraints on possible counterparts to the local high velocity cloud population in an external group, reaching a $3\\sigma$ column density of $10^{19}$~atom~cm$^{-2}$ and a $6\\sigma$ limiting mass of $1.5\\times10^5$~$M_{\\mbox{\\scriptsize \\sun}}$. ",
        "introduction": "\\label{sec:intro} Most galaxies in the local universe are located in groups, where dwarfs dominate by number. Understanding the influence of this environment on (dwarf) members is therefore essential to our comprehension of galaxy formation and evolution. Nearby groups are particularly attractive for this purpose since both internal (galactic) and external (group-related) evolution processes can be studied in details. Bureau \\& Carignan (2002) presented neutral hydrogen Very Large Array (VLA) D-array observations of the dwarf irregular galaxy HoII, a prototype galaxy for studies of self-propagating star formation (Puche et al.\\ 1992). HI was detected to radii over $16\\arcmin$ or $4R_{25}$, with a total mass $M_{\\mbox{\\tiny HI}}=6.44\\times10^8$~$M_{\\scriptsize \\sun}$. Most importantly, the integrated HI map revealed a characteristic comet-like appearance, with a large but faint component extending to the northwest and the HI appearing compressed on the opposite side, suggesting that HoII is affected by ram pressure from an intragroup medium (IGM).  HoII lies roughly $0.5$~Mpc northeast of the core of the nearby M81 Group ($D=3.2$~Mpc), along with the dwarf irregular galaxy M81dwA (Kar52) and the blue compact dwarf galaxy UGC4483. No obvious signs of interaction were detected, but the three galaxies are most likely part of the NGC2403 subgroup, infalling toward M81 (Karachentsev et al.\\ 2000). With large uncertainties, ram pressure stripping of the outer parts of the disk would require an IGM density $n_{\\mbox{\\tiny IGM}}\\ga4.0\\times10^{-6}$~atoms~cm$^{-3}$ near HoII, or about $1$\\% of the virial mass of the group uniformly distributed over a volume just enclosing it, consistent with the known X-ray properties of small groups. The presence of an IGM in such a poor group, lacking any early-type galaxy, would have important consequences for the evolution of group members, and it could explain why many HI holes in HoII are located in low surface density regions of the disk, where no star formation is expected or observed (Rhode et al.\\ 1999; Bureau \\& Carignan 2002). ",
        "conclusions": "\\label{sec:conclusions} We have presented preliminary results from a wide-field but sensitive HI synthesis survey of the outskirts of the M81 Group, in an area encompassing the dwarf galaxies HoII, M81dwA, and UGC4483. No signs of large-scale tidal features are present, and no intergalactic HI clouds are detected, this to a level better than $1.5\\times10^{19}$~atoms~cm$^{-2}$ or $1.25\\times10^5$~$M_{\\scriptsize \\sun}$ (per channel and beam), which appears to rule out counterparts to the Local Group HVCs if located at cosmologically relevant distances. The HI distribution and kinematics of HoII are consistent with that expected from tidal interactions, especially M51-like events, but ram pressure models including gas cooling are also possible. It thus appears premature to rule out either possibilities, and thus the presence of a hot extended IGM in the M81 Group. It is also possible that M81dwA is a tidal dwarf. A search for a (faint) stellar counterpart to the northwest HI arms in HoII should resolve the issue, given that ram pressure is only expected to affect the gas."
    },
    "0311/hep-ph0311240_arXiv.txt": {
        "abstract": "We show that a moduli space of the form predicted by string theory, lifted by supersymmetry breaking, gives rise to successful inflation for large regions of parameter space without any modification or fine tuning. This natural realization of inflation relies crucially on the complex nature of the moduli fields and the multiple points of enhanced symmetry, which are generic features of moduli space but not usually considered in inflationary model building.  Our scenario predicts cosmic perturbations with an almost exactly flat spectrum for a wide range of scales with running on smaller, possibly observable, scales. The running takes the form of either an increasingly steep drop off of the spectrum, or a rise to a bump in the spectrum before an increasingly steep drop off. ",
        "introduction": "\\label{intro}  While inflation \\cite{Gliner,guth1} solves many cosmological problems, the realization of such an accelerated expansion phase in the early Universe in particle theory still remains an open question. Modular cosmology has also attracted a lot of attention \\cite{example,towards} as string theory generically predicts ubiquitous moduli fields, though the moduli tend to cause cosmological problems \\cite{bs} rather than solve them. We introduced a successful modular cosmology scenario in Ref.~\\cite{smc} where the problems in modular cosmology, in addition to the persistent fine-tuning problems in inflationary model building \\cite{vpp}, are solved in a simple and self-consistent manner. In this paper we show that the unresolved problem of the simplest single modulus model of Ref.~\\cite{smc} is naturally resolved by a more careful analysis of the expected form of the modulus' potential (see Section~\\ref{ntheta}), and give a more complete analysis of the cosmic perturbations produced in the model.  The outline of our paper is as follows. In Section~\\ref{modulispace} we review the expected form of a modulus' potential. In Section~\\ref{inflation} we describe our scenario of inflation on moduli space and present our detailed potential, including a toy model which we use for numerical calculations. In Section~\\ref{perts} we calculate the cosmic perturbations produced in our model. In Section~\\ref{obs} we summarize the observational predictions of our model. In Section~\\ref{conc} we discuss and conclude. ",
        "conclusions": "\\label{conc}  We investigated inflation on a simple one complex dimensional moduli space and found that the inflationary dynamics is surprisingly rich. Both real dimensions as well as multiple points of enhanced symmetry play an essential role. Simplifying to a single real dimensional model or neglecting the effects of the points of enhanced symmetry would lead to completely different expectations for the spectrum of curvature perturbations produced during the inflation. Interestingly, while such over simplified models tend to produce observationally inviable spectra, we have shown that, for large regions of parameter space, inflation on a simple one complex dimensional moduli space produces an almost exactly flat spectrum for a wide range of scales consistent with current observations. On smaller, possibly observable, scales running of the spectrum becomes significant providing the opportunity for the model to be tested by future observations.   \\subsection*{Acknowledgements} E.D.S. thanks Misao Sasaki and Takahiro Tanaka for helpful discussions, and K.K. thanks Joanne D. Cohn for helpful advice and continuous encouragement. We also thank the SF03 Cosmology Summer Workshop for hospitality while this work was in progress. E.D.S. was supported in part by the Astrophysical Research Center for the Structure and Evolution of the Cosmos funded by the Korea Science and Engineering Foundation and the Korean Ministry of Science, and by the Korea Research Foundation grant KRF PBRG 2002-070-C00022. K.K. was supported in part by NSF under grant AST-0205935."
    },
    "0311/astro-ph0311296_arXiv.txt": {
        "abstract": "Relativistic phenomena (orbital precession, Shapiro delay, and/or orbital decay) have been measured in Arecibo timing observations of four pulsar--white dwarf binaries, leading to constraints on the neutron star masses.  We have detected the decay of the PSR J0751+1807 orbit due to gravitational radiation emission, the first such measurement in a binary with a low mass ratio ($m_2/m_1\\ll 1$).  Timing data constrains the mass of this pulsar to bet between 1.6 and 2.8\\,M$_{\\sun}$.  Masses of the other pulsars are in marginal agreement with the canonical pulsar mass of 1.35\\,M$_{\\sun}$, but higher values are preferred. ",
        "introduction": "The conventional wisdom is that all pulsars have mass close to the Chandrasekhar value, $\\sim 1.35\\,{\\rm M}_{\\sun}$.  This belief has arisen from high precision measurements of neutron star--neutron star binary systems;  in all cases, the pulsar and the secondary neutron star have masses within a narrow range, 1.33 to 1.45\\,M$_{\\sun}$ (Stairs et al. 2002, Weisberg \\& Taylor 2003, Deich \\& Kulkarni 1996). Perhaps all neutron stars are born within this narrow mass range, or perhaps the particular evolution of double neutron star leads to this result (Brown 1995). Even if all neutron stars are born near the Chandrasekhar mass, they may achieve higher masses through accretion of matter from orbital companions.  Millisecond pulsars in orbits with white dwarfs, in particular, are known to have accreted mass during their low-mass X-ray binary phases, and hence they are prime candidates to be relatively massive pulsars.  There has been some observational evidence for heavy pulsars in these systems (Kaspi, Taylor, \\& Ryba 1994; van Straten et al. 2001), although a statistical study argued that they are no heavier than pulsars in double neutron star binaries (Thorsett \\& Chakrabarty 1999).  In this status report, we summarize observations of four millisecond pulsars in pulsar--white dwarf binaries made with the Arecibo telescope. This is an update of Nice, Splaver, \\& Stairs (2003), incorporating significant new timing data from PSR~J0751+1807 and an improved timing model (including orbital-annual parallax) for PSR J1713+0747.   \\begin{table}[t] \\caption{Mass Measurements from Timing Analysis} \\begin{tabular}{l@{}lcccc} \\tableline \\tableline \\multicolumn{2}{c}{Pulsar} & Orbital  & Eccentricity & Mass         & Pulsar Mass \\\\ &             & Period   &              &  Function    &  (95\\% confidence) \\\\ &             & $P_b$ (days)   &   $e$  & $f_1$ (M$_{\\sun}$) & $m_1$ (M$_{\\sun}$) \\\\ \\tableline J&0621+1002 & \\phantom{0}8.32 & 0.002\\,457 & 0.0270 & 1.1$-$2.3 \\\\ J&0751+1807 & \\phantom{0}0.26 & 0.000\\,003 & 0.0010 & 1.6$-$2.8 \\\\ J&1713+0747 &           67.83 & 0.000\\,075 & 0.0079 & 1.2$-$2.1 \\\\ B&1855+09   &           12.33 & 0.000\\,022 & 0.0056 & 1.4$-$1.8 \\\\ \\tableline \\tableline \\end{tabular} \\end{table} ",
        "conclusions": "The sources individually provide some evidence for pulsars substantially more massive than the Chandrasekhar value, and the collection of constraints shown in Figure 1 provides a compelling case for heavy neutron stars."
    },
    "0311/astro-ph0311069_arXiv.txt": {
        "abstract": "We model the infrared (IR) emission from the ring-like dust disk around the main-sequence (MS) star $\\epsilon$ Eridani, a young analog to our solar system, in terms of a porous dust model previously developed for the extended wedge-shaped disk around the MS star $\\beta$ Pictoris and the sharply truncated ring-like disks around the Herbig Ae/Be stars $\\hra$ and $\\hda$. It is shown that the porous dust model with a porosity of $\\simali$90\\% is also successful in reproducing the IR to submillimeter dust emission spectral energy distribution as well as the 850$\\mum$ flux radial profile of the dust ring around the more evolved MS star $\\eri$. Predictions are made for future {\\it SIRTF} observations which may allow a direct test of the porous dust model. ",
        "introduction": "} Together with Vega, Fomalhaut ($\\alpha$ PsA), and $\\beta$ Pictoris, the nearby (distance to the Earth $d \\approx 3.2\\pc$) low-mass ($\\mstar \\approx 0.8\\msun$) K2V main-sequence (MS) star $\\epsilon$ Eridani constitutes the ``Big-Four'' family of prototypical ``Vega-excess'' stars, which are MS stars with infrared (IR) radiation in excess of what is expected from their stellar photospheres (Gillett 1986). The IR excess associated with ``Vega-excess'' stars is generally attributed to the thermal emission of dust grains orbiting the central star in the form of a disk or ring and heated by its stellar radiation. The dust is thought to be ``second generational'' in nature rather than ``primordial''; i.e., the dust is considered as {\\it debris} generated by collisions of larger objects such as cometary and/or asteroidal bodies rather than {\\it remnants} left over from the star formation process (Backman \\& Paresce 1993; Lagrange, Backman, \\& Artymowicz 2000; Zuckerman 2001).  In recent years, $\\eri$ has received much attention since it is thought to be analogous to the solar system at an age about 10--20\\% of that of the Sun: (1) in comparison with the other three ``Big-Four'' prototypes which all are much more luminous A-type stars, $\\eri$ is {\\it closer} in type, mass, and luminosity ($L_\\star \\approx 0.33 L_\\odot$; Soderblom \\& D\\\"appen 1989) to the Sun (see Greaves \\& Holland 2000); (2) $\\eri$ has a ring-like structure of $\\simali$55$\\AU$ in radius (which is very similar to the Kuiper Belt in size) and a central cavity inside $\\simali$30$\\AU$ (which is equivalent in scale to Neptune's orbit) as revealed by its 850$\\mum$ image obtained using the SCUBA ({\\it Submillimetre Common-User Bolometer Array}) camera at the {\\it James Clerk Maxwell Telescope} (JCMT; Greaves et al.\\ 1998). (3) as the oldest (with an age of $\\approx 800\\Myr$; Henry et al.\\ 1996) ``Big-Four'' Vega-type star, the formation of planets in its disk may have completed since planet formation is thought to be complete by about 10 to 100$\\Myr$ (Lissauer 1993); (4) a roughly Jupiter-mass planet with a period of $\\simali$6.9$\\yr$ appears to be present inside the central cleared region of the $\\eri$ disk, as suggested by the radial velocity measurements spanning the years 1980.8--2000.0 (Hatzes et al.\\ 2000).  To better understand the origin and evolution of protoplanetary dust disks and the formation process of planetary systems, we need to know the physical, chemical, and dynamical properties of their constituent grains. In previous papers (Li \\& Greenberg 1998, Li \\& Lunine 2003a,b), we have shown that a simple porous dust model is successful in reproducing the spectral energy distributions from the near-IR to millimeter wavelengths of the disks around the Herbig Ae/Be stars $\\hra$ and $\\hda$ and the Vega-type star $\\beta$ Pictoris, including the 9.7$\\mum$ amorphous and the 11.3$\\mum$ crystalline silicate features or the 7.7$\\mum$ and 11.3$\\mum$ polycyclic aromatic hydrocarbon (PAH) features. In this {\\it Letter} we will study the dust IR to submillimeter emission of the disk around the more evolved star $\\eri$. The objectives of this {\\it Letter} are four-fold: (1) we wish to infer the physical and chemical properties of the dust in the $\\eri$ disk and to infer the disk structure; (2) we wish to gain insights into the evolution of dust mass, size, morphology, and mineralogy at different stellar ages; (3) we wish to know how widely applicable is the porous dust model; if it can be shown that the porous dust model is ``universally'' valid for disks at different evolutionary stages and with different geometrical structures, it will be a valuable guide for interpreting future data sets obtained by the upcoming {\\it Space Infrared Telescope Facility} (SIRTF); (4) we wish to make mid-IR spectral, imaging and broadband photometry predictions for the $\\eri$ disk; these predictions can be compared to future SIRTF observations in order to test further the validity of the porous dust model. ",
        "conclusions": "} The porous dust model, previously developed for the extended wedge-like disk around the $\\simali$15$\\Myr$-old MS star $\\beta$ Pictoris and the narrowly confined ring-like disks around the Herbig Ae/Be stars $\\hra$ ($\\simali$8$\\Myr$ old) and $\\hda$ ($\\simali$5$\\Myr$ old), is shown also applicable to the dust disk around the $\\simali$800$\\Myr$-old MS star $\\eri$, the oldest ``Big-Four'' Vega-excess star. Modeled as a ring peaking at $\\simali$55$\\AU$ from the star with a FWHM $\\simali$30$\\AU$, analogous to the young solar system Kuiper Belt, the observed IR to submillimeter emission as well as the 850$\\mum$ flux density radial distribution of the $\\eri$ disk is closely reproduced by porous grains of $\\simali$$7.1\\times 10^{-3}\\mearth$ mass consisting of $\\simali$90\\% vacuum in volume. Predictions are made for future SIRTF/MIPS imaging observations that will offer a useful test of the porous dust model."
    },
    "0311/astro-ph0311543_arXiv.txt": {
        "abstract": "Planar flames are intrinsically unstable in open domains due to the thermal expansion across the burning front---the Landau-Darrieus instability.  This instability leads to wrinkling and growth of the flame surface, and corresponding acceleration of the flame, until it is stabilized by cusp formation.  We look at the Landau-Darrieus instability for C/O thermonuclear flames at conditions relevant to the late stages of a Type Ia supernova explosion.  Two-dimensional direct numerical simulations of both single-mode and multi-mode perturbations using a low Mach number hydrodynamics code are presented.  We show the effect of the instability on the flame speed as a function of both the density and domain size, demonstrate the existence of the small scale cutoff to the growth of the instability, and look for the proposed breakdown of the non-linear stabilization at low densities.  The effects of curvature on the flame as quantified through measurements of the growth rate and computation of the corresponding Markstein number.  While accelerations of a few percent are observed, they are too small to have any direct outcome on the supernova explosion. ",
        "introduction": "\\label{sec:intro}  A carbon/oxygen flame propagating outward from the center of a white dwarf is subjected to a number of instabilities that wrinkle and accelerate the flame.  If the flame can accelerate to a significant fraction of the speed of sound, a deflagration alone can account for the explosion (see \\citealt{hillebrandtniemeyer2000} and references therein; \\citealt{reinecke2002b, gamezo2003}).  One such instability, the Landau-Darrieus (LD) instability \\citep{darrieus:1938,landau:1944} affects a planar flame front, even in the absence of gravity, driven by the thermal expansion across the flame.  The LD unstable flame will wrinkle, eventually forming cusps in the nonlinear regime that stabilize the flame.  The growth rate for a LD unstable flame, including the effects of the finite thickness of the flame, $l_f$, is \\begin{equation} \\label{eq:ld} \\omega = k U_l \\frac{\\alpha}{\\alpha + 1} \\left [ \\sqrt{\\alpha - \\frac{1}{\\alpha} + 1 + k \\; l_f \\Ma (k \\;l_f \\Ma + 2 \\alpha)} - 1 + k \\; l_f \\Ma \\right ] \\end{equation} \\citep{zeldovich85}, where $\\alpha = \\rho_{\\mathrm{fuel}}/\\rho_{\\mathrm{ash}}$ is the density ratio across the flame, $k$ is the wavenumber, $U_l$ is the laminar flame speed, and $\\Ma$ is the Markstein number.  The Markstein number is a measure of the response of the flame speed, $U$, to curvature, \\begin{equation} U = U_l \\left ( 1 + \\Ma \\; l_f \\frac{\\partial^2 y_f}{\\partial x^2} \\right ) \\enskip , \\end{equation} where $y_f$ is the position of the flame interface and $x$ is the transverse coordinate.  We note that our sign convention for $\\Ma$ is opposite that used in \\citet{zeldovich85} (see \\citealt{flame-curvature} for a discussion of the Markstein number for astrophysical flames).  This growth rate has been confirmed experimentally \\citep{clanet1998} for chemical flames.  In this paper, we present two-dimensional direct numerical simulations (DNS) of the LD instability for low-density C/O flames in white dwarf interiors. We start by briefly reviewing previous LD simulations and results relevant for thermonuclear flames.  \\citet{sivashinsky1977} and \\citet{michelson1977} studied the non-linear regime of the LD instability analytically and numerically, using an integro-differential equation to describe the deformation of a discontinuous flame front.  This was extended later by \\citet{gutman1990}, who concentrated on the regime where thermal conduction dominates over mass diffusion (this is the regime appropriate for astrophysical flames).  They experiment with boundary conditions (both adiabatic and periodic) and discuss the influence of the domain size on the final cusp behavior.  Their simulations show a multi-mode perturbation cusping, and slowly, the cusps merging until only a single one survives.  In the widest domains, a superimposed cellular structure appears on the single remaining cusp.  The mechanism behind this cellular instability does not, at this moment, seem to be well understood (although see \\citealt{kupervasser}).  A similar approach, using a more general equation for the front, the Frankel equation, considering spherically expanding flames, was performed by \\citet{blinnikovsasorov}.  They observe a fractal structure for the front through a range of spatial scales.  The unstable flame exhibits cell-splitting, owing to the expanding spherical geometry.  \\citet{helenbrooklaw1999} presented simulations that treat the fuel and ash states as incompressible, and link them with jump conditions across the discontinuous front.  In contrast to other methods, feedback between the flame and fluid flow is included in this model.  They found that the flow dominates the LD instability in setting the scale for the wrinkling.  This has important consequences for the flame in Type Ia supernova (SNe Ia), as the medium will be turbulent, and some evidence suggests that a large-scale dipole flow is setup during the ignition phase of the explosion \\citep{kuhlen2003}.  In the real star, it may be likely that the fluid motions swamp any effects of the LD instability.  Furthermore, the LD instability will be competing with the Rayleigh-Taylor instability \\citep{SNrt}.  Recently, front tracking methods treating a discontinuous flame coupled to full hydrodynamics have been used for both terrestrial \\citep{qian1998} and astrophysical flames \\citep{rhn2003} to investigate LD unstable flames.  The work of \\citet{rhn2003} used a level set method coupled to fully compressible PPM \\citep{ppm} to follow a discontinuous flame front, using flame speeds from \\citet{timmeswoosley1992}.  In their highest resolution run, they report seeing a cellular structure distorting the single cusp, perhaps arising through the same process as described in \\citet{gutman1990}. \\citet{rhn2003} find the flame speed rising to $1.3\\times$ the laminar speed.  The work described above all treat the flame as a discontinuous front---some model is required to describe the behavior of the flame on scales smaller than the grid resolution.  The finite thickness of the flame is known to have some influence on the LD instability, such as setting the small scale cutoff for the growth, so simulations that resolve the thermal structure are important.  Direct numerical simulations of the LD instability at various densities in the astrophysical context were performed by \\citet{niemeyerhillebrandt1995a}, showing cusp formation, but it was unclear from their studies whether there was any acceleration over the laminar speed.  Their lowest density flame ($5\\times 10^7~\\gcc$) showed significant growth beyond simple cusping, but it is likely that this was either numerical in origin, or due to the interaction of the transient perturbations left over from the initialization with the flame.  LD unstable flames in laminar flows are unlikely alone to lead to the accelerations required of the SNe Ia models.  An interesting question is whether the non-linear stabilization can break down, whether through some property intrinsic to the flames or through external forces. \\citet{kerstein1996} discusses one such mechanism, active turbulent combustion, involving the interaction of the LD unstable flame with turbulence.  It has been proposed that this process could occur in Type Ia supernovae \\citep{niemeyerwoosley1997}, since large amounts of the thermonuclear energy released goes into the expansion of the star, and it is this thermal expansion that drives the LD instability.  \\citet{niemeyerwoosley1997} argue that this feedback between the flame and its own turbulence could lead to large accelerations of the flame, although, at present, this has not been demonstrated in numerical simulations.  We focus on DNS of the Landau-Darrieus instability at low densities ($\\rho \\sim 2\\times 10^7~\\gcc$ -- $8\\times 10^7~\\gcc$) using a low Mach number numerical formulation, as described in \\citet{SNeCodePaper}.  Here we are able to extend to lower densities than previous studies, and also, through the use of adaptive mesh refinement, use finer resolution and larger domains.  For the present study, we restrict ourselves to two-dimensional calculations.  These direct numerical simulations are truly parameter free (i.e., no laminar flame speed is input), and could serve to test\\slash calibrate different methods of representing unresolved flames on stellar-sized domains.  Corresponding DNS calculations of the Rayleigh-Taylor instability will be presented in a second paper \\citep{SNrt}. Three-dimensional results will be presented in a future paper. Additionally, we do not look to seed the flow with turbulence in the present study, but rather choose to focus solely on the pure LD instability.  We provide a brief description of the numerics and parameters used for the present simulations in \\S\\ref{sec:numerics}. In \\S\\ref{sec:results}, we discuss the results, and finally, in \\S\\ref{sec:conclusions} we conclude. ",
        "conclusions": "\\label{sec:conclusions}  We presented direct numerical simulations of the Landau-Darrieus instability at low densities in conditions relevant for Type Ia supernova explosions.  We employed a low Mach number hydrodynamics method---new in the application to astrophysics---to advance the flows and enable us to follow the development of the LD instability from the linear regime through the nonlinear regime.  Resolving the flame ensures that we account for the full coupling between the flame and the flow and the effects of curvature.  The consistency of our results with the theoretical predictions validate this approximation for astrophysical flames.  Fully compressible versions of the same simulations would not be possible, or prohibitively expensive and error prone.  Acceleration of the flame was seen, due to the increase of flame surface, until the nonlinear cusp formation sets in, at which point the velocity stabilized a few percent higher than the laminar speed. The maximum increase in flame speed observed was $\\sim 2\\%$---this is in contrast to the results presented in \\cite{rhn2003}, where accelerations of $\\sim 30\\%$ were observed.  There are several possible reasons to explain these differences.  Most of the numerical results for the LD instability treat the flame as infinitely thin. Here we've included the effects of the finite thickness of the flame, and therefore, implicitly included curvature effects.  At the densities we consider, the burning is expected to terminate at the intermediate mass elements, so we burn to magnesium whereas they burn to nickel.  Furthermore we use half carbon/half oxygen fuel instead of pure carbon.  This will affect the expansion factor across the flame. Finally, fully understanding the role of the numerics requires a detailed comparison of the different methods, using the same input physics, and could potentially be a research project in the future.  The speedup increases with decreasing density (or increasing expansion parameter), as expected.  With time, it seems as if the velocities asymptote to the same value, for a given density.  When a multimode perturbation is considered, it is observed that the cusps eventually merge into a single cusp.  We were able to perform fits of the growth of the cusp amplitude with time to compute the growth rates.  Using the theoretical prediction (Equation~[\\ref{eq:ld}]), we were able to compute the Markstein number for these flames, finding values $< -2.0$, and growing in magnitude with decreasing density.  These values and the trend with density are consistent with the results of \\citet{flame-curvature}.  For the $4\\times 10^7~\\gcc$ density run, the large number of domain widths allowed us to compare to the predictions of the fastest growing mode and small scale cutoff from the theoretical dispersion relation, and we found strong agreement.  Furthermore, this demonstration of the small scale cutoff and the deviation from the infinitely thin flame ($\\Ma = 0$) estimates of the growth rate, as well as the difference in the growth of the flame speed and area emphasize the importance of curvature in small-scale flame dynamics.  As we zoom out from these fully resolved flames, several steps will be needed at intermediate scales between the fully resolved flame, where curvature effects appear important, and the full star, with validation of the flame model on each scale through direct comparison to the immediately smaller scale simulations.  The present direct numerical simulations provide the anchor for this subgrid model validation.  These curvature effects will be explored more fully in paper II, focusing on the reactive Rayleigh-Taylor instability.  We saw no breakdown of the nonlinear stabilization mechanism as hinted to in the calculations presented by \\citet{niemeyerhillebrandt1995a}. Furthermore, none of the flames we studied showed a superposed cellular structure, although, there is some evidence \\citep{gutman1990} that this phenomena is present only on very large scales, and \\citet{blinnikovsasorov} admit that it may actually be seeded by numerical noise.  The widest domain that we were able to simulate was $\\sim 350$ flame widths.  \\citet{helenbrooklaw1999} only saw this instability when they increased their domain size to 20~times the wavelength where the LD growth rate is the maximum.  This would translate into a domain size of $230$~cm for the $4\\times 10^7~\\gcc$ flame, which would require a 23,000 zone wide simulation if we are to resolve the flame.  The growth rate for a domain that wide would be very small, requiring a much longer simulation.  This is very expensive, and furthermore, it is not known if this factor of 20 is universal.  We are not aware of any resolved calculations of the LD instability that have shown this secondary instability.  We note that all of our calculations considered planar geometries only.  On the scales where the spherical geometry of the star is likely to make a difference, the flame will be disturbed by the larger scale motions in the star, other instabilities as well (notably the Rayleigh-Taylor instability, presented in \\citealt{SNrt}) and subjected to the interaction with flame-generated turbulence.  All of these processes will dominate the pure LD instability, considered in the present paper."
    },
    "0311/astro-ph0311319_arXiv.txt": {
        "abstract": "We review the main results from recent numerical simulations of turbulent fragmentation and star formation. Specifically, we discuss the observed scaling relationships, the ``quiescent'' (subsonic) nature of many star-forming cores, their energy balance, their synthesized polarized dust emission, the ages of stars associated with the molecular gas from which they have formed, the mass spectra of clumps, and the density and column density probability distribution function of the gas. We then give a critical discussion on recent attempts to explain and/or predict the star formation efficiency and the stellar initial mass function from the statistical nature of turbulent fields. Finally, it appears that turbulent fragmentation alone cannot account for the final stages of fragmentation: although the turbulent velocity field is able to produce filaments, the spatial distribution of cores in such filaments is better explained in terms of gravitational fragmentation. ",
        "introduction": "It is well known that stars form in dense cores within molecular clouds, but the physical processes that control the formation of low-mass stars within these cores are not well understood yet. The ``standard'' scenario assumes that low-mass protostellar cores are in quasi-static equilibrium supported against gravitational collapse by a combination of magnetic and thermal pressures (see, e.g., \\opencite{Shu_etal87}). In this scenario, a core forms stars once magnetic support is lost through the process known as ``ambipolar diffusion'': neutral gas particles slowly drift through the ions, which are held up by the magnetic field, allowing the core to eventually attain the ``critical'' mass-to-flux ratio, beyond which the gravitational energy exceeds the magnetic energy and collapse sets in from the inside-out.  The more recent turbulent scenario, on the other hand, suggests that cores are formed by compressible motions in the turbulent velocity field of the cloud \\cite{Sasao73, Elmegreen93, Padoan95, BVS99, KHM00}. Those cores with an excess of gravitational energy collapse rapidly to form stars, while the others with sufficiently large internal or kinetic energies re-expand once the turbulent compression subsides \\cite{Taylor_etal96, VSB02}  Although the standard scenario has explained many observational features of protostellar cores, it does not address the problem of how the quasistatic cores can form and persist within a turbulent environment. On the other hand, the turbulent scenario naturally incorporates the process of core formation, and is increasingly succesful in reproducing the observed core properties. In this paper, we will review the main results of the turbulent scenario, the observational facts that numerical works have successfully reproduced, the simulations' predictions concerning the star formation efficiency and the attempts to explain the initial mass function from statistical considerations of a turbulent field. Finally, we discuss the final stages of fragmentation, which appear to be gravitational rather than turbulent. ",
        "conclusions": "This review shows that numerical simulations of turbulent fragmentation have been quite successful in reproducing a number of observational properties of molecular cloud cores, while simultaneously making a number of predictions about the behavior of turbulence in molecular clouds and their cores, as well as providing an explanation for another set of observed features. Although much work remains to be done, and a full theory is still elusive, the progress made is highly encouraging, and makes the dynamical picture of star formation a plausible alternative to the ambipolar-diffusion-mediated magnetostatic-equilibrium standard model.   \\vskip -1 cm"
    },
    "0311/astro-ph0311633_arXiv.txt": {
        "abstract": "I present two-band HST STIS imaging, and WIYN spectral mapping, of ongoing mass transfer in the interacting galaxy pair NGC 1409/10 (where NGC 1410 is the Seyfert galaxy also catalogued as III Zw 55). Archival snapshot WFPC2 imaging from the survey by Malkan et al. showed a dust feature stretching between the galaxies, apparently being captured by NGC 1409. The new images allow estimates of the mass being transferred and rate of transfer. An absorption lane typically 0.25\" (100 pc) wide with a representative optical depth $\\tau_B = 0.2$ cuts across the spiral structure of NGC 1410, crosses the 7-kpc projected space between the nuclei, wraps in front of and, at the limits of detection, behind NGC 1409, and becomes a denser ($\\tau_B = 0.4$) polar feature around the core of NGC 1409. Combination of extinction data in two passbands allows a crude three-dimensional recovery of the dust structure, supporting the front/back geometry derived from colors and extinction estimates. The whole feature contains of order $10^6$ solar masses in dust, implying about $2 \\times 10^7$ solar masses of gas, requiring a mass transfer rate averaging $\\approx 1$ solar mass per year unless we are particularly unlucky in viewing angle. Curiously, this demonstrable case of mass transfer seems to be independent of the occurrence of a Seyfert nucleus, since the Seyfert galaxy in this pair is the donor of the material. Likewise, the recipient shows no signs of recent star formation from incoming gas, although NGC 1410 has numerous luminous young star clusters and widespread H$\\alpha$ emission. ",
        "introduction": "Galaxy interactions are clearly linked to bursts of star formation, and in more restricted ways, to triggering of nuclear activity. What physical mechanisms mediate these connections remains unclear, since there are multiple well-motivated and plausible processes. Among these, mass transfer between the galaxies in an interacting system has been long discussed, but evidence for its occurrence has remained largely circumstantial. Such transfer could play a particular role in dumping gas into the inner regions of a galaxy, since the angular-momentum barrier which restricts the inward transport of disk gas may be greatly reduced if the gas comes from another galaxy with the appropriate encounter geometry.  The conditions for mass transfer have been examined analytically as well as numerically. Sotnikova (1990) considered the fate of dense ISM clouds crossing between galaxies and heated by both conduction and radiation in a hot IGM, finding that the clouds could survive long enough to reach the second galaxy, and that a stream of clouds from galaxies in low-eccentricity orbits could reach a quasi-steady state configuration. Her numerical simulations (Sotnikova 1988) indicated that this situation is reached much faster when the spiral spin and orbit are parallel than the antiparallel case, in roughly an orbital timescale of the binary system.  Wallin \\& Stuart (1992) presented an extensive study of mass transfer from restricted $n$-body simulations, including over 1000 encounter geometries to map the systematics of mass transfer. The largest amounts finally captured by the second galaxy are found for encounters with the orbital plane close to the donor's disk plane, with this favoritism enhanced by the strong dependence on orientation of periapsis for non-planar encounters. The intuitive increases in mass transfer for closer periapsis and larger mass ratio also hold. These simulations do not give much detail on how the companion captures the mass, but should be robust as to the totals bound in its potential well. These results fit with patterns in the 86 simulations described by Howard et al. (1993), where the particles were tracked but there was no modelled companion structure to interact in detail with the incoming particles.  The observational signatures of mass transfer can be hidden in the complex kinematics and star-formation behavior of gas-rich interactions and mergers. Thus, it has long been recognized that such effects are easiest to see in pairs of mixed morphology, since large amounts of cool gas in the early-type E/S0 members of these systems very likely originated in the gas-rich companion (as set forth by Domingue et al. 2003). This has been particularly discussed in the context of emission lines, far-infrared emission, and blue (post-starburst) optical colors for these galaxies, as possible evidence of past gas transfer (de Mello et al. 1995, 1997). This would also account for the Holmberg effect, a correlation between colors of pair members, even for early-type members (Holmberg 1958, Demin et al. 1984, Madore 1986, Reduzzi \\& Rampazzo 1995). Marziani et al. (2003) present kinematic evidence of ongoing or recent dumping of gas in the pair Arp 194, and stress that mass transfer itself is necessary but not sufficient for cross-fuelling of nuclear activity or starbursts.  An unusually clear example of current mass transfer appeared in the HST WFPC2 snapshot survey of AGN and starburst galaxies by Malkan, Gorjian, \\& Tam (1998). Their data for NGC 1409 showed a dust lane crossing between NGC 1409 and its close companion (and Seyfert galaxy) NGC 1410, also known as III Zw 55 (although there has been some ambuguity on the literature as to which member of the pair is which). This feature extends even beyond the region shown in Fig. 1 of Malkan et al. (where the galaxy identified there as the Seyfert galaxy NGC 1410 is in fact its earlier-type and more quiescent companion NGC 1409). Such cases highlight the value of dust absorption as a tracer of the interstellar medium, also exploited by several groups analyzing HST imagery of the nuclei of early-type galaxies (Martel. et al. 2000, Tran et al. 2001, Martini et al. 2003).  The extensive tidal envelope of stars around NGC 1409/10 provides backlighting to trace this feature and estimate the total mass and transfer rate involved. To allow such a study, two-color STIS imagery was obtained with higher signal-to-noise in each band. The results are described here, along with mapping in H$\\alpha$ and [N II] which constrains the encounter geometry and timing through the velocity field.  We see the NGC 1409/10 pair deep into a strong encounter, as shown by the presence of a substantial stellar envelope extending well beyond the two galaxy cores, to a radius at least 15 kpc. The nuclei are projected at a separation of only 7 kpc, suggesting that there may have been significant tidal loss of stars from both galaxies. The spiral NGC 1410 in particular is rather small in linear size, with the nearly face-on spiral pattern spanning a diameter of 5.3 kpc (for $\\rm H_0=72$ km s$^{-1}$ Mpc $^{-1}$ following the WMAP analysis of Spergel et al. 2003). This galaxy hosts a Seyfert nucleus and active disk star formation. In contrast, NGC 1409 is a barred early-type system, best described as type SB0. Because of the small angular extent of these structures, these types are not obvious from typical ground-based images (such as shown by Keel 1996). ",
        "conclusions": "HST STIS images in two broad passbands have been used to trace a dust lane marking mass transfer in the interacting galaxy pair NGC 1409/10. The combined color and extinction behavior support the impression that there is a single feature crossing the 7-kpc gap between NGC 1410 and NGC 1409, then wrapping behind NGC 1409 and reappearing on the opposite side to become a small lane crossing its disk over the pole. The depth into the starlight derived from the two-color data and the residual intensity  yield estimates of the dust mass, totalling about $2 \\times 10^6$ solar masses and likely accompanied by $3 \\times 10^8$ solar masses of gas. At the characteristic orbital velocities in this system, this suggests a mean rate of mass transfer slightly above one solar mass per year onto NGC 1409.  This mass-transfer calculation may be too simple, since the feature may be more exactly described as an isochrone than a pipeline. It is striking that only in a small region has material been launched away from NGC 1410, perhaps indicating that a special location in the system was needed to start such transfer. This fits with the fact that the dust lane appears to cross the spiral pattern of NGC 1410.  In light of the mechanisms discussed to accont for star formation and nuclear activity in interacting galaxies, it is ironic that the Seyfert galaxy NGC 1410 is the donor, rather than the recipient, of mass transfer. The recipeint NGC 1409 shows no signs of any fate of the gas reaching its disk. Limits to its rate of star formation stand at a few per cent of the inflow rate, and its nuclear activity is limited to a modest LINER. Either the flow has yet to actually reach the dense inner regions, or some additional trigger must set in after a critical density or mass of infalling gas is achieved. If flows such as these drive starbursts in interacting systems, such a cycle is needed to drive star-formation rates an order of magnitude greater than the estimated flow rate in NGC 1409/10."
    },
    "0311/astro-ph0311405_arXiv.txt": {
        "abstract": "\\noindent Studies of massive black holes (BHs) in nearby galactic centers have revealed a tight correlation between BH mass and galactic velocity dispersion. In this paper we investigate how the BH mass versus velocity dispersion relation and the nuclear luminosity versus velocity dispersion relation in QSOs/active galactic nuclei (AGNs) are connected with the BH mass versus velocity dispersion relation in local galaxies, through the nuclear luminosity evolution of individual QSOs/AGNs and the mass growth of individual BHs. In the study we ignore the effects of BH mergers and assume that the velocity dispersion does not change significantly during and after the nuclear activity phase. Using the observed correlation in local galaxies and an assumed form of the QSO/AGN luminosity evolution and BH growth, we obtain the simulated observational appearance of the BH mass versus velocity dispersion relation in QSOs/AGNs. The simulation results illustrate how the BH accretion history (e.g., the lifetime of nuclear activity and the possibility that QSOs/AGNs accrete at a super-Eddington accretion rate at the early evolutionary stage) can be inferred from the difference between the relation in QSOs/AGNs and that in local galaxies. We also show how the difference may be weakened by the flux limit of telescopes. We expect that a large complete sample of QSOs/AGNs with accurate BH mass and velocity dispersion measurements will help to quantitatively constrain QSO/AGN luminosity evolution and BH growth models. ",
        "introduction": "\\label{sec:intro} \\noindent The existence of massive black holes (BHs) in nearby galactic centers, as a prediction of the widely accepted QSO model that QSOs are powered by gas accretion onto massive BHs (e.g., \\citealt{Lynden69,S82,R84}), is now believed to be confirmed \\citep[e.g.,][]{KR95,Magorrian98,Gebhardt03,Pinkney03}. Studies of the BHs in nearby galaxies have revealed that BH mass in nearby galactic centers is tightly correlated with galactic velocity dispersion \\citep{FM00,Geb00a,Tremaine02} and is also (less tightly) correlated with the luminosity (or mass) of elliptical galaxies or bulges of spiral/S0 galaxies (e.g., \\citealt{KG01}). These correlations suggest a close link between the formation and evolution of BHs and their host galaxies. However, why these correlations exist and whether they also exist in distant galaxies have not yet had definite answers.  The purpose of this paper is to investigate the relationship between BH mass and velocity dispersion in QSOs/active galactic nuclei (AGNs). It is important to study this in QSOs/AGNs at least for the following two reasons: \\begin{enumerate} \\item Various BH growth models have been proposed to explain the origin of the tight correlation between BH mass and velocity dispersion in nearby galaxies \\citep[e.g.,][]{SR98,Fabian99,Blandford99,Ostriker00,HK00,BS01,AGR01,ITS03,King03}, including accretion of either baryonic gas or non-baryonic dark matter onto seed BHs or hierarchical mergers of intermediate-mass BHs (which might be end products of the first generation of stars, or Population III stars) with masses of typically a few hundred $M\\sun$ \\citep[e.g.,][]{Schetal02}. How can these BH growth models be tested by observations? Among these BH growth processes, currently BH mergers are unlikely to be observed directly although detecting gravitational wave signals emitted from BH mergers might be possible in the future. The possible observational features of accretion of non-baryonic material are still unclear; only accretion of baryonic gas, which appears as QSO/AGN phenomena, is detectable and has been extensively studied. We expect that the investigation of the BH mass versus velocity dispersion relation in QSOs/AGNs will provide an observationally achievable way to understand BH growth and the origin of the correlation of local BHs with their host galaxies. As a matter of fact, some observational investigations on the relation in QSOs/AGNs have been performed in the past several years \\citep[e.g.,][]{laor98,wandel99,Geb00b,Ferrarese01,Shields03}.  \\item Furthermore, according to current observations, the total local BH mass density is consistent with the total mass density accreted onto BHs during QSO/AGN phases, which suggests that BH mass growth comes mainly from gas accretion during QSO/AGN phases, rather than from accretion of non-baryonic material or mergers of intermediate-mass BHs \\citep{YT02,AR02,Fab03}.  Since QSOs/AGNs represent the population of galaxies housing growing BHs, the study of the relationship between BH mass and velocity dispersion in QSOs/AGNs can provide valuable information on the BH growth history, the evolution of galaxies, and the origin of the tight correlation between BH mass and velocity dispersion in local galaxies. \\end{enumerate}  In this paper we show how the BH mass (or nuclear luminosity) versus velocity dispersion relation in QSOs/AGNs is connected with the BH mass versus velocity dispersion relation in local galaxies through the nuclear luminosity evolution of individual QSOs/AGNs and mass growth of individual BHs. The basic model assumptions are as follows: \\begin{enumerate} \\item BH mergers are assumed to not be important for BH growth. In principle, the mass growth of a BH may come from both gas accretion due to QSO/AGN phases and mergers with other BHs. However, currently, the BH merger process and rate are very uncertain. In addition, comparison of the mass density distribution in nearby galaxies with that accreted due to QSO phases has shown that BH mergers are not necessarily required at least for growth of high-mass ($\\ga 10^8\\msun$) BHs \\citep{YT02}. \\item The velocity dispersion of host galaxies is assumed to not change significantly during and after the nuclear activity phase.  In principle, both nuclear luminosity evolution/BH mass growth and the evolution of host galaxy velocity dispersions contribute to the evolution of the nuclear luminosity/BH mass versus velocity dispersion relation. However, ignoring the velocity dispersion evolution may greatly simplify our analysis and highlight the effects of nuclear luminosity evolution/BH mass growth. Comparison of observations with the predictions obtained by ignoring velocity dispersion evolution might also provide clues on the effect of velocity dispersion evolution. In addition, this assumption has also been adopted in many other physical models to explain the correlation between BH mass and velocity dispersion in nearby galaxies (e.g., \\citealt{SR98,Fabian99,Blandford99,BS01,AGR01,King03}). \\end{enumerate} The model includes the following three basic components: the BH mass distribution of local galaxies at a given host galaxy velocity dispersion, the curve of the nuclear luminosity evolution/mass growth of individual BHs, and the nuclear luminosity/BH mass distribution in QSOs/AGNs at the given velocity dispersion. In principle, given any two of them, the last one can be constrained by the model. In this paper we use the first two components as model inputs and the third one as the model output.  The paper is organized as follows: In \\S~\\ref{sec:m0sigma} we review current observational results on the BH mass versus velocity dispersion relation in local galaxies.  In \\S~\\ref{sec:msigmaqso} we formulate the connection of the nuclear luminosity/BH mass versus velocity dispersion relation in QSOs/AGNs with the BH mass versus velocity dispersion relation in nearby galaxies through the nuclear luminosity evolution/mass growth of individual BHs. Detailed models on the nuclear luminosity evolution/mass growth of individual BHs are presented in \\S~\\ref{sec:calLM}. Applying the observational results of the local BH demography (in \\S~\\ref{sec:m0sigma}) and the assumed QSO luminosity evolution and BH growth curves (in \\S~\\ref{sec:calLM}) to the analysis made in \\S~\\ref{sec:msigmaqso}, we predict the BH mass versus galactic velocity dispersion relation in QSOs/AGNs and show the simulation results of its observational appearance in \\S~\\ref{sec:results}. Discussion on current observational results of the relation in QSOs/AGNs is given in \\S~\\ref{sec:dis}. Finally, conclusions are summarized in \\S~\\ref{sec:con}.  In \\S~\\ref{sec:msigmaqso} readers who are primarily interested in the results of the formulation might want to skip the detailed mathematical manipulations and move straight to the model predictions given by equations (\\ref{eq:calPLsigma1}), (\\ref{eq:calPMsigma}), and (\\ref{eq:calPMsigmasimp}) (with the aid of Table~\\ref{tab:tab1}). The QSO luminosity evolution or BH growth is incorporated into those equations through the lifetime of the nuclear activity $\\tau\\life$ and a distribution function of the probability that the progenitor of a local BH with mass $M_{\\rm BH,0}$ (where the subscript ``0'' represents the current cosmic time $t_0$) had a mass $M_{\\rm BH}$ or nuclear luminosity $L$ in its nuclear activity history [i.e., $P(M_{\\rm BH}|M_{\\rm BH,0})$ or $P(L|M_{\\rm BH,0})$ in Table~\\ref{tab:tab1}]. The observational counterparts of the model predictions are given by equations (\\ref{eq:calPLsigmaQSO}) and (\\ref{eq:calPMsigmaQSO}). Comparison of the observations with the model predictions may help to strengthen the existing constraints or provide new information on the QSO/AGN luminosity evolution and BH growth.  In this paper we set the Hubble constant to $H_0=100h\\kms$, and the cosmological model used is ($\\Omega_{\\rm M}$,$\\Omega_{\\Lambda}$,$h$)$=(0.3,0.7,0.65)$. ",
        "conclusions": "\\label{sec:con} \\noindent In this paper we have investigated how the BH mass/nuclear luminosity versus velocity dispersion relation in QSOs/AGNs is connected with the local BH mass versus velocity dispersion relation through the mass growth/nuclear luminosity evolution of individual BHs (see eqs.~\\ref{eq:calPMsigma} and \\ref{eq:calPLsigma1}), by ignoring BH mergers and assuming that the velocity dispersion does not significantly change during and after the nuclear activity phase. Circumventing the uncertain history of triggering the accretion onto seed BHs, the relation in QSOs/AGNs considered in this paper is for a sample of QSOs/AGNs at all redshifts, rather than for QSOs/AGNs at a given cosmic time.  Using the observed local BH mass versus velocity dispersion relation and the assumed form of the QSO/AGN luminosity evolution and BH growth, we simulate the observational appearance of the BH mass versus velocity dispersion relation in QSOs/AGNs. The simulation results quantitatively show how the BH accretion history (e.g., characterized by the timescales $\\tau\\I$ and $\\tau\\D$ in this paper; see \\S~\\ref{sec:calLM}) affects the difference between the relation in QSOs/AGNs and that in local galaxies. A simple example to illustrate this is that: if the BH mass increases by a factor of $\\ga 10$, mainly via Eddington accretion, the relation in QSOs/AGNs will significantly deviate from the relation in nearby galaxies, with the mean logarithm of BH mass at a given velocity dispersion being smaller than that of local BHs by an absolute difference of $\\ga 0.3$~\\dex (see Figs.~\\ref{fig:f1}a and 1b). We quantitatively show that if QSOs/AGNs are not always accreting material at the Eddington rate (e.g., accreting at super-Eddington rates at their early evolutionary stages or at sub-Eddington rates at their late evolutionary stages), the distribution function of BH mass at a given velocity dispersion in QSOs/AGNs can be skewed (i.e., the probability of observing QSOs/AGNs with low BH mass will be relatively decreased compared to that for QSOs/AGNs with high BH mass). We also show that the observational difference between the relation in QSOs/AGNs and that in local galaxies may be weakened by the selection effect due to the flux limit of telescopes, since QSOs/AGNs housing small BHs may be missing from observations. To further constrain the BH accretion history from the BH mass versus velocity dispersion relation in QSOs/AGNs, a large, complete QSO/AGN sample with accurate BH mass and velocity dispersion measurements from observations is needed.  In addition, the observational appearance of the nuclear luminosity versus velocity dispersion relation in QSOs/AGNs can also be predicted by the method described in this paper, and it would also be useful to compare the expectation with future observations.  We are grateful to the referee for a careful reading of the manuscript and helpful comments. Q.Y.  acknowledges support provided by NASA through Hubble Fellowship grant HF-01169.01-A awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA, under contract NAS-5-26555."
    },
    "0311/astro-ph0311411_arXiv.txt": {
        "abstract": "We model the broad-band spectral energy distribution of the innermost ``core-jet'' region of the redshift z=1.187 quasar PKS 1127-145.  We propose a scenario where the high energy photons are produced via the Compton scattering of thermal IR radiation by the relativistic particles in a parsec-scale jet. The  high energy spectrum, together with the observed radio variability and superluminal expansion, suggest that PKS 1127-145 may be a blazar, despite the fact that its optical/UV component is likely dominated by thermal radiation from an accretion disk.  The relation of PKS 1127-145 to MeV - blazars is discussed. ",
        "introduction": "X-ray emission from  the central regions of Active Galactic Nuclei (AGN) is associated with accretion flows or jets. The latter dominate the observed X-ray fluxes in blazars, the objects with powerful, relativistic jets oriented close to the line of sight. The X-ray spectra of these objects extend to high energy $\\gamma$-rays and  origin of both depends on the total luminosity of a source (e.g. Fossati et al. 1998). In the least luminous sources, called High Energy Peaked BL Lacs, the X-rays are produced by synchrotron emission, while the $\\gamma$-rays -- presumably by  Comptonization of the synchrotron radiation (SSC). In intermediate luminosity objects -- Low Energy Peaked BL Lacs -- both X-rays and $\\gamma$-rays are likely produced by the SSC process. Finally, in the most luminous sources -- Flat Spectrum Radio Quasars (FSRQ) -- the high energy spectra are dominated by so-called external radiation Comptonization (ERC) process, but with a possible contribution to the soft/mid X-ray bands from the SSC radiation. Seed photons for the ERC process may come directly from the accretion disk (Dermer \\& Schlickeiser 1993), from the broad line region (BLR) (Sikora, Begelman \\& Rees 1994), and from an optically thick molecular torus  in the form of thermal IR radiation from hot dust (B{\\l}a{\\.z}ejowski et al. 2000). The luminosity of FSRQ detected with the Compton Gamma Ray Observatory ({\\it CGRO}) is usually dominated by GeV radiation. However, there is a subclass of FSRQ with their luminosity peaked in the MeV band.  Hereafter we call the former 'GeV-blazars', and the latter 'MeV-blazars'.  A possible explanation for this division could be related to the location of the radiation site in the jet.  As discussed by Sikora et al. (2002), in GeV-blazars radiation is produced closer to the central black hole, while in MeV-blazars, the active region is significantly farther from the center (see Sect. 3 for details). Currently, there are about a dozen known MeV-blazars  or candidates for such a class; unfortunately any searches for them are limited by the low sensitivity of $\\gamma$-ray detectors in the $1 - 30$ MeV band, where their spectra often have a peak.   PKS 1127-145 is a radio loud quasar at redshift $z=$1.187 with a radio spectrum peaked around 1~GHz, characteristic of the GigaHertz-Peaked (GPS) radio-sources (Stanghellini et al. 1998; O'Dea 1998). {\\it Chandra} observations have revealed a large scale (at least $\\sim$300 kpc) X-ray jet associated with weak radio emission (Siemiginowska et al. 2002). The extended radio structure is much weaker (nearly 3 orders of magnitude) than the core emission. This value is significantly lower in comparison to the typical FRI and FRII sources (Kellerman \\& Owen 1988).  VLBI observations show a double component with a parsec scale jet on one side (Wehrle et al. 1992; Stanghellini et al. 1998). The core emission dominates also in X-rays and the total kpc-scale X-ray jet emission is $\\sim$ 60 times fainter than the core emission (Siemiginowska et al. 2002).  The strong X-ray core has also been detected by BeppoSax (Giommi et al. 2002), with the spectrum extending up to $\\sim 300$ keV. Finally, a very likely counterpart to that core is an EGRET source 2 EG J1134-1515 (Thompson et al. 1995).  We propose that the high energy emission originating in the central parsec-scale region of the quasar PKS~1127-145 can be described in terms of the blazar phenomenon. This can be justified by the following:  \\begin{itemize}  \\item  as is the case for many $\\gamma$-ray detected blazars,  but contrary to steep spectrum radio quasars (SSRQ), the X-ray and $\\gamma$-ray luminosities of PKS~1127-145 exceed the optical/UV luminosity by a very large factor (for X-rays in SSRQ see, e.g., Sanders, et al. 1989; Laor et al. 1997);  \\item the X-ray spectrum  of PKS~1127-145 is much harder than in SSRQ ($\\alpha_x \\simeq 0.3-0.5$ vs. $0.6 - 0.8$) (e.g., Reeves \\& Turner 2000);   \\item  in contrast to SSRQ, the luminosity is peaked in the $\\gamma$-ray band (Thompson et al. 1995);   \\item the radio spectrum is flat and variable (Wehrle et al. 1992), and superluminal motion is observed (Jorstad et al. 2001).  \\end{itemize}  At the same time, unlike a typical blazar, PKS 1127-145 has a very low optical polarization (Impey and Tapia 1990) and optical variability (Bozyan et al. 1990). This suggests a thermal origin of the optical/UV radiation. Such co-existence of the strong, non-thermal $\\gamma$-ray component with a prominent thermal UV bump is a characteristic of the MeV-blazars.  We show that the jet  emission model developed by Sikora et al. (2002) for MeV-blazars can explain these spectral properties and make prediction that PKS 1127-145 can be variable  in X-rays and $\\gamma$-rays on the time scale of the order of a month. ",
        "conclusions": "Our model shows that the X-ray and gamma-ray spectrum of PKS~1127-145 may originate in the parsec scale jet. This emission can be described entirely by the Comptonization of the IR radiation with a possible contribution from ERC(BEL) at energies $>1$GeV. This object is similar to the other MeV-blazars where the SSC component is very weak and, therefore, the X-ray spectrum has a very hard slope down to the lowest energies.   Another interesting feature of PKS~1127-145 is the co-existence of a prominent UV thermal bump (as seen in radio-lobe dominated quasars), with the extremely luminous $\\gamma$-ray component. Both components are very common for the MeV-blazars (see e.g. Tavecchio et al. 2000). In our model this co-existence can be associated with the distance between the shock and the jet apex which affects the synchrotron emission. The ratio $u_{ext}'/u_B'$ is higher for larger shock distances in MeV-blazars than for smaller shock distances in GeV-blazars, where in the former case the external radiation is dominated by infrared radiation of dust, while in the latter case it is dominated  by broad emission lines, and $u_B' = B'^2/8 \\pi$ is the magnetic energy density in the shocked plasma in a jet. This results in a weaker synchrotron spectrum in MeV-blazars than in GeV-blazars assuming that both emit comparable bolometric luminosity.  Secondly, the entire synchrotron component is shifted into lower frequencies in MeV-blazars (because the average frequency of the synchrotron radiation is $\\nu_{SYN} \\propto B^{'}$ and in our model magnetic field scales with distance like $B^{'} \\propto 1/r$).  Both effects diminish the synchrotron contribution to the UV-band, thus the thermal UV-bump becomes visible. Lower synchrotron luminosities in MeV-blazars explain  also why in these objects the SSC component is weak ($L_{SSC} \\propto {L_{SYN}}^2$) and, therefore, its contribution to the X-ray band is negligible (see Sikora et al. 2002).   The available EGRET data do not constrain the location of the luminosity peak and do not provide a spectral slope for the PKS~1127-145 $\\gamma$-ray spectrum. Therefore, in order to minimize the number of free model parameters we adopted a single-power law electron injection function, instead a double-one assumed in modeling of other MeV-blazars (Sikora et al. 2002). This approach leads to a different location of the luminosity peak in our model spectrum of PKS 1127-145 ($h\\nu \\sim 100$MeV) than in the models of the other MeV-blazars ($1$MeV $<h\\nu < 30$MeV). Better $\\gamma$-ray data are needed in order to apply the same model to both MeV-blazars and PKS~1127-145 and to constrain better the model parameters. However, based on the current observations of PKS~1127-145, we conclude that it is quite similar to other MeV-blazars.  Since in the model, high energy radiation is produced at a few parsecs, we can  predict that PKS~1127-145 should be variable in the X-ray and $\\gamma$-ray bands on a time scale of the order of a month. Monitoring of the source with gamma-ray instruments such as Integral or GLAST may provide an  opportunity to verify this prediction.  PKS~1127-145 has been included in samples of GigaHertz Peaked Spectrum (GPS) radio sources (Stanghellini et al. 1998). GPS sources are considered to be at an early stage of their expansion into large scale radio sources.  However, GPS samples are heterogenous and may contain many blazars  (Lister 2003; Siemiginowska et al. 2003; Stanghellini 2003). The X-ray data provide a way to identify blazars and separate them from the ``true'' GPS sources. It is important to understand a number of blazars in the GPS samples, so we can test whether they are consistent with blazars being ``young'' GPS sources, but seen along the jet, with the emission from the GPS double radio component suppressed by the parsec scale jet emission. This is critical to our understanding of the nature and evolution of radio sources.   Finally, it is worth noting the similarity of the PKS 1127-145 with other high-z X-ray luminous quasars with very hard X-ray spectra, e.g.: 1508+5714 (Moran \\& Helfand 1997); GB 1428+4217 (Fabian et al. 1998; Fabian et al. 2001); PKS 2149-306 (Elvis et al. 2000); PMN J0525-3343 (Fabian et al. 2000). All of them are FSRQ and with their peculiar spectra, similar to those of MeV-blazars, they can represent those FSRQ which have nuclei embedded in a very dusty environment. Such hypothesis can be further explored by detailed infrared spectroscopy of those sources and by observations of $\\gamma$-ray spectra with GLAST which are expected to be relatively soft."
    },
    "0311/astro-ph0311557_arXiv.txt": {
        "abstract": "{ We present broad-band (0.1--90\\,keV) spectral and temporal properties of the three Intermediate Polars, RE\\,0751$+$144\\,(PQ\\,Gem), RX\\,J0558.0$+$5353\\,(V405\\,Aur) and RX\\,J1712.6$-$2414\\,(V2400\\,Oph) based on simultaneous soft and hard X-ray  observations with the BeppoSAX satellite. The analysis of their spectra over the wide energy range of BeppoSAX instruments allows us to identify the soft and hard X-ray components and to determine simultaneously their temperatures. The black--body temperatures of the irradiated poles of the white dwarf atmosphere are found to be 60--100\\,eV, much higher than those found in their synchronous analogues, the Polars. The temperature of the optically thin post--shock plasma  is well constrained in RX\\,J1712.6$-$2414 and in RE\\,0751$+$144 (13 and 17\\,keV) and less precisely determined in RX\\,J0558.0$+$5353. In the first two systems evidence of subsolar abundances is found, similarly to that estimated in other magnetic Cataclysmic Variables. A Compton reflection component is present in RX\\,J0558.0$+$5353 and in RE\\,0751$+$144 and it is favoured in RX\\,J1712.6$-$2414. Its origin is likely at the irradiated white dwarf surface. Although these systems share common properties (a soft X-ray component and optical polarized radiation), their X-ray power spectra and light curves at different energies suggest accretion geometries that  cannot be reconciled with a single and simple configuration. ",
        "introduction": "Intermediate Polars (IPs) are the hardest X-ray emitting Cataclysmic Variables (CVs) containing a magnetized white dwarf (WD) that is accreting material from a late-type Main Sequence star.  Unlike the strongly magnetized (B$\\sim$10-230\\,MG) orbitally phase--locked Polars for which $\\rm P_{orb}=P_{rot,WD}$, the magnetic field of the white dwarf is not strong enough ($\\rm B<$10\\,MG) to lock its rotation to the binary orbit and thus $\\rm P_{rot,WD} < P_{orb}$. IPs typically populate the long period distribution of magnetic CVs, whilst Polars are generally found at short orbital periods. Hence, the wide orbits and the low magnetic fields in IPs can allow the formation of an accretion disc, which is truncated at the magnetospheric radius. However, other accretion configurations can occurr in IPs, such as the complex disc-overflow, where the stream of matter from the secondary star overpasses the accretion disc and is  channeled towards  the magnetic poles (Hellier 1991) or the direct (or stream) accretion without an intervening disc. This latter configuration is encountered in Polars, where the accretion material directly flows towards the magnetic poles of the white dwarf.  \\begin{table*}[ht!]     % \\centering \\caption{BeppoSAX observations of RX\\,J0558, RE\\,0751 and RX\\,J1712} \\vspace{0.05in} \\begin{tabular}{llccccccc} \\hline \\hline ~ &   ~ &   ~ &     ~ &    ~ & \\cr  Object & Date & MECS &  LECS & PDS & Flux$^{5}$ \\cr  &      &  $\\rm T_{expo}$(ks)$^{1}$ ~~~~ Count Rates$^{2}$ & $\\rm T_{expo}$(ks)$^{1}$ ~~~~   Count Rates$^{2}$ &  $\\rm T_{expo}$(ks)$^{3}$  ~~~~   Count Rates$^{4}$   \\cr \\noalign {\\hrule} ~ & ~ & ~ & ~ & & \\cr RX\\,J0558 & Oct. 7-8, 1996 & 44.8 ~~~~~~~~~~~~~~~~~ 0.27 & 13.7 ~~~~~~~~~~~~~~~ 0.09 & 84.4 ~~~~~~~~~~~~~~~~ 0.32 & 2.79 \\cr  RE\\,0751  & Nov. 9-12, 1996 & 114.7 ~~~~~~~~~~~~~~~ 0.32 & 27.5 ~~~~~~~~~~~~~~~ 0.15 & 110.4 ~~~~~~~~~~~~~~~~0.38 & 5.30 \\cr  RX\\,J1712  & Aug. 17-19, 1998 & 82.6 ~~~~~~~~~~~~~~~  ~0.50 & 33.0 ~~~~~~~~~~~~~~ 0.28 & --  ~~~~~~~ ~~~~~~~~~~~~~~~ -- & 3.12  \\cr  & ~ &    ~ &             ~ &      &   \\cr \\noalign {\\hrule} \\hline \\hline \\end{tabular} ~\\par \\begin{flushleft} $^{1}$: Total on-source exposure time.\\par $^{2}$: Net count rates in units of cts\\,s$^{-1}$.\\par $^{3}$: PDS total exposure times relative to two units (see text).\\par $^{4}$: PDS net count rates in 13-90\\,keV range.\\par $^{5}$: 0.1-10\\,keV phase averaged absorbed flux in units of 10$^{-11}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ as derived from best fit parameters in Table\\,2.\\par  \\end{flushleft} \\end{table*}   \\noindent The X-ray emission in IPs is generally strongly pulsed at the spin $\\rm P_{rot} = 1/\\omega$ period of the white dwarf, indicating that material accretes onto the magnetic poles via a disc, as it loses memory of the orbital motion (disc-fed systems). This material flows onto the poles in an arc--shaped accretion curtain (Rosen et al. 1988). However, not all systems are found to pulse only  at the white dwarf rotation. The presence of additional frequencies in the X-ray power spectra, such as the orbital $\\Omega$ and the beat (or synodic) $\\omega -\\Omega$,  indicate a disc-overflow accretion whose relative proportion can vary from system to system and with time (Hellier 1991; Wynn \\& King 1992; Norton et al. 1992). Only one system, RX\\,J1712$-$24\\,(V2400\\,Oph) is  known to date to be pulsed at the beat period (Buckley et al. 1997), thus being an unique example of a discless IP.  \\noindent The X-ray emission of IPs is hard and optically thin ($\\rm kT_{Brems.} \\sim  5-30\\,keV$), originating from the post-shock region above the magnetic pole. These hard X-rays are highly absorbed by cold material (column densities up to $\\rm 10^{23}\\,cm^{-2}$) within the accretion flow (Ishida et al. 1994) and are expected to be reflected by the WD atmosphere (Beardmore et al. 1995; Done et al. 1995; Matt 1999; Ezuka \\& Ishida 1999; de Martino et al. 2001). Unlike   Polars, which also show a strong soft X-ray black-body component (20-50\\,eV), only three systems, RE\\,0751$+$144\\,(PQ\\,Gem), RX\\,J0558.0$+$5353\\,(V405\\,Aur) and the faint RX\\,J0512.2$-$3241\\,(UU\\,Col), were discovered as soft X-ray IPs (Mason et al. 1992; Haberl et al.  1994; Burwitz \\& Reinsch 2001). The first two are also known to show circular polarized optical emission, suggesting a close link of soft X-ray IPs to the Polars. Indeed the soft X-ray emission in Polars is due to reprocessing, in the white dwarf atmosphere, of hard X-rays and cyclotron radiation, which are emitted from the post-shock region. In particular, the relative proportion of hard-to-soft X-ray emissions was found to depend on the magnetic field strength and local mass accretion rate (Voelk \\& Beuermann 1998). Hence, the soft IPs might be considered as true non-synchronous Polars and probably as their progenitors. Why most IPs do not show such a soft X-ray component is still an open question.   In this paper we present BeppoSAX observations of RE\\,0751$+$144\\,(PQ\\,Gem) (henceforth RE\\,0751), RX\\,J0558.0$+$5353\\,(V405\\,Aur) (henceforth RX\\,J0558) and RX\\,J1712.6$-$2414\\,(V2400\\,Oph) (henceforth RX\\,J1712), which together with RX\\,J0028.8$+$5917\\,(V709\\,Cas) (de Martino et al. 2001) are so far the only IPs for which a simultaneous broad-band X-ray coverage has been performed. All three IPs are characterized by  circular polarization in their optical emission. Also, while the first two were discovered as soft X-ray IPs by ROSAT, RX\\,J1712 was only recently found by XMM-Newton (Haberl 2002) to possess a soft, but hot (80-100\\,eV) X-ray component. Here we discuss their temporal and spectral properties over the wide energy range of the BeppoSAX instruments. ",
        "conclusions": "We have presented the first simultaneous soft and hard X-ray data of the three soft X-ray IPs, RE\\,J0751, RX\\,J0558 and RX\\,J1712, over the wide energy range of the BeppoSAX satellite.  All three sources have strong X-ray modulation at pe\\-riods identified as the white dwarf rotations in RX\\,J0558 and RE\\,0751, and as the synodic period in RX\\,J1712. We have inferred the temperatures of both post--shock flow and of the irradiated polar region of the white dwarf. The temperatures of the black--body component are very high, up to 80-100\\,eV in RX\\,J0558 and RX\\,J1712, much higher than those found in the Polars. The temperature of the hard X-ray emitting region is found between 13-30\\,keV. Furthermore, a Compton reflection component is definitely present in RX\\,J0558 and RE\\,0751, while in RX\\,J1712 it is favoured. The reflection  likely originates at the white dwarf surface.  \\noindent In the three systems, we find differences in the periodic behaviour of the spectral components. We find that the absorbing material partially covering the X-ray source is phase--dependent in RE\\,0751  and RX\\,J1712 but not in RX\\,J0558. The column densities can be as high as 10$^{23}\\,\\rm cm^{-2}$, as  in RX\\,J1712, which can heavily hamper the detection of soft X-rays in other hard X-ray IPs.  \\noindent Although these soft X-ray IPs share common properties, their X-ray temporal and spectral characteristics  suggest accretion patterns that cannot be reconciled with a single and simple configuration."
    },
    "0311/hep-ph0311312_arXiv.txt": {
        "abstract": "It has been suggested that a scalar field with negative kinetic energy, or ``ghost,'' could be the source of the observed late-time cosmological acceleration. Naively, such theories should be ruled out by the catastrophic quantum instability of the vacuum.  We derive phenomenological bounds on the Lorentz-violating ultraviolet cutoff $\\Lambda$ which must apply to low-energy effective theories of ghosts, in order to keep the instability at unobservable levels.  Assuming only that ghosts interact at least gravitationally, we show that $\\Lambda \\lsim 3$ MeV for consistency with the cosmic gamma ray background.  We also show that theories of ghosts with a Lorentz-conserving cutoff are completely excluded. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311007_arXiv.txt": {
        "abstract": "{In this paper we present new near infrared (NIR) imaging and spectroscopic data of the quasar \\object{3C 48} and its host galaxy. The data were obtained with the ESO-VLT camera ISAAC. We report the first detection of the apparent second nucleus 3C 48A about 1\\arcsec NE of the bright QSO nucleus in the NIR bands $J$, $H$, and $Ks$. 3C 48A is highly reddened with respect to the host, which could be due to warm dust, heated by enhanced star formation or by interstellar material intercepting the radio jet. In fact, all colors on the host galaxy are reddened by several magnitudes of visual extinction. Imaging and initial spectroscopy also reveal a stellar content of about 30\\% to the overall QSO-light in the NIR. These results are important input parameters for future models of the stellar populations by taking extinction into account. ",
        "introduction": "Revealing the physical properties of quasi stellar objects (QSOs) and their host galaxies is one key to understand how the overall galaxy structure affects the extreme activity in active galactic nuclei (AGN) and how in turn the QSO and its black hole influence the surrounding galaxy structure.  The host galaxy of the radio-loud quasar 3C 48 is an unusually large and bright object with excess far-infrared emission, $L_\\mathrm{FIR}=5\\times 10^{12}L_\\odot$ \\citep{1985ApJ...295L..27N}, showing a one-sided radio jet extending about 0.5\\arcsec to the north, fanning out to the east to about 1$''$ \\citep{1991Natur.352..313W}. At its redshift of $z\\approx 0.3695$ \\citep{1984ApJ...281..535B} \\object{3C 48} is one rare member of the class of compact steep spectrum (CSS) radio sources with a radio-spectral index $\\alpha\\sim 0.9$ \\citep{2001A&A...373..381T}. The quasar is rich in molecular gas, $M_{\\mathrm{H}_2}\\approx 2.7\\times 10^{10} M_\\odot$, as CO(1-0) observations show \\citep{1993ApJ...415L..75S,1997A&A...322..427W}.  There is also morphological evidence for a recent merger event, like a possible double nucleus and a tidal tail extending several arcseconds to the northwest \\citep[~and references therein]{1991AJ....102..488S, 1999A&A...343..407C, 2000ApJ...528..201C}. But it is still unclear, whether the brightness peak $\\sim 1''$ northeast (NE) of the QSO is really the nucleus of the companion galaxy, as it could also be a region of interaction of the radio jet with the interstellar medium of the host galaxy, inducing star formation \\citep{2000ApJ...528..201C} or even both. \\cite{2000ApJ...528..201C} found unusually high stellar velocity dispersions as well as very young stellar populations, possibly due to the interaction of the jet with the ambient medium \\citep[see also][]{1999A&A...343..407C}. The merger morphology, its high FIR luminosity, and its large content of molecular gas put 3C 48 in an evolutionary scheme envisioned by \\citet{1988ApJ...325...74S}. It fits nicely as an transition object between ultra-luminous infrared galaxies (ULIRGs) and pure QSOs, being still dominated by its FIR excess but also displaying QSO features like its luminosity and rather flat spectral energy distribution.  The 3C 48 nucleus is an ideal target for high-resolution near infrared studies of AGN and their hosts, in order to study the influence of the nuclear activity as well as the ongoing merger on the host environment. 3C 48 also offers access to a unique set of diagnostic lines (Pa$\\beta$, Pa$\\gamma$, \\ion{Si}{I}, CO(6-3) and see Table \\ref{3C48DiagnosticLines}) to study stellar populations, effects of extinction, varying metallicity and initial mass function.  A major advantage of the NIR is that contamination due to scattered light is much smaller than in the visible. Previous work did not take extinction due to molecular material into account, which must play a role in the dusty and gas rich host galaxy. Therefore it is likely that the patchy velocity and age distribution found by \\cite{2000ApJ...528..201C} may be in part due to variations in extinction.  Section 2 of this article describes details of our new NIR observations, the data reduction procedures, and calibration in both imaging and spectroscopy. In Section 3 we present and discuss the results of photometry of the nuclear region, the 3C 48A potential second nucleus and the tidal tail, and we present the initial spectroscopic results. Section 4 gives a final discussion and summary. ",
        "conclusions": "\\subsection{3C 48 Surface Photometry} \\label{surfacePhotometry} To place the QSO and the host galaxy into a two-color diagram, photometry was carried out at several positions on the galaxy in the three bands $J$, $H$, and $Ks$. \\paragraph{Overall Colors:} Overall magnitudes of 3C 48 were determined using an aperture of a diameter of 10\\arcsec (= 68 pixels) centered on the centroid of the brightness distribution. In this case one arcsecond corresponds to a linear scale of 4~kpc. These magnitudes are dominated by the unresolved active nucleus. Additionally two smaller apertures with diameters of 5\\arcsec (= 34 pixels) and 1\\arcsec (= 7 pixels) respectively have been applied, centered on the same position. A comparison allowed to estimate the stellar contribution to the AGN luminosity (see below).  The IRAF task \\emph{qphot} (centering and integrating the flux) has been used to determine the flux in ADUs in these apertures in the individual bands. Table \\ref{3c48OverallMag} displays the resulting magnitudes using the above determined zero-points. The error is dominated by sky-noise, determined from the standard deviation within 5x5 pixel boxes at several positions on the sky. This yields a precision in magnitudes of about $\\Delta \\mathrm{m} \\approx \\pm 0.1$ in the individual bands.  From these magnitudes the $H-K$ and $J-H$ colors for the three apertures have been calculated (Table \\ref{3c48OverallColor}).  \\paragraph{Colors of the Host Galaxy:} In order to obtain information on the nature of the host galaxy, it is necessary to carry out photometry at several positions on the host. Using for example models for color evolution \\citep{1993ApJ...405..538B, 1997AJ....113..550H}, one can estimate the age of the stellar content residing in the disk of the host galaxy. With surface photometry one can additionally investigate the importance of reddening due to (warm) dust. We therefore measured the flux at several positions across the galaxy, which are indicated in Figure \\ref{3c48ManyAp}. Altogether we choose 12 apertures of 1.2\\arcsec (8 pixels) diameter at representative positions. Three apertures are lying in the south and two in the tidal tail region to the northwest. Five apertures lie around the active nucleus and cover the location of the potential second nucleus $\\sim 1''$ to the northeast. Aperture number 12 lies at the position within an apparent counter tidal \\citep{2000ApJ...528..201C} reaching out to about 7\\arcsec to the southeast (see Fig. \\ref{3c48ManyAp}). In our NIR images we do not see a tidal tail feature reaching out to about 8\\farcs 5 as proposed in \\cite{1999MNRAS.302L..39B}.  \\paragraph{$J-H/H-K$ two Color Diagram of 3C 48:} The two-color diagram in Fig. \\ref{3c48TwoColors} shows (1) the overall colors, (2) the colors derived from the 12 distinct positions on interesting locations in the 3C 48 host (as discussed above) and (3) a model for color-evolution at a redshift of $z=0.35$ (of about the distance to 3C 48) \\citep{1997AJ....113..550H} and (4) the stellar contribution to the light measured within an aperture centered on the nucleus (see also Fig. \\ref{3c48CoreSubtracted}).  (1) \\emph{Overall colors}. Considering only the overall colors (green dots), there is a trend of redder colors in $H-K$ as well as in $J-H$ for smaller apertures (10\\arcsec down to 1\\arcsec). This is in agreement with the idea, that in a large aperture there is a contribution of both the host galaxy (filled circles) and the active nucleus. In a smaller aperture, in contrast there only remains the nuclear contribution which is usually redder due to gas and dust the active core is embedded in. The dotted line in the two color diagram (Fig. \\ref{3c48TwoColors}) presents the influence of mixing the nuclear color, determined in the 1\\arcsec aperture, with a stellar contribution, taking typical stellar colors at the 10~Gyr point\\footnote{We have chosen this age, because the dominant, underlying stellar population is of about 10~Gyr \\citep{2000ApJ...528..201C} and is also typical for 'inactive' stellar dominated galaxies.} of the evolutionary curve and holding the flux in the Ks-band constant, so that the total flux is conserved for the start and end points. Starting with no stellar contribution at the point of the 1\\arcsec (4~kpc) aperture and arriving at 100\\% stellar contribution at the 10~Gyr color point, the data points of the two larger apertures lie on this curve within the errors and correspond to a stellar contribution of about 15\\% for the 5\\arcsec (20~kpc) aperture and about 30\\% for the 10\\arcsec (40~kpc) aperture (see Fig. \\ref{3c48CoreSubtracted}). The stellar contribution originates in the disk of the host galaxy. The 10\\arcsec aperture covers essentially the whole galaxy. Therefore the stellar contribution of the underlying host galaxy to the light from the active nucleus is of the order of 30\\% at 2.2~$\\mu$m. A consistent check for the stellar content can be the CO(6-3) absorption, a tracer for the stellar content \\citep{1993A&A...280..536O}. One has to note that for the 1\\arcsec aperture for the nuclear region a stellar contribution still has to be taken into account.  \\begin{table}[ht] \\caption{Photometry of 3C 48 with three different apertures centered on the QSO} \\label{3c48OverallMag} \\begin{center}  \\begin{tabular}{@{}p{0.8cm}@{}p{1.1cm}@{}p{1.2cm}@{}p{0.8cm}@{}p{1.5cm}@{}p{1.5cm}@{}p{1.4cm}@{}}  \\hline Band & DIT~[s] &  Airmass & k    & ZP$_\\mathrm{c}$~[mag] & Diam.~[\\arcsec]  & Mag\\\\ \\hline $J$    & 20       & 1.90     & 0.11 & 25.13        & 10          & 14.4$\\pm 0.1$\\\\ &          &          &      &              & 5           & 14.4\\\\ &          &          &      &              & 1           & 15.1\\\\ \\hline $H$    & 15       & 2.01     & 0.06 & 24.77        & 10          & 13.5$\\pm 0.1$\\\\ &          &          &      &              & 5           & 13.6\\\\ &          &          &      &              & 1           & 14.3\\\\ \\hline $Ks$   & 10       & 2.01     & 0.07 & 24.26        & 10          & 12.8$\\pm 0.1$\\\\ &          &          &      &              & 5           & 12.8\\\\ &          &          &      &              & 1           & 13.4\\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\begin{table} \\caption{Colors for individual apertures centered on the QSO} \\label{3c48OverallColor} \\begin{center} \\begin{tabular}{c c c} \\hline Diameter~[\\arcsec] & $H-K$ & $J-H$\\\\ \\hline 10 & 0.7$\\pm 0.1$ & 0.9$\\pm 0.1$\\\\ 5 & 0.8$\\pm 0.1$ & 0.9$\\pm 0.1$\\\\ 1& 0.8$\\pm 0.1$ & 0.9$\\pm 0.1$\\\\ \\hline \\end{tabular} \\end{center} \\end{table}  (2) \\emph{Colors from single apertures}. Colors derived from measurements of aperture 12 are unusually blue compared to the colors derived from other apertures on the host galaxy. Its colors are more typical for normal galaxies at redshifts close to zero \\citep{1978ApJ...220...75F}. This suggests that this feature ($\\sim 6.5''$ to the southeast of the quasar nucleus), while being aligned with the apparent counter-tidal tail to the southeast, in fact is a foreground galaxy. This contravenes with the redshift derived in \\cite{2000ApJ...528..201C} ($z\\approx 0.8112$). As both results assume that this feature is not associated with 3C 48, further investigation is necessary. Furthermore recent multi-particle simulations \\citep{scharw} indicate, that starting from an Antennae-like setup \\citep[e.g.~][]{1972ApJ...178..623T} one can arrive at a 3C 48 look-alike due to different projection parameters. In the case of \\cite{scharw}, the counter-tidal tail is extending from the SW towards a NE direction in front of the galaxy host. Based on their HST images \\cite{2000ApJ...528..201C} already speculated on the presence of a tidal tail in a SE-SW direction. This extension of the host towards the SW is also visible in our J- and H-band maps (Fig. \\ref{3c48jhk}). But it could also belong to the tidal tail resulting from the multi particle simulations.  \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics{H4724F2.eps}} \\caption{Positions of apertures with radii of 1.6\\arcsec on 3C 48. The apertures are labeled with numbers. The size of the map is 26$\\times$26 arcsec$^2$ corresponding to 104$\\times$104~pc$^2$.} \\label{3c48ManyAp} \\end{center} \\end{figure}  (2) and (3) \\emph{Color-evolution model}. The two crosses on the evolutionary curve in Fig. \\ref{3c48TwoColors} show the loci of stellar populations at an age of 2 and 10~Gyr, respectively. Noticing that the color of the upper right end of the model curve corresponds to an age of 20~Gyr for the stellar population, a striking result is that all derived colors are considerably reddened in comparison to the model shown. Indicated in the figure is also the effect of reddening for a visual extinction of $A_V=1$ (as discussed in \\citealt{1999hia..book.....G}). This indicates, that extinction might play an important role in deriving the physical properties of the host galaxy. As already mentioned, \\cite{2000ApJ...528..201C} did not take extinction into account in their calculations. They state that the underlying old population has an age of about 10~Gyr. Taking the median color of all our measurements on the host \\begin{eqnarray*} \\mathrm{Median}(H-K)&=& 0.6\\\\ \\mathrm{Median}(J-H)&=& 0.8 \\end{eqnarray*} we find that within the errors and including a visual extinction of $A_V=1$ \\citep{1985ApJ...288..618R} this result is consistent with an age of 10 Gyr. But we also sample regions of active star formation (as indicated by \\citealt{2000ApJ...528..201C}), i.e. regions with much younger stellar components. The authors find e.g. a young starburst population within the tidal tail at the position of our aperture 11. The $J-H$ color is about 0.25~mag bluer than the median color of the host galaxy. Its position is shifted in $H-K$ by about 0.2~mag to the red with respect to the evolutionary track shown in Fig. \\ref{3c48TwoColors}. This indicates a younger stellar population which may be associated with warm dust that contributes to the red $H-K$ color in emission. \\cite{2000ApJ...528..201C} quote an age of the corresponding stellar population of about 10~Myrs mixed with an underlying old stellar population of about 10~Gyr. Starburst populations with those ages lie at the lower (bluer) end of the color-evolution curve in the two color diagram. The presence of warm dust and the location of the aperture in the two color diagram, however, may indicate that a considerable reddening is at work. In this case the optical radiation would only be indicative of the properties of a fraction of the underlying stellar population.  \\begin{figure*} \\sidecaption \\includegraphics[width=12cm]{H4724F3.eps} \\caption{ $J-H$ vs. $H-K$ two color diagram showing the overall colors (filled circles with aperture diameter given in arcsec) and colors derived from apertures on distinct positions on the galaxy (crosses with aperture number as label (see Fig. \\ref{3c48ManyAp})). The median color of all apertures without the nuclear one is shown as a triangle. The solid line is an evolutionary model for colors \\citep{1997AJ....113..550H}. Evolution proceeds from the lower left corner at 0~Gyr to the upper right corner at 20~Gyr. The initial population is a 1~Gy starburst followed by passive evolution. The dotted color line represents the stellar contribution to the QSO light when increasing the aperture (cf Fig. \\ref{3c48CoreSubtracted}). The arrow indicates reddening in colors for a visual extinction of $A_V=1$ \\citep{1985ApJ...288..618R}.} \\label{3c48TwoColors} \\end{figure*}  (4) \\emph{The circum-nuclear region}. The flux from central regions of the host galaxy is still dominated by emission from the active nucleus. In order to get a more detailed view on the underlying host we subtracted the contribution of the active nucleus. Here we used the fact that the QSO-nucleus is unresolved, i.e. star like in extent. We have done the subtraction by taking a star within a distance of 1\\arcmin to 3C 48 shifting its centroid with subpixel-accuracy onto the position of the 3C 48-centroid and subtracting it scaled to the flux of the active nucleus. One has to take care that no 'holes' appear in the galaxy after this process. Ideally this should give a smooth brightness distribution in the central region after subtraction and primarily the host contribution should remain.  \\begin{figure*}% \\sidecaption \\includegraphics[width=12cm]{H4724F4.eps} \\caption{ Two color diagram showing the NIR properties of the brightness-excess region 3C 48A after subtraction of the active nucleus contribution (stars) together with the median color of the host (triangle) and the overall colors (dots) already shown in figure \\ref{3c48TwoColors}. Also shown is the stellar contribution to the overall color, when increasing the aperture (dotted line).} \\label{3c48CoreSubtracted} \\end{figure*}  \\begin{table}[ht] \\caption{Nuclear Flux Contribution} \\label{3c48NuclearContribution} \\begin{center} \\begin{tabular}{c c} \\hline Band & Mag  \\\\ \\hline $J$    & $15.1\\pm 0.1$\\\\ $H$    & $14.2\\pm 0.1$\\\\ $Ks$    &$13.5\\pm 0.1$\\\\ \\hline \\end{tabular} \\end{center} \\end{table}  Figure \\ref{3c48CoreSubtractedImages} displays the core-subtracted images in the $J$-, $H$- and $K$-bands and Table \\ref{3c48NuclearContribution} lists the nuclear contributions which have been subtracted. After the subtraction there remains no really smooth brightness distribution but certain residuals. This is most probable due to psf-mismatching. It would be favorable to take a composite psf by measuring several stars, giving a more realistic sampling of the point-spread function, impossible in this case due to the restricted FOV of the science exposures. But a brightness excess region about 1\\arcsec northeast of the QSO is clearly visible in the residual images as indicated in Figure \\ref{3c48CoreSubtractedImages}. \\begin{table}[ht] \\caption{Parameters of Gaussian Profiles for 3C 48A. $\\Delta$mag$_\\mathrm{Peak}$ is the difference in magnitudes between the host contribution and host plus 3C 48A contribution.} \\label{3c48AGaussian} \\begin{center} \\begin{tabular}{c c c} \\hline Band & FWHM~[\\arcsec] & $\\Delta$mag$_\\mathrm{Peak}$\\\\ \\hline J    & 2.1           & 0.42\\\\ H    & 2.1           & 0.49\\\\ Ks   & 1.2           & 2.84\\\\ \\hline \\end{tabular} \\end{center} \\end{table} \\begin{figure*}[ht] \\resizebox{\\hsize}{!}{\\includegraphics{H4724F5a.eps} \\includegraphics{H4724F5b.eps} \\includegraphics{H4724F5c.eps}} \\caption{$JHK$ images with subtracted nucleus (see text for details). (a) $J$, (b) $H$ and (c) $Ks$. The circles of $1''$ radius centered on the active nucleus indicate the location of 3C 48A northeast of the QSO. In all three bands an additional flux component at the position of 3C 48A becomes visible after the subtraction of the QSO nucleus.} \\label{3c48CoreSubtractedImages} \\end{figure*} Its peak position in all three bands is $\\sim 1''$ north-east of the QSO-location, having a position angle of about 40$^\\circ$ (as measured from north to east). This feature is most prominent in the Ks image whereas it gets fainter in the $H$ and $J$ images.  In order to estimate the amount of magnitudes with which 3C 48A exceeds the underlying host galaxy emission, we fitted a Gaussian to this emission region. After subtraction of a Gaussian with this scaling we recognized a circular symmetric residual emission from the central region as expected. On the opposite side of 3C 48 at the same distance as 3C 48A we determined the flux value of the host galaxy. The difference between these two values then gives an estimate of the 3C 48A contribution presented in Table \\ref{3c48AGaussian}.  $J-H$ and $H-K$ colors were determined by measuring the flux in $J$, $H$ and $Ks$ in apertures of the size of the FWHM of the Gaussian profile. Figure \\ref{3c48CoreSubtracted} displays the data in a two color diagram. The determined colors of 3C 48A have to be taken with care, because of the psf mismatch mentioned above. Additionally, the $Ks$ band image is quite noisy which makes the measurement process more complicated. Together this leads to a high degree of freedom in choosing an appropriate Gaussian fit of this region. The apertures have been chosen to be of the sizes listed in Table \\ref{3c48AGaussian}, i.e. rather small, in order to reduce the impact of the imperfect QSO-subtraction process, which did not result in a smooth brightness distribution and produced some 'holes'.  Figure \\ref{3c48CoreSubtracted} shows that the 3C 48A area is heavily reddened in the $H-K$ direction. This indicates the presence of hot dust which usually shifts colors exclusively in the $H-K$ direction as it influences predominantly the $K$ band \\citep{1985MNRAS.214..429G}. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{H4724F6.eps}} \\caption{Overlay of the central region of 3C 48 in the $H$-band (centered on the QSO) with a 1.66~GHz radio continuum map from \\citet{1991Natur.352..313W}, showing the expansion of the radio jet towards the location of 3C 48A. } \\label{3c48Merlin} \\end{figure}  An overlay of the NIR images with a radio continuum map of Merlin and European VLBI data at 1.66~GHz \\citep{1991Natur.352..313W} shows an apparent alignment of the radio jet and the position of 3C 48A (Fig. \\ref{3c48Merlin}). Such an alignment was also found by \\citet{1991AJ....102..488S}. Together with the results of \\citet{scharw} there is still no definite answer on the origin of the excess brightness of 3C 48A. Probably it is a combination of both, a second nucleus and a jet-ISM interaction.  \\subsection{The $H$-band Spectrum} The most prominent features in the $H$ spectrum (Fig.~\\ref{3c48SHSpec}) are the Pa$\\beta\\ \\lambda$17561 and Pa$\\gamma\\ \\lambda$14985 lines (cf. Table \\ref{3C48DiagnosticLines}). \\begin{table} \\centering \\caption{NIR diagnostic Lines in 3C 48} \\begin{tabular}{c c c} \\hline Line   & $\\lambda_\\mathrm{rest}$~[$\\mu$m] & $\\lambda_\\mathrm{obs}$~[$\\mu$m]\\\\ \\hline \\multicolumn{3}{c}{$H$-band}\\\\ Pa$\\beta$  & 1.2818 & 1.7561\\\\ Pa$\\gamma$ & 1.0938 & 1.4985\\\\ \\multicolumn{3}{c}{$K$-band}\\\\ CO(6-3)    & 1.6185 & 2.2173\\\\ \\ion{Si}{I}        & 1.5892 & 2.1772\\\\ H$_2$(1-0) S(7) & 1.7475 & 2.3941\\\\ \\ion{Fe}{II}       & 1.5335 & 2.1009\\\\ \\ion{Fe}{II}       & 1.6435 & 2.2516\\\\ \\ion{Fe}{II}       & 1.7449 & 2.3905\\\\ \\hline \\end{tabular} \\label{3C48DiagnosticLines} \\end{table}  The baseline is subtracted using a straight line. A Gaussian fit to the Pa$\\beta$ and Pa$\\gamma$ lines gives line fluxes of \\begin{eqnarray} f_{Pa\\beta} &=&1.0\\times 10^{-24}\\ (\\pm 20\\%)\\ \\mathrm{W~cm}^{-2}\\\\ f_{Pa\\gamma}&=&5.6\\times 10^{-25}\\ (\\pm 20\\%)\\ \\mathrm{W~cm}^{-2} \\end{eqnarray}  The line ratio Pa$\\beta$$/$Pa$\\gamma$ can then be calculated: \\begin{equation} \\frac{\\mathrm{Pa}\\beta}{\\mathrm{Pa}\\gamma}=1.8(\\pm 30\\%) \\end{equation} When using line ratios for {\\it Menzel's case B} (see e.g. \\citeauthor{1999hia..book.....G} \\citeyear{1999hia..book.....G} or~ \\citeauthor{osterbrock89:_astrop_gaseous_nebul_activ_galac_nuclei} \\citeyear{osterbrock89:_astrop_gaseous_nebul_activ_galac_nuclei}), this result is consistent within its errors with temperatures of $T=10^4$~K and electron densities of $\\rho=10^4$~cm$^{-3}$, which are typically found for the atomic interstellar medium in QSO nuclei.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{H4724F7.eps}} \\caption{H band spectrum of the nuclear region of 3C 48. The dotted line is the original spectrum, whereas the solid line shows the spectrum after a 10 point smoothing. Indicated are the redshifted positions of the Pa$\\beta$ and Pa$\\gamma$ lines. } \\label{3c48SHSpec} \\end{figure}  \\subsection{The $K$-band Spectrum} The $K$ spectrum (Fig.~\\ref{3c48SKSpec}) is less noisy but there is hardly any feature to be recognized at its original spectral resolution of $\\Delta\\lambda=0.005\\ \\mu$m. Therefore we smoothed the spectrum with a boxcar of a width of 4 pixels. There are two features appearing in the spectrum, one left of the indicated \\ion{Si}{I} line and one at a wavelength of about 2.3~$\\mu$m. The latter one could not be identified. At this location there seems to be no prominent emission line. We suggest this to be a noise peak. The absorption feature shortward of the \\ion{Si}{I} line is most probably an atmospheric or stellar residual from the telluric-line removal and can therefore not be attributed to 3C 48. The noise per pixel of the smoothed spectrum is about half of the noise in the original spectrum \\begin{equation} \\sigma_\\mathrm{pixel}=3\\times 10^{-24}\\ \\mathrm{W~cm}^{-2}\\ \\mu\\mathrm{m}^{-1}. \\end{equation} \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{H4724F8.eps}} \\caption{$K$ band spectrum of the nuclear region of 3C 48. The dotted line is the original spectrum, whereas the solid line shows the spectrum after a 4 point smoothing. Indicated are the redshifted positions of the \\ion{Si}{I} and CO(6-3) absorption lines. } \\label{3c48SKSpec} \\end{figure}  We investigated the noise features at the CO(6-3) line at a redshifted wavelength of 2.217~$\\mu$m (cf. Table \\ref{3C48DiagnosticLines}). The CO(6-3)~$\\lambda 1.6185$ line is a tracer of the stellar content of a galaxy, as this line arises in the atmospheres of cool stars \\citep{1993A&A...280..536O}. It is weak in early K- but strong in late M-type stars. We determined the line flux by using the triangular area above the absorption line. The flux and noise integrated over a wavelength interval of $\\Delta\\lambda=10$~nm is \\begin{eqnarray} f_\\mathrm{CO(6-3)} &\\le& 2.43\\times 10^{-26}\\ \\mathrm{W~cm}^{-2} \\\\ \\sigma_\\mathrm{integrated} &=& 2.46\\times 10^{-26}\\ \\mathrm{W~cm}^{-2}\\nonumber\\\\ \\end{eqnarray} where the noise was determined by measuring the fluctuations throughout the spectrum except at the CO(6-3) wavelength. We thus have a $1\\sigma$ detection of the line. The depth of the line is about 3\\% of the continuum level at that wavelength. According to \\cite{1993A&A...280..536O} the typical depth of this line is about 15\\% of the continuum for a pure stellar contribution. Taking into account our photometrically determined stellar contribution of about 30\\% of the galaxy light, this leads to an expected CO(6-3) line depth of $\\sim 4\\%$. Accordingly our spectroscopic result seems to be consistent with the photometric results within the errors."
    },
    "0311/astro-ph0311231_arXiv.txt": {
        "abstract": "We investigate the mass profile of $\\Lambda$CDM halos using a suite of numerical simulations spanning five decades in halo mass, from dwarf galaxies to rich galaxy clusters. These halos typically have a few million particles within the virial radius ($r_{200}$), allowing robust mass profile estimates down to radii below $1\\%$ of $r_{200}$. Our analysis confirms the proposal of Navarro, Frenk \\& White (NFW) that the shape of $\\Lambda$CDM halo mass profiles differs strongly from a power law and depends little on mass.  The logarithmic slope of the spherically-averaged density profile, as measured by $\\beta=-d\\ln \\rho/d\\ln r$, decreases monotonically towards the center and becomes shallower than isothermal ($\\beta < 2$) inside a characteristic radius, $r_{-2}$. The fitting formula proposed by NFW provides a reasonably good approximation to the density and circular velocity profiles of individual halos; circular velocities typically deviate from best NFW fits by less than $10\\%$ over the radial range which is well resolved numerically.  On the other hand, systematic deviations from the best NFW fits are also noticeable. Inside $r_{-2}$, the profile of simulated halos becomes shallower with radius more gradually than predicted and, as a result, NFW fits tend to underestimate the dark matter density in these regions. This discrepancy has been interpreted as indicating a steeply divergent cusp with asymptotic inner slope, $\\beta_{0}\\equiv\\beta(r=0)\\sim 1.5$. Our results suggest a different interpretation. We use the density and enclosed mass at our innermost resolved radii to place strong constraints on $\\beta_{0}$: density cusps as steep as $r^{-1.5}$ are inconsistent with most of our simulations, although $\\beta_{0}=1$ is still consistent with our data. Our density profiles show no sign of converging to a well-defined asymptotic inner power law. We propose a simple formula that reproduces the radial dependence of the slope better than the NFW profile, and so may minimize errors when extrapolating our results inward to radii not yet reliably probed by numerical simulations. ",
        "introduction": "\\label{sec:intro}  Disk galaxy rotation curves; strong gravitational lensing by galaxies and clusters; the dynamics of stars in elliptical galaxies and of gas and galaxies in clusters; these are just examples of the various luminous tracers that probe the inner structure of dark matter halos. Such observations place strong constraints on the distribution of dark matter in these highly non-linear regions that may be contrasted directly with theoretical predictions for halo structure. Given the sensitivity of such predictions to the nature of the dark matter, these observational constraints constitute provide tests that may question or even rule out particular models of dark matter.  Robust prediction of the inner structure of cold dark matter (CDM) halos is a quintessential N-body problem, albeit one of considerable complexity due to the large overdensities and, consequently, the short crossing times involved. Indeed, only recently have computational capabilities improved to the point of allowing realistic simulation of the regions which house the luminous components of individual galaxies.  This work builds upon the pioneering efforts of Frenk et al (1988), Dubinski and Carlberg (1991), and Crone et al (1994), among others, which led to the identification of a number of key features of the structure of dark matter halos assembled by hierarchical clustering. One important result of this early work concerns the absence of a well defined central ``core'' of constant density in virialized CDM halos.  In this sense, dark matter halos are ``cuspy'': the dark matter density increases apparently without bounds toward the center of the halo.  A second result concerns the remarkable similarity (``universality'') in the structure of dark matter halos of widely different mass. This was first proposed by Navarro, Frenk \\& White (1996, 1997; hereafter NFW), who suggested a simple fitting formula to describe the spherically-averaged density profiles of dark matter halos, \\begin{equation} \\rho(r)={\\rho_s \\over (r/r_s) (1+(r/r_s))^{2}}, \\label{eq:nfw} \\end{equation} where $\\rho_s$ and $r_s$ are a characteristic density and radius, respectively. The larger the mass of a halo, the lower its characteristic density, reflecting the lower density of the universe at the (later) assembly time of more massive systems.  Further simulation work of similar numerical resolution (see, e.g., Cole \\& Lacey 1996, Huss et al 1999) provided support for the NFW conclusions, although small but systematic differences began to emerge as the numerical resolution of the simulations improved (see, e.g., Moore et al 1999, hereafter M99, Ghigna et al 2000, Fukushige \\& Makino 1997, 2001, 2003). These authors reported deviations from eq.~\\ref{eq:nfw} that increase systematically inward, and thus are particularly noticeable in high-resolution simulations.  In particular, Fukushige \\& Makino (2001) reported that NFW fits tend to underestimate the dark matter density within the scale radius $r_s$. M99 reached a similar conclusion and interpreted this result as indicating a density cusp steeper than that of the NFW profile. These authors preferred a modified fitting function which diverges as $r^{-1.5}$ near the center, \\begin{equation} \\rho(r)={\\rho_M \\over (r/r_M)^{1.5} (1+(r/r_M)^{1.5})}. \\label{eq:mooreetal} \\end{equation} One should note, however, that there is no consensus amongst N-body practitioners for such modified profile (see, e.g., Klypin et al 2001 and Power et al 2003, hereafter P03), and that there is some work in the literature suggesting that the central density cusp might actually be {\\it shallower} than $r^{-1}$ (Subramanian et al 2000; Taylor \\& Navarro 2001; Ricotti 2003).  This unsettled state of affairs illustrates the difficulties associated with simulating the innermost structure of CDM halos in a reliable and reproducible manner. The high density of dark matter in such regions demands large numbers of particles and fine time resolution, pushing to the limit even the largest supercomputers available at present. As a result, many of the studies mentioned above are either of inadequate resolution to be conclusive or are based on results from a handful of simulations where computational cost precludes a systematic assessment of numerical convergence.  Establishing the detailed properties of the central density cusp, as well as deriving the value of its asymptotic central slope, are important for a number of reasons. For example, steeper cusps place larger amounts of dark matter at the center, exacerbating the disagreement with observations that suggest the presence of a constant density core in low surface brightness galaxies or in strongly barred galaxies (Moore 1994, Flores \\& Primack 1994, McGaugh \\& de Blok 1998, Debattista and Sellwood 1998, van den Bosch et al 2000). Steep cusps would also be important for direct detection experiments for dark matter, as a possible gamma-ray annihilation signal of WIMPS at the Galactic center would be particularly strong for $r^{-1.5}$ cusps (Calc\\`aneo-Rold\\'an \\& Moore 2000; Taylor \\& Silk 2003; Stoehr et al 2002).  Finally, the detailed structure of the central cusp is not the only focus of contrasting claims in the literature. For example, the ``universality'' of CDM halo structure has been questioned by Jing \\& Suto (2000), who find that the slope of the density profile at a fixed fraction of the virial radius steepens towards lower halo masses. Klypin et al (2001), on the other hand, point out that such a systematic trend is entirely consistent with universality as originally claimed by NFW, and just reflects the mass dependence of halo characteristic density.  We address these conflicting issues here using a suite of nineteen high-resolution simulations of the formation of halos in the standard $\\Lambda$CDM cosmogony. Halo masses are chosen in three main groups: ``dwarf'' halos with $M_{200} \\sim  10^{10} \\, h^{-1} \\, \\,M_{\\odot}$, ``galaxy'' halos with $M_{200}\\sim 10^{12} \\, h^{-1} \\,M_{\\odot}$ and ``cluster'' halos of mass $M_{200}\\sim 10^{15} \\, h^{-1} \\, M_{\\odot}$. This allows us to gain insight into the effects of cosmic variance at each mass scale, as well as to explore the mass dependence of the structure of $\\Lambda$CDM halos. We define the mass of a halo to be that contained within its virial radius, that is, within a sphere of mean density contrast $200${\\footnote{We use the term `density contrast' to denote densities expressed in units of the critical density for closure, $\\rho_{\\rm crit}=3H^2/8\\pi G$. We express the present value of Hubble's constant as $H(z=0)=H_0=100\\, h$ km s$^{-1}$ Mpc$^{-1}$}}.  This paper is organized as follows. Section~\\ref{sec:numexp} describes briefly the numerical simulations; \\S~\\ref{sec:res} discusses our main results; and we summarize our conclusions in \\S~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have analyzed the mass profile of $\\Lambda$CDM halos in a series of simulations of high mass, spatial, and temporal resolution. Our series targets halos spanning five decades in mass: ``dwarf'' galaxy halos with virial circular velocities of order $V_{200}\\sim 30$ km s$^{-1}$; ``galaxy''-sized halos with $V_{200} \\sim 200$ km s$^{-1}$; and ``cluster'' halos with $V_{200} \\sim 1200$ km s$^{-1}$.  Each of the nineteen halos in our series was simulated with comparable numerical resolution: they have between $8\\times 10^5$ and $4\\times 10^6$ million particles within the virial radius, and have been simulated following the ``optimal'' prescription for time-stepping and gravitational softening laid down in the numerical convergence study of P03.  The high resolution of our simulations allows us to probe the inner properties of the mass profiles of $\\Lambda$CDM halos, down to $\\sim 0.5\\%$ of $r_{200}$ in our best resolved runs. These results have important implications for the structure of the inner cusp in the density profile and resolve some of the disagreements arising from earlier simulation work. Our main conclusions may be summarized as follows.  \\begin{itemize}  \\item $\\Lambda$CDM halo density profiles are ``universal'' in shape: i.e., a simple fitting formula reproduces the structure of all simulated halos, regardless of mass. Both the NFW profile and the profile proposed by M99 describe the density and circular velocity profiles of simulated halos reasonably well. Best NFW fits to the circular velocity profiles deviate by less than $10\\%$ over the region which is well resolved numerically. Best M99 fits reproduce circular velocity profiles to better than $25\\%$ over the same region. It should be noted, however, that the deviations increase inwards and are typically maximal at the innermost resolved radius, a result that warns against extrapolating to smaller radii with these fitting formulae.  \\item $\\Lambda$CDM halos appear to be ``cuspy'': i.e., the dark matter density increases monotonically towards the center with no evidence for a well-defined ``core'' of constant density. We find no evidence, however, for a central asymptotic power-law bin the density profiles. These become progressively shallow inwards and are significantly shallower than isothermal at the innermost resolved radius, $r_{\\rm conv}$. At $r\\sim 0.01 \\, r_{200}$, the average slope of ``cluster'', ``galaxy'' and ``dwarf'' halos halos is $\\beta(r_{\\rm conv}) \\sim 1.1$, $\\sim 1.2$, and $\\sim 1.35$, respectively. This is steeper than predicted by the NFW profile but shallower than the asymptotic slope of the M99 profile.  \\item The density and enclosed mass at $r_{\\rm conv}$ may be used to derive an upper limit on any asymptotic value of the inner slope. Cusps as steep as $\\beta_0=1.5$ are confidently ruled out in essentially all cases; the asymptotic slope of the NFW profile ($\\beta_0=1$) is still consistent with our data.  The radial dependence of $\\beta(r)$ differs from that of the NFW profile, however, decreasing more slowly with decreasing radius than is predicted. For some scalings of the NFW fitting formula to the numerical data, this shape difference appears as a dark matter ``excess'' near the center which has (erroneously) been interpreted indicating a steeply divergent density cusp.  \\item A simple formula where $\\beta(r)$ is a power law of radius reproduces the gradual radial variation of the logarithmic slope and its apparent failure to converge to any specific asymptotic value (eq.~\\ref{eq:newprof}). This formula leads to much improved fits to the density profiles of simulated halos, and may prove a safer choice when comparison with observation demands extrapolation below the innermost converged radii of the simulations.  \\end{itemize}  Our study demonstrates that, although simple fitting formulae such as NFW are quite accurate in describing the global structure of $\\Lambda$CDM halos, one should be aware of the limitations of these formulae when interpreting observational constraints. Extrapolation beyond the radial range where these formulae have been validated is likely to produce substantial errors. Proper account of the substantial scatter in halo properties at a given halo mass also appears necessary when assessing the consistency of observations with a particular cosmological model. Direct comparison between observations and simulations (rather than with fitting formulae) is clearly preferable whenever possible. Given the computational challenge involved in providing consistent, robust, and reproducible theoretical predictions for the inner structure of CDM halos it is likely that observational constraints will exercise to the limit our hardware and software capabilities for some time to come.   This work has been supported by computing time generously provided by the High Performance Computing Facility at the University of Victoria, as well as by the Edinburgh Parallel Computing Centre and by the Institute for Computational Cosmology at the University of Durham. Expert assistance by Colin Leavett-Brown in Victoria and Lydia Heck in Durham is gratefully acknowledged. JFN is supported by the Alexander von Humboldt Foundation, the Natural Sciences and Engineering Research Council of Canada, and the Canadian Foundation for Innovation."
    },
    "0311/astro-ph0311141_arXiv.txt": {
        "abstract": "We study the soft X-ray emission (0.1 -- 1 keV) from the Warm-Hot Intergalactic Medium (WHIM) in a hydrodynamic simulation of  a Cold Dark Matter universe. Our main goal is to investigate how such emission can be explored with a combination of imaging and spectroscopy, and to motivate future X-ray missions.  We first present high resolution images of the X-ray  emission in several energy bands, in which emission from different ion species dominates.  We pick three different areas  to study the high resolution  spectra of X-rays from the warm-hot IGM: (a) a galaxy group; (b) a filament and (c) an underluminous region. By taking into account the background X-ray emission from AGNs and foreground emission from the Galaxy, we compute composite X-ray spectra of the selected regions. We  briefly investigate angular clustering of the soft-X-ray emission, finding a strong signal. Most interestingly, the combination of high spectral resolution and angular information allows us to map the emission from the WHIM in 3 dimensions. We cross-correlate the positions of galaxies in the simulation with this redshift map of  emission and detect the presence of six different ion species (Ne IX, Fe XVII, O VII, O VIII, N VII, C VI) in  the large-scale structure traced by the galaxies. Finally we show how such emission can be detected and studied with future X-ray satellites, with particular attention to a proposed mission, the Missing Baryon Explorer, or $\\sl MBE$. We present simulated observations of the WHIM gas with $\\sl MBE$. ",
        "introduction": "Cosmological hydrodynamic simulations provide us with  quantitative predictions for the state of baryonic matter in the universe (see, e.g, \\citealp{cos99,dav01}). At high redshift ($z\\gtrsim 3$), most baryons are diffuse, with a density close to the cosmic mean, and are located in the intergalactic medium (IGM).  They can be probed readily using Ly$\\alpha$ absorption systems in the optical spectra of high redshift quasars (see e.g.,  \\citealp{cen94,zan95,hkw96,wmhk97}). At low redshift, the process of gravitational  collapse has caused a minority of baryons to condense into structures such as stars, galaxies, groups and clusters of galaxies. While a  substantial fraction ($\\sim$ 30\\%) of baryons are trapped within the gravitational potential wells of these large scale structures in the form of hot intracluster or intragroup gas, the remaining baryons are left in intergalactic space and are predicted to form filamentary structures seen in cosmological hydrodynamic simulations. Some of these structures contain baryons with relatively low temperatures ($T \\lesssim 10^4$ K), and therefore enough residual neutral hydrogen to be detectable as the low-redshift Ly$\\alpha$ forest (see, e.g., \\citealp{dhkw99,pss00}). However,  a larger fraction have been shock-heated to temperatures between $10^5 - 10^7$ K. Given these temperatures, these baryons, which form the so-called ``warm-hot intergalactic medium,'' or ``WHIM,'' are best probed using UV/X-ray observations.  Recently, important progress has been made in the detection of absorption features produced by highly ionized metals in this WHIM gas in the UV/X-ray spectra of background quasars. In the UV band, a significant number of \\ion{O}{6} absorption lines have been seen with the Far Ultraviolet Spectroscopic Explorer ($\\sl FUSE$) and the Hubble Space Telescope ($\\sl HST$) (see, e.g., \\citealp{stl98,tsj00,tsa00,ssr02}). The distribution and derived properties of these \\ion{O}{6} lines are consistent with predictions from simulations (see, e.g., \\citealp{fbr01,cto01,che02}). In the X-ray band, \\citet{fan02}, \\citet{mwc02}, \\citet{cag02} and \\citet{mck03} reported on the detection of intervening \\ion{O}{7} and/or \\ion{O}{8} absorption lines with {\\sl Chandra} and {\\sl XMM}-Newton. \\citet{nic02}, \\citet{fsc03} and \\citet{ras03} also reported the detection of $z \\approx 0$ X-ray absorption lines. These low redshift lines may be attributable, at least in part, to the WHIM gas in our Local Group.  The absorption line method does have the advantage of directly probing the absorbing ions, as the detected line strength is proportional to the column density and so to the ion number density. However, it suffers the limitation of being only able to probe one dimensional information, along limited lines of sight, and detectability is also constrained by the flux of background sources. In order to fully reveal the three dimensional structure and physical properties of the WHIM gas, we need to study its {\\it emission} with imaging/spectroscopic methods. In this paper, we will focus on the X-ray emission from the WHIM gas. In general, X-ray emission from a hot, diffuse plasma contains two parts, a continuum which is contributed by various emission mechanisms, and line emission from metals in the IGM. While we will briefly discuss the continuum emission from the WHIM, this paper will be devoted mostly to studying the line emission.  Several authors have extensively studied the broad band soft X-ray emission (0.5 -- 2 keV) from the diffuse background, particularly the X-ray emission from the WHIM gas, using hydrodynamical simulations (see, e.g., \\citealp{cko95,poc01,cro01}). Much effort has been put into investigating the overall intensity of the emission, and how to distinguish it from the two dominant sources of the soft X-ray background: X-ray emission from  active galactic nuclei (AGNs) and from the hot gas in our Galaxy (see, e.g, \\citealp{ksn00,ksm01}). Simulations predict that the mean intensity of emission from the WHIM gas is between (2 -- 4) $\\times 10^{-13}\\rm\\ ergs\\ s^{-1}cm^{-2}deg^{-2}$, only about 5 -- 15\\% of the total extragalactic X-ray emission between 0.5 -- 2 keV (\\citealp{poc01,cro01}). While it seems to be a difficult task to detect most of this WHIM emission with current X-ray telescopes,  it appears feasible  to  detect the high intensity tail of the WHIM X-ray intensity distribution.  For instance, \\citet{sch00} reported on the presence of a possible filament with intensity $\\approx 6 \\times 10^{-13}\\rm\\ ergs\\ s^{-1}cm^{-2}deg^{-2}$ with the {\\sl ROSAT} PSPC; and \\citet{mar03} and \\citet{kaa03} reported detecting \\ion{O}{7} and/or \\ion{O}{8} from an extended, diffuse structure, possibly from WHIM gas; and \\citet{mcc02} recently performed a large-field ($\\sim$ 1 sr), high spectral resolution (5 -- 12 eV) observation with a  sounding rocket and detected \\ion{C}{5}/\\ion{O}{7}/\\ion{O}{8}  lines from the diffuse background, some fraction of which may be from the WHIM at zero redshift.  We expect that future X-ray missions will enable high resolution imaging/spectroscopic observations of the emission from the WHIM gas. We explore such a scenario with a high resolution hydrodynamic simulation in this paper. Our purpose is to address two important questions: with high spectral/spatial resolution, what does the X-ray emission from the diffuse, hot gas in intergalactic space look like; and are proposed future X-ray missions capable of conducting useful observations? Particularly, we will focus on a proposed mission called the Missing Baryon Explorer, or $\\sl MBE$. Given its large field-of-view ($29.5\\arcmin\\times29.5\\arcmin$), high spectral ($\\sim 4$ eV) and moderate angular ($\\sim 5\\arcmin\\times5\\arcmin$) resolution, we find that $\\sl MBE$ will be capable of revealing the X-ray emission from the WHIM gas. In addition, we will explore the clustering properties of the WHIM gas, and its correlation with nearby large scale structures traced by galaxies.  Recently, \\citet{yos03} examined the detectability of the X-ray emission from the WHIM gas independently with a similar numerical approach. They presented a detailed study of \\ion{O}{7} and \\ion{O}{8} emission lines from WHIM. They also investigated a dedicated X-ray mission, {\\sl DIOS} (Diffuse Intergalactic Oxygen Surveyor). It turns out that {\\sl DIOS} and {\\sl MBE} very similar properties.  While there are some differences between our work and theirs, such as the selection of the energy bands (\\citealp{yos03} concentrated on 0.5 -- 0.7 keV while we look at a broader band between 0.1 -- 1 keV) and the choice of the IGM metallicity models, where they can be compared our results agree with each other reasonably well.  Our plan for this paper is as follows.  We give a brief explanation of the simulation method and our treatment of the simulation data in \\S\\ref{simu}. A full study of the emission from the WHIM gas is presented in \\S\\ref{soft}, followed by measurement of its auto- and cross-correlation properties, making use of the galaxy distribution in the simulation for the latter.  We investigate the detectability of emission with the proposed {\\sl MBE} X-ray telescope in \\S\\ref{dete}, and \\S\\ref{disc} is our discussion and summary. ",
        "conclusions": "\\label{disc}  We have presented a detailed study of predictions for the  soft X-ray emission from the Warm-Hot Intergalactic Medium.  Our main conclusions can be summarized as follows:  \\begin{enumerate}  \\item We have examined images of soft X-ray emission from the hot IGM with a wide field of view ($200\\arcmin\\times200\\arcmin$). Having spectral information for each pixel allowed us to view the sky in several different soft X-ray energy bands, where emission lines from different ion species dominate. The low energy emission is markedly more diffuse than X-rays in the higher bands.  \\item We selected three representative simulated patches of sky to study with high resolution X-ray spectra: (a) a galaxy group; (b) a filament; (c) a void-like, underluminous region. In each case, we saw numerous emission lines from elements such as C, N, O, Fe, Ne, and identified their redshifts.  \\item By taking into account the background X-ray emission from the AGNs and foreground emission from the Galaxy, we obtained composite X-ray spectra of the selected regions. We found that in areas where the galaxy group is present, the X-ray spectrum is dominated by the hot intragroup medium, while in the void-like areas, the spectrum is completely dominated by the AGNs plus the Galactic foreground. In the filament case, although the majority emission comes from the AGNs plus the Galactic foreground, it is still possible to detect strong emission lines from the WHIM.  \\item We found that the emission in the  (0.5-0.8 keV) band (covering many of the emission lines)  shows strong small scale angular clustering, as measured by the angular autocorrelation function. Also, spectral information for each pixel allowed us to measure three dimensional clustering. In particular, we computed the cross-correlation between eight different ion  species and the simulated galaxy distribution.  We found that \\ion{O}{8} showed the strongest correlation signal, being measurable on scales as large as $\\sim 4-5 h^{-1}$ Mpc, while low energy species such as \\ion{C}{5} and \\ion{N}{6} showed no correlation at all. This variation of the correlation strengths  can be explained by ($\\epsilon_iZ_i$), the combination of line emissivity $\\epsilon_i$ and elemental abundance $Z_i$: the larger ($\\epsilon_iZ_i$), the stronger the correlation signal.  \\item Finally, we studied the detectability of the WHIM with several proposed X-ray missions. By comparing the effective area, resolving power and instrumental field-of-view, we found that the Missing Baryon Explorer ($\\sl MBE$) shows the most promise for detecting weak emission lines from the WHIM gas. We also showed simulated MBE instrument-folded spectra of our three selected patches of sky.  \\end{enumerate}  Accurate prediction of the soft X-ray emission depends on several important factors, including X-ray emissivity, ionization fraction, metal abundances, etc. The largest uncertainty in our predictions of the soft X-ray emission comes from the metal abundance in the IGM. Theoretically, simulations are beginning to address the formation and distribution of metals in a self consistent way (see e.g., \\citealp{cos99,she03}; Aguirre et al. 2001a,b,c). The metallicity-density relationship we adopted in this paper was motivated by this work,  and future studies of UV and soft X-ray emission lines should examine the question of metal enrichment in more detail (see, e.g., \\citealp{fur03}).  Observationally, metals are known to be distributed in a variety of systems with large dispersion in metallicity $Z$. For example, the metallicity can be as high as $10Z_{\\odot}$ in the central regions of some active galactic nuclei (AGNs) (\\citealp{mdp93,tls97}), and as low as $10^{-3}Z_{\\odot}$ in some halo stars in our galaxy \\citep{bee99}. When we focus on the low redshift IGM, the main evidence for an enriched IGM comes from the observations of the X-ray-emitting gas within galaxy clusters (\\citealp{mcc75,mus96,mlo97,dwh00}). With {\\sl ASCA}, \\citet{mus96} found that the metallicities in four clusters vary between $0.3-0.5Z_{\\odot}$. Studying $\\sim 40$ clusters at $z>0.14$, \\citet{mlo97} concluded that there is no metallicity evolution in clusters up to $z\\sim 0.3$. The observations of the Ly$\\alpha$ systems at low redshifts give metallicities as high as $\\sim 0.05Z_{\\odot}$ \\citep{bty98} indicating a metal enrichment in even lower density systems at low redshifts. The observations of damped Ly$\\alpha$ systems [N(\\ion{H}{1}) $> 10^{20}\\rm\\ cm^{-2}$] show a similar result \\citep{del00}. Due to the lack of observations, the metallicity of the ``WHIM'' is unknown.  Another uncertainty comes from the sensitive dependence of the X-ray emissivity and ionization fraction on physical properties of the WHIM gas, such as its temperature and density. For instance, both  the emissivity and ionization fraction peak over a rather narrow range of temperatures, and drop rapidly at low and high temperatures (see Figure~\\ref{f1} and Figure 2, 3, \\& 4 of \\citealp{che02}), which means a small change in temperature will cause order-of-magnitude changes in both quantities. This will in turn cause a dramatic change in observables such as intensity and emission line ion species. Because of this the actual observed spectrum depends sensitively on the underlying physical properties of the WHIM gas.  In this paper, we adopted several approximations to simply the calculations. First we did not take into account photoionization -- we used collisional ionization when calculating the ionization fractions. It is important to be aware that photoionization from background radiation may substantially alter the metal ionization fractions for low density, low temperature gas (see, e.g., \\citealp{che02}), which in turn will change the observed spectrum. However, since most of the emission spectra that can be observed with {\\sl MBE} or other telescopes are produced by hot gas in the high density, high temperature tail of the WHIM distribution, where collisional ionization dominates, we expect that photoionization will not substantially change those observables.  Metal cooling is ignored in our simulation: the cooling is primarily through H and He. Metal line cooling could be very important in some cases, especially for high density, warm/hot gas (see, e.g., \\citealp{sdo93}). However, since our results are very sensitive to the temperature variation, it is important to quantify the possible effect of metal cooling on WHIM gas temperature. To cool some hot gas, giving it a small drop in temperature $\\delta T$, the characteristic time scale $\\delta t$ is \\begin{equation} \\delta t = \\alpha \\frac{k}{n\\Lambda (T)} \\delta T. \\end{equation} Here $\\alpha$ is $3/2$ or $5/2$, depending on whether the cooling is isochoric or isobaric; $k$ is Boltzmann's constant; $n$ is the electron density; and $\\Lambda (T)$ is the cooling function. Assuming that  we are interested in how cooling acts on timescale of order the Hubble time, we can estimate the effect of metal cooling on the temperature decrement $\\delta T$ in terms of the physical parameters of the WHIM gas: \\begin{equation} \\delta \\log T \\approx 0.1\\ n_{-5} \\Lambda_{-23} T^{-1}_6 \\delta t_{\\rm Hubble}, \\end{equation} where $n_{-5}=n/(10^{-5}\\ \\rm cm^{-3})$; $T_6=T/(10^6\\ \\rm K)$; $\\delta t_{\\rm Hubble}$ is the cooling time in units of Hubble time scale. Here $\\Lambda_{-23}$ is cooling rate in units of $10^{-23}\\ \\rm ergs\\ cm^3s^{-1}$, typical of the WHIM gas with 0.1 solar abundance \\citep{sdo93}. So in general, metal cooling has only a minor effect on our results; however, since $\\Lambda (T) \\propto T^{-1}$ in the temperature range we are interested in, the metal cooling process is more efficient for low temperature gas, especially so when $\\Lambda (T)$ reaches its peak at $T\\approx 2\\times 10^5$ K.  It has been known for quite some time that non-gravitational heating, and/or cooling processes  can change the overall level of cosmological X-ray emission substantially (\\citealp{pen99,wfn01,vbr01a}).  The energy deposited by supernovae in our simulations has little direct  effect on the IGM. It is added as thermal energy to particles in high density regions and is radiated away quickly. Although the direct impact of feedback is therefore minimal, the cooling included serves to  change the entropy distribution of gas near the centers of virialized objects and lower the overall level of emission.  Other treatments of starburst energy such as those which involve giving wind particles kinetic energy so that they escape beyond galaxy and group virial radii into the IGM (e.g., \\citealp{she03}) are also likely to change the distribution of X-ray  emitting gas. Different treatments of feedback have been shown  by \\citet{bvo01} and \\citet{vbr01b} to change the pdf of X-ray surface brightnesses in the range which would affect the WHIM emission lines we focus on here. Varying the feedback model would be useful in the future in order to assess the range of effects on the WHIM emission.  Finally, the ionization state of the gas was calculated assuming ionization equilibrium. Non-equilibrium evolution of the ionization state can significantly change the ionization fraction (see, e.g., \\citealp{ech86,hca01}). The characteristic time scale for hot gas to reach ionization equilibrium, or the recombination time scale is $\\sim (n\\alpha_{rec})^{-1}$, where $\\alpha_{rec}$ is the recombination rate. At a temperature of $\\sim 10^6$ K, the recombination rate for \\ion{O}{7}, for example, is $\\sim 1.4\\times 10^{-12}\\ \\rm cm^3s^{-1}$ \\citep{svs82}. Here we include contributions from both radiative and dielectronic recombination. Assuming $n \\sim 10^{-5}\\ \\rm cm^{-3}$, typical of WHIM gas, the recombination time scale is $\\sim$ 2 Gyrs, less than the Hubble time. However, we must keep in mind that most of the hot gas was shock-heated at low redshift, so it is very likely that a significant portion of the WHIM gas has not reached its recombination time scale and is in a non-equilibrium ionization state.  A reasonable question to ask is: given detected X-ray emission lines, how do we identify lines with ion species and redshifts? It is impossible to carry out line identification with a single line; however, by simultaneously observing multiple lines such as the \\ion{O}{7} triplet and/or \\ion{O}{8} redshifted to the same degree, which is within the capability of the energy resolution of proposed X-ray missions such as {\\sl MBE}, we can easily identify lines with redshifts. In addition, measurements of the line ratios can provide direct temperature/density diagnostics for the ionized plasma (see, e.g., \\citealp{vcm86}). Observations of soft X-ray emission lines have the potential to not only reveal the presence of the missing baryons but to give us detailed information on their clustering, temperature, and  metallicity.  X-ray emission provides a direct probe of the WHIM gas. Unlike the X-ray absorption method, which can yield only one dimensional information, by combining imaging/spectroscopic observations, X-ray emission can show us the fully three-dimension distribution of the WHIM gas. This is essential not only for closure of the local cosmic baryon budget, but also if we are to understand the distribution of the WHIM gas, and how structures form and evolve. Moreover, by applying line diagnostic methods, X-ray emission lines can be used to probe the metallicity, temperature and density of X-ray emitting gas. This in turn can give us important information on galaxy formation and evolution, given that most of these metals which are responsible for X-ray emission lines are produced by supernova explosions and propagate from the interstellar medium into intergalactic medium. Although the relatively low density and temperature make it hard to observe, currently proposed X-ray missions will for the first time provide us with detectable signals from the X-ray emission of the WHIM gas.  This work is supported by Carnegie Mellon University. T.F. thanks members of the astrophysics group at CMU for their help and stimulated discussions. T.F. also thanks to the hospitality of the Kavli Institute for the Theoretical Physics at the University of California, Santa Barbara, where part of the work was conducted. This research was supported in part by the National Science Foundation under Grant No. PHY99-0794."
    },
    "0311/astro-ph0311188_arXiv.txt": {
        "abstract": "{We study the globular cluster system (GCS) of the giant elliptical \\object{NGC\\,4374} (\\object{M84}) in the Virgo cluster using B and R photometry. The colour distribution is bimodal with peaks at B-R=1.11 and B-R=1.36, fitting well to those found in other early-type galaxies. The radial profile of the cluster number density is flatter than the galaxy light. Using the luminosity function we derive a distance modulus of $\\mu=31.61\\pm0.2$, which within the uncertainty agrees with the distance from surface brightness fluctuations. Blue and red clusters show similar radial concentrations and azimuthal distributions. The total number of clusters is  $N=1775\\pm150$, which together with our distance modulus leads to a specific frequency of $S_{N}=1.6\\pm0.3$. This value is surprisingly low for a giant elliptical, but resembles the case of merger remnants like \\object{NGC\\,1316}, where the low specific frequency is probably caused by the luminosity contribution of an intermediate-age population. A further common property  is the high rate of type Ia supernovae which also may indicate the existence of a younger population. However, unlike in the case of \\object{NGC\\,1316}, one cannot find any further evidence that \\object{NGC\\,4374} indeed hosts younger populations. The low specific frequency would also fit to a S0 galaxy seen face-on. ",
        "introduction": "The study of globular clusters in early type galaxies has reached a state, where surprises have become rare when only individual galaxies are studied. Progress in understanding the relation between the morphology (i.e. the specific frequency of clusters, their spatial as well as their colour distribution) of a GCS and the host galaxy properties normally emerges from analyzing larger galaxy samples (e.g. Kundu \\& Whitmore \\cite{kundu01}; Larsen et~al. \\cite{soren01}) However, from time to time, one encounters galaxies which exhibit some peculiarity in their GCS, which one would like to understand in a more general framework.  For example, it is well known that ``normal'' elliptical galaxies have specific frequencies higher than about 3 (see Sect.~\\ref{sec:SpecFreq} for the definition and Elmegreen \\cite{elmegreen99} for a review of specific frequencies). \\object{NGC\\,1316}, the brightest galaxy in the Fornax cluster, has nevertheless a low specific frequency of only $\\sim0.9$ (Grillmair et~al. \\cite{grillmair99}; G\\'omez et~al. \\cite{gomez01}). Since it is known as a merger remnant, it is tempting to seek an explanation not in the low number of globular clusters, but in the high luminosity due to the presence of an intermediate-age stellar population which formed in the merger a few Gyrs ago. Indeed, Goudfrooij et~al. (\\cite{goud01a}), by means of spectroscopy of the brightest clusters, identified several intermediate-age globular clusters among them. Also the apparently high SN Ia rate (\\object{NGC\\,1316} hosted SN 1980N and SN 1981D) could indicate a strong intermediate-age population, given that the progenitor population of SNe Ia is suspected to have an age of a few Gyrs (Yungelson et~al. \\cite{yungelson95}; Yoshii et~al. \\cite{yoshii96}; McMillan \\& Ciardullo \\cite{mcmillan96}.)  The question arises whether the combination of a low specific frequency and a high SN Ia rate is ubiquitous or whether more examples can be found pointing to the possibility to use the specific frequency as an indicator for the presence of a younger population.  Early-type galaxies, which hosted more than one SN Ia, are rare. Besides \\object{NGC\\,1316}, \\object{NGC\\,4374} in Virgo is one of the few ellipticals with two or even three SNe events: 1957B, 1980I (also labeled ``intergalactic'' due to its location between \\object{NGC\\,4374} and \\object{NGC\\,4406}) and 1991bg.  The GCS of \\object{NGC\\,4374}  is not well investigated. Ajhar et~al. (\\cite{ajhar94}) presented VRI photometry of globular clusters for 10 galaxies in Virgo and Leo ellipticals, among them \\object{NGC\\,4374}. Due to the small field, properties like the specific frequency or the spatial distribution could not be addressed.  In the light of a possible intermediate-age population, it is worthwhile to investigate in detail the GCS of \\object{NGC\\,4374} with a larger field. ",
        "conclusions": "We analysed the GCS of the giant elliptical \\object{NGC\\,4374}, host of two or even three SNe Ia. This may lead to the suspicion that an intermediate-age population may be present in \\object{NGC\\,4374}, and perhaps manifest in the GCS properties. However, the colour distribution does not show anomalous features. It is best represented by a bimodal distribution  with peaks at B-R=1.11 and B-R=1.36, resembling the colour distribution found in other early-type galaxies.  A unimodal distribution is  ruled out by a KMM test.  The  spatial distributions  of the red and blue subpopulations do not show significant differences.  However, the number of clusters in \\object{NGC\\,4374} turns out to be surprisingly lower than what is expected for its high luminosity. This could indicate  the presence of an intermediate-age population, as  has been the case for other early-type galaxies with low specific frequencies. It is intriguing that these examples (\\object{NGC\\,1316} and \\object{NGC\\,5018})  also have been hosts of type Ia SNe.  The other possibility is that \\object{NGC\\,4374} is a S0 galaxy seen face-on, in which case the low specific frequency of 1.6 would fit to other objects of this class. One then would expect a disk-like kinematic behaviour of the  metal-rich globular clusters or of a subsample . Measuring radial velocities of many globular clusters could provide further clues."
    },
    "0311/astro-ph0311377_arXiv.txt": {
        "abstract": "We study numerically the collapse of rotating, magnetized molecular cloud cores, focusing on rotation and magnetic braking during the main accretion phase of isolated star formation. Motivated by previous numerical work and analytic considerations, we idealize the pre-collapse core as a magnetized singular isothermal toroid, with a constant rotational speed everywhere. The collapse starts from the center, and propagates outwards in an inside-out fashion, satisfying exact self-similarity in space and time. For rotation rates and field strengths typical of dense low-mass cores, the main feature remains the flattening of the mass distribution along field lines -- the formation of a pseudodisk, as in the nonrotating cases. The density distribution of the pseudodisk is little affected by rotation. On the other hand, the rotation rate is strongly modified by pseudodisk formation. Most of the centrally accreted material reaches the vicinity of the protostar through the pseudodisk. The specific angular momentum can be greatly reduced on the way, by an order of magnitude or more, even when the pre-collapse field strength is substantially below the critical value for dominant cloud support. The efficient magnetic braking is due to the pinched geometry of the magnetic field in the pseudodisk, which strengthens the magnetic field and lengthens the level arm for braking. Both effects enhance the magnetic transport of angular momentum from inside to outside. The excess angular momentum is carried away in a low-speed outflow that has, despite claims made by other workers, little in common with observed bipolar molecular outflows.  We discuss the implications of our calculations for the formation of true disks that are supported against gravity by rotation. ",
        "introduction": "Rotation is observed ubiquitously in star-forming cores. On the scale of $\\sim$0.1~pc, the detected rotation rate is small, with a typical value corresponding to a ratio of rotational to gravitational energy of $0.02$ (Goodman et al. 1993). The rotation rate is much lower than expected if the core is condensed out of a low density background medium with angular momentum conserved. The slow core rotation is thought to be due to magnetic braking (Mestel 1985; Mouschovias \\& Ciolek 1999), which is particularly efficient during the long, magnetically subcritical phase of cloud evolution driven by ambipolar diffusion (Basu \\& Mouschovias 1994).  This assumes that the clouds are strongly magnetized to begin with, as envisioned in the standard picture of isolated low-mass star formation (Shu, Adams \\& Lizano 1987). Once a magnetically supercritical core develops, it begins to collapse dynamically. It is often assumed that magnetic braking then becomes inefficient, with most of the angular momentum of the core material carried by the collapsing flow into the central star-plus-disk or binary system. However, the typical value of specific angular momentum inferred for dense cores is an order of magnitude higher than the typical binary value, at least in the Taurus cloud complex (see Fig.~8 of Simon et al. 1995 and discussion in \\S~8.5 of Myers 1995). This suggests that a large fraction of the core angular momentum may still be removed {\\it during collapse}. The problem of angular momentum redistribution in a collapsing magnetized core thus warrants a closer examination.  The efficiency of magnetic braking depends sensitively on the magnetic field geometry (Mestel 1985). The exact field geometry of a magnetized core on the verge of collapse is difficult to determine without following the cloud evolution in three dimensions. Simplified numerical calculations  of cloud evolution driven by ambipolar diffusion have shown, however, that by the time a central density cusp is formed the mass-to-flux ratio is approximately spatially constant in the supercritical core, and the (volume) density decreases with radius roughly as a power-law of $r^{-2}$ (Lizano \\& Shu 1989; Basu \\& Mouschovias 1994). This self-similar behavior motivated Li \\& Shu (1996) to idealize the end state of the ambipolar diffusion-driven core formation as a scale-free configuration supported in part by a large scale magnetic field and in part by the thermal pressure gradient, with a fixed mass-to-flux ratio and an exact $r^{-2}$ density profile. Such a configuration turns out to be a singular isothermal toroid, for which the  magnetic field geometry can be computed exactly (see also Baureis, Ebert, \\& Schmitz 1989).  The field geometry is distorted by fluid motions during the subsequent inside-out collapse.  In a companion paper to this journal (Allen, Shu \\& Li 2003; hereafter Paper I), we have studied the collapse of magnetized, singular isothermal toroids without rotation. The main features of the collapse solution are governed by the anisotropy in the magnetic forces, which induces in the toroid an evacuated region near the axis, where the plasma $\\beta$ parameter (the ratio of thermal to magnetic energy density) drops below unity (see Fig.~4a of Paper I). As the core collapses from inside out, matter in the magnetically dominated region slides along the field lines towards the equatorial region, creating a dense, flattened structure -- a pseudodisk (Galli \\& Shu 1993a,b). The material in the pseudodisk collapses radially towards the central point mass in a magnetically diluted free-fall, dragging the footpoints of field lines along with it (see Fig.~4b of Paper I). A highly pinched field configuration (``split monopole\") is produced, with important implications for the magnetic transport of angular momentum from inside to outside. In the presence of rotation, the collapse would lead to a spinup of the pseudodisk, which is linked magnetically to a more slowly rotating envelope of infalling and nearly static material.  We will show in this paper that the twisting of magnetic field lines actually drives a slow outflow during the core collapse, which can remove much of the angular momentum remaining in the infalling pseudodisk.  This removal poses a substantial brake against the eventual formation of a centrifugally supported disk unless the trapped magnetic field is lost in the realistic situation by nonideal MHD effects.  As the first step toward the dynamical calculations, we will generalize the magnetized, singular isothermal toroid solutions of Li \\& Shu (1996) to include rotation in \\S~\\ref{initial}. These magnetized rotating toroids are then induced to collapse from inside out in \\S~\\ref{collapse}, and the collapse is followed numerically in a manner similar to that described in Paper I. The numerical solutions are discussed in \\S~\\ref{discussion}, with an emphasis on the efficiency of magnetic braking during core collapse and the implications for the difficulty of forming rotationally supported disks unless significant flux loss appears in the realistic problem at small scales, as is expected theoretically for high-density situations (e.g., Nishi, Nakano \\& Umebayashi 1991; Desch \\& Mouschovias 2001). We end with a brief summary and conclusion in \\S~\\ref{conclusion}. ",
        "conclusions": "\\label{discussion}  \\subsection{Magnetic Braking and Keplerian Disk Formation} \\label{nonideal}  Formation of centrifugally supported circumstellar disks is a central problem in star and planet formation.  The initial distribution of angular momentum in the dense core, which collapses to form the star and disk, and the angular momentum redistribution during the collapse and disk accretion both affect the final outcome. We have shown that almost all of the core angular momentum can be removed in a low-speed outflow driven by magnetic braking during the collapse, provided that (1) the core is significantly magnetized and (2) the ideal MHD approximation holds. We now examine these two conditions from a practical point of view.  The degree of magnetization of dense cores of molecular clouds prior to collapse is somewhat uncertain. In the standard scenario for isolated low-mass star formation (Shu et al. 1997; Mouschovias \\& Ciolek 1999), the cores are envisioned to form out of a magnetically subcritical background cloud (bounded at its surface by significant thermal or ram or turbulent pressure), with the evolution driven by ambipolar diffusion. The cores remain well magnetized, even during the supercritical phase of evolution, when their dimensionless mass-to-flux ratio $\\lambda$ is about 2 (Basu \\& Mouschovias 1994; Nakamura \\& Li 2003). This ratio, corresponding to $H_0 -v_0^2/2 \\approx 0.5$ (Li \\& Shu 1996, this paper), is consistent with the field strength obtained for the well studied cores L1544 (Crutcher \\& Troland 2000) and B1 (Goodman et al. 1989) from Zeeman measurements, after likely geometric corrections (Crutcher et al. 1994; Shu et al. 1999; Ciolek \\& Basu 2000). Polarization has been detected in the dust continuum emission from a number of cores (e.g., Ward-Thompson et al. 2000; Matthews \\& Wilson 2002), which also points to a significant degree of magnetization. In addition, the cores are known to be elongated in general, with a typical aspect ratio of 2:1 (Myers et al. 1991); the elongation is too large to be induced by rotation, and is consistent with the presence of a dynamically important magnetic field (Li \\& Shu 1996). Even if the molecular clouds are magnetically supercritical to begin with and cores form in them dynamically, it would still be difficult to avoid having an $H_0-v_0^2/2$ at least as high as 0.2 to 0.3, which corresponds to a weak field of $\\sim 3~\\mu$G threading a slab of interstellar medium of $A_{\\rm V}=1$. Based on these arguments, we believe that the dense cores formed in present-day star-forming clouds should trap enough magnetic flux for magnetic braking to operate efficiently during gravitational collapse, as long as the field remains well coupled to the bulk of the neutral core matter.  Good coupling between the magnetic fields and neutral matter is maintained through collisions with ions in relatively low-density regions of a dark, molecular cloud core. In such regions, the primary source of ionization, and the time scale for ambipolar diffusion is typically an order of magnitude longer than the dynamic time scale\\footnote{In the more turbulent, lower density envelope surrounding a dense core, the time scale for ambipolar diffusion can be shortened considerably through field fluctuations (Fatuzzo \\& Adams 2002). The faster ambipolar diffusion may affect the angular momentum evolution during the early phase of core formation. }, which means that there is little time for magnetic flux to be redistributed in a dynamically collapsing flow. Thus, ideal MHD is a reasonable first approximation in collapse calculations. Nishi et al. (1991; see also Desch \\& Mouschovias 2001) claim that good coupling is lost above a volume density of order $10^{11}$~cm$^{-3}$.  Magnetic decoupling occurs for densities higher than this critical value, which reduces the efficiency of the magnetic braking. The decoupling density well exceeds the density in most of the pseudodisk, which is typically less than $\\sim 10^7$~cm$^{-3}$ (see the density contours in Fig.~\\ref{fig:n1r} for an example), except in one or two zones right next to the central cell. Therefore, most of the angular momentum of the collapsing core material is already removed by magnetic braking before the braking is rendered ineffective by decoupling.  Let us now examine the supposition that the gas volume density will eventually exceed the critical value for decoupling. While true in rotationally supported protoplanetary disks with similar physical properties to the so-called ``minimum solar nebula,'' it is not necessarily the case for the dynamically collapsing pseudodisks, whose surface densities are much lower and provide less shielding against cosmic ray ionization. It is an easy matter to estimate the surface density $\\Sigma$ of a flattened configuration accreting axisymmertically in quasi-steady state $\\dot M$ at radial speeds that approach free-fall values $\\sqrt{ 2GM/\\varpi}$ onto a central mass $M$: \\begin{equation} \\Sigma = {\\dot M\\over 2\\pi \\varpi \\sqrt{2GM/\\varpi} } = 0.16 \\; {\\rm g \\, cm}^{-2} \\left( {{\\dot M}\\over 10^{-6} \\;  M_\\odot {\\rm yr}^{-1}}\\right) \\left( {M\\ \\varpi \\over M_\\odot \\; {\\rm AU}}\\right)^{-1/2}. \\label{surfdenff} \\end{equation} Within a factor of 2 or so, the inner regions of pseudodisks obtained by numerical simulation in this paper satisfy the estimate of equation (\\ref{surfdenff}).  Galactic cosmic rays can keep a layer with a surface density of about 100 g cm$^{-2}$ sufficiently ionized as to couple well to magnetic fields (Nishi et al. 1991).  Thus, in the absence of a magnetized YSO wind or jet that can sweep out cosmic rays before they penetrate pseudodisks, Galactic cosmic rays will make a pseudodisk with a surface density profile given by equation (\\ref{surfdenff}) sufficiently well-ionized all the way to the surface of a protostar with characteristic radius $\\sim 10^{-2}$ AU. Even after a YSO jet or wind turns on, if the event is accompanied by stellar magnetic activity and the emission of X-rays, the observed levels of X-rays in low-mass protostars are sufficient to keep a layer with the surface-density profile of equation (\\ref{surfdenff}) marginally well-coupled (Glassgold, Najita, \\& Igea 1997).  Why then, should centrifugally supported disk ever arise?  Because the coupling to magnetic fields in the pseudodisks, while good, is not perfect.  Nonideal MHD effects of two types have been described in the literature that can substantially modify the discussions considered so far in this paper.  Li \\& McKee (1996) focused on slippage of magnetic flux in the radial direction by the effects of ambipolar diffusion.  They proposed that the magnetic flux decoupled this way from the mass that enters the star can drive a hydromagnetic shock against the collapsing inflow. Behind the shock, the accretion flow can be slowed down by a large factor, and become magnetically dominated. It is conceivable that magnetic braking would still remove a large fraction of the remaining angular momentum in the slowly-contracting post-shock region (Krasnopolsky \\& K\\\"onigl 2002).  Galli \\& Shu (1993b) focused on the vertical stratification produced by the highly pinched magnetic field in the pseudodisk as a condition whereby magnetic flux can be lost by reconnection. As we also noted in this paper and paper I, infall in the pseudodisk tends to produce a split-monopole geometry with reversed poloidal fields across the midplane. In an ideal MHD calculation, the reversed magnetic fields across the midplane are supported by current sheets. In reality, the large electric current density near the midplane will dissipate by the presence of nonzero electrical resistivity.  The dissipation of these currents corresponds to annihilation of the reversed fields above and below the midplane. The higher magnetic pressures of the surrounding regions will then press toward the middle, leading to further field annihilation.  Except for the continued introduction of new interstellar flux by infall into the pseudodisk, the entire magnetic flux threading the pseudodisk could eventually be lost by this reconnection process.  In quasi-steady state, we therefore expect a ``sandwich-like'' structure to develop. The ``bread'' of the sandwich is the well-coupled, magnetized, regions where angular momentum can be removed at a sufficient pace as to allow the matter to spiral essentially unimpeded at dynamical rates toward the central protostar.  This bread is the radially inward extension of the original pseudodisk, and it may already have been detected in the case of IRS L1489 by the observations of Boogert, Hogerheijde, \\& Blake (2002).  The ``meat'' of the sandwich is the part which is sufficiently shielded by the bread against external ionizing radiation as to become essentially decoupled from any magnetic fields.  This decoupling causes eventual spinup of the inwardly spiraling material, because of the lack of magnetic braking, and the eventual formation of a centrifugally supported disk.  This ``true'' disk is bounded above and below by the dynamically infalling remnant of the pseudodisk.  How massive the meat becomes relative to the bread is, at this point, entirely a matter of speculation.  It could vary considerably between regions of isolated star formation, such as Taurus-Auriga, and regions of crowded star formation, such as Orion, because of the difference of ionizing radiation (particulates and photons) associated with star-birth activity.  Like many examples in astrophysics, moreover, such as the Sun or other magnetically active stars, the reconnection process may occur not quiescently but explosively and intermittently in flares.  It is interesting to speculate whether such intermittent behavior has anything to do with FU Orionis outbursts. Other effects may enter to further complicate the situation.  The flow could become turbulent because of Kelvin-Helmholtz instabilities that arise from the velocity shear between the bread and meat parts of the sandwich, or because of interchange instabilities in 3D that manifest themselves in the pinched magnetic configuration of the split-monopole. Under laminar conditions, resistive effects become important in the same density regime where ambipolar diffusion becomes competitive with the dynamics (Nishi et al. 1991), so a combination of the effects discussed by Galli \\& Shu (1993b) and Li \\& McKee (1996) may be operative.  All these effects operating together in lightly ionized media, under turbulent conditions, may considerably increase the size scales over which nonideal effects must be considered (Zweibel \\& Brandenburg 1997a, b). Challenging nonideal 3D simulations with adaptive mesh refinement may be needed to quantify such possibilities.  \\subsection{Comparison with Previous Work}  Magnetic braking in star formation has been studied previously by many authors, in connection with the long-standing ``angular momentum problem''. Most of the early studies were concerned with the relatively low-density, initial stage of star formation using semi-analytic treatments (e.g., Mestel \\& Paris 1984; Nakano 1989; see Mouschovias 1990 for a review). Recent numerical calculations have extended the treatment to higher densities associated with the runaway core contraction either prior to star formation (Basu \\& Mouschovias 1994) or slightly beyond the formation of a central object (Tomisaka 2002). These calculations showed that, during the short runaway phase when the central density increases by several orders of magnitude, the angular momentum is nearly conserved, because the density increase involves little positional change of fluid elements in both radial and azimuthal directions (and thus little pinching and twisting of field lines). Our work deals exclusively with magnetic braking in the infall or accretion phase, {\\it after} the formation of a compact central object. One may naively think that, once a molecular-cloud core starts to collapse dynamically, there should be little time for magnetic braking to operate. And yet we find that the bulk of the angular momentum, if not the total, can be removed during the dynamical collapse, even with a magnetic field of relatively moderate strength. Why is there this difference between previous expectations and current calculations?  There are two key ingredients behind the efficient braking during the accretion phase: (1) the domination of magnetic energy over the rotational energy, and (2) the distinct geometry of the magnetic field. The foot points of the magnetic field lines associated with the matter already landed on the central object are pinned near the origin by accretion, creating a split monopole for the magnetic geometry (see Fig.~\\ref{fig:n1r}). The split-monopole field lines push on the field lines outside (see also Li \\& Shu 1997), bending them into an hour-glass shape that is conducive to efficient magnetic braking. The enhancement comes from both the greater field strengths due to pinching and the radial fan-out (as opposed to vertical collimation) which lengthens the level arm for magnetic braking. These two effects are not taken into account in previous simple estimates. They are analogous to magneto-centrifugally driven winds (K\\\"onigl \\& Pudritz 2000; Shu et al. 2000). Indeed, the slowly moving outflow we find in the collapse solution is fundamentally similar to these winds, except for scale and for the ultimate origin of the field (interstellar, disk dynamo, or stellar dynamo).  Magneto-rotationally driven outflows were also found in the collapse calculations of Tomisaka (1998, 2002). These calculations start from infinitely long cylinders with magnetic field lines parallel to the axis. The magnetized cylinder is induced to break up into dense cores by a sinusoidal density perturbation along its length. The subsequent core evolution prior to the formation of a central object is more dynamic than that driven by ambipolar diffusion. It produces a pivotal state at time $t=0$ (when a {\\it hydrostatic} core starts to form in the calculation) that differs substantially from the singular toroids that we adopted based on ambipolar diffusion calculations. Nevertheless, soon after the appearance of the hydrostatic core, an outflow develops, which Tomisaka attributed to core formation. A more precise cause for the outflow, we believe, is the trapping of a finite amount of magnetic flux (rather than mass) in a small (core) region, which gives rise to the central magnetic monopole responsible for the pinched magnetic geometry that is conducive to outflow driving. For numerical reasons, Tomisaka's calculations were stopped a short time into the accretion phase, typically $4\\times 10^{-3}$ times the free-fall time at the center of the initial cylinder beyond the moment of hydrostatic core formation (t=0). During this brief period of time, the outflow appears to have removed most of the angular momentum of the material close to the origin (Tomisaka 2000). In particular, no rotationally supported disk is able to form, consistent with our findings which, because of their self-similar nature, apply to all times $t>0$.  Formation of rotationally supported disks during the collapse of dense cores formed in strongly magnetized molecular clouds was recently examined by Krasnopolsky \\& K\\\"onigl (2002) under thin-disk approximation. The material in the disk is connected through field lines to an external medium, which brakes the disk rotation with a parameterized efficiency. They find that, in the ideal MHD limit, a rotationally supported disk can form, provided that the braking efficiency is low enough (see their Fig.~3). If the braking is more efficient, the disk angular momentum can be radiated away completely, as shown in their Fig.~6. Our calculations and those of Tomisaka (1998, 2002), which treat the dynamics of the material above and below the disk self-consistently, suggest that the braking efficiency is high enough to prevent the rotationally supported disk from forming in a strongly magnetized core, as long as ideal MHD holds. In the presence of ambipolar diffusion, Krasnopolsky \\& K\\\"onigl (2002) find that the bulk neutral material can retain a small rotation even in the strong braking case, because the neutral component must rotate relative to the charged component in order to feel the magnetic braking torque (see their Fig.~10). Whether the residual angular momentum is large enough in realistic situations to form a rotationally supported disk well outside the forming star remains to be seen.  To conclude this subsection, we note an interesting feature of regarding numerical errors, which are always present, of course, in any simulation. Simulations of self-similar configurations such as ours confer the advantage that although one does not know (in analytically intractable situations) what the exact details of the solution should be, one does know from dimensional considerations how each quantity should scale.  This knowledge provides a powerful check on the present simulations, and therefore we have a good handle on the numerical reliability of the results quoted in this paper.  \\subsection{Observational Implications: The Case of IRAM 04191}  Rotation introduces a new feature into our collapse solution - the slow outflow (see also Tomisaka 1998, 2002). Despite its  crucial role in angular momentum redistribution, this outflow is unlikely to be observed soon: it is too easily masked by the faster-moving bipolar molecular outflows driven by a jet/wind from near the central object. Another prominent feature of the collapse solution -- the flattened, high-density pseudodisk in the equatorial region -- should be less affected by the bipolar outflow. Indeed, there are many examples of dense structures surrounding deeply embedded YSOs that are elongated perpendicular to the outflows. They can be plausibly interpreted as the pseudodisks formed during the collapse of magnetized toroids or, on a larger scale, the more static toroids themselves, although detailed kinematic studies and direct measurements of the strength and morphology of magnetic fields are needed to strengthen the case. The best studied circumstellar structure of this type is perhaps the one around the source IRAM 04191 (Belloche et al. 2002), which we discuss in some depth below.  IRAM 04191 is a low-luminosity ($L_{\\rm bol}\\sim 0.15 L_\\odot$) protostar in the nearby Taurus molecular cloud, surrounded by a massive envelope ($M_{\\rm env}\\sim 1.5 M_\\odot$; Andr{\\'e, Motte, \\& Bacmann 1999). This Class 0 source drives a well-developed molecular outflow, as all Class 0 sources do. Perpendicular to the outflow lies a flattened structure that is mapped in both dust continuum and several molecular lines (Belloche et al. 2002). The line study revealed that the flattened structure is infalling at a velocity between half and one isothermal sound speed, with a substantial amount of rotation. The (linear) rotational speed increases from $\\sim 0.2\\ a$ at a radius of $\\sim 11,000$~AU to $\\sim a$ at $\\sim 3,500$~AU, where $a\\approx 0.2$~km~s$^{-1}$ is the isothermal sound speed at 10~K (see Fig.~12c of Belloche et al. 2002). We identify the increase in rotational speed with the spinup in the outer part of the pseudodisk as a result of radial contraction under the approximate conservation of specific angular momentum, before significant magnetic braking sets in. In our model, the rotational speed peaks at some radius, and then decreases inward in a region of strongly pinched magnetic field that efficiently removes the angular momentum of the accreting flow, as shown in Fig.~\\ref{fig:n1r}c. If the flattened structure of IRAM 04191 is indeed a pseudodisk, then its rotational speed should decrease inward inside $\\sim$ 3,500 AU according to our model. This is consistent with, but not required by, the available spectral data, particularly the ``S''-shaped position-velocity diagrams of several molecular species. Higher resolution observations are needed to test the model.   The test may be hampered by the freeze-out of molecules onto dust grains in the high-density region near the protostar, as evidenced by the ``hole'' in the NH$_3$ distribution around this source (Wootten et al. 2003).  Interpreting the flattened structure of IRAM 04191 strictly in terms of the models of this paper is not without difficulties. In the fiducial case where $v_0=0.25$ (comparable to the $0.2 \\ a$ seen at 11,000 AU in IRAM 04191), the rotational speed achieves a peak near the observed value $\\sim a$ at a reduced radius of $x_{\\rm p}\\sim 0.25$. If this corresponds to the radius of $3,500$~AU for IRAM 04191, then the collapse would have been initiated $\\sim 3.3\\times 10^5$~yrs ago, and the central star would have a mass of $\\sim 0.8 M_\\odot$. But the dynamical time for the molecular outflow of IRAM 04191 is only $\\sim 8\\times 10^3$~yrs, much shorter than the above estimate. The low luminosity 0.15 $L_\\odot$ of the central source also suggests a fairly small mass for the central star, accreting gas at not much higher than the standard rates for Taurus (see, e.g., Adams, Lada, \\& Shu 1987).  A flawed way to reconcile theory and observation, within the context of quasi-steady flow, is to postulate a significant temporal offset between the pivotal instant $t=0$ and the appearance of a powerful outflow in the system.  Such an offset requires, for example, some way to store infalling material (from a pseudodisk) in a rotationally supported disk without producing either much accretion luminosity or an outflow for a duration approaching $10^5$ yr.  This task is not easy to do without inducing fierce nonaxisymmetric gravitational instabilities that would undermine the assumption of little mass accretion into the central regions.  Moreover, the observational limit claimed for the mass or size of any disk in IRAM 04191 is less than $\\sim 10^{-3}$ $M_\\odot$ or 10 AU.  Efficient magnetic braking has evidently prevented a more massive/larger disk from forming close to the star, and such braking works against the idea of appreciable mass storage in a disk.  A better solution might be to suppose that, in addition to an offset between the pivotal instant and the onset of a wind, accretion onto the central source is time-variable, with FU Orionis outbursts punctuating quiescent periods of relative calm (Hartmann, Kenyon \\& Hartigan 1993).  In such a picture, IRAM 04191 has just recently experienced such an outburst (accounting for the associated powerful molecular outflow) and has now faded considerably in its accretion luminosity (accounting for the low luminosity currently). Oscillatory inflow into the central regions in the presence of steady infall from the pseudodisk is a common bane of our simulations (see the discussion in Paper I). It remains to be seen whether the numerical artifact of field pinching and unpinching observed in the simulations might not have their physical analogues in more realistic calculations performed with nonideal MHD.  Another plausible option invokes a pre-existing contracting and rotating toroid to explain the observations of the regions at a few thousand AU and beyond. Subsonic cloud contraction prior to the formation of a formal density singularity for $t < 0$ is consistent with the observed motions in both IRAM 04191 and the well-studied starless core L1544 (Tafalla et al. 1998). As discussed already in Paper I, subsonic contraction of the requisite magnitude might be initiated by the dissipation of turbulent support in dense molecular-cloud cores (Myers 1999) or resulted from ambipolar diffusion (e.g., Ciolek \\& Basu 2000). We have performed simulations of the fiducial (rotating and magnetized) case with an initial uniform radial inflow $u_r$ equal to half the sound speed $a$.  As we found in Paper I for the example without rotation, the results look like the configuration started without initial radial motion, except that the density profiles next to the axis are partially filled in by the additional inflow and the central accretion rate is increased by about a factor of 2.  \\subsection{Conclusion} \\label{conclusion}  To summarize, we have studied numerically the inside-out collapse of rotating, magnetized singular isothermal toroids. We find that a pseudodisk forms, just as in the nonrotating case, through which most of the core material is accreted. The pinched magnetic field in the pseudodisk removes most of the angular momentum of the accreted matter, and deposits it in a low-speed wind. The efficient magnetic braking during the protostellar accretion phase has important implications for the formation of rotationally supported disks that need to be pursued with nonideal MHD calculations.  The complementary side to this phenomenon is the discovery of an extremely simple and efficient mechanism of angular momentum transport that allows direct accretion of extended rotating material onto relatively compact objects without the necessary concomitant formation of a standard viscous accretion disk. Whether this mechanism has anything to do, for example, with the mini-spiral that seems to be feeding ionized gas, with low radiative efficiency, into the massive black hole at the center of our own Galaxy (Lo \\& Claussen 1983) remains to be seen. We have speculated in this paper that the process, with the inclusion of nonideal MHD effects, opens new vistas by which the two classical conundrums of star formation discussed long ago by Mestel and Spitzer -- the angular-momentum problem and the magnetic-flux problem - might be resolved, perhaps simultaneously, in the disks and pseudodisks surrounding newly formed stars.}"
    },
    "0311/astro-ph0311606_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\hspace{18pt}Recent studies of low-mass stars in the star-forming regions Taurus \\citep{Gullbring98} and Orion \\citep{Robberto03} indicate surprisingly low accretion rates between $\\dot{M}\\simeq 10^{-7} - 10^{-10} M_{\\odot} yr^{-1}$. These accretion rates are far too low to build up solar-type stars on timescales of the order of 10$^{6}$ yr, as it is expected in the current picture of low-mass star formation (see \\citealt{Lada99} and references therein). An elegant alternative to the quiescent, steady accretion can be the occurence of periodic, short events of intense accretion ($\\dot{M}\\simeq 10^{-4} M_{\\odot} yr^{-1}$), during which a significant fraction of the disk mass can be accreted \\citep{Hartmann96}. The well-known FUor objects \\citep{Herbig66,Herbig77} may represent this intense accretion phase of otherwise quiescent T Tauri stars \\citep{Hartmann96}. Thus, if FUor outbursts are typical to {\\it all} T Tauri stars, these outbursts might be the dominant way of mass accretion.  Models based on the enhanced accretion rates are in general successful in explaining the observed properties of the FUor outbursts \\citep{Hartmann85,Kenyon88}, but the causes of the onset of the rapid accretion are still disputed. Suggested possibilities include thermal instabilities in the disk (see, e.g., \\citealt{Lin85}; \\citealt{Kawazoe93}; \\citealt{Bell94}) and perturbations from close companions \\citep{Bonnell92,Clarke96}, but up to now none of these explanations could be confirmed observationally.  However, earlier high-angular resolution observations led to the detection of companions for the FUor objects L1551 IRS 5 \\citep{Rodriguez98} and Z CMa \\citep{Koresko91}, with a separation of 45 and 93 AU respectively. Moreover, \\citet{Kenyon93} found RNO 1B and 1C to be a FUor binary. Infrared long-baseline interferometry shows that FU Ori might have a close companion with a separation as small as 0.35 $\\pm$ 0.05 AU, although these observations could be explained more naturally with the presence of a circumstellar disk \\citep{Malbet98}.  In this letter we present adaptive optics imaging of the young pre-main sequence star FU Ori and the detection of a faint red star in its apparent vicinity. Further results from our imaging campaign to search for companions of young stars were reported by \\citet{Apai03}.  In order to avoid confusion, throughout the letter we use the name FU Ori for the prototype star ($\\alpha = 05^h45^m22.^s6, \\delta = +09^\\circ04^{\\prime}12^{\\prime\\prime}$), while the term FUor stands for the class of stars named after FU Ori. ",
        "conclusions": "\\hspace{18pt}Our AO-assisted near-infrared imaging led to the detection of a red point source south of the well-known outburst star FU Ori. We discuss the possible nature of this object and conclude that the observed properties are consistent with either a pre-MS star at the distance of FU Ori or -- with much smaller probability -- a red background giant star. We argue that mid-resolution optical spectroscopy can discriminate between these cases.  As the FUor outbursts may be of central importance for star formation, a possible companion to the prototype object FU Ori can have a major impact on the outburst models. In particular, the close fly-by of the companion could lead to disk perturbations, therefore, providing an external triggering of the outbursts of FUor stars.  We would like to thank the staff members of the Calar Alto Observatory for their support during the observations. We are grateful to Ch. Leinert for his useful comments. H. Wang acknowledges the support by NSFC grants 10243004 and 10073021."
    },
    "0311/astro-ph0311095_arXiv.txt": {
        "abstract": "{We study  the  hard X-ray  emission of two samples of Seyfert 2 galaxies \\textit{with} and \\textit{without} Polarized Broad Lines (PBL). In the hard X-ray domain, absorption effects do not significantly modify  the intrinsic emission allowing us a direct access to the central engine.\\\\ The purpose of this study is  to compare the primary  emission of the two Seyfert 2  subclasses in order to investigate the nature of their central engine   and to test  unified models according to which they both have a hidden Seyfert 1 nucleus.  We compute the average hard X-ray spectra of Seyfert 2 galaxies with and without PBL observed with BeppoSAX/PDS (15-136 keV). The two spectra have a  common general behavior at  first sight, but investigating deeper we find  differences in the  intrinsic  properties  of  the two categories of Seyfert 2 galaxies. Sy 2 with polarized broad lines  have physical conditions close to those of Sy 1 galaxies whereas Sy 2 without PBL differ substantially, suggesting that they may  have a particular  place in the  scheme of Seyfert galaxies. ",
        "introduction": "Seyfert galaxies are classified into two spectroscopic groups based on the presence (type 1) or absence (type 2) of broad permitted optical emission lines. The first discovery of Polarized Broad Lines (PBL hereafter) in NGC 1068 \\citep{Antonucci} and later in  several other Seyfert 2 galaxies \\citep{MillerGood,Tran95,Young,Heisler,Awaki,Moran2000,Alexander,GuPbl,Lumsden} suggests that  Seyfert 2 galaxies harbor a bright Seyfert 1 nucleus hidden from our view by an optically and geometrically thick obscuring torus (``Unified models'', see the review of \\citet{Antonucci93}). In this frame, the difference between type 1 and 2 Seyfert galaxies is only due to the viewing angle. X-ray  observations  support this model showing the presence of large column densities along the line of sight of Seyfert 2 galaxies \\citep {Risaliti, Guainazzi2001}. However, other observations are inconsistent with   unified models, in particular in the hard X-ray domain where \\citet{Z95} and \\citet{Deluit1} (D03 hereafter)  show differences in the spectral index, energy cutoff and amount of reflection between Sy 1 and Sy 2 galaxies. Furthermore, \\citet{Alexander} refutes the absorption argument to explain the absence of broad lines in some Sy 2 galaxies  showing that the density of the absorbing medium is not linked to the detectability of broad lines.\\\\ The polarization is presumed to  be due to nuclear emission scattered by  a gas of electrons located further out than the obscuring material. According to current  models, the detection of polarized broad emission lines in type 2 Seyfert  galaxies implies that these objects have a Hidden Broad Line Region (HBLR henceforth). However, not all Sy 2 galaxies that have been observed in polarized light show the presence of  broad lines. This might imply an intrinsic difference between Seyfert 2 with PBL (PBLs hereafter) and those with  No PBL (NPBLs hereafter). We propose to test this by studying their hard X-ray intrinsic emission in a spectral domain where absorption does not modify significantly the primary emission.\\\\ The existence of Sy 2 without PBL could be due to  the absence of a broad line region  in some Sy 2 galaxies or to the absence of a scattering medium or different physical conditions within the inner regions.  \\\\ To investigate these hypotheses, we intend to describe the general emission properties of various  classes of Seyfert 2 galaxies by computing their average hard X-ray spectrum.\\\\ In D03 we computed the average spectra of Sy 1, Sy 1.5 and Sy 2 galaxies. We found that  the Sy 1 and Sy 2 galaxies  of our sample have a different behavior,  in particular concerning  the presence of a cutoff which is required in the Sy 1 emission and absent in the Sy 2 average spectrum. Another point of disagreement was the role played by  Compton reflection possibly influencing  the Sy 2 emission, unlike Sy 1. \\\\ We  pursue our investigations  within the Seyfert 2 class of our sample to search for common properties between Sy 1 galaxies and the two Sy 2 subclasses. To test the argument that the absorption is the only reason for which we cannot see the HBLR, we consider Compton-thin Sy 2 galaxies. After presenting  the average spectra obtained, we study physical processes responsible for the observed X$\\gamma$ emission. ",
        "conclusions": "Far from having definitively proven that NPBLs have a different central engine hence that they do  not have a hidden Sy 1 nucleus, we show that their primary emission  presents  substantial differences compared to PBLs and Sy 1 galaxies. \\subsection{The high-energy cutoff} Sy 1 and PBLs galaxies  possess  a cutoff of the same order, indicating that the temperature of the electron gas is similar for both classes (T$_{\\textrm{\\tiny{Sy1}}}$=E$_{c}$/k$_{\\textrm{\\tiny{B}}}$=8.27$\\cdot$10$^{8}$ K and T$_{\\textrm{\\tiny{PBLs}}}$=9.88$\\cdot$10$^{8}$ K). The presence of similar cutoffs in the PBLs and Sy 1 emission reveals that the same global physical conditions occur in the inner regions of their source. This supports unified models. \\\\ The absence of a  constrained cutoff in NPBLs emphasizes differences in the primary emission of various classes of Sy 2 galaxies. However, a cutoff could arise at higher energies thus produced by very high electron gas temperatures. We therefore have to extend the energy domain of our investigation. This lack could also be explained by a wide  distribution of NPBLs  cutoff energies. \\\\ The genuine absence of a cutoff would suppose that different emission processes arise within the NPBLs class.\\\\ Nevertheless, the cutoff argument cannot  explain the absence of polarized broad lines in the NPBLs sources. Another way to explain the existence of two Sy 2 types is by the study of the central engine activity,  explored in the next subsection. \\subsection{The central engine activity} We show in this study that  the NPBLs countrate at 15 keV is two times weaker than PBLs (section 4),  the NPBLs spectrum is softer than Sy 1 and PBLs (see 4.3.1 and 4.3.2) and in particular the NPBLs of our sample are less luminous (see 2.3.2). Therefore, the central engine  of NPBLs appears weaker than PBLs. \\\\ We speculate that Sy 2 galaxies without polarized broad lines correspond to a state in which the accretion flow is weak and devoid of  line emitting clouds. This flow may be insufficient  to produce a large UV flux which in turn would explain why Sy 2 without PBL cannot cool through Compton processes between the electron population and the UV photons, thus explaining the absence of cutoff in the high energy emission. \\subsection{The absorption/environment role in Seyfert 2 classification} The absence of  polarized broad lines and of a cutoff in NPBLs galaxies cannot both be  explained by very high column densities. Indeed, we consider only Compton-\\textit{thin} Sy 2 galaxies  thus sources slightly influenced by absorption. The  argument  according  to  which  the  column density of our NPBLs would be underestimated is improbable here, since this would imply a large divergence between the average emission of Sy 1 and Sy 2 without PBL in the 15-40 keV domain, that is not observed.\\\\ The absence  of polarized broad lines in NPBLs may be due to either the lack of a scattering medium above the absorbing gas or to the genuine absence of line emitting clouds in the nuclear regions. \\\\ However, the absence of a scattering medium  alone cannot explain all the differences seen in the continuum emission of the two Sy 2 types. \\subsection{Implications} We summarize in Table 5 the characteristics of   Seyfert 2 galaxies with and without polarized broad lines. \\begin{table*}[!hbt] \\begin{center} \\caption{Emission properties of  PBLs and NPBLs of our sample} \\begin{tabular}[!hbt]{l | l } \\hline \\hline Sy 2 Class & Characteristics  \\\\ \\hline PBLs & Cutoff at $\\sim$ 80 keV \\\\ & lower spectral index than NPBLs \\\\ & Spectral similarities with Sy 1 galaxies  \\\\ \\hline NPBLs & no constrained cutoff  \\\\ & softness of the emission\\\\ & central engine activity weaker than PBLs  \\\\ & not more absorbed than PBLs   \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} The PBLs and  Sy 1 galaxies of our sample have similar physical conditions in the inner regions of their sources. Consequently,  the  observed  differences  in  the line properties and in X-ray emission can be due to a greater absorption in PBLs compared to  Sy 1 galaxies, which agrees with unified models.\\\\ The NPBLs, however,  are not only hidden Sy 1 nucleus since the absorption argument is unable to reproduce the different spectral properties found in this study. To confirm this assertion and to  allow  a meaningful analysis of the central engine signatures of the different types of Sy galaxies, we  need a larger sample of Sy 2 having both spectropolarimetric observations and high quality X-ray data. \\\\ We conclude that while Seyfert 2 galaxies with polarized broad lines have  continuum emission properties that match those of Seyfert 1 galaxies, Seyfert 2 galaxies without  PBL do seem to present genuine differences in their nuclear properties."
    },
    "0311/astro-ph0311576_arXiv.txt": {
        "abstract": " ",
        "introduction": "The Michelson-Morley experiment \\cite{mm} was designed to detect the relative motion of the Earth with respect to a preferred reference frame, the ether, by measuring the shifts of the fringes in an optical interferometer. These shifts, that should have been proportional to the square of the Earth's velocity, were found to be much smaller than expected. Thus, that experiment was taken as an evidence that there is no ether and, as such, represented the essential ingredient for deciding between Lorentzian Relativity and Einstein's Special Relativity.  However, according to some authors, the fringe shifts observed by Michelson and Morley, while certainly smaller than the classical prediction corresponding to the orbital velocity of the Earth, were {\\it not} negligibly small. This point was clearly expressed by Hicks, see page 36 of Ref.\\cite{hicks} \"..the numerical data published in the Michelson-Morley paper, instead of giving a null result, show a distinct evidence of an effect of the kind to be expected\" and also by Miller, see Fig.4 of Ref.\\cite{miller}. In the latter case, Miller's refined analysis of the half-period, second-harmonic effect observed in the original experiment, and in the subsequent ones by Morley and Miller \\cite{morley}, showed that all data were consistent with an effective, {\\it observable} velocity lying in the range 7-10 km/s. For comparison, the Michelson-Morley experiment gave a value $v_{\\rm obs} \\sim 8.8 $ km/s for the noon observations and a value $v_{\\rm obs} \\sim 8.0 $ km/s for the evening observations.  Due to the importance of the issue, we have decided to re-analyze the original data obtained by Michelson and Morley and re-calculate the values of $v_{\\rm obs}$ for their experiment. Our findings completely confirm Miller's indication of an observable velocity $v_{\\rm obs}\\sim 8.4$ km/s in their data.  In addition assuming, as in the pre-relativistic physics, the existence of a preferred reference frame, but using Lorentz transformations to connect with the Earth's reference frame, it turns out that this $v_{\\rm obs}$ corresponds to a {\\it real} Earth's velocity, in the plane of the interferometer, $v_{\\rm earth} \\sim 201 \\pm 12$ km/s. This value, which is remarkably consistent with Miller's kinematically  calculated value $v_{\\rm earth} \\sim 203 \\pm 8$ km/s \\cite{miller}, suggests that the magnitude of the fringe shifts is determined by the typical velocity of the Solar System within our galaxy (and not, for instance, by its velocity $v_{\\rm earth} \\sim 336$ km/s with respect to the centroid of the Local Group).  We emphasize that the use of Lorentz transformations is absolutely crucial. In fact, in this case, differently from the classical predictions, the fringe shifts measured with an interferometer filled with a dielectric medium of refractive index ${\\cal N}_{\\rm medium}$ are proportional to the Fresnel's drag coefficient $1- 1/{\\cal N}^2_{\\rm medium}$. For this reason, a large `kinematical' velocity $\\sim 200$ km/s is seen, in an in-air-operating optical system, as a small `observable' velocity $\\sim 8.4$ km/s. At the same time, without using Lorentz transformation, there was no hope to understand why, for the same value of $v_{\\rm earth}$, the effective $v_{\\rm obs}$ had to be $\\sim 3$ km/s for the Illingworth experiment (performed in an apparatus filled with helium) or $\\sim 1$ km/s for the Joos experiment (performed in an evacuated housing).  \\vskip 10 pt ",
        "conclusions": "In this paper we have presented our re-analysis of the original data obtained by Michelson and Morley \\cite{mm}. Contrary to the generally accepted ideas, but in agreement with the point of view expressed by Hicks in 1902, Miller in 1933 and Munera in 1998, the results of that experiment cannot be considered null. The observed fringe shifts, although smaller than the classical prediction corresponding to the orbital motion of the Earth, point to an effective observable velocity $v_{\\rm obs}=8.4 \\pm 0.5$ km/s. As emphasized at the end of Sect.2 and at the end of Sect.5, this value is exactly the same average value obtained by Miller in his observations at Mt.Wilson (after the critical re-analysis of Shankland et al.\\cite{shankland}).  Therefore, it becomes natural to explore the existence of a preferred reference frame and use Lorentz transformations to extract the real kinematical velocity corresponding to this $v_{\\rm obs}$. In this case, we find a value (in the plane of the interferometer) $v_{\\rm earth}=201 \\pm 12 $ km/s. This value is in excellent agreement with the value that was kinematically calculated by Miller on the base of the observed variations of the ether-drift effect in different epochs of the year. As a consequence, we conclude that the magnitude of the fringe shifts is determined by the typical velocity of the Solar System within our galaxy. This conclusion is also consistent with the most recent data for the anisotropy of the two-way speed of light in the vacuum.  It is somewhat surprising that the crucial role of Lorentz transformations has been overlooked for so long time in the analysis of the Michelson-Morley experiment. This is, probably, due to the wrong impression that there should be very little difference from the classical predictions since the velocities in the game (30, 200, 300,..) km/s are so much smaller than the speed of light. The point is, however, that Lorentz transformations preserve exactly the value of the speed of light in the vacuum. Therefore, in a medium any anisotropy has to be proportional to the Fresnel drag coefficient $k_{\\rm medium} \\ll 1$. For an in-air-operating optical interferometer, this effect is responsible for the huge reduction of the Earth's velocity from its relevant kinematical value $v_{\\rm earth}\\sim 200 $ km/s down to its observable value $v_{\\rm obs}\\sim 8.4$ km/s governing the magnitude of the fringe shifts in the Michelson-Morley and Miller's experiments. At the same time, without using Lorentz transformation, there was no hope to understand why, for the same value of $v_{\\rm earth}$, the effective $v_{\\rm obs}$ had to be $\\sim 3$ km/s for the Illingworth experiment  or $\\sim 1$ km/s for the Joos experiment.  This set of puzzling experimental results can now be described within a single consistent framework. Therefore, we conclude that an alternative interpretation of the Michelson-Morley type of experiments is now emerging. Contrary to the point of view that was generally accepted in the past century, they provide experimental evidence for the existence of a preferred reference frame.  \\vskip 60 pt"
    },
    "0311/astro-ph0311556.txt": {
        "abstract": "{We present results from the \\xmm\\ observation of the binary cluster A1750 at $z = 0.086$. We have performed a detailed study of the surface brightness, temperature and entropy distribution and confirm that the two main clusters of the system (A1750 N and A1750 C) have just started to interact. From the temperature distribution, we calculate that they are likely to merge sometime in the next 1 Gyr.  The more massive cluster, A1750 C, displays a more complicated temperature structure than expected. We detect a hot region associated with a density jump $\\sim$ 450 kpc east of the cluster centre, which appears to be a shock wave. This shock is not connected with the binary merger, but it is {\\em intrinsic} to A1750 C itself. From simple physical arguments and comparison with numerical simulations, we argue that this shock is related to a merging event that A1750 C has suffered in the past 1-2 Gyr. The larger scale structure around A1750 suggests that the system belongs to a rich supercluster, which would presumably increase the likelihood of merger events. These new \\xmm\\ data thus show that A1750 is a complex system, where two clusters are starting to interact before having re-established equilibrium after a previous merger. This merger within a merger indicates that  the present day morphology of clusters may depend not only on on-going interactions or the last major merging event, but also on the more ancient merger history, especially in dense environments. ",
        "introduction": "Clusters of galaxies are the largest gravitationally bound structures in the Universe, and because of their observable properties and  sheer number, they represent invaluable cosmological probes.  The hierarchical scenario of structure formation in the Universe predicts the  growth of clusters of galaxies through accretion of smaller units by gravitational infall and mergers. During a merger event more than 10$^{63}$ ergs are dissipated in the intra-cluster medium (ICM) by shock heating, leading to strong variations of physical characteristics such as the temperature, density and entropy. Numerical N-body simulations have shown that mergers produce substructure which is detected in both the density and temperature distributions (Schindler \\& M\\\"uller~\\cite{schindlermuller93}; Roettiger et al~\\cite{roettiger97}; Ricker~\\cite{ricker98}). The temperature substructures survive for $\\sim 4$--6 times longer than density substructures, thus the temperature distribution is a strong indicator of the cluster history and dynamical state. X-ray spectro-imaging observations are well suited to investigate in a deeper detail the observable characteristics of the merger event.  Density substructure in X-ray observations of some low-redshift clusters (Forman \\& Jones~\\cite{formanjones82}; Mohr, Fabricant \\& Geller~\\cite{MFG93}), together with the detection of other substructure in optical observations (Escalera et al. \\cite{escalera}), suggests that these clusters are dynamically young (Buote \\& Tsai \\cite{BuTs})  The first quantitative X-ray temperature maps were obtained with ASCA (e.g., Markevitch et al~\\cite{marketal98}) but, as with the later BeppoSax observations (e.g., De Grandi \\& Molendi~\\cite{DeGM99}), resolution was low due to the energy dependent Point Spread Function (PSF) of these telescopes. Thanks to the high effective area and spatial resolution now available from \\xmm\\ and Chandra, precise spatially resolved temperature maps are now possible. Results confirm the increase in temperature in the merging regions (e.g., Arnaud et al. \\cite{arnaud01a}; Markevitch \\& Vihklinin \\cite{MarkVihk01}; Markevitch et al.~\\cite{mark02}; Neumann et al.~\\cite{neumann03}) but the new observations have also detected new phenomena, such as the cold fronts (Markevitch et al. \\cite{mark00}).  Substructures in clusters contain a fossil record of the merger history. Statistical studies of  cluster morphology can thus provide an important test of cosmological models of structure formation.  In practise, attempts to compare the prediction of numerical simulations of structure formation to quantitative statistical studies of nearby cluster morphology have been hampered by our current poor knowledge of the effect of cluster growth on morphology. Despite the large amount of data available, the dynamical evolution of the ICM and the relation between the galaxies and gas during a merger event is still poorly understood (see Buote \\cite{buote02} for a review).  A better understanding of the physics of merger events, in particular the relaxation time of substructures (Nakamura \\cite{nakamura95}) is required.  Hence a systematic study of merging clusters of galaxies would help a great deal in deepening our understanding of the process of cluster formation and evolution. Our \\xmm\\ Guaranteed  Time (GT) time program was established for this purpose. We selected a small sample of clusters displaying the signature of substructure (on the basis of ROSAT imaging analysis), and which could be considered as being in different {\\em epochs} of a merger event.  In this paper we present the first object in our sample: the binary cluster A1750. In optical, A1750 shows a multi-peaked galaxy distribution (Beers et al. ~\\cite{beersetal91}; Ramirez \\& Quintana ~\\cite{RQ90}; Donnelly et al.~\\cite{donnelly01}); the two major peaks have a radial velocity difference of $\\sim 1300$ km s$^{-1}$. A1750 was observed in X-ray with {\\em Einstein} (Forman et al. \\cite{forman81}), and subsequently identified as a canonical binary cluster, though three peaks are clearly visible in the {\\em Einstein} X-ray image. The cluster was later observed with ROSAT  and ASCA (Novicki et al.~\\cite{Noveetal98}; Donnelly et al.~\\cite{donnelly01}); three peaks are also visible in the $2^{\\circ}$ ROSAT image. The latter authors combine ROSAT and ASCA observations, as well as the galaxy distribution, giving the first kinematic and dynamical description of the cluster. They detected an enhancement of order $30\\%$ in the temperature between the two main peaks in the X-ray emission, and suggested that a shock region is developing in the gas by compression. However, Donnelly et al.~(\\cite{donnelly01}) concluded that the two clusters could effectively be considered as isolated objects, as the hot gas in the interacting region did not have a strong impact in the global temperature estimate, and its contribution to the total density was weak. Donnelly et al.'s mass analysis indicated a mass ratio of order 1.3, but their dynamical analysis was not conclusive as to whether or not the system is bound.  In this paper we use the high sensitivity and spatial resolution of \\xmm\\ to gain new insights into this system. Taking advantage of the large field of view of \\xmm, we observed the two main subclusters in a single pointing. The exceptional sensitivity of \\xmm\\ allows us to produce temperature and entropy maps with an accuracy which considerably surpasses previous attempts. We detect new features in all maps and relate these to the dynamical state of the system, arguing that, while the clusters are just beginning to interact, at least one of them (A1750 C), and possibly both, may not yet have relaxed from  a previous merger event.  Throughout the paper we assume $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$, q=0.5. In this cosmology 1 arcmin corresponds to 130 kpc at the mean redshift of the cluster ($z=0.086$).  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "We have reported the \\xmm\\  GT observation of the merging cluster A1750. The main conclusions of this work may be summarised as follows:  \\begin{itemize}  \\item In the morphological analysis with 2D $\\beta$-models, we detect excess and off-centre emission in the cores of both A1750 N and C. We do not find any significant substructure in the region between the two clusters after subtraction of the $\\beta$-models.  However, on close examination of the surface brightness in the region between the two clusters, we detect an increase of the same order as expected if the region between the clusters is a weak shock region (Sect.~\\ref{sec:current}). We further detect excess residual emission to the north of A1750 N, which is likely (from comparison the \\xmm\\ and PSPC images) related to the larger scale structure of the system.  \\item We have produced temperature maps by applying two different algorithms. We  measured a temperature increase of order $30\\%$ in the region between A1750 N and C. A1750 N exhibits  a relatively smooth, uniform temperature distribution, but there are significant temperature variations within A1750 C, which appear intrinsic to A1750 C and not connected to the merger between the two clusters.  \\item We use the 2D information to select interesting regions for a 1D analysis. For the regions corresponding to the relaxed, isothermal-looking parts of each cluster (R1 and R2), the 1D and 2D \\betamod\\ parameters are in excellent agreement with each other and with previous ROSAT analysis, although the \\xmm\\ data require an additional central component properly to fit the profiles. Both temperature profiles appear to decline modestly ($< 10\\%$) with radius.  \\item The density profile of the disturbed side of A1750 C (R3) is better described by a $\\beta$-model and a power law, separated by a jump of order $20\\%$. The temperature profile of this region is very irregular, but appears constant across the jump.  \\item Compared to the scaled entropy profiles of relaxed clusters, the scaled entropy profiles of both regions of A1750 C exhibit a high scaled entropy value in the centre, similar to what is observed in merging cluster simulations. Only beyond $\\sim 0.2 \\rv$ does the profile of R2 attain the $S\\propto r^{1.1}$ expected from shock heating.  \\item In the assumption of an head-on merger between A1750 C and N we calculated the Mach number and velocity of the merger from the measured temperatures. With the resulting values (${\\cal M}$ = 1.64 and $v_{merg} \\sim$1400 km s$^{-1}$), we estimate that the two clusters will reach core passage in --- roughly --- less than 1 Gyr.  \\item By applying the Rankine-Hugoniot jump condition to the measured density jump in A1750 C (R3) we confirm the detection of a weak shock of ${\\cal M}$ = 1.2, in agreement with the measured temperature of the pre-shock and post-shock gas. Under simplified assumptions and by comparison with numerical simulations, we estimate that a merger --- intrinsic to A1750 C --- has occurred sometime in the past 1-2 Gyrs.  \\item An inspection of the larger scale X-ray image suggests that A1750 N and C are the two main clusters within a filament which includes at least another small cluster at the same redshift of A1750 C (A1750 S). At still larger scale, A1750 lies in a filamentary supercluster containing, in total, 13 clusters.  \\end{itemize}  We thus confirm that A1750 is a merging cluster in an early phase, when the two units have just started to interact and stand at a distance slightly lower than their virial radii. However, the global dynamical status of the system is far more complicated than expected. We conclude that the main cluster (A1750 C) has already suffered a previous merger and is now in the phase of re-establishing equilibrium. This phase will be interrupted by the current merger with A1750 N.  As far as we know, this is the first time such complex dynamical signatures have been observed in clusters at this merger phase.  The present day morphology of clusters may thus depend not only on on-going mergers (or the last merger) but also on the more ancient merging history, especially in dense environments. This has to be taken into account in the interpretation of statistical studies of cluster morphology."
    },
    "0311/astro-ph0311210_arXiv.txt": {
        "abstract": "In the course of a comprehensive mm/submm survey of massive star-forming regions, the vicinities of a sample of 47 luminous IRAS sources were closely investigated with SCUBA and IRAM bolometers in order to search for massive protostellar candidates. A particularly interesting object has been found in the surroundings of the bright FIR source IRAS 07029-1215. Follow-up line observations show that the object is cold, that it has a massive envelope, and that it is associated with an energetic molecular outflow. No infrared point source has been detected at its position. Therefore, it is a very good candidate for a member of the long searched-for group of massive protostars. ",
        "introduction": "The study of the earliest stages of massive star formation (MSF) is one of the most exciting subjects of present research in this field. Whether the dominant formation process for massive stars is disk accretion (see e.g. \\citealp{jia96}) or coalescence \\citep{bon98} is still an open question. The presence of outflows in regions of massive star formation favours, however, the first scenario. Similarly to previous surveys (e.g. \\citealp{beu02b}, \\citealp{mue02} or \\citealp{hun00}), we have searched for MSF candidates located near bright IRAS sources using SCUBA and IRAM bolometers (\\citealp{pos03} and Klein et al. 2003, {\\sl in prep.}). Our focus was on objects unrelated to the IRAS sources themselves, looking for the earliest evolutionary stages not yet seen in the near- to mid-infrared wavelength regime. Even if we do not yet know what the observational features of massive protostars are \\citep{eva02}, their invisibility at NIR/MIR wavelengths is a defining criterion. Near IRAS~07029-1215, an object with an IRAS luminosity of $L(IRAS)=1700$~L$_{\\odot}$ \\citep{hen92} at its distance of $d=1$~kpc, we discovered a deeply embedded object (UYSO~1, for \\textsl{Unidentified Young Stellar Object~1}), powering a high-velocity bipolar CO outflow. The IRAS source is located in the bright \\HII{} region S~297, illuminated by the B1II/III star HD~53623 \\citep{hou88}. UYSO~1 is located very close to the edge of a dark cloud. ",
        "conclusions": "The emerging picture of UYSO~1 drawn from both mass estimates and subsequent empirical relations as well as SED considerations thus is that of an early B star surrounded by an envelope of about 30-40~M$_{\\odot}$. The object seems to be in a particularly early evolutionary state, but already drives a bipolar outflow with a total mass of M$_{\\rm outflow} = 5.4$~M$_{\\odot}$. To our knowledge, there are only very few comparable sources that are undetected at infrared wavelengths (see e.g. \\citealp{hun98}, \\citealp{bdm99} and \\citealp{fon03}). With additional high-resolution observations, we hope to be able to specify these conclusions and learn more about this evolutionary state. Possibly, parameters of a putative disk around UYSO~1 could be derived then."
    },
    "0311/astro-ph0311026_arXiv.txt": {
        "abstract": "{We analyse high-resolution VLT+UVES spectra of the low-ionization intrinsic absorber observed in the BAL QSO SDSS~J001130.56$+$005550.7.  Two narrow absorption systems at velocities $-$600 km s$^{-1}$ and $-$22000 km s$^{-1}$ are detected. The low-velocity system is part of the broad absorption line (BAL), while the high-velocity one is well detached.  While most narrow absorption components are only detected in the high-ionization species, the lowest velocity component is detected in both high- and low-ionization species, including in the excited \\siiie\\ and \\ciie\\ lines. From the analysis of doublet lines, we find that the narrow absorption lines at the low-velocity end of the BAL trough are completely saturated but do not reach zero flux, their profiles being dominated by a velocity-dependent covering factor. The covering factor is significantly smaller for \\mgii\\ than for \\siiv\\ and \\nv , which demonstrates the intrinsic nature of absorber. From the analysis of the excited \\siiie\\ and \\ciie\\ lines in the lowest velocity component, we find an electron density $\\simeq$ 10$^{3}$ cm$^{-3}$. Assuming photoionization equilibrium, we derive a distance $\\simeq$ 20 kpc between the low-ionization region and the quasar core.  The correspondence in velocity of the high- and low-ionization features suggests that all these species must be closely associated, hence formed at the same distance of $\\sim$ 20 kpc, much higher than the distance usually assumed for BAL absorbers. ",
        "introduction": "Intrinsic absorption lines in quasars are usually classified as broad absorption lines (BALs) or narrow absorption lines (NALs). These absorption line systems are to be distinguished from cosmologically ``intervening'' systems unrelated to the quasar environment (Barlow et al. \\cite{bar1997b}).  Broad (velocity width FWHM $>$ 2000 \\ks ) troughs (BALs) are detected in roughly 15\\% of optically selected quasars (Hewett \\& Foltz \\cite{hf2003}, Reichard et al.  \\cite{rei2003b}).  They are blueshifted with respect to the QSO emission lines.  BAL outflows occur at velocities of typically 0.1$c$ (Weymann et al. \\cite{wey1991}). Most BAL QSOs have absorption in high-ionization species like \\civ\\ $\\lambda$1549, \\siiv\\ $\\lambda$1397 and \\nv\\ $\\lambda$1240. A minority of them also show absorption due to lower ionization species (LoBAL) such as \\mgii\\ $\\lambda$2798 or \\aliii\\ $\\lambda$1857.  NALs have velocity widths of at most a few hundred \\ks . NALs are not only observed at redshifts $z_{\\rm abs} \\simeq z_{\\rm em}$ but also at blueshifted velocities comparable to those seen in BAL QSOs (Barlow et al. \\cite{bar1997b}, Hamann et al. \\cite{ham1997a}). NALs can also appear redshifted up to $\\sim$ 2000 \\ks even though they are frequently blueshifted.  Since common doublet transitions are resolved in NALs using high-resolution spectroscopy, they constitute useful diagnostics of the quasar environment.  The so-called mini-BALs (e.g. Churchill et al. \\cite{chu1999}) have intermediate absorption widths, i.e. FWHM between a few hundred and 2000 \\ks .  BALs and NALs indicate that outflows from Active Galactic Nuclei span a large range of velocity widths.  While it is known that NALs may form in various environments (Hamann et al. \\cite{ham2001} and references therein), it is not clear whether at least some of them are directly related to the BAL phenomenon. In some rare cases, NALs and BALs may be observed in the same quasar, providing an opportunity to directly investigate this issue.  In this paper we report high-resolution spectroscopy of SDSS~J001130.56$+$005550.7 (hereafter \\sd ; Schneider et al. \\cite{sch2002}) discovered in the Sloan Digital Sky Survey (York et al. \\cite{yor2000}).  This object is one of the $\\sim$10$^5$ quasar candidates (Richards et al. \\cite{ric2002}) for which the survey is obtaining redshifts, in addition to the $\\sim10^6$ galaxies which comprise the bulk of the spectroscopic targets (Blanton et al. \\cite{bla2003}), selected from astrometrically calibrated drift-scanned imaging data (Gunn et al. \\cite{gun1998}, Pier et al. \\cite{pie2003}) on the SDSS $ugriz$ AB asinh magnitude system (Fukugita et al. \\cite{fuk1996}, Lupton et al. \\cite{lgs1999}, Hogg et al. \\cite{hog2001}, Stoughton et al. \\cite{sto2002}, Smith et al. \\cite{smi2002}).  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics*{H4512f1.eps}} \\resizebox{\\hsize}{!}{\\includegraphics*{H4512f2.eps}} \\caption{SDSS spectrum ($R \\simeq 2000$) of \\sd .Top: the full spectrum. $F_{\\lambda}$ is in units of 10$^{-17}$ ergs cm$^{-2}$ s$^{-1}$ \\AA$^{-1}$. Bottom: a portion of the spectrum illustrating the \\siiv\\ $\\lambda\\lambda$ 1393.75, 1402.77 and \\civ\\ $\\lambda\\lambda$ 1548.20, 1550.78 doublets. Absorption lines at $z$ = 2.29071 (part of system A) and at $z$ = 2.05844 (part of system B) are indicated} \\label{fig:sdss} \\end{figure}  \\sd\\ has strong intrinsic absorption extending up to a velocity of $\\sim$5500 \\ks\\ in \\civ .  It just misses formal classification as a BAL QSO: the balnicity index defined by Weymann et al. (\\cite{wey1991}) is measured to be zero using the method of Reichard et al.  (\\cite{rei2003a}).  However, it has a value of 2040 \\ks\\ on the absorption index scale designed by Hall et al. (\\cite{hal2002}) to include troughs too narrow or close to the quasar redshift for consideration by the balnicity index.  Since the formal classification is rather arbitrary and non-physical, we consider \\sd\\ as a BAL QSO in the following.  \\sd\\ shows both high- and low-ionization absorptions resolved into narrow components, together with a well-detached high-velocity $z_{\\rm abs} \\ll z_{\\rm em}$ narrow line system. Because some absorption arises from excited levels, \\sd\\ also provides a rare opportunity to derive the electron gas density and to constrain the distance to the absorber.  Observations are reported in Sect. \\ref{sec:Obs} and the characteristics of the spectrum in Sect. \\ref{sec:Over}. The narrow lines are analysed in Sect. \\ref{sec:Ana} in order to establish their intrinsic nature, to evaluate the covering factor of the different ions, and to estimate the electron density in the flow using excited lines. Discussion and conclusions form the last section. ",
        "conclusions": "\\label{sec:Dis}  One of the main results of our study is that the narrow absorption components seen at the low-velocity end of the BAL troughs are completely saturated, their profiles being dominated by the covering factor. Velocity and ionization dependent partial covering is often seen in NALs (Barlow \\& Sargent \\cite{bar1997a}, Hamann et al. \\cite{ham2001}), and has also been reported for mini-BAL and BAL QSOs (Arav et al. \\cite{ara1999b}, Srianand \\& Petitjean \\cite{sri2000}, de Kool et al. \\cite{dek2001}, \\cite{dek2002b}, Hall et al. \\cite{hal2003}).  The partial covering proves the intrinsic nature of the absorbers.  However, our most intriguing result is the distance between the ionizing source and the absorber that we derived on the basis of the excited lines.  BALs are usually thought to be formed much closer to the central engine, at distances $\\sim$ 0.1 -- 1 pc, roughly four to five orders of magnitude smaller than the distance of $\\sim$ 20 kpc we have measured. Our estimate is clearly uncertain, but the errors cannot explain such a large difference.  Note that the \\ciie\\ and \\siiie\\ must be related to the other absorption features.  The correspondence in velocity of the high- and low-ionization features in component A1 (Figs.~\\ref{fig:vel} and \\ref{fig:velz}) clearly indicate that all species must form at roughly the same location.  Also, the continuity of the opacity and covering factor properties towards the higher velocity components A2 and A3 (Fig. \\ref{fig:cova}) suggests a common formation for these components too.  The components of the absorbing system A are definitely part of the BAL outflow because of their association with the high-ionization features seen in the wider \\civ\\ BAL (cf. the SDSS low-resolution spectrum in Fig.~\\ref{fig:sdss}).  It is therefore likely that the BAL formation region (BALR) in \\sd\\ --- at least for component A --- is at much higher distance than commonly thought.  In fact, small BALR distances are based more on theoretical considerations than on direct observational evidence. Here we summarize the few direct estimates that have been made of the distances to confirmed intrinsic outflows; all distances have been converted to our cosmology (Sect. 4.2).  Two NALs that have undergone time variability have thereby confirmed themselves as intrinsic and yielded upper limits to their distances in the kiloparsec range: $<$ 1 kpc in UM~675 (Hamann et al. \\cite{ham1997b}), and $<$ 2 kpc in QSO 2343+125 (Hamann et al. \\cite{ham1997c}).  Large distances have been inferred from the presence of \\ciie\\ or \\siiie\\ absorption in two other NALs which are known to be intrinsic absorbers due to partial covering: 20 kpc for 3C~191 (Hamann et al. \\cite{ham2001})\\footnote{Note that 3C 191 is a radio-loud quasar, whereas \\sd\\ is radio-quiet by virtue of its non-detection in the FIRST survey (Becker et al. \\cite{bwh1995}). Absorbers apparently can be found at large distances from the ionizing source in both RLQs and RQQs.},  and $>$ 570 pc for APM 08279+5255 (Srianand \\& Petitjean \\cite{sri2000}).  A related but distinct technique is the use of numerous \\ion{Fe}{ii} absorption lines to estimate absorber distances.  Using this method, $r$ $\\sim$ 310 pc has been measured for the low-velocity system in FIRST J084044.5+363328, whereas $r$ $\\sim$ 1.35 pc is found for the high-velocity system in the same object using different diagnostic (de Kool et al. \\cite{dek2002b}).  However, analysis of the similar NAL in FIRST J121442.3+280329 yielded a small distance range $\\sim$ 1 -- 30 pc for the entire outflow (de Kool et al. \\cite{dek2002a}), similar to that found for LBQS 0059-2735 (Wampler et al. \\cite{wam1995}). The largest distance found by this method has been for the low-ionization BAL in FIRST J104459.6+365605 (de Kool et al.  \\cite{dek2001}).  On the basis of excited \\ion{Fe}{ii} lines and velocity correspondence between \\ion{Fe}{ii} and \\mgii\\ lines, $r$ $\\sim$ 630 pc was found.  We emphasize that -- again based on velocity correspondence -- our high distance estimate in \\sd\\ also applies to the high-ionization absorbers and not only to the low-ionization ones, as found in FIRST J104459.6+365605.  However, Everett et al. (\\cite{eka2002}) model a multiphase outflow at only $r$ $\\sim$ 4 pc that can reproduce the observations of FIRST J104459.6+365605, including the density measurements used by de Kool et al. to infer a much larger distance.  In their model the intrinsic outflow does not have a constant density, but instead consists of a relatively low-density, high-ionization wind with embedded higher-density, lower-ionization clouds.  Absorption in the high-ionization wind modifies the spectrum seen by the outer region of the wind and the low-ionization clouds, effectively reducing the value of $L_{\\rm LL}$ in Eq.~(4) and thereby decreasing the inferred distance. The ionization and density in the wind decrease with increasing distance until species such as \\ion{Fe}{ii} and \\mgii\\ are present at the required $n_e$.  Denser clouds embedded in the flow at that distance are invoked to produce \\ion{Mg}{i} absorption at the same velocities as \\ion{Fe}{ii} and \\mgii .  Everett et al. claim that a multiphase outflow at a small distance from the ionizing source could also explain the absorption in 3C 191. Hamann et al. do in fact state that the outflow in 3C 191 must span a range of densities (or distances) to explain the presence of \\nv\\ and \\ion{Mg}{i} at the same velocities.  A range of densities could also be present in \\sd\\ but is not required, because \\ion{Mg}{i} is not detected.  In any case, the density diagnostics used to infer large absorber distances in \\sd\\ and 3C 191 are different than those used in FIRST J104459.6+365605.  Therefore, detailed modeling should be done to determine whether or not a multiphase outflow at small distances can indeed reproduce the observed densities, column densities and velocity structure of all ions observed in these objects.  A further observational test of a multiphase model for \\sd\\ and 3C 191 could come from spectra extending shortward of Ly$\\alpha$ past the Lyman limit, directly measuring $L_{\\rm LL}$ and constraining any modification of the spectrum by the inner wind.  It would also be worth monitoring objects with large inferred distances, to search for time variable NALs which could require much smaller distances.  If they are confirmed, large distances from the continuum source and the broad emission line region may be difficult to reconcile with the evidence of velocity and ionization dependent partial covering in these absorbers.  On the one hand, partial covering implies only that the projected size scale of the absorber is less than or comparable to that of the emitting region.  {\\em That inferred size for the absorber is independent of the distance from the source,} as long as the distance is small compared to the angular diameter distance to the quasar.  Partial covering by outflows at kpc-scale distances is therefore possible, at least in principle.  However, it is difficult to understand how such distant outflows could be common, for several reasons.  The usual problems of cloud (or density inhomogeneity) survival and confinement (e.g. Hamann et al. \\cite{ham2001}, de Kool et al.  \\cite{dek2001}) are exacerbated by the requirement that such structures must survive for the time needed to reach such large distances.  Also, if such outflows are seen in $\\sim$10\\% of quasars, they cover somewhere between $\\sim$10\\% of the unobscured lines of sight around all quasars and all such sightlines in $\\sim$10\\% of quasars at distances of $\\sim$10~kpc, implying outflows of very large masses ($\\sim 10^8-10^9 M_{\\sun}$; Hamann et al. \\cite{ham2001}).  It is worth noting that the observed partial covering may in fact be due to additional emission from an extended region of scatterers comparable in size to the BALR.  Such a region has already been suggested to explain spectropolarimetric measurements (Cohen et al. \\cite{coh1995}, Goodrich \\& Miller \\cite{goo1995}) but its size and location have never been directly measured.  Resonance scattering in a roughly axially symmetric BALR could also explain abnormal doublet line ratios usually interpreted as partial covering (e.g. Branch et al. \\cite{bra2002}).  Spectropolarimetry could provide some tests of these hypotheses.  Whether our results are generic to BAL QSOs is not clear. The objects for which BALR sizes have been measured so far show absorption somewhat intermediate between mini-BALs and BALs, rather than the very wide and deep troughs often seen in BAL QSOs. Perhaps these intermediate objects have different outflows, or represent an older evolutionary stage of the BAL QSO phenomenon when most of the material has dissipated, preferentially leaving behind denser clumps.  Clearly these results are puzzling and raise many questions.  To determine the range of distances spanned by intrinsic outflows will require detailed study of other, similar objects with \\ciie\\ and \\siiie\\ absorption, including photoionization modeling to determine if multiphase models can explain such outflows."
    },
    "0311/astro-ph0311356_arXiv.txt": {
        "abstract": "We present the discovery of a high redshift, X-ray selected AGN by the Calan-Yale Deep Extragalactic Research (CYDER) survey: CXOCY J033716.7-050153, located at $z=4.61$, the second high redshift AGN discovered by this survey. Here, we present its optical, near-IR and X-ray properties and compare it with other optical and X-ray selected high redshift AGN. The optical luminosity of this object is significantly lower than most optically selected high redshift quasars. It also has a lower rest frame UV to X-ray emission ratio than most known quasars at this redshift. This mild deviation can be explained either by dust obscuring the UV radiation of a normal radio quiet AGN emitting at $10\\%$ of its Eddington luminosity or because this is intrinsically a low luminosity radio loud AGN, with a super-massive Black Hole of $\\sim 10^8M_\\sun$ emitting at 1\\% of its Eddington luminosity. Deep radio observations can discriminate between these two hypotheses. ",
        "introduction": "Due to its large intrinsic luminosity, an AGN (Active Galactic Nucleus) can be observed up to very high redshifts, and AGNs can therefore be used to probe the early epochs of the universe when the formation of large structures began. Recent observations have found a considerable number of quasars at high redshift \\citep{anderson01,schneider02}, with some of them at redshift $>6$ \\citep{fan03}, when the universe was about 1 Gyr old.  While AGNs were first noticed to be a class of objects worthy of detailed follow-up in other bands on the basis of their radio emission, e.g. \\citep{schmidt63}, more recent work has focused on surveys at shorter wavelengths.  Due to their low spatial density, most AGNs at high redshift (defined here to be $z>4$) were discovered by large area, shallow optical surveys. An example of this is the Sloan Digital Sky Survey (SDSS; \\citealt{york00}), which in its first data release \\citep{schneider03} presented 238 high redshift quasars. However, these quasars were selected based on their optical properties, and optical surveys are subject to bias against the discovery of quasars which are either intrinsically dim at these wavelengths or are obscured by dust like the type 2 AGN that are believed to comprise most of the AGN background \\citep{hasinger02}.  The other main AGN selection technique is X-ray detection, as X-ray emission appears to be a universal characteristic of AGNs at all observed redshifts \\citep{kaspi00}. However, given the relative faintness of the majority of the sources, X-ray surveys need to rely on optical spectroscopy for most of the identifications. With new X-ray missions like the Chandra X-ray Observatory and its superb angular resolution, the identification of X-ray sources at other wavelengths has become much easier, allowing for successful multiwavelength follow-up of X-ray sources. The nature of the X-ray emission also ensures that the optical selection bias due to obscuration is strongly reduced (although the X-ray selection may introduce other biases). This is especially true at high redshift, where the observed X-ray photons were emitted at higher energies and can penetrate even considerable amounts of obscuring material.  According to a recent review by \\citet{brandt03} only seven of the high redshift AGNs known were discovered by their X-ray emission. Three of them were found with ROSAT  \\citep{henry94,zickgraf97,schneider98}, while the other four were detected by Chandra \\citep{silverman02,barger02,castander03a}. One of them was detected in the radio and classified as radio loud \\citep{zickgraf97}.  In this paper, we present the discovery of the eighth high redshift X-ray selected AGN, and the second discovered by the Calan Yale Deep Extragalactic Research (CYDER) survey, at a redshift $z=4.61$. The CYDER survey is a collaborative effort between Universidad de Chile and Yale University to study faint stellar and extragalactic populations in detail. One of the key aims is the characterization of the population of faint X-ray sources. In order to do this, some of the fields of the survey were selected to overlap with moderately deep ($\\sim$ 50 ks) Chandra pointings. Multiwavelength follow-up in these fields includes optical imaging using the CTIO 4m telescope, near IR images obtained using the DuPont telescope at Las Campanas Observatory (LCO) and optical spectroscopy from VLT at Cerro Paranal and Magellan at LCO. A more complete review of the optical and spectroscopic properties of the X-ray sources in the first two fields can be found in \\citet{castander03b}.   In \\S 2 we describe the X-ray, optical and near-IR observations that led to the discovery of the AGN presented here.  In \\S 3 we discuss its observed properties and compare it with other high redshift quasars, both optical and X-ray selected.  We present our conclusions in \\S 4. Throughout this paper we assume $H_0=70h_{70}$ km s $^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.3$ and $\\Omega_\\Lambda =0.7$. These cosmological parameters are consistent with the recent findings by \\citet{spergel03}.  We define the photon index $\\Gamma$ as the exponent giving a photon flux density X-ray spectrum $dN/dE\\propto E^{-\\Gamma}$ in photons cm$^{-2}$ s$^{-1}$ keV$^{-1}$. ",
        "conclusions": "We present the discovery of the CXOCY J033716.7-050153, the second high redshift X-ray selected AGN discovered by the CYDER survey.  This is a faint object in the optical. In fact there are only two AGNs fainter at these high redshifts. The X-ray flux relative to its optical emission is unusual, having a larger value of $\\alpha_{ox}$ than the typical AGN at high redshift. There are two effects that could bring this object into better accordance with the observed correlations for the general high redshift population:  The first possibility is that this is a radio loud low luminosity AGN with a Black Hole of $\\sim 3\\times 10^8M_\\sun$ , emitting at $\\sim 1\\%$ of its Eddington Luminosity. In this case the accretion disk is cooler than a typical AGN at high redshift, making the optical colors look redder.  Also, the high value of $\\alpha_{ox}$, even taking into account the correlation with optical luminosity, can be explained by the presence of second emission component, probably coronal, which contributes to the X-ray flux measurement.  A second possibility is that this is a radio quiet AGN emitting at an Eddington ratio of $\\sim 10\\%$, in which the optical emisssion is obscured by a moderate amount of dust which would imply a neutral hydrogen column density of $\\sim 10^{21}$ cm$^{-2}$, causing the observed optical colors to be redder than the normal AGN. If this is the case, the observed value of $\\alpha_{ox}$ is affected by the obscuration of the optical light, and after correcting for it, the corrected value agrees very well with the measurements for other high redshift AGNs, including the observed correlation between $f_{X}$ and $f_{UV}$.  Detection of a flat radio spectrum in this source above about 10 $\\mu$Jy would confirm the Low Luminosity AGN hypothesis, while a non-detection in the radio with a sensitivity limit of a few $\\mu$Jy would strongly support the radio quiet, mildly obscured AGN hypothesis."
    },
    "0311/astro-ph0311483_arXiv.txt": {
        "abstract": "I review the present status of ground-based $\\gamma$-ray astronomy, concentrating on the population of galactic TeV sources. A number of new telescope systems are now being completed, and promise to yield exciting new discoveries, expanding rapidly the number of sources. The TeV galactic sources today include a number of plerions, shell-type SNR, an X-ray binary, and also one unidentified candidate. Their present status, and our understanding of their TeV $\\gamma$-ray emission processes are summarised and some motivation driving development of the field is outlined. ",
        "introduction": "Ground-based $\\gamma$-ray astronomy (at energies $E>100$ GeV) has been developing rapidly in the last 10 years. The advent of a number of {\\em next-generation} instruments in 2003/2004 promises exciting new results and answers to some long-standing questions. Strongly motivating the development of new detectors and telescopes, are the questions surrounding the acceleration of particles to multi-TeV energies in many galactic environments. It has been known for some time that various evolutionary paths following a supernova explosion can lead to the production of high energy cosmic ray (CR) electrons and hadrons to TeV energies and above, and with them, $\\gamma$-radiation as the most accessible {\\em tracer} of this process. The main motivation driving the search for galactic TeV sources has been concerned with the origin of these CR. In particular the shell type SNR have long been thought responsible for the hadronic CR up to the 'knee' energy $E\\sim 10^{15}$eV. From the first convincing discovery of TeV $\\gamma$-rays from the Crab plerion $\\sim$20 years ago (Weekes et al. 1989), galactic objects now comprise a large fraction of catalogued TeV $\\gamma$-ray emitters.  In this review, I will outline the present and future developments in ground-based $\\gamma$-ray astronomy, emphasising the imaging atmospheric \\C technique, and also review the various galactic objects identified as sources of TeV $\\gamma$-rays. Some present and future scientific questions are also posed. Other reviews of the field, some including summaries of experiments using other techniques to detect TeV $\\gamma$-rays may be found in Aharonian 1999, Weekes 2000, Rowell 2001 and Ong 2003. ",
        "conclusions": ""
    },
    "0311/astro-ph0311174_arXiv.txt": {
        "abstract": "We present results from {\\sl Chandra} and {\\sl XMM-Newton} X-ray observations of NGC\\,4214, a nearby dwarf starburst galaxy containing several young regions of very active star-formation. Starburst regions are known to be associated with diffuse X-ray emission, and in this case the X-ray emission from the galaxy shows an interesting morphological structure within the galaxy, clearly associated with the central regions of active star-formation. Of the two main regions of star formation in this galaxy, X-ray emission associated with the older is identified whereas little is detected from the younger, providing an insight into the evolutionary process of the formation of superbubbles around young stellar clusters. The spectra of the diffuse emission from the galaxy can be fitted with a two temperature component thermal model with $kT=0.14$\\,keV and 0.52\\,keV, and analysis of this emission suggests that NGC\\,4214 will suffer a blow-out in the future.  The point source population of the galaxy has an X-ray luminosity function with a slope of $-0.76$. This result, together with those for other dwarf starburst galaxies (NGC\\,4449 and NGC\\,5253), was added to a sample of luminosity functions for spiral and starburst galaxies. The slope of the luminosity function of dwarf starbursts is seen to be similar to that of their larger counterparts and clearly flatter than those seen in spirals. Further comparisons between the luminosity functions of starbursts and spiral galaxies are also made. ",
        "introduction": "Hierarchical models of structure formation in the Universe suggest that dwarf galaxies were some of the first objects to be formed \\citep*{Navarro_95}, and hence were amongst the earliest sites of star formation. If dwarf galaxies in the local Universe are analogues of those formed at high redshift then studying the structure and evolution of local dwarf galaxies will be useful in understanding star formation at high redshifts.  When starburst galaxies undergo intense periods of star formation many supernovae occur in a relatively small region of the galaxy. The ejecta from these combine with the winds and ejected material from the massive stars in OB associations and Super Star Clusters (SSCs) in the surrounding area to form a superbubble around the cluster which expands, sweeping up surrounding medium to produce large volumes of hot, shocked gas with $T\\sim 10^8$K. The simulations of \\citet{Suchkov_94} show that X-ray emission from starburst galaxies is most likely to be due to the shocks formed from the interaction between the starburst outflow and the surrounding medium, rather than from the material within the winds. The bubbles of hot gas expand along the steepest density gradient -- usually the minor axis of the galaxy, and will eventually expand out of the galaxy's gravitational potential well. Indeed, H$\\alpha$ images of starburst galaxies show filaments extending along the minor axis (\\citealt{Marlowe_95}, \\citealt{Lehnert_96a}). The low mass and hence small gravitational potential well of dwarf galaxies means that it more likely than in other galaxies for material to escape into the surrounding intergalactic medium and form galactic superwinds. In some models these starburst-driven winds result in enough interstellar medium (ISM) being ejected from the galaxy via galactic superwinds to prevent any further star formation from being able to occur \\citep{De_Young_94}. \\citet{Lehnert_96b} find that superwinds are most commonly formed in galaxies which have extreme IR properties, such as high IR luminosities and warm far-IR flux. Dwarf starburst galaxies can therefore play an important role in enriching their surroundings with heavy elements produced during starburst episodes (\\citealt{Ferrara_00}, \\citealt{Dekel_86}).  NGC\\,4214 is a nearby dwarf starburst galaxy with an almost face-on orientation. Using the tip of the red giant branch method the distance to NGC\\,4214 has been estimated by \\citet*{Maiz-Apellaniz_02} as $D=2.94$Mpc, and this is the value used throughout this paper. We note that a distance of 2.7Mpc was derived by \\cite{Drozdovsky_02} using the same method. A few $\\times10^7$ years ago the galaxy underwent a burst of star formation during which around $5\\%$ of the stellar mass of the galaxy was formed \\citep{Huchra_83}. Colour magnitude diagrams for the field stars outside the main star-forming regions show evidence for both old and young populations, and are typical for dwarf galaxies. The diagrams are dominated by the ``red tangle'' -- an area of the diagram that contains RGB, AGB and blue-loop stars -- which implies that there is an underlying population of old stars \\citep{Drozdovsky_02}. The basic properties of NGC\\,4214 are given in Table~\\ref{tab:4214}.  NGC\\,4214 has been studied at many wavelengths, each of which provide an insight into its structure and composition. These observations appear to be show ambiguous results when considering whether NGC\\,4214 has undergone a recent interaction with a nearby galaxy. $I$~band observations of NGC\\,4214 show a smooth, symmetric disk \\citep{Fanelli_97}. The outer contours of the $I$~band emission are regular, which suggest that there is no evidence that the galaxy has undergone a recent merger. NGC\\,4214 is H{\\small I}-rich like many other dwarf starbursts. In contrast, observations of H{\\small I}~emission by \\citet{Allsopp_79} show an extension towards a nearby irregular galaxy NGC\\,4190 suggesting a possible tidal interaction between the two galaxies could have triggered the starburst in NGC\\,4214. Observations in the far ultra-violet (FUV) shows the presence of knots of emission at the centre of the galaxy embedded in diffuse emission \\citep{Fanelli_97}. These knots denote the presence of OB associations which have also been observed at optical wavelengths ({\\sl e.g.} \\citealt{Huchra_83}, \\citealt{Leitherer_96}, \\citealt{MacKenty_00}).  More recent observations of NGC\\,4214 using the Hubble Space Telescope ({\\sl HST}) have shown evidence for a two-stage starburst \\citep{MacKenty_00}, producing SSCs and OB associations. These observations show that the galaxy has a complicated morphology and is dominated by two large H{\\small II} complexes often referred to as the North West (NW) and South East (SE) complexes (these structures are also referred to as NGC\\,4214-I and NGC\\,4214-II in the literature as well). In addition many knots of H$\\alpha$ emission are observed along with structurally varying diffuse gas surrounding the main regions of emission. Recent CO observations \\citep{Walter_01} show three main sites of molecular emission, two of which correspond to the NW and SE complexes, with an additional site at which star formation has yet to commence.  In optical appearance NGC\\,4214 is similar to other dwarf irregular galaxies, however \\citet{Maiz-Apellaniz_99} suggest that it has an unusually thin galactic disk with a thickness of around 200pc; outflows from the galaxy extending this by a further 200pc or so. Flatter galaxies preferentially undergo blow-out which is when supernovae explosions blow through the disk of the galaxy, forming a way for gas from subsequent explosions to escape without having much affect on the surrounding gas (\\citealt{De_Young_94}, \\citealt{Mac_Low_99}).    \\begin{table} \\begin{center} \\caption{Properties of NGC\\,4214} \\begin{tabular}{ c c c }  \\hline \\hline Parameter & Value & Ref. \\\\ \\hline  $\\alpha$[2000] & $12^{h}15^{m}39.26^{s}$ & \\\\ $\\delta$[2000] & $+36^{\\circ}19{'}33.4{''}$ & \\\\ Distance & 2.94Mpc & 1\\\\ 1 arcsec & 855pc & \\\\ $D_{25}$ & $8.91'\\times5.62'$ & 2 \\\\ Position angle & 42.2$^{\\circ}$ & 2 \\\\ Metallicity & $\\sim0.25$ & 3 \\\\ m$_B$ & 10.20\\,mag & 4 \\\\ S$_{60\\mu m}$ & 17.9\\,Jy & 5 \\\\ S$_{100\\mu m}$ & 29.0\\,Jy & 5 \\\\ $L_{FIR}$ & 9.1$\\times10^{41}\\rm{erg\\,s}^{-1}$ & 6 \\\\ SFR & 0.04\\,$M_{\\odot}\\,\\rm{yr}^{-1}$ & 7 \\\\ \\hline \\hline \\end{tabular} \\end{center} References: (1) \\citet{Maiz-Apellaniz_02}; (2) LEDA database (available at http://leda.univ-lyon1.fr); (3) \\citet{Kobulnicky_96}; (4) NED (NASA Extragalactic Database, available at http://nedwww.ipac.caltech.edu); (5) \\citet{Soifer_89}; (6) see Section~\\ref{sect:disc_ptsrcs}, Table~\\ref{table:lumfunc} for more details; (7) using the $L_{FIR}$ relation from \\citet{Kennicutt_98b}.  \\label{tab:4214}   \\end{table}    The launch of the {\\sl Chandra} Observatory has revolutionised X-ray astronomy, providing an unprecedented spatial resolution of $0.5''$. This resolution enables the identification of point sources down to fainter luminosities than those seen with the previous X-ray telescopes such as {\\sl ROSAT} and {\\sl Einstein}. It is now possible to study the diffuse X-ray emission from galaxies without severe contamination from unresolved point sources, allowing a more accurate definition of its properties. The {\\sl XMM-Newton} telescope launched in 1999 combines both reasonable spatial resolution with very high sensitivity. This allows the temperature and the elemental abundances of hot, diffuse gas to be studied in great detail. The combination of data from both instruments promises a much deeper understanding of starburst galaxies.  Prior to the launch of {\\sl Chandra} and {\\sl XMM-Newton} X-ray data from several dwarf starburst galaxies had been obtained from {\\sl ROSAT} observations (examples of {\\sl ROSAT} data on dwarf starburst galaxies can be found in \\citet{Stevens_98a} and \\citet{Stevens_98b}). However, {\\sl ROSAT} lacked both the sensitivity and resolution to be able to disentangle the point source emission from X-ray binaries from the diffuse emission, which severely limits the interpretation of the data.  This analysis of NGC\\,4214 adds to a small sample of published results from {\\sl Chandra} and {\\sl XMM-Newton} observations of dwarf starbursts. One significant example is the analysis of a {\\sl Chandra} observation of NGC\\,1569 \\citep*{Martin_02}. With the resolution of {\\sl Chandra} it is possible to study not just the global properties of the X-ray emission but also the properties of different spatial components. \\citet{Martin_02} found a strong correlation between the X-ray and H$\\alpha$ emission, and found a high $\\alpha$-element to Iron abundance ratio ($\\alpha$/Fe), which is consistent with a large amount of Type {\\small II} supernova ejecta contribution \\citep{Martin_02}. It is concluded from this study that elements can be released from dwarf galaxies into the surrounding intergalactic medium. Further comparison is presented in Section~\\ref{sect:xray+ir}. In addition to NGC\\,1569 published data is available from {\\sl Chandra} observations of the dwarf starbursts NGC\\,4449 \\citep{Summers_03a}, NGC\\,3077 \\citep{Ott_03} and NGC\\,5253 \\citep{Summers_03b}.   In Section~\\ref{sect:obs} we describe the {\\sl Chandra} and {\\sl XMM-Newton} observations of NGC\\,4214 and the data reduction methods used. In Section~\\ref{sect:ptsrcs} the properties of the point sources are analysed, and in Section~\\ref{sect:diff} the properties of the diffuse X-ray emission are investigated. The morphology of this emission is discussed and compared the H$\\alpha$ emission from the central regions of the galaxy. In Section~\\ref{sect:discussion} the luminosity function of NGC\\,4214 is compared to those of spiral and other starburst galaxies and the properties of the diffuse emission and their implications for the formation of superbubbles are discussed. Finally in Section~\\ref{sect:concl} the main conclusions are summarised.  \\begin{figure*} \\centering \\includegraphics[scale=1.0]{figure1.ps} \\caption{The NGC\\,4214 field, as seen with {\\sl Chandra}, showing the whole of the ACIS-S3 chip ($8.5'\\times8.5'$), binned up by a factor of 2. NGC\\,4214 can be seen as a region of diffuse emission at the centre of the image. The large ellipse represents the $D_{25}$ ellipse of the galaxy; the smaller ellipses show the location and extent of the discrete sources as found by {\\sl wavdetect}.} \\label{fig:whole_chip} \\end{figure*} ",
        "conclusions": "\\label{sect:concl}  We have presented the results from both {\\sl Chandra} and {\\sl XMM-Newton} observations of NGC\\,4214, which show the diffuse X-ray emission from this dwarf starburst galaxy. The high spatial resolution of {\\sl Chandra} has allowed the detailed structure of the X-ray emission of NGC\\,4214 to be observed for the first time. 17 point sources were observed within the $D_{25}$ ellipse of the galaxy, and further analysis of these suggest that the majority are HMXBs associated with the galaxy as opposed to being background AGN.  The luminosity function of the point sources is similar to those seen in other starbursts such as NGC\\,5253 and NGC\\,4449. Indeed, when combined with those of NGC\\,4449 and NGC\\,5253 and added to the sample of \\citet{Kilgard_02} this data strengthens the argument that the luminosity functions of starburst galaxies are flatter than those seen in normal spirals, and indeed extends this result to include dwarf starburst galaxies. It is seen that the starburst galaxies tend to have a higher 60$\\mu$m to 100$\\mu$m flux ratio than spirals, a criteria for being more likely to produce a galactic superwind \\citep{Lehnert_96b}. However, there is no clear segregation between the spirals and the starbursts when the luminosity function slopes are compared to the ratio of $L_{FIR}/L_B$. \\citet{Lehnert_96b} state that for a superwind to form a large excess of IR luminosities is needed, and although this criteria is not met in our sample overall the starbursts show a larger $L_{FIR}/L_B$ ratio than the spirals, as would be expected. In general the starbursts have a larger amount of star formation per unit area (using the $D_{25}$ ellipse area as a scaling parameter).  The diffuse emission from the galaxy is observed in both the {\\sl Chandra} and {\\sl XMM-Newton} observation, and is seen to almost twice the extension in {\\sl XMM-Newton} (radius $\\sim1.4'$) than in {\\sl Chandra}. The data suggest the presence of gas at two separate temperatures, with {\\sl Chandra} analysis providing values for these of around 0.3 and 4.6\\,keV and a combined absorption corrected flux of $3.4\\times10^{-13}$erg~s$^{-1}$cm$^{-2}$ corresponding to an unabsorbed luminosity of $L_X=3.5\\times10^{38}$erg~s$^{-1}$. Using derived X-ray properties the SFR of NGC\\,4214 is calculated as 0.07$M_{\\odot}\\,\\rm{yr}^{-1}$, in reasonable agreement with the value of 0.04$M_{\\odot}\\,\\rm{yr}^{-1}$ derived using $L_{FIR}$.  The X-ray emission coincides with regions of H$\\alpha$ emission and hence is coincident with star forming regions. The presence of X-ray emission coincident with the NW complex, but not the SE complex of H$\\alpha$ emission agrees with previous speculation that there is a time delay before the wind from a cluster of stars can break out and start interacting with the surrounding medium. Using some basic assumptions about the properties of the X-ray emitting gas is has been calculated that it is likely that a blow-out will occur from the disk of NGC\\,4214, enriching the surrounding medium with materials from this starburst."
    },
    "0311/astro-ph0311497_arXiv.txt": {
        "abstract": "{We present the results of the dynamical analysis of the rich, hot, and X--ray very luminous galaxy cluster A2219, containing a powerful diffuse radio--halo.  Our analysis is based on new redshift data for 27 galaxies in the cluster region, measured from spectra obtained at the TNG, with the addition of other 105 galaxies recovered from reduction of CFHT archive data in a cluster region of $\\sim 5\\arcmin$ radius ($\\sim 0.8$ \\h at the cluster distance) centered on the cD galaxy.  The investigation of the dynamical status is also performed by using X--ray data stored in the Chandra archive. Further, valuable information comes from other bands -- optical photometric, infrared, and radio data -- which are analyzed and/or discussed, too.  We find that A2219 appears as a peak in the velocity space at $z=0.225$, and select 113 cluster members. We compute a high value for the line--of--sight velocity dispersion, $\\sigma_{\\rm v}= 1438^{+109}_{-86}$ \\kss, consistent with the high average X--ray temperature of 10.3 keV. If dynamical equilibrium is assumed, the virial theorem leads to $M\\sim2.8\\times 10^{15}$\\msun for the global mass within the virial region.  However, further investigation based on both optical and X--ray data shows significant signs of a young dynamical status. In fact, we find strong evidence for the elongation of the cluster in the SE--NW direction coupled with a significant velocity gradient, as well as for the presence of substructure both in optical data and X--ray data.  Moreover, we point out the presence of several active galaxies. We discuss the results of our multi--wavelength investigation suggesting a complex merging scenario where the main, original structure is subject to an ongoing merger with few clumps aligned in a filament in the foreground oriented along an oblique direction with respect to the line--of--sight. Our conclusion supports the view of the connection between extended radio emission and merging phenomena in galaxy clusters. ",
        "introduction": "In the hierarchical scenario for large--scale--structure formation, mergers are an essential ingredient of galaxy cluster evolution (e.g., White \\cite{whi97}; Evrard \\& Gioia \\cite{evr02}).  One of the most recent aspects of the merging phenomenology is the possible connection of cluster mergers with the presence of extended, diffuse radio sources, halos and relics (Feretti \\cite{fer99}; Buote \\cite{buo02}; Giovannini \\& Feretti \\cite{gio02}).  Radio--halos are located at the cluster center, with typical sizes of 0.5 $h^{-1}$ Mpc, regular shape, steep radio spectra and no significant polarization; while relics are located in peripheral regions of the clusters, with irregular shape and generally highly polarized spectra (e.g., Feretti \\& Giovannini \\cite{fer96}; Giovannini \\& Feretti \\cite{gio02}). The synchrotron radio emission of halos and relics demonstrates the existence of large scale cluster magnetic fields, of the order of 0.1--1 $\\mu$G, and of widespread relativistic particles of energy density 10$^{-14}$ -- 10$^{-13}$ erg cm$^{-3}$. Although the properties of halos and relics are not well understood, several suggestions for the mechanism transferring energy to the relativistic electrons have been made (e.g., Ensslin \\cite{ens00}).  Cluster mergers were suggested to provide the large amount of energy necessary for electron re--acceleration and magnetic field amplification (Feretti \\cite{fer99}; Feretti \\cite{fer02}; Sarazin \\cite{sar02}). However, the precise radio--halo formation scenario is still debated (e.g., Ensslin \\& R\\\"ottgering \\cite{ens02}).  In fact, the diffuse radio sources are quite uncommon and only recently we can study these phenomena on the basis of a sufficient statistics ($\\sim 30$ clusters up to $z\\sim 0.3$, e.g., Giovannini et al. \\cite{gio99}; see also Giovannini \\& Feretti \\cite{gio02}).  Presently, growing evidence of connection between diffuse emission and cluster merging is based on X--ray data (e.g., B\\\"ohringer \\& Schuecker \\cite{boh02}; Buote \\cite{buo02}). Studies based on a large number of clusters have found a significant relation between the radio and the X--ray surface brightness (Govoni et al. \\cite{gov01}) and connections between the presence of radio--halos/relics and irregular and bimodal X--ray surface brightness distribution (Schuecker et al. \\cite{sch01}).  However, we are far from understanding the whole picture. Unfortunately, to date optical data is lacking or poorly exploited. Sparse literature concerns some few individual clusters (e.g., Colless \\& Dunn \\cite{col96}; G\\'omez et al. \\cite{gom00}; Barrena et al. \\cite{bar02}; Mercurio et al. \\cite{mer03}).  Optical data are a powerful way to investigate the presence and the dynamics of cluster mergers (e.g., Girardi \\& Biviano \\cite{gir02}), too.  The spatial and kinematic analysis of member galaxies allow us to detect and measure the amount of substructure, to identify and analyze possible pre--merging clumps or merger remnants.  This optical information is really complementary to X--ray information since galaxies and ICM react on different time scales during a merger (see numerical simulations by Roettiger et al. \\cite{roe97}). Moreover, additional information comes from spectral types of member galaxies, since cluster mergers could stimulate star formation in cluster galaxies; thus the spectral signatures of past activity are useful to determine the relevant time--scales (e.g., Bekki \\cite{bek99}; Terlevich et al. \\cite{ter99}).  To investigate the connection between diffuse radio sources and cluster mergers, we have performed spectroscopic observations at TNG and a multi--wavelength analysis of the Abell cluster A2219 (Abell et al. \\cite{abe89}). This cluster shows the presence of a powerful radio--halo (Giovannini et al. \\cite{gio99}; Bacchi et al. \\cite{bac03}) and is a rich, X--ray luminous, hot cluster at moderate redshift $z\\sim0.22$ -- Abell richness $=3$ (Abell et al. \\cite{abe89}); $L_\\mathrm{X}$(0.1--2.4 keV)$\\sim$5$\\times 10^{44} \\ h^{-2}$ erg\\ s$^{-1}$ (Ebeling et al. \\cite{ebe96}); $T_\\mathrm{X}\\sim$10 keV (Rizza et al. \\cite{riz98}). Possible evidence of the young dynamical status of A2219 comes from the ROSAT/HRI X--ray data analysis and strong/weak lensing features (Rizza et al. \\cite{riz98}; Smail et al. \\cite{sma95}; Dahle et al. \\cite{dah02}).  This paper is organized as follows.  We present the new redshift data of A2219 in Sect.~2 and the analysis of optical data in Sect.~3. Our analysis of recent X--ray Chandra archival data is shown in Sect.~4. Active galaxies are analyzed and discussed on the basis of multi--wavelength data using also mid--IR and radio data in Sect.~5. Finally, in Sect.~6, we summarize and discuss our results suggesting a tentative picture of the dynamical status of A2219.  Unless otherwise stated, we give errors at the 68\\% confidence level (hereafter c.l.)  Throughout the paper, we assume a flat cosmology with $\\Omega_{\\rm m}=0.3$ and $\\Omega_{\\Lambda}=0.7$.  For the sake of simplicity in rescaling we adopt a Hubble constant of 100 \\ks \\hh.  In this context, 1$\\arcmin$ corresponds to $\\sim 0.152$ \\hh. ",
        "conclusions": "We analyze the internal dynamics of A2219 on the basis of spectroscopic data for 132 galaxies in a cluster region of $\\sim 5\\arcmin$ ($\\sim 0.8$ \\h at the cluster distance) centered on the cD galaxy. Out of 132 spectra, 27 galaxies come from new observations carried out at the TNG.  We find that A2219 appears as a peak in the velocity space at $z=0.225$, and select 113 cluster members.  The investigation of the dynamical status is also performed by using X--ray data stored in the Chandra archive. Further valuable information comes from other bands -- optical photometric, infrared, and radio data -- which are analyzed and/or discussed, too.  The value we obtain for the line--of--sight velocity dispersion $\\sigma_{\\rm v}\\sim 1400$ \\ks is in the high--tail of the distribution of cluster velocity dispersions (see Fadda et al. \\cite{fad96}; Mazure et al. \\cite{maz96}; Girardi \\& Mezzetti \\cite{gir01}).  This global value of $\\sigma_{\\rm v}$ is consistent with the average value of $T_{\\rm X} \\sim 10$ keV coming from the X--ray analysis ($\\beta_{\\rm{spec}}$=1.10$^{+0.15}_{-0.12}$, see also Fig.~\\ref{figprof}) and with the value of $L_{\\rm X}(0.1-2.4\\ \\rm{keV})\\sim 5\\;h^{-2}\\;10^{44}$ erg s$^{-1}$ (Ebeling et al. \\cite{ebe96}), converted in $L_{\\rm{X,bol}}\\sim2\\cdot 10^{45} \\ h^{-2}$ erg s$^{-1}$ (see $L_{\\rm X}$--$\\sigma_{\\rm v}$ relation by, e.g., Wu et al. \\cite{wu99}, and Girardi \\& Mezzetti \\cite{gir01}).  Therefore, on the basis of global properties only, we would conclude that A2219 is not far from the dynamical equilibrium and would trust the large virial mass estimate $M(\\sim 2.2$\\,\\hh$)\\sim 2.8\\times10^{15}$ \\msun.  However, the high values of $\\sigma_{\\rm v}$, $T_{\\rm X}$, and $L_{\\rm X}$ could be due to the expected increase during a cluster--merging phase (e.g., Schindler \\& B\\\"ohringer \\cite{sb93}; Schindler \\& M\\\"uller \\cite{sm93}). Their agreement could be a mere coincidence, and our virial mass estimate an overestimate of the true value.  We find observational evidence that the cluster is not dynamically relaxed.  Our optical and X--ray analyses show that A2219 is elongated in the SE--NW direction. This elongation is seen in: the spatial distribution of the colour--selected likely cluster members; the shape of the cD galaxy; the X--ray contour levels of the Chandra image; the gradient in the velocity dispersion and in the X--ray temperature. In particular the ellipticity of X--ray isophotes is significantly higher than 0 and corresponds to an axis-ratio=0.66, which is much lower than the usual value for galaxy clusters (median axis-ratio with 99\\% c.l. errors=$0.82^{+0.04}_{-0.02}$ from Mohr et al. \\cite{moh95}).  Note that Dahle et al. (\\cite{dah02}) found that the mass distribution of this cluster recovered from a weak lensing analysis appears even more elongated than the light and galaxy number density distribution.  Very interestingly, the optical data show a somewhat smaller PA than X--ray data ($\\sim 110\\degr$ and $\\sim 130\\degr$, respectively).  A different direction in the elongation of light and hot gas distribution could be the result of a cluster merger in an advanced phase as shown by numerical simulations where collisional and collisionless components react in different way (e.g., Ricker \\& Sarazin \\cite{ric01}; Schindler \\cite{sch02}). In fact, the evidence of an ongoing merger in the cluster core was first suggested by Smail et al. (\\cite{sma95}) to explain the difference they find in the orientation of the mass and gas distributions (PA=115 and 128 degrees), as obtained from gravitational lensing analysis and ROSAT/HRI X--ray data, respectively.  As for the presence of the velocity gradient roughly in the E--W direction, this suggests that the structures are moving into the plane of the sky at $\\sim 45\\degr$ (e.g., Quintana et al. \\cite{qui96}; Roettiger \\& Flores \\cite{roe00}).  In fact, intermediate angles are the most suitable to allow detection of the cluster elongation and contemporary the velocity gradient during a merging phase.  The hypothesis of an advanced--phase merging is also supported by the lack of a cool center in the cluster (see also Allen \\& Fabian \\cite{all98}), as expected since strong cluster mergers disrupt the cool cores (Peres et al. \\cite{per98}; Roettiger et al. \\cite{roe96}), and the existence in the core of several star--forming or post--starburst galaxies, whose activity could be produced by the cluster merger (e.g., Bekki \\cite{bek99}). In particular, three of these active galaxies are part of the filament at NW of the cD, also visible as a cool component in the X--ray softness ratio map (see Fig.~\\ref{figxy} and Fig.~\\ref{figsoftness}), suggesting that this structure might be already involved in a merger event.  Moreover, the shape of the radio tail of the radio galaxy no. 77 indicates that the velocity of this galaxy was originally pointing towards the cluster center (cD galaxy), but has rather suddenly changed the motion direction, probably because of the cluster merger.  The structure we find in the cluster core could be related to the cluster merger, too. Considering optical data, the Dressler--Schectman test shows the possible existence of a substructure located at SE, very close to the cD galaxy. As for X--ray data, our wavelet multi--scale analysis puts in evidence three significant emission peaks located within 1$\\arcmin$ from the cluster center. Of these, two peaks are possibly connected to intrinsic X--ray emission from the cD and the active galaxy no. 81, but no optical counterpart is found for the third, intermediate peak. Our wavelet X--ray analysis also detects an external diffuse substructure located 2$\\arcmin$ SE of the cD galaxy.  We suggest that this structure is a premerging clump.  In fact, it coincides with the cold SE X--ray region and thus could be well explained by the two--dimensional superposition of a cold clump in the foreground of the main cluster.  The presence of this foreground clump could also explain why, on average, the velocity at the East is lower than in the West since at least for some galaxies the low velocity is of cosmological rather than of kinematic origin.  These anomalous low velocity galaxies could also produce a spurious enhancement of the observed velocity dispersion in the SE region (see Table~2).  A population of very close foreground galaxies could also explain why {\\em very red} galaxies strangely differ from {\\em not--very red} galaxies both in kinematic and spatial properties (see Sect.~3.3). In fact, while the former are really associated with the main structure of A2219, the latter could be mixed with the clump members, i.e. a population of foreground -- and thus less red -- galaxies. Unfortunately, we have very poor redshift information in the SE region to confirm our hypothesis and identify clump members.  On a much larger scale, the cluster is not so well isolated in the redshift space, and a next possible merging clump might be the one at $z\\sim 0.19$ (see Fig.~\\ref{figden}).  Summarizing the above evidence, we suggest a complex merging scenario for A2219 involving many clumps, possibly in different dynamical states. We suggest that the main, original structure (hosting the cD galaxy) is subject to an ongoing merger with few clumps aligned in a foreground filament obliquely oriented with respect to the line--of--sight; the filament projected on the sky lies in the SE--NW direction. In this scenario, the merging with one or more clumps are already in a very advanced phase, well after the first core passage, while one or more clumps are still in a pre--merging phase.  In the context of the above scenario, we argue that the presence of a radio--halo in A2219 is related to the existence of an ongoing, very advanced phase of merging between the main cluster and one or more groups. In fact, the elongation of radio map contours (see Fig.~\\ref{figisofote}) has the same direction of the X--ray surface brightness contours. Moreover, the time necessary to re--accelerate the electrons producing the radio--halo as a consequence of the cluster merger ($\\sim$1 Gyr) is comparable to that needed to obtain post--starburst spectral signatures in galaxies of which A2219 is rich (see Sect.~5).  In conclusion, our scenario supports the view of the connection between extended radio emission and merging phenomena in galaxy clusters."
    },
    "0311/astro-ph0311342_arXiv.txt": {
        "abstract": "We analyze the filamentarity in the Las Campanas redshift survey (LCRS) and determine the length scale at which filaments are statistically significant.  The largest length-scale at which filaments are statistically significant, real objects, is between $70$ to $80 \\, h^{-1}$Mpc, for the LCRS $-3^o$ slice. Filamentary features longer than $80 \\, h^{-1}$Mpc, though identified, are not statistically significant; they arise from chance alignments.  For the five other LCRS slices, filaments of lengths $50 \\, h^{-1}$Mpc to $70 \\, h^{-1}$Mpc are statistically significant, but not beyond. These results indicate that while individual filaments up to $80 \\, h^{-1}$Mpc are statistically significant, the impression of structure on larger scales is a visual effect. On scales larger than $80 \\, h^{-1}$Mpc the filaments interconnect by statistical chance to form the the filament-void network. The reality of the $80 \\, h^{-1}$Mpc features in the $-3^o$ slice make them the longest coherent features in the LCRS.  While filaments are a natural outcome of gravitational instability, any numerical model attempting to describe the formation of large scale structure in the universe must produce coherent structures on scales that match these observations. ",
        "introduction": "One of the most striking visual features in the distribution of galaxies in the Las Campanas redshift Survey (LCRS, Shectman et al. 1996) is that they appear to be distributed along filaments.  These filaments are interconnected and form a network, with voids largely devoid of galaxies comprising the region in between the filaments. This network of interconnected filaments encircling voids extends across the entire survey and may be referred to as the ``Cosmic Web''. Similar networks of filaments and voids are also visible in other galaxy surveys, e.g., CfA (Geller \\& Huchra 1989), 2dFGRS (Colless et al. 2001, Colless et al. 2003) and SDSS (EDR) (Stoughton et al. 2002, Abazajian et al. 2003). These, if they are a genuine feature of the galaxy distribution, represent the largest structural elements in the hierarchy of structures observed in the universe, namely galaxies, clusters, superclusters, filaments and the cosmic web.  The analysis of filamentary patterns in the galaxy distribution has a long history dating back to papers by Zel'dovich, Einasto and Shandarin (1982), Shandarin and Zel'dovich (1983) and Einasto et al. (1984). In the last paper the authors analyze the distribution of galaxies in the Local Supercluster. They use the Friend-of-Friend algorithm with varying neighborhood radius to identify connected systems of galaxies referred to as ``clusters''.  As they increase the neighborhood radius, they find that the clusters which are initially spherical become multi-branched with multiple filaments of lengths up to several tens of $h^{-1}$Mpc extending out in different directions. Finally, as the radius is increased further, they find that the filaments get interconnected and join neighboring superclusters into an infinite network of superclusters and voids.  A later study (Shandarin \\& Yess 2000) used percolation analysis to arrive at a similar conclusion for the distribution of the LCRS galaxies. The large-scale and super large-scale structures in the distribution of the LCRS galaxies have also been studied by Doroshkevich et al. (2001) and Doroshkevich et al. (1996) who find evidence for a network of sheet like structures which surround underdense regions (voids) and are criss-crossed by filaments.  The distribution of voids in the LCRS has been studied by M{\\\" u}ller, Arbabi-Bidgoli, Einasto \\& Tucker (2000) and the topology of the LCRS by Trac, Mitsouras, Hickson \\& Brandenberger (2002) and Colley (1997).  A recent analysis (Einasto et al. 2003a) indicates a supercluster-void network in the Sloan Digital Sky Survey also.   Traditionally, correlation functions (Peebles 1993) have been used to quantify the statistical properties of the galaxy distribution. For the LCRS, the two-point correlation function is a power law, \\b \\xi(r) = \\left({r\\over r_o}\\right)^{-1.52} \\e with the correlation length $r_o = 6.28$ $h^{-1}$Mpc on scales $2.0 \\, h^{-1}$Mpc to $16.4 \\, h^{-1}$Mpc.  On scales larger than $30-40 \\, h^{-1}$Mpc, $\\xi$(r) fluctuates closely around zero indicating a statistically homogeneous galaxy distribution at and beyond these scales (Tucker et al. 1997). This raises the question whether the filamentary features which appear to span scales larger than $100 \\, h^{-1}$Mpc are statistically significant features of the galaxy distribution or if they are mere artifacts arising from chance alignment of the galaxies.  A quantitative estimator of filamentary structure, Shapefinder, was defined (Bharadwaj et al. 2000) to provide a measure of the average filamentarity for a point distribution in 2D. (See Sahni, Sathyaprakash \\& Shandarin 1998 for general introduction to Shapefinders and Sheth et al. 2003 for the application of Shapefinders to 3D simulations of structure formation.)  The Shapefinder statistic was used to demonstrate that the galaxy distribution in the LCRS exhibits a high degree of filamentarity compared to a random Poisson distribution having the same geometry and selection effects as the survey. This analysis provides objective confirmation of the visual impression that the galaxies {\\em are} distributed along filaments. This, however, does not establish the statistical significance of the filaments. The features identified as filaments are essentially chains of galaxies, a crucial requirement being that the spacing between any two successive galaxies along a chain is significantly smaller than the mean inter-galaxy separation.  A chain runs as long as it is possible to find another nearby galaxy which is not yet a member of the chain, and breaks when no such galaxy is to be found. The fact that the LCRS galaxies are highly clustered on small scales increases the probability of finding pairs of galaxies at small separations. This enhances the occurrence of long chains of galaxies, and we expect to find a higher degree of filamentarity arising just from chance alignments in the LCRS compared to a Poisson distribution.  To establish whether the observed filaments are statistically significant or if they are a result of chance alignments of smaller structural elements, it is necessary to compare the sample of galaxies (here, the LCRS slices) with a distribution of points which has the same small scale clustering properties as the original sample and for which we know that all large-scale filamentary features are solely due to chance alignments. This is achieved using a statistical technique called Shuffle (Bhavsar \\& Ling 1988) whereby the statistical significance of the filamentarity in a clustered dataset can be assessed.  Shuffle generates fake data-sets, practically identical in their clustering properties to the original data up to a length scale $L$, but in which all structures longer than $L$, both real and chance, of the original data, have been eliminated. In these Shuffled data, filaments spanning length-scales larger than $L$ are visually evident, even expected to be identified as a signal by the statistics used to quantify the filamentarity, but all filaments spanning length-scales larger than L have formed accidentally. The measure of the occurrence of filaments spanning length-scales larger than $L$ in the Shuffled data gives us a statistical estimate of the level at which chance filaments spanning length-scales larger than $L$ occur in the original data. Here we use Shuffle to estimate the degree of filamentarity expected from chance alignments in the LCRS and use this to determine the statistical significance of the observed filamentarity.  We present the method of analysis and our findings in Section 2.  In Section 3 we discuss our results and present conclusions. ",
        "conclusions": "A look at the filamentary features at different levels of coarse-graining (Figure 1) reveals that the size of the largest filamentary feature increases monotonically with successive iterations of coarse-graining until it spans the entire survey (Bharadwaj et al. 2000). As coarse-graining proceeds, individual filaments form and then interconnect to form the supercluster-void network, in keeping with the earlier analysis (Einasto et al. 1984; Shandarin \\& Yess 2000; Einasto et al. 2003a) discussed in the introduction.  Although, the length of the interconnected network of filaments increases monotonically, the ratio of the length to the number of holes (Genus) stabilizes and then decreases (Bharadwaj et al. 2000). For the $-3^{o}$ slice, this ratio stabilizes around $140 \\, h^{-1}$Mpc at $FF \\sim$ 0.4. This ratio may be interpreted as the perimeter of the typical void in the network. This leads to a picture where there are voids of diameter $60 \\, h^{-1}$Mpc encircled by filaments of thickness $10 \\, h^{-1}$Mpc (Peebles 1993, Sheth et al. 2003, Sheth 2003) interconnected to form a large web. A void of this size, along with the filament at its perimeter, would span a length-scale $\\sim 80 \\, h^{-1}$Mpc. The results of this paper show that such voids encircled by filaments are statistically significant features. Although our analysis also finds a web of interconnected filaments which spans length-scales larger than $80 \\, h^{-1}$Mpc and runs across the entire survey, this is not statistically significant. The web arises from chance interconnections between the filaments encircling different voids.  Studies of the distribution of Abell superclusters (Einasto et al. 1997b) show that the mean distance between neighboring superclusters is about $50 \\, h^{-1}$Mpc for poor superclusters and about $100 \\, h^{-1}$Mpc for rich superclusters. The distribution of the SDSS superclusters (Einasto et al. 2003a) and the LCRS superclusters (Einasto et al.  2003b) shows a similar behavior.  Visualizing the superclusters as being randomly distributed, we would expect filaments joining the superclusters to develop as the density field is progressively smoothened.  Such filaments will arise from the chance alignments of shorter, genuine, statistically significant filaments. The filaments joining superclusters will span length-scales comparable to the mean inter-supercluster separation and the statistical properties of filaments would be stable to shuffling {\\it i.e.}, it would not change if the superclusters were rearranged randomly. Our results may be interpreted as being indicative of the superclusters being randomly distributed on scales larger than $80 \\, h^{-1}$Mpc with the mean inter supercluster separation also being of this order {\\footnote{ It is interesting to note that a study of the SDSS(EDR) superclusters conducted by Doroshkevich et al. (2003) using Minimal Spanning Trees concludes that the large-scale filaments appear to randomly connect the sheet-like structures in more denser environments.}}.  The presence of statistically significant features on scales as large as $70$ to $80 \\, h^{-1}$Mpc may seem surprising given the fact that the correlation analysis fails to detect any clustering on scales beyond $30$ to $40 \\, h^{-1}$Mpc. This is due to the inability of the two point (and higher order) correlation functions in detecting coherence at large-scales.  Pattern specific methods (like Shapefinders) are necessary to detect and quantify coherent large scale features in the galaxy distribution. It is interesting to note that the two dimensional power spectrum for the LCRS (Landy et al. 1996) exhibits strong excess power at wavelengths $\\sim 100 \\, h^{-1}$Mpc, a feature which may possibly be related to the filamentary patterns studied here.  The analysis of the three dimensional distribution of Abell clusters (Einasto et al. 1997a, 1997b) reveals a bump at $k=0.05 \\, h \\, {\\rm Mpc}^{-1}$ in the power spectrum.  Also, the recent analysis of the SDSS shows a bump at $k=0.05 \\, h \\, {\\rm Mpc}^{-1}$ in the power spectrum (Tegmark et al.  2003). While it is interesting to conjecture that these features may be related to the presence of filaments, we should also note that the filaments are non-Gaussian features and cannot be characterized by the power spectrum alone.   A point which should be noted is that the filaments quite often run in a zig-zag fashion (Figure 1), and the length of a filament which spans a length-scale of $80 \\, h^{-1}$Mpc may be significantly larger than $80 \\, h^{-1}$Mpc. Also, the present analysis is two dimensional whereas the filaments actually extend in all three dimensions. The length of the filaments may be somewhat larger in three dimensions, and a little bit of caution may be advocated in generalizing our results {\\footnote{In this context, it is interesting to note that in a recent analysis of mock SDSS catalogs based on $\\Lambda$CDM model, Sheth (2003) finds the length-scales of the largest superclusters to be $\\sim 60 \\, h^{-1}$Mpc. This indicates as to what might be anticipated in extending this work to 3D redshift surveys.}}.  In the gravitational instability picture, small disturbances in an initially uniform matter distribution grow to produce the large-scale structures presently observed in the universe. It is possible to interpret the filaments in terms of the coherence of the deformation (or strain) tensor (Bond, Kofman \\& Pogosyan 1996) of the smoothened map from the initial to the present positions of the particles which constitutes the matter. Our analysis shows that the deformation tensor has correlation to length-scales up to $80 \\, h^{-1}$Mpc and is uncorrelated on scales larger than this. The ability to produce statistically significant filamentarity on scales up to $80 \\, h^{-1}$Mpc will be a crucial quantitative test of the different models for the formation of Large Scale Structure in the universe.   We next address the question of the length-scale beyond which the distribution of galaxies in the LCRS may be considered to be homogeneous. The analysis of Kurokawa, Morikawa \\& Mouri (2001) shows this to occur at a length-scale of $\\sim 30 \\, h^{-1}\\,{\\rm Mpc}$, whereas Best (2000) fails to find a transition to homogeneity even on the largest scale analyzed.  The analysis of Amendola \\& Palladino (1999) shows a fractal behavior on scales less than $\\sim$ $30 h^{-1}\\,{\\rm Mpc}$ but is inconclusive about the transition to homogeneity. The results presented in this paper set a lower limit to this length-scale at around $80 \\, h^{-1}$Mpc, in keeping with estimates based on the multi-fractal analysis of LCRS (Bharadwaj, Gupta \\& Seshadri 1999) who find that the LCRS exhibits homogeneity on the scales $80$ to $200 \\, h^{-1}\\, {\\rm Mpc}$. In a separate approach based on the analysis of the two point correlation applied to actual data and simulations Einasto and Gramann (1993) find that the transition to homogeneity occurs at about $130 \\, h^{-1} {\\rm Mpc}$.  For the LCRS, the scale of the largest coherent structure is at least twice the length-scale at which the two-point correlation function becomes zero. Beyond this scale the filaments interconnect statistically to form a percolating network.  This filament-void network of galaxies is not distinguishable, in a statistical sense, beyond scales of $80 \\, h^{-1}$Mpc. If the LCRS slices can be considered a fair sample of the universe then this suggests the scale of homogeneity for the universe."
    },
    "0311/astro-ph0311518_arXiv.txt": {
        "abstract": "I present a review by epoch of baryons in the intergalactic medium (IGM), from the first star until today.  Recent observations indicate a protracted period of reionization, suggesting multiple populations of reionizers; detection of these $z\\ga 6$ sources is a key goal that is now coming within reach.  The optical Lyman alpha forest ($2\\la z\\la 4$) is well-described by the Fluctuating Gunn-Peterson Approximation, but puzzling observations of associated metal lines and nearby galaxies may provide insights into galactic feedback processes. IGM studies at $z\\la 1.5$ are progressing rapidly thanks to ultraviolet absorption line studies of both Ly$\\alpha$ absorbers and the warm-hot intergalactic medium. Cosmological hydrodynamic simulations have played an integral part in these advances, and have helped to shape our understanding of the IGM at each epoch.  Working jointly, observations and theory continue to expand our knowledge of the IGM as the earliest stages of galaxy formation and the dominant reservoir of baryons at all redshifts. ",
        "introduction": "All baryons in the Universe were originally in the intergalactic medium (IGM). Today, the majority of baryons remain there.  Understanding the evolution of the IGM is therefore a critical component for a complete theory of how baryons evolve in the Universe, from the smooth Dark Ages to present-day galaxies and large-scale structure.  In this review I will highlight some recent work towards understanding the evolution of the IGM's various components, their physical state and connection to galaxies, with particular emphasis on the contributions and insights provided by numerical simulations.  This review is divided into four sections.  \\S2 presents a simulation-based overview of the various baryonic components in the Universe.  \\S3 discusses some of the latest theoretical and observational results probing the epoch of reionization.  \\S4 presents a short history of optical Ly$\\alpha$ forest studies, and expounds on two of its more puzzling aspects: The presence of metals in the diffuse IGM, and the effects of galaxies on their surrounding gas.  \\S5 highlights recent developments in probing the $z\\la 1.5$ IGM, including prospects for detecting the so-called missing baryons. \\S6 provides a short summary. ",
        "conclusions": "Observations of the IGM have progressed rapidly in recent years, which in turn has spurred dramatic advances in our theoretical understanding of the various baryonic components in the Universe.  Still, many questions crucial to our understanding of galaxy and structure formation remain unanswered, including: \\begin{itemize} \\item What are the sources that reionize the IGM, and how can we detect them? \\item How and when did metals get into the IGM? \\item What effects do galactic feedback have on surrounding IGM gas, and how does it change with redshift? \\item How do we obtain a complete observational census of baryons today? \\end{itemize} Here I have given a status report on these topics, which currently present more questions than answers.  Major advances will only come through joint progress in observational capabilities, particularly in the ultraviolet, along with improving simulations for the evolution of baryons in the Universe."
    },
    "0311/astro-ph0311032_arXiv.txt": {
        "abstract": "How does turbulence\\index{turbulence} contribute to the formation and structure of the dense interstellar medium (ISM)? Molecular clouds are dense, high-pressure objects.  It is usually argued that gravitational confinement\\index{gravitational confinement} causes the high pressures, and that the clouds are (magneto)hydrostatic\\index{hydrostatic} objects supported by a balance between magnetic and turbulent pressures and gravity. However, magnetic pressures\\index{pressure, magnetic}\\index{pressure}\\index{pressure, turbulent} appear too weak, and MHD turbulent support not only requires driving, but also results in continuing gravitational collapse, as has now been demonstrated in simulations reaching $512^3$ zones. Models of supernova-driven, magnetized turbulence readily form transient, high-pressure, dense regions that may form molecular clouds. They are contained not by self-gravity, but by turbulent ram pressures from the larger flow.  Apparent virialization may actually be a geometrical effect.  Turbulent clouds are unlikely to be in hydrostatic equilibrium, instead either collapsing or expanding, although they may appear well-fit by projected equilibrium Bonnor-Ebert spheres.  Collapsing clouds probably form stars efficiently, while expanding ones can still form stars by turbulent compression, but rather inefficiently. ",
        "introduction": "\\label{sec:equil}  The very high Reynolds numbers of the interstellar medium (ISM) ensure that it is turbulent and not laminar.  How can we model a turbulent flow that is intrinsically non-steady and chaotic?  Most models of the structure of the ISM rely on some notion of equilibrium: (magneto)hydrostatic equilibrium, virial equilibrium, energy equipartition, or at least statistical equilibrium (or stationarity). In this talk I use the structure of the dense ISM as a case study in the problems with these different assumptions.  In this case, making such assumptions leads to the conclusion that molecular clouds are gravitationally bound, which may not generally be the case. ",
        "conclusions": ""
    },
    "0311/astro-ph0311204_arXiv.txt": {
        "abstract": "{We present a summary of high-spatial resolution follow-up observations of the elliptical (E) and lenticular (S0) galaxies in the \\sauron\\/ survey using the \\oasis\\/ integral field spectrograph. The \\oasis\\/ observations explore the central 8\\arcsec$\\times$10\\arcsec regions of these galaxies using a spatial sampling four times higher than \\sauron, often revealing previously undiscovered features. Around 75\\% (31/48) of the \\sauron\\/ E/S0s with central velocity dispersion $\\gsim 120$~\\kms\\/ were observed with \\oasis, covering well the original \\sauron\\/ representative sample. We present here an overview of this follow-up survey, and some preliminary results on individual objects, including a previously unreported counter-rotating core in NGC\\,4382; the decoupled stellar and gas velocity fields of NGC\\,2768; and the strong age gradient towards the centre of NGC\\,3489. ",
        "introduction": "The \\sauron\\/ project (de Zeeuw et al. 2002) is a study of the structure of 72 representative nearby early-type galaxies and spiral bulges based on measurements of the two-dimensional stellar kinematics and distribution of absorption line-strengths, together with the kinematics and distribution of ionized-gas, measured with \\sauron, a panoramic integral-field spectrograph (IFS) mounted at the William Herschel Telescope (WHT), La Palma (Bacon et al. 2001). The aims of this survey and the observed galaxy sample are described in de Zeeuw et al. (2002).  The stellar kinematic maps for the sub-sample of 48 E/S0s is presented in Emsellem et al. (2003).  To provide a large field of view, the spatial sampling of \\sauron\\/ was set to 0\\farcsec94$\\times$0\\farcsec94 (per lenslet) for the survey, therefore often undersampling the typical seeing at La Palma (0\\farcsec8 FWHM). This does not, in principle, affect the results for the main-body of early-type galaxies, as they generally exhibit smoothly-varying structure on scales larger than the seeing. Towards the galaxy nucleus, however, there are often sharp, localised features in the kinematics, such as decoupled cores or central disks, as well as distinct stellar populations and ionized-gas distributions. Such features may be only partially resolved in the \\sauron\\/ data, and in some cases are not even visible.  In addition, at HST resolution, all elliptical galaxies exhibit power-law central luminosity profiles: $\\Sigma (r) \\propto r^{-\\gamma}$, with clear trends between the nuclear luminosity profiles and certain global properties, such as the degree of rotational support, isophotal shape, and stellar populations. It is therefore crucial to fill the gap between the medium (few 100s of pc) to large-scale (few kpc) structures probed with \\sauron\\/ and the inner ($< 200$~pc) components probed by HST. We have thus conducted a complimentary study on a subset of the \\sauron\\/ sample using the \\oasis\\/ (Optically Adaptive System for Imaging Spectroscopy) spectrograph, mounted on the Canada-France-Hawaii Telescope (CFHT), Hawaii. ",
        "conclusions": "We have conducted extensive high-spatial resolution observations of the E/S0 sub-sample of the \\sauron\\/ survey using the \\oasis\\/ IFS. This provides a unique data set on the centres of these galaxies, complimenting the panoramic view delivered by \\sauron. We have presented here some preliminary results to illustrate the quality of the data and the analysis techniques we use. The \\oasis\\/ data set will be published as a complimentary project to \\sauron, relating the inner properties of these galaxies with those of the outer regions, as well as providing a high-resolution catalogue of a substantial sample of early-type galaxies. The data will be made publicly available when the final reduction and analysis is complete."
    },
    "0311/astro-ph0311387.txt": {
        "abstract": "{% We present a unique set of nested stellar kinematical maps of \\NGC{3377} obtained with the integral-field spectrographs \\oasis and \\sauron. We then construct general axisymmetric dynamical models for this galaxy, based on the \\schw numerical orbit superposition technique applied to these complementary measurements. We show how these two datasets constrain the mass of the central massive object and the overall mass-to-light ratio of the galaxy by probing the inner and outer regions respectively. The simultaneous use of both datasets leads us to confirm the presence of a massive black hole with a mass of $\\Mbh = 7_{-5}^{+4}\\,10^{7}$~\\Msun{} (99.7\\% confidence level), with a best-fit stellar mass-to-light ratio $\\MLI = 2.1 \\pm 0.2$ (for an assumed edge-on inclination). % ",
        "introduction": "\\NGC{3377} is a prototypical `disky' E5-6 galaxy with `boxy' outer isophotes in the Leo~I group \\citep[e.g.][]{1990AJ....100.1091P} at an assumed distance of $D = 9.9$~Mpc. It has a power-law central luminosity profile \\citep{HST-centersIV, 2001AJ....121.2431R}, and its total absolute magnitude of $\\sim -19$ ($B$) is intermediate between that of the classical `boxy' giant ellipticals and the `disky' lower-luminosity objects \\citep[e.g.][]{KB96}. Previous dynamical models of this galaxy suggest the presence of a central massive black hole (BH) of $\\sim 10^8$~\\Msun{} (\\citealt[][hereafter K+98]{ngc3377-K98}; \\citealt[][hereafter G+03]{G3I-03}), while the $\\Mbh-\\sigma$ relation \\citep[and references therein]{2002ApJ...574..740T} predict a BH of $\\sim 4\\,10^7$~\\Msun.  In this paper, we present a unique combined set of nested integral-field spectroscopic observations of \\NGC{3377}, from \\sauron/William Herschel Telescope (WHT) and \\oasis/Canada-France-Hawaii Telescope (CFHT), on which we will base our dynamical modeling.  \\citet{sauron-paper1} presented the large-scale two-dimensional stellar kinematics of \\NGC{3377}, obtained with the panoramic integral-field spectrograph \\sauron as part of a representative survey of nearby early-type galaxies \\citep{sauron-paper2}. The resulting kinematic maps cover $32\\arcsec\\times 43\\arcsec$, with an effective spatial resolution of $\\sim 2\\arcsec$~FWHM, and reveal modest but significant deviations from axisymmetry, not only in the stellar motions, but also in the morphology and kinematics of the emission-line gas. This is unexpected, as lower-luminosity steep-cusped systems such as \\NGC{3377} were assumed to be axisymmetric \\citep[e.g.][but see \\citealt{2002ApJ...567..817H}]{1983ApJ...266...41D, HST-centersIII, Valluri98}.  Here we also present for the first time observations of the inner $5\\arcsec\\times 6\\arcsec$ of \\NGC{3377} with a spatial resolution of $\\sim 0\\farcs6$~FWHM, obtained at the CFHT with \\oasis in its adaptive-optics-assisted mode. The resulting nested set of high-quality integral-field maps make \\NGC{3377} a nearly ideal case for detailed dynamical modeling, aimed at determining the stellar mass-to-light ratio $M/L$, the internal orbital structure, and also the BH mass \\Mbh.  In this paper, we construct axisymmetric dynamical models which incorporate all the integral-field kinematic data. Our approach is similar to that followed by \\citet{sauron-M32} for \\objM{32}, but we use an independent version of the modeling software, include all the available integral-field datasets and study the full line-of-sight velocity distributions (LOSVDs). By construction, we have to ignore any signature of non-axisymmetry, but we will include them in triaxial models in a follow-up work, in the way done by \\cite{2003astro.ph..1070V} for similar nested observations of \\objNGC{4365}. The comparison of the axisymmetric and triaxial models will help clarify which (if any) of the parameters based on the axisymmetric modeling are robust. This is relevant because many of the nearly twenty galaxies for which BH masses have been determined by means of (edge-on) axisymmetric models \\citep{2002ApJ...574..740T, G3I-03} show similar signs of non-axisymmetry.  This paper is organized as follows. \\refsec{sec:photometry} presents the photometric data and the mass model adopted for \\NGC{3377}, and \\refsec{sec:kinematics} describes the kinematic data. The dynamical models are detailed in \\refsec{sec:dynamical-models}, and discussed in \\refsec{sec:discussions}. \\refsec{sec:concl} presents the conclusions.  % ============================================================ % Photometry & mass model % ============================================================ ",
        "conclusions": "\\label{sec:concl}  We have presented a unique set of nested integral-field spectroscopic observations of the E5-6 `disky/boxy' cuspy galaxy \\NGC{3377}, obtained with the \\oasis and \\sauron instruments. Both sets display regular stellar kinematics, with indications for an inner disk, coherent with the diskyness of the light distribution up to $\\sim 30\\arcsec$. However, the IFS observations also reveal a significant twist of the line of zero velocity of $\\sim 10^{\\circ}$ up to $r \\approx 10\\arcsec$ --~with no corresponding isophotal twist~-- indicative of a moderately triaxial intrinsic shape or of a central bar.  Based on this complementary and intrinsically coherent dataset, we have constructed general (three-integral) axisymmetric dynamical models for this galaxy using the orbit superposition technique of \\schw, non-parametrically fitting the full LOSVDs delivered by the 3D spectrographs. In order to properly take into account the observed non-axisymmetry of the kinematic maps, we have added a systematic component to the Monte-Carlo-computed statistical errors on the LOSVDs used in the fit. The best-fit model has a BH mass of $\\Mbh = 7_{-5}^{+4}\\,10^7$~\\Msun{} and a mass-to-light ratio $\\MLI = 2.1 \\pm 0.2$ (99.7\\% confidence level) for an assumed distance of $D = 9.9$~Mpc.  The use of fully coherent IFS measurements on the observation side, and of proper errors and regularization constraints on the modeling side allowed us to derive this significantly improved constraints on both \\ML[]{} and \\Mbh. However, some assumptions used in the modeling --~the most important presumably being the axisymmetry~-- are in noticeable contradiction with our IFS observations. Therefore, BH mass estimates for \\NGC{3377} based on axisymmetric models should be considered with caution.  % Ackowledgements ------------------------------------------------------------"
    },
    "0311/astro-ph0311562_arXiv.txt": {
        "abstract": "We study the evolution of the cosmological perturbations after inflation in curvaton models where the non-relativistic curvaton decays into both radiation and a cold dark matter component. We calculate the primordial curvature and correlated isocurvature perturbations inherited by the radiation and cold dark matter after the curvaton has decayed. We give the transfer coefficient in terms of the initial curvaton density relative to the curvaton decay rate. ",
        "introduction": "Observational results obtained by a variety of recent cosmic microwave background and large scale structure experiments have greatly improved our picture of the Universe on the largest observable scales. These observations are consistent with a primordial spectrum of adiabatic, Gaussian, scale-independent perturbations~\\cite{trotta,amen,lewis,wmap,gauss_test,valvi,julien,gm,sdss}. An attractive explanation is that these primordial perturbations in the early radiation-dominated era originate from quantum fluctuations in an earlier inflationary era where small scale vacuum fluctuations can be stretched up to arbitrarily large scales \\cite{LLBook}. Although it is often assumed that primordial perturbations originate from adiabatic fluctuations of a slowly-rolling scalar field driving the inflationary expansion (the inflaton), {\\em any} light scalar field (with effective mass less than the Hubble scale) will acquire an almost scale-invariant spectrum during an almost exponential expansion. Whether or not these isocurvature fluctuations in non-inflaton fields produce a primordial curvature perturbation depends upon whether or not the local value of that scalar field subsequently affects the local radiation density \\cite{Mollerach,LM}.  In the curvaton scenario \\cite{curvaton,MT,Enqvist} a light, weakly coupled scalar field remains decoupled from the inflaton during inflation and remains decoupled after the end of inflation. Once the Hubble scale drops below the effective mass of the curvaton, the field oscillates and eventually decays. But if the decay is late enough then the density of the non-relativistic curvaton may have grown large enough to contribute significantly (more than about 1\\%) to the total energy density of the universe. Then the local radiation density after curvaton decay depends upon the local value of the curvaton at the end of inflation and it is the initial isocurvature perturbation of the curvaton field that determines the primordial density perturbation. This scenario leads to new possibilities for building particle physics models of inflation \\cite{modelbuilding} and qualitatively different constraints upon the inflationary parameters \\cite{modelconstraints}.  In the conventional slow-roll inflaton scenario there is a remarkable simplification when calculating observational constraints: one can calculate the primordial perturbations directly in terms of the inflaton field fluctuations during inflation, independently of the detailed physics of the end of inflation, reheating and the subsequent radiation-dominated era.  This is because the inflaton fluctuations describe {\\em adiabatic} perturbations about the background homogeneous cosmology on large scales and one can define a conserved curvature perturbation $\\zeta$ directly from energy conservation \\cite{WMLL,conserved}.  By contrast, fluctuations in the curvaton field describe {\\em non-adiabatic} perturbations about the background cosmology and the total curvature perturbation $\\zeta$ is no longer conserved on large scales. The non-adiabatic nature of the initial perturbations leaves open the possibility of residual isocurvature perturbations and detectable non-Gaussianity in the curvaton scenario \\cite{curvaton,MT,LUW,cdm}. One can define individual curvature perturbations for the radiation and curvaton fluids that remain approximately conserved so long as the fluids can be assumed to be non-interacting \\cite{WMLL,conserved}. This has been used to derive an approximate calculation of the resulting primordial perturbations in the curvaton scenario in terms of the curvaton density at the decay time in what has been called the sudden-decay approximation \\cite{curvaton,LUW}.  However an accurate calculation of the resulting primordial perturbations requires numerical solution of the perturbation equations. This was performed in \\cite{MWU} for a two-fluid model where non-relativistic curvaton particles decay into radiation. In this paper we extend this analysis to a three-fluid system to include the possibility that the curvaton also decays into non-relativistic cold dark matter (CDM). This allows us to calculate the residual isocurvature perturbation that may be inherited by the matter in this scenario. See Ref.~\\cite{MT} for a numerical solution of the perturbed curvaton equations in a field theory description. Analytic estimates of the amplitude of the residual isocurvature perturbation in other models for the CDM were given in Ref.~\\cite{cdm}.  In section~\\ref{sec:back} we set up the coupled evolution equations for a system of dimensionless variables describing the evolution of the background homogeneous cosmology. We then give the coupled evolution equations for linear density perturbations in each fluid on uniform-curvature hypersurfaces in section \\ref{sec:perteq}. In an appendix we present equations describing the evolution of curvature and isocurvature perturbations in manifestly gauge-invariant form and discuss the appearance of singularities for some gauge-invariant quantities in a system such as this that includes energy transfer between fluids. This is seen to be due to a singular choice of hypersurface rather than a breakdown of perturbation theory.  The primordial perturbations that result from an initial curvaton perturbation can be characterised by a transfer coefficient that is a unique function of a single dimensionless parameter that specifies the initial density of the curvaton relative to the curvaton decay rate. We give an accurate analytic approximation (to better than 1\\%) for the curvature transfer coefficient. The same coefficient also determines non-Gaussianity of the primordial perturbations \\cite{curvaton} and, in our three-fluid model, the amplitude of the residual isocurvature perturbation. We are able to derive, as an exact result, a simple consistency relation between the ratio of isocurvature to curvature perturbations and the non-Gaussianity of the primordial perturbations, which was previously derived only using the sudden-decay approximation.  We present our conclusions in section \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}   In this paper we have studied the coupled evolution of background densities and linear perturbations of three interacting fluid components: radiation, a curvaton field and cold dark matter.  We solve numerically the evolution equations for density perturbations on spatially flat hypersurfaces to calculate the spectrum of curvature and isocurvature (or relative entropy) perturbations that result from initial perturbations in the curvaton field.  Whereas the resulting radiation curvature perturbation, $\\zeta_{\\gamma,{\\rm out}}$, is dependent on the initial values of background parameters (such as curvaton density), we find the simple result that the final CDM curvature perturbation, $\\zeta_{\\rm m,out}$, is always exactly equal to the initial curvaton perturbation $\\zeta_{\\sigma,{\\rm in}}$, in a model where the CDM is directly produced by curvaton decay. We are able to demonstrate this analytically by constructing a conserved quantity, $\\zeta_{\\mbox{\\scriptsize comp}}$, that at early times is equal to the curvaton perturbation and at late times is equal to the CDM curvature perturbation. In effect this solves the evolution of one degree of freedom and allows us to solve numerically only for $\\zeta_\\gamma$ as a function of input parameters.  The background evolution of the fluids is described by a trajectories in phase-space distinguished by a dimensionless parameter $p=[\\Omega(H/\\Gamma_1)^{1/2}]_{\\rm in}$. The ratio $r$ between $\\zeta_{\\gamma,{\\rm out}}$ and $\\zeta_{\\sigma,{\\rm in}}$ is a function solely of this trajectory parameter $p$ and is well-fit by a simple numerical function given in Eq.~(\\ref{eq:defrfit}). This coincides with the ``sudden-decay approximation'' \\cite{curvaton,MWU} for $p\\ll1$.  The resulting radiation and CDM perturbations after curvaton decay can then be used to predict the spectrum of CMB anisotropies in a given model where the CDM is produced directly as a curvaton decay product. CMB constraints are usually quoted in terms of the primordial curvature $\\zeta_\\p$ and an isocurvature perturbation $\\S_\\p$ around the epoch of primordial nucleosynthesis, leading to \\begin{equation} \\label{thematrix} \\left( \\begin{array}{c} \\zeta \\\\ \\S \\end{array} \\right)_\\p = \\left( \\begin{array}{cc} 1 & r/3 \\\\ 0 & (1-r) \\end{array} \\right) \\left ( \\begin{array}{c} \\zeta \\\\ \\S \\end{array} \\right)_{\\rm in} \\,. \\end{equation} In the curvaton scenario $\\zeta_{\\rm in}=\\zeta_{\\gamma,{\\rm in}}=0$ and $\\S_{\\rm in}=3(\\zeta_{\\sigma,{\\rm in}}-\\zeta_{\\gamma,{\\rm in}})=3\\zeta_{\\sigma,{\\rm in}}$. Although the model we have studied describes the production of CDM directly from the curvaton decay, one would expect a similar result to hold for models where the net baryon or lepton number is directly produced from curvaton decay.  The matrix equation~(\\ref{thematrix}) is a particular case of the general form of the transfer matrix for curvature and isocurvature perturbations presented in Ref.~\\cite{amen,wbmr} where the contribution of the initial isocurvature perturbation to the resulting primordial curvature and isocurvature perturbation is model dependent. For instance in a curvaton scenario where the CDM is a thermal relic which decouples from radiation sometime after the curvaton has decayed we have instead \\begin{equation} \\label{theadmatrix} \\left( \\begin{array}{c} \\zeta \\\\ \\S \\end{array} \\right)_\\p = \\left( \\begin{array}{cc} 1 & r/3 \\\\ 0 & 0 \\end{array} \\right) \\left ( \\begin{array}{c} \\zeta \\\\ \\S \\end{array} \\right)_{\\rm in} \\,. \\end{equation}  The matrix equations (\\ref{thematrix}) and (\\ref{theadmatrix}) using the numerically determined function $r(p)$ [or its approximate form $r_{\\rm fit}(p)$] enable one to relate CMB observations to curvaton model parameters such as the curvaton decay rate, $\\Gamma$, and the curvaton density, $\\rho_\\sigma$, at some early time, characterised by the value of the Hubble rate $H$, and the initial spectrum of curvaton perturbations $\\zeta_\\sigma=(1/3)\\delta\\rho_\\sigma/\\rho_\\sigma$.  For example, in the model where the CDM is directly produced by curvaton decay we confirm \\cite{curvaton} that the relative amplitude of primordial isocurvature to curvature perturbations is a fixed ratio $3(1-r)/r$ which is a function solely of the background parameter $p$ defined in Eq.~(\\ref{eq:defp}). This holds independently of the ``sudden-decay approximation'' previously used.  In any curvaton model the transfer coefficient $r$ is independent of scale in the large-scale limit (by definition) so that the scale-dependence of the primordial perturbation spectra are inherited directly from the scale dependence of the curvaton perturbation at the end of inflation \\cite{curvaton,wbmr} \\begin{equation} \\label{nsigma} \\Delta n_\\sigma \\equiv \\frac{d\\ln{\\cal P}_{\\delta\\sigma}}{d\\ln k} = - 2\\epsilon + 2\\eta_\\sigma \\,. \\end{equation} In this expression $\\epsilon\\equiv -\\dot{H}/H^2$ is the usual dimensionless slow-roll parameter during inflation and $\\eta_\\sigma\\equiv m^2_\\sigma/3H^2$ is the effective mass of the curvaton relative to the Hubble scale during inflation. Both the curvature and isocurvature perturbation (if any) originate from initial curvaton perturbation and are thus completely correlated, independent of the physics of the curvaton decay.  In the absence of any additional source for the isocurvature perturbations, the curvature and any isocurvature perturbations must share the same spectral tilt, given by Eq.~(\\ref{nsigma}).  Note that the curvature transfer coefficient $r$ can also be related to the non-Gaussianity of the primordial perturbations. If $r$ is small then the initial curvaton perturbation $\\zeta_{\\sigma,{\\rm in}}$ must be correspondingly large in order to produce a given amplitude of primordial curvature perturbation. This leads to a larger contribution from second-order perturbations in the curvaton density, which shows up in the non-linearity parameter \\cite{komsper,gauss_test} \\begin{equation} f_{\\rm NL} \\approx \\frac5{4r} \\,, \\end{equation} for $r\\ll 1$ \\cite{curvaton}. (See \\cite{bmr-nongauss} for corrections when $r\\sim1$.)  There is thus a possible consistency test of this curvaton model using two observables.  Existing constraints on the relative amplitude of primordial isocurvature to curvature perturbations \\cite{lewis} constrain $r>0.9$ in models where the CDM is directly produced by curvaton decay, and hence limit the allowed non-Gaussianity. However in curvaton models which produce purely adiabatic perturbations ($\\S_\\p=0$), determinations of the non-linear parameter $f_{\\rm NL}$ may be the only way to determine independently both the transfer parameter $r$ (and hence the model parameter $p$) and the initial amplitude of the curvaton perturbation $\\zeta_\\sigma$."
    },
    "0311/astro-ph0311081_arXiv.txt": {
        "abstract": "{ In radiative layers of rotating stars the luminosity carried by circulation currents through a surface of constant entropy (circulation luminosity) is shown to be positive. The corresponding decrease in the temperature gradient is important in the secondary of contact binaries. This result removes the deadlock in the theory of contact binaries.  The resulting treatment of contact binaries is investigated, assuming thermal equilibrium. The sources of the circulation luminosity in the secondary can be written as the product of a circulation function (a normalized non-negative function of the fractional mass) and an amplitude. If the amplitude is adjusted to give a prescribed temperature difference $\\Delta T_{\\rm e}=T_{{\\rm e}1}-T_{{\\rm e}2}$, the choice of the circulation function is (in a broad range) unimportant. This invariance extends in a close approximation to all observable properties as well as to the internal structure. The temperature difference is bound to be positive. The fractional extent of radiative regions is larger in the secondary than in the primary. In the course of evolution the period increases and the mass ratio decreases. Comparing thermodynamic quantites on level surfaces, pressure and density are larger in the secondary than in the primary. The specific entropy is larger in the primary than in the secondary. The temperature difference is remarkably small and almost vanishing when averaged over the level surfaces occupied in both components.  The only free parameter (apart from $\\Delta T_{\\rm e}$) is the efficiency $f_{\\rm E}$ of the energy transfer from the primary to the secondary. Using standard values for the parameters, a survey of unevolved and evolved contact configurations is presented. Observational tests are passed. In stable systems the degree of contact is small. Stable systems in the period-colour diagram, unevolved and evolved, cover the strip (and only the strip) of observed systems in this diagram. Lower limits for period and effective temperature, compatible with the observed limits, are caused by the requirement of thermal stability. Stable systems with mass ratios very close to unity are possible, in accordance with recent observations. Since stability considerations are essential in these observational tests the results support the assumption of thermal equilibrium as well as the treatment of the stability problem.  Models for individual observed systems with reliable data are well-determined (apart from some freedom in $\\Delta T_{\\rm e}$) and can be used  to calibrate the efficiency and to determine metallicity and age. All models obtained so far are stable. This again supports the assumption of thermal stability. The results show that evolutionary effects are important and that the efficiency is very small ($f_{\\rm E}=10^{-3}\\ldots 10^{-5}$).  Arguments are presented that the velocity field in the common envelope has a reversing layer, with motions from the secondary to the primary in the layers just above the critical surface and from the primary to the secondary in the surface layers. ",
        "introduction": "\\label{intro}  This paper is the third one in a systematic attempt to overcome the deadlock in the theoretial treatment of contact binaries. In the first paper (K\\\"ahler \\cite{k02a}, hereafter K1) we discussed the uncertainties in the structure equations. Serious uncertainties were found to concern only the energy sources and sinks associated with internal mass motions. In the second paper (K\\\"ahler \\cite{k02b}, hereafter K2) we made the usual assumption (hereafter assumption A) that sources and sinks occur only in the common envelope. Investigating a typical system we found a variety of solutions which are reasonable from an observational point of view, but all of them suffer from theoretical deficiencies. In particular, all solutions with good light curves (i.e. a small temperature difference between the components) are of the type proposed by Moss \\& Whelan (\\cite{mwh73}), i.e. the energy sources in the secondary's envelope occur in the extreme superadiabatic layers. As shown by Hazlehurst (\\cite{ha74}) this is impossible since the heat capacity of these layers is too small.  The results in K2 show that assumption A is too restrictive. When thermal equilibrium is assumed this conclusion is cogent since it does not depend on details, e.g. the choice of a transport equation for the energy transfer from the primary to the secondary. In the case of thermal cycles there is more freedom since they depend on the transport equation as well as on the equation governing the rate of mass exchange. Nevertheless the experience of decades shows that the conclusion is justified also in this case. Associated to a cycle is a configuration in unstable thermal equilibrium. If this configuration is in shallow contact (and if extreme superadiabatic transfer is excluded) the temperature difference is large for a considerable part of the cycle. If the configuration is in good contact the temperature difference is always large. In both cases the light curve is bad.  Assumption A turns out to be a deadlock in the treatment of contact binaries. Energy sources or sinks in the interior of at least one component and corresponding changes in the temperature gradient are necessary. According to Lucy (\\cite{lucy68}) there is no freedom in the structure of radiative regions since the luminosity carried by circulation currents through a level surface vanishes (hereafter Lucy's theorem). The only remaining freedom concerns the extent of turbulent regions. A decrease of a convective envelope is in conflict with Lucy's theorem. An increase of the turbulent envelope in the secondary is in conflict with Roche geometry. An increase in the primary is compatible with Roche geometry, but the resulting model (K\\\"ahler \\cite{k89}, K\\\"ahler \\& Fehlberg \\cite{kf91}) is in conflict with the second law of thermodynamics (Hazlehurst \\cite{ha93}) since (in a large region) internal mass motions transfer energy from cooler to hotter layers.  Accordingly, the deadlock boils down to Lucy's theorem. This theorem has been derived --- also by Mestel (\\cite{mes}) and Roxburgh et al. (\\cite{rox}) --- assuming that all thermodynamic quantities are constant on equipotential surfaces, i.e. assuming pseudo-barotropic layers in strict hydrostatic equilibrium. Although the theorem is usually taken for granted (e.g. Maeder \\& Zahn \\cite{maeza}), it is valid only in the zero-order approximation and not in higher approximations (Mestel \\cite{mes}). In contact binaries departures from strict hydrostatic equilibrium (and thus baroclinic effects) are necessary for the energy transfer from the primary to the secondary, and baroclinic effects are expected to be important also in the interior of the components since the variation of the gravity on a level surface is large. For these reasons it is likely that Lucy's theorem does not apply. As observed also by Tassoul \\& Tassoul (\\cite{tass95}) in a discussion of rotating stars there is no reason to claim that the net flux of energy carried by the circulation currents through each level surface vanishes.  Removing the deadlock and assuming thermal equilibrium we are able to predict the modification of the original Lucy (\\cite{lucy68}) model which solves the contact binary problem. Roche geometry requires that (in comparison with the Lucy model) either the (averaged) temperature gradient in the primary be increased or that the gradient in the secondary be decreased. An enhanced gradient is in conflict with the second law. Accordingly we need a reduced gradient in the secondary and thus a reduced radiative luminosity and/or a reduced extent of convective layers. The reduction is possible if some part of the luminosity in the secondary's interior is carried by circulation currents to the common envelope.  In this paper we investigate the resulting treatment of contact binaries, assuming that thermal equilibrium is enabled by circulation currents in the secondary's interior. We discuss the effects of circulation, present a survey of unevolved and evolved theoretical configurations, and perform observational tests. Having obtained models for individual observed systems we shall finally return to the hydrodynamic problem.      \\begin{figure} \\psscalefirst \\psfig{figure=H4648.f1,width=8.8cm} \\caption[]{The degree of contact $F$ in dependence on the  amplitude $f$ of the circulation in System~1, for different values of the parameter $a$ (see text).} \\label{fig1} \\end{figure}  \\begin{figure} \\psscalefirst \\psfig{figure=H4648.f2,width=8.8cm} \\caption[]{The temperature difference $\\Delta T_{\\rm e}$ in dependence on the  amplitude $f$ in System~1, for different values of the parameter $a$.} \\label{fig2} \\end{figure}    \\begin{figure} \\psscalefirst \\psfig{figure=H4648.f3,width=8.8cm} \\caption[]{$\\Delta T_{\\rm e}$ as a function of $f$ for $a=0.95$ in System~1. Branches~1 and 3 (solid lines) of the linear series are stable, while branch~2 (dashed line) is unstable (see text). } \\label{fig3} \\end{figure} ",
        "conclusions": "\\label{concl}  The present investigation was based on the assumption that the circulation luminosity in the secondary's interior is positive and large enough to enable thermal equilibrium. The first part of this assumption can be abandoned. As shown in the preceding section the circulation luminosity is certainly positive in the radiative interior, not only in the secondary but also in the primary. In systems with intermediate or small mass ratios the effects in the primary are small while the effects in the secondary are important since the core luminosity is small. In systems with large mass ratios (e.g. OO~Aql) the neglect of the circulation luminosity in the primary introduces a slight error.  The resulting treatment of contact binaries is simple. Sources of the circulation luminosity deep in the secondary's interior are necessary, but details are unimportant. Free parameters are the temperature difference between the components $\\Delta T_{\\rm e}$, bound to be positive and small, and the efficiency $f_{\\rm E}$ of the energy transfer between the components. The fractional extent of radiative regions is larger in the secondary than in the primary. In the course of evolution the period increases and the mass ratio decreases.  Adopting standard values for the parameters, a survey of unevolved and evolved contact configuration has been obtained. The results are successful in explaining the basic observational facts, in particular the preference for shallow contact, the minimum period and the minimum temperature, the existence of stable configurations with mass ratios very close to unity, and the period-colour relation. Since stability considerations are essential in these observational tests we consider the results as strong arguments not only in favour of thermal equilibrium but also in favour of the treatment of the stability problem.  Assuming thermal equilibrium, models of individual observed systems with reliable data are well-determined apart from some freedom in $\\Delta T_{\\rm e}$. When stable they can be used to to calibrate the efficiency and to determine metallicity and age. The models obtained so far are stable, which lends again support to the assumption of thermal equilibrium. Evolutionary effects are important. The efficiency is very small ($f_{\\rm E}=10^{-3}\\ldots 10^{-5}$).  Concerning the hydrodynamic problem, arguments have been presented that the velocity field in the common envelope has a reversing layer with motions from the secondary to the primary in the layers just above the critical surface and from the primary to the secondary in the surface layers.  Details of the circulation, in particular the circulation function, can be determined only in a three-dimensional discussion. In comparison with rotating stars the problem is more difficult since the geometry is more complex. In other respects the problem is simpler since rotation law and geometry are known in a close approximation. This problem will be investigated in a forthcoming paper, again assuming thermal equilibrium.  Although we have strong arguments in favour of thermal equilibrium we have no proof, and it is uncertain whether a proof is possible."
    },
    "0311/astro-ph0311048_arXiv.txt": {
        "abstract": "We present the results of high resolution SPH simulations of the evolution of gravitationally unstable protoplanetary disks. We report on calculations in which the disk is evolved using a locally isothermal or adiabatic equation of state (with shock heating), and also on new simulations in which cooling and heating by radiation are explicitly modeled. We find that disks with a minimum Toomre parameter $< 1.4$ fragment into several gravitationally bound protoplanets with masses from below to a few Jupiter masses.  This is confirmed also in runs where the disk is given a quiet start, growing gradually in mass over several orbital times. A cooling time comparable to the orbital time is needed to achieve fragmentation, for disk masses in the range $0.08-0.1 M_{\\odot}$. After about 30 orbital times, merging between the bound condensations always leads to 2-3 protoplanets on quite eccentric orbits. ",
        "introduction": "The 3D disk models initially extend between 4 and 20 AU. The central star is modeled as a softened point mass with a mass of $1 M_{\\odot}$. There are no fixed boundaries; the disk is free to expand and contract and both the central star and the disk can respond to their mutual gravitational field. The disk surface density profile is of the type $\\Sigma \\sim r^{-3/2}$. Disk models have masses between $0.075$ and $0.125 M_{\\odot}$ and initial temperature profiles as those used in MA02. The minimum temperatures are reached at the outer edge of the disk and range from 36 to 60 K. More details on the setup of the initial conditions are explained in MA03. Models are evolved with a locally isothermal or with an adiabatic equation of state; the new runs adopt an adiabatic equation of state with the addition of radiative cooling and heating. In all cases we include a standard Monaghan artificial viscosity to model shocks, with typical values of the coefficients $\\alpha=1$ and $\\beta=2$. The artificial viscosity term appears in both the momentum and the energy equation, hence irreversible shock heating is included, except in the locally isothermal equation of state, where by construction the thermal energy of any given particle is assumed to be constant.  Radiative cooling is implemented using a  prescription similar to that used by Rice et al. (2003a) (see also Pickett et al. 2003); the cooling timescale is simply assumed to be proportional to the local orbital frequency of particles, $t_{cool}(r)=A{\\Omega(r)}^{-1}$. In addition we introduce a density dependent criterion, so that when a region grows beyond a specified threshold, radiative cooling is completely switched off. In the runs presented here the density threshold is fixed at $5 \\times 10^{-10}$ g/cm$^{3}$ --- this is a conservative choice based on the recent calculations by Boss (2001) with radiative transfer, which show that at such densities the temperature evolves in a nearly adiabatic fashion. In runs that are evolved using a locally isothermal equation of state we simply switch to an adiabatic equation of state throughout the disk once such density threshold is reached somewhere (see also MA02).  In the runs with radiative cooling we heat the inner part of the disk by means of another radially dependent term (this goes exponentially to zero at $R=10$ AU) so that a gravitationally stable disk ($Q >2$) develops a temperature profile similar to that used in the initial conditions of the locally isothermal runs (the latter was indeed motivated by the results of the radiative transfer models of Boss (1996) that include irradiation from the central star and compressional heating due to material infalling onto the disk from the molecular cloud, see MA03).  The simulations are run with GASOLINE, a parallel tree-based gravity code with SPH hydrodynamics which uses multistepping to evolve efficiently a wide density range (Wadsley, Stadel \\& Quinn 2003). ",
        "conclusions": "Gravitational instability continues to remain a very attractive mechanism to explain the origin of gas giants, especially those found in extrasolar planetary systems. With our new runs we showed that the mass range of bound condensations arising along the spiral arms is significantly broad once the relevant disk parameters are changed; the outcome is not necessarily ``SuperJupiters'', protoplanets even as small as Saturn can be formed. The general picture is confirmed with runs that directly model radiative heating and cooling; the only necessary requirement for the instability to proceed is that the cooling time must be comparable or only slightly larger than the orbital time at some point of  disk evolution. Once protoplanets are formed, mergers during the first few orbital times and accretion of disk gas over longer timescales are the two ways by which they can grow in mass. How much they can grow will depend a lot on the details of heating and cooling. Our results suggest that, even if the cooling time remains comparable to the orbital time in the outer disk, a mass of $\\sim 10 M_J$ might be an upper limit if gravitational torques cause the dissipation of most of the disk on timescales of about $10^4$ years (see also MA03)."
    },
    "0311/astro-ph0311424_arXiv.txt": {
        "abstract": "The X-ray cluster C0337-2522 at redshift $z=0.59$ hosts in its core a group of five elliptical galaxies. Using N-body simulations we show that a multiple merging event among the five galaxies is expected to take place in the next few Gyrs, forming a central brightest cluster galaxy. We also find indications that dynamical friction heating associated with this event is likely to modify the central slope of the cluster dark matter density profile. ",
        "introduction": "We have identified a group of five ellipticals (Es) located in the core of the X-ray cluster C0337-2522 at redshift $z=0.59$ (ROSAT Deep Cluster Survey~\\cite{rosati}). This system represents a strong candidate for the initial stages of brightest cluster galaxy (BCG) formation through galactic cannibalism~\\cite{ostriker}.  Here, we explore how many of the five galaxies under consideration are expected to merge before $z=0$ and we study the properties of the merger remnant.  We address these questions by using N-body simulations, exploring several initial conditions compatible with the imaging and kinematic information from our ESO-VLT data. The details of observations and numerical simulations are given in~\\cite{nipoti}.  Recent observational studies~\\cite{sand1,sand2} of a few galaxy clusters find central dark matter (DM) density distributions $(\\rho_{\\rm DM}~\\propto~r^{-\\beta})$ flatter ($\\beta \\sim 0.5$) than predicted by cold dark matter simulations~\\cite{navarro,moore} ($\\beta \\sim 1-1.5$). A possible interpretation of this discrepancy is that dynamical friction heating~\\cite{bertin} is effective in flattening the DM cusp. We investigate this hypothesis using our simulations, where the initial DM profile is cuspy ($\\beta=1$), galaxies are deformable, and initial conditions correspond to an observed cluster.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=.4\\textwidth]{nipotiF1.eps} \\end{center} \\caption[]{Final DM, stellar, and total density profiles for a representative simulation.  In this case, the best--fitting inner slope of the final DM profile is $\\beta=0.68$. The solid curve is the initial ($\\beta=1$) cluster DM profile.} \\label{eps1} \\end{figure} ",
        "conclusions": ""
    },
    "0311/astro-ph0311612_arXiv.txt": {
        "abstract": "We have measured the proper motion (PM) and kinematics of the ansae in NGC 7009 using high dispersion echelle spectra and archive narrow band HST images.  Assuming that the ansae are moving at equal and opposite velocities from the central star we obtain a system radial velocity of $-53 \\pm 2$ km.s$^{-1}$, the eastern ansa approaching and the western ansa receding at $v_r=5.3 \\pm 1$ km.s$^{-1}$ with respect to this value.  The PM of the eastern ansa is $28 \\pm 8$ mas.yr$^{-1}$, which with our weighted distance to NGC 7009 of $0.86\\pm 0.34$ kpc gives $V_{exp}=112 \\pm 32$ km.s$^{-1}$. The electron temperature and density in both ansae were determined to be $T_e \\sim 9000 \\pm 400$ K and $n_e \\sim 2300 \\pm 400$ cm$^{-3}$. The dynamic age of the ansae is $\\sim 925 \\pm 260$ yrs. and the implied PM of the central star is $\\mu_{CS}$ = 1 $\\pm$ 0.5 mas.yr$^{-1}$. This is in qualitative but not quantitative agreement with previous work. ",
        "introduction": "Most low and intermediate mass stars become White Dwarfs after a brief phase as Planetary Nebulae (PNe) caused by heavy mass loss on the AGB \\citep{ibr83}. Symmetrical and mildly asymmetrical PNe are explained by versions of the interacting wind model \\citep{kpf78} but the mechanisms causing the strong asymmetries observed in the majority of PNe are still unexplained. That both binaries and magnetic fields probably play a role is now generally accepted. For an overview of formation mechanisms, models, and asymmetries see \\citet{kas00}.  The wealth of different PNe shapes is shown in \\citet{bal87} and \\citet{scm92}, and features have been described in detail by many authors: e.g. point symmetry \\citep{sch93}, bipolarity (e.g. \\citealt{csc95}), multipolarity \\citep{sat98}, BRETS \\citep{lop97} and Fast, Low Ionization Emission Regions or FLIERS \\citep{bal94} are all used to indicate the various observed morphologies.  NGC7009 has long been known to have a complex morphology. \\citet{all41} described the outer features as ansae (handles in Latin), and work by \\citet{rea85}, and \\citet{bpi87} show that these ansae are expanding near the plane of the sky at high velocities ($\\sim 10^2$ km/s) relative to the nucleus.  \\citet{lil65} used photographic plate material to analyze the expansion of NGC7009, including the ansae, for which they obtained 16 $\\pm$ 3 mas.yr$^{-1}$.  In this paper we present the kinematics and proper motion of the ansae using observational material from the HST and CTIO 4m echelle spectrograph, and show that the previous results are qualitatively but not quantitatively correct. ",
        "conclusions": "We conclude that the ansae in NGC7009 --which is at a distance of 860$\\pm$\\,340pc-- move outward at the supersonic velocity of $\\pm$\\,112\\,km.s$^{-1}$ near the plane of the sky (2.6\\,$\\pm$\\,1\\,degrees), have a dynamic age of 920$\\pm$\\,260 yrs, a proper motion of 28\\,$\\pm$\\,8\\,mas.yr$^{-1}$, and that the CS likely moves at 1\\,$\\pm$\\,0.5\\,mas.yr$^{-1}$."
    },
    "0311/astro-ph0311338_arXiv.txt": {
        "abstract": "The most intense monitoring observations yet made were carried out on the Seyfert 1 galaxy NGC 4151 in the optical and near-infrared wave-bands. A lag from the optical light curve to the near-infrared light curve was measured. The lag-time between the $V$ and $K$ light curves at the flux minimum in 2001 was precisely $48^{+2}_{-3}$ days, as determined by a cross-correlation analysis. The correlation between the optical luminosity of an active galactic nucleus (AGN) and the lag-time between the UV/optical and the near-infrared light curves is presented for NGC 4151 in combination with previous lag-time measurements of NGC 4151 and other AGNs in the literature. This correlation is interpreted as thermal dust reverberation in an AGN, where the near-infrared emission from an AGN is expected to be the thermal re-radiation from hot dust surrounding the central engine at a radius where the temperature equals to that of the dust sublimation temperature. We find that the inner radius of the dust torus in NGC 4151 is $\\sim 0.04$ pc corresponding to the measured lag-time, well outside the broad line region (BLR) determined by other reverberation studies of the emission lines. ",
        "introduction": "One attractive idea of unifying active galactic nuclei (AGNs) of types Seyfert 1 and 2, depending on whether strong, broad emission lines are visible or not, hinges upon the existence of a dust torus that surrounds the broad-line region (BLR) \\citep{Ant93}. A direct view of the BLR through the inner hole of the dust torus seen face-on yields a Seyfert 1 spectrum, whereas an obscuration of the BLR by the dust torus seen edge-on results in a Seyfert 2 spectrum. Although this unified model has been supported in many ways already, the strongest justification would come from confirming the existence of the dust torus surrounding the BLR. Since modern technology has not yet reached a point where one might spatially resolve the innermost structure of AGNs, the only method available to uncover it is a lag analysis in light curves, known as reverberation mapping.  The near-infrared continuum ($\\lambda \\ga 2\\ \\mu$m) emission of Seyfert galaxies is considered to be dominated by thermal radiation from hot dust surrounding the central engine, based on the decomposition of spectral energy distribution (SED) of AGNs (e.g. Kobayashi et al. 1993). The dust closest to the central engine has higher temperatures, but cannot survive inside some critical radius at which it begins to sublimate. Accordingly, the distribution of dust should be torus-like, with a hole at the critical radius, and the dust at the innermost point of the torus has the maximum temperature equal to the sublimation temperature. This dust produces thermal emission which has a peak luminosity at the corresponding near-infrared wavelength.  Clear evidence that the hot dust is located some distance from the center of the central engine in AGNs is the existence of a lag time between the UV or optical light curve and the near-infrared light curve of an AGN. It is interpreted in terms of light travel time from the central engine, which almost simultaneously emits UV and optical radiation to the surrounding hot dust, which reprocesses the radiation from the central engine to thermal near-infrared radiation. Previous studies reported such lags of $10-1000$ days for at most ten AGNs \\citep{CWG89,glass92,sitko93,okn93,nelson96,okn99,bari92,nelson01,okn01}, however, the number of targets and the measurement errors were somewhat limited.  To make it possible to carry out much more intense monitoring observations of many AGNs in optical and the near-infrared the MAGNUM (\\underline{M}ulticolor \\underline{A}ctive \\underline{G}alactic \\underline{Nu}clei \\underline{M}onitoring) project installed a 2-meter telescope at the Haleakala Observatories in Hawaii \\citep{kob98a,yoshii02,yoshii03}, and started its operation in early 2001.  In this {\\it Letter}, as an early result of the MAGNUM project, we present the $V(0.55\\ \\mu\\rm{m})$ and $K(2.2\\ \\mu\\rm{m})$ light curves of a bright Seyfert 1 galaxy NGC 4151 for the period including the epoch of flux minimum in 2001, and report a measurement of lag-time between them based on a cross-correlation analysis. Due to the high frequency and high signal-to-noise ratio of our monitoring observations, the measured lag is the most precise ever made, not only for NGC 4151 but also for any other AGNs. Then it was compared with the previously measured lag-times of other AGNs in literature, and interpreted by thermal dust reverberation in NGC 4151. The lag-time between the $V$ and $K$ light curves corresponds to the light travel-time between the central engine and the inner radius of the dust torus, and with this result, we argue in favor of the unified model that NGC 4151 has a dust torus whose innermost radius is well outside the BLR. ",
        "conclusions": "An optical and near-infrared monitoring observation of the Seyfert 1 galaxy NGC 4151 was carried out, and a lag-time between the $V$ and $K$ light curves at the flux minimum in 2001 was estimated at $\\Delta t =48^{+2}_{-3}$ days by a cross-correlation analysis. We present that the lag-time between the UV or optical and the $K$ light curves is correlated with the absolute $V$ magnitude of AGN, and conclude that the $K$ band emission of NGC 4151 is dominated by the thermal radiation of innermost hot dust whose temperature is comparable to dust sublimation, and that the inner radius of the dust torus of NGC 4151 is $\\sim 0.04$ pc, which is well outside the BLR.  The ongoing intense multicolor monitoring observations of NGC 4151 and many other AGNs in the MAGNUM project will not only study the detailed geometry of the dust tori and further examine the unified model, but also constrain the real-time change in size of the dust tori and the time-scales of dust formation and destruction under extreme circumstances in AGNs from measurements of $\\Delta t$ at different epochs along the light curves and the near-infrared color with flux variation."
    },
    "0311/hep-ph0311342_arXiv.txt": {
        "abstract": "Thermal emission rates and mean free paths of photons in a color-flavor locked (CFL) phase of quark matter at high densities and moderate temperatures are evaluated. Our calculations are based on a low-energy effective theory for CFL matter describing Goldstone boson excitations and their electromagnetic as well as strong interactions. In-medium coupling strengths of vector mesons are estimated to be smaller than in vacuum. As a consequence of in-medium modified pion dispersion relations, novel processes such as $\\pi^+ \\pi^- \\rightarrow \\gamma$ and $\\gamma \\to \\pi^+ \\pi^-$ become possible. The total photon emissivity is found to be very large, exceeding contributions from thermal $e^+e^-$ annihilation above temperatures of about 5~\\mev. At the same time, the corresponding mean free paths become very small. Our results imply that the photon flux from the surface of a (hypothetical) CFL star in its early hot stages saturates the black-body limit.  Schematic estimates for the early thermal evolution of the star are also presented. ",
        "introduction": "The discovery of asymptotic freedom, leading to the formulation of quantum chromodynamics (QCD) as the theory of strong interactions, was soon followed by the suggestion that matter at sufficiently high densities consists of a deconfined phase of quarks~\\cite{col:1974}. Only shortly afterwards it was pointed out~\\cite{bar:1977} that the true ground state of cold dense quark matter exhibits color superconductivity (CSC), characterized by diquark condensation with an estimated energy gap $\\Delta$ of the order of 1~\\mev\\ between the highest occupied and the unperturbed quark state at the Fermi surface. Since this magnitude of the gap is rather small for phenomenological applications, CSC subsequently received little attention. The situation changed when reinvestigations \\cite{alf:1997,rapp:1997} using nonperturbative forces (\\eg, instanton-induced) showed that the gap can be substantially larger, $\\Delta\\simeq100~\\mev$ for moderate quark chemical potentials, $\\mu_q\\simeq$~350~\\mev. Similarly large values are obtained from estimates based on perturbative calculations at asymptotically high densities~\\cite{pis:1998,son:1998}. Thus, from the practical point of view, the existence of color superconductivity in compact stars has (re-) emerged as an exciting possibility.  The detailed properties of CSC matter relevant to astrophysical applications depend on the interplay of the quark chemical potential, the $q$-$q$ interaction strength, and the bare masses of the (light) quarks $u$, $d$ and $s$. In particular, for $\\mu_q$ below the (constituent) strange quark mass, only $u$ and $d$ quarks are subject to BCS pairing. The corresponding phase is known as 2-flavor CSC (2SC). In the idealized case where the quark chemical potential is much larger than the strange quark mass ($m_s$), the latter becomes negligible and all three flavors exhibit likewise pairing. The preferred symmetry (breaking) pattern in this phase corresponds to the so-called color-flavor locking (CFL)~\\cite{alf:1998}, since the underlying diquark condensate is invariant only under simultaneous color and flavor transformations. In the present work, we will focus on the CFL phase (for a recent review and a more exhaustive list of references, cf.~\\cite{sch:2003}).  The investigation of compact stellar objects largely relies on electroweak emission spectra over a broad range of frequencies. Thus, emissivities and mean free paths ({\\em mfp}'s) of neutrinos and photons are important quantities to assess bulk properties of matter within compact stars.  It is therefore vital to address the question in how far the existence of CSC inside compact stars affects their properties, such as neutrino and photon emission as well as thermal evolution. Pertinent studies have recently been carried out for CSC phases in various scenarios~\\cite{jai:2002,red:2002,jai:2001}. Ref.~\\cite{jai:2002} mostly addressed neutrino emission for CFL matter at temperatures in the sub-MeV range ($\\sim$0.1~MeV or $10^9$~Kelvin) relevant to the long-term evolution of a compact star. In  Ref.~\\cite{red:2002} the CFL neutrino problem has been investigated for a temperature range of up to 30~\\mev, corresponding to the hottest phases expected to occur during a supernova explosion or in the early stages of a proto-neutron star (see, \\eg, Ref.~\\cite{sha:1983}). Photon emission has been calculated in Ref.~\\cite{jai:2001} in a weak-coupling expansion for (gapped) quarks. The objective of the present article is to evaluate both photon emissivities and \\mfp's in the strong coupling regime of CFL matter in the (tens of) MeV temperature region making use of an underlying effective theory description. The main processes of interest therefore involve scattering and decay processes of \\gb, taking into account their in-medium modified dispersion relations.  For comparison, the effects of thermal electrons and positrons are assessed, which prevail towards lower temperatures. Furthermore, we incorporate generalized vector mesons, relevant at the high-temperature end of the CFL phase. As a lower limit, and in order to minimize the number of poorly known parameters, such as meson masses, we restrict the effective theory description to the $SU(2)$ flavor sector, \\ie, neglect contributions from mesons involving strange quarks (cf.~Ref.~\\cite{Turbide:2003si} for a recent analysis in hot hadronic matter). We will comment on the effects of kaons and other strangeness-bearing mesons in the summary.  The outline of the article is as follows. In Sect.~\\ref{sec:effL} we will recall the basic elements of the effective low-energy Lagrangian which describes the electromagnetic and strong interactions of pseudoscalar \\gb. In Sect.~\\ref{sec:phot} the evaluation of photon emission rates and mean free paths will be presented, including novel aspects specific to the CFL setting. Sect.~\\ref{sec:evo} is devoted to an estimate of the photon-driven cooling behaviour in the early stages of the thermal evolution of a compact star which accommodates CFL matter throughout. Finally, Sect.~\\ref{sec:sum} contains our summary and conclusions. Explicit expressions for medium-modified Feynman rules and process amplitudes have been relegated to the Appendix. ",
        "conclusions": "\\label{sec:sum}  The main objective of this work has been an order of magnitude assessment of photon emission rates and mean free paths in a CFL phase of high-density quark matter. Our approach is based on a low-energy effective theory that describes the prevailing degrees of freedom at large densities and moderate temperatures, $T \\sim \\morder{10 \\mev}$. These are generalized \\gb\\ that appear due to breaking of chiral and baryon number symmetries. We have followed the hidden local symmetry approach to introduce vector-meson excitations as well as the electromagnetic field in our analysis. The assumption of the KSRF relation in medium enabled us to relate the vector-meson self-interaction strength $\\widetilde{g}$ and two-pion coupling $\\grho$ to the characteristic scales of the CFL phase, \\ie, the gap $\\Delta$ and the quark chemical potential $\\mu_q$, according to $\\grho=\\wt{g}\\simeq \\sqrt2\\, \\Delta/f_\\pi$, where $f_\\pi \\simeq 0.21 \\, \\mu_q$. In the presence of in-medium modified pion dispersion relations, electromagnetic gauge invariance could be maintained by requiring both the $\\rho\\,$-$\\,\\gamma$ vertex as well as the rho propagator (with mass $m_\\rho\\simeq 2\\Delta$) in its vacuum form (which emerged as a consequence of neglecting the distinction between temporal and spatial components of the rho mass).  The resulting photon production and scattering amplitudes to lowest order have been convoluted over available phase space in CFL matter using standard kinetic theory expressions. Our results suggest that photon emissivities might become very large, up to $\\morder{10^{53} \\erg \\cm^{-3} \\s^{-1}}$ at the highest temperatures expected to occur during (or immediately following) a supernova explosion, $T\\simeq 30 \\mev$~\\cite{sha:1983}. Over a large range of CSC gap values, $\\Delta_0=$ 50 -- 150~\\mev, we find that at temperatures below $\\sim$5~\\mev\\ the dominant process is $e^+ e^- \\to \\gamma \\, \\gamma$, whereas for larger $T$ electromagnetic pion annihilation, $\\pi^+ \\pi^- \\to \\gamma \\, \\gamma$, is most important. Processes invol\\-ving strongly interacting rho mesons could also play a major role. The temperature where the strong processes begin to outshine electromagnetic ones could not be accurately determined (although it might be well below 30~MeV), mostly due to uncertainties in the vector-meson mass spectrum. An interesting feature of the CFL phase is the appearance of the novel processes $\\pi^+ \\pi^- \\leftrightarrow \\gamma$, which are triggered by the in-medium pion dispersion relation due to the breaking of Lorentz invariance.  We furthermore have found that the photon mean free path is extremely small over a wide range of temperatures. In particular, it is far below the typical radius of a compact star as long as the temperature is above about a few tens of keV. With a gap value of 50 (150)~\\mev, the processes providing the largest opacity for $T$ up to 5 (15)~\\mev\\ are Compton scattering off thermal electrons and positrons, $\\gamma \\, e^\\pm \\to \\gamma \\, e^\\pm$, as well as photon annihilation, $\\gamma \\, \\gamma \\to e^+ \\, e^-$, in agreement with previous calculations in a low-temperature approximation~\\cite{sho:2002b}. At larger $T$, Compton scattering off pions and photon annihilation into pion pairs are dominant.  The comparison with analogous calculations of neutrino properties in CFL matter~\\cite{red:2002} revealed that the photon emissivity is much larger than the neutrino one, whereas the mean free path is substantially shorter. A schematic estimate of the photon flux from a hypothetical CFL star with uniform temperature profile indicates that the photon flux could even exceed the volume-driven neutrino flux (temperature gradients will favor, however, the latter).  We have also presented an order of magnitude analysis of the cooling time of the outermost surface layers of a (hypothetical) CFL star, in the very early stages of its evolution. Since, at the highest temperatures considered here, the thermal conductivity of CFL matter is very small, one may neglect heat transport in the interior of the star. The thus found cooling time is of the order $10^{-16} \\s$. This is to be compared with the diffusion time of a massless Goldstone boson, $t_D \\sim 10^{-14} \\s$, corresponding to the distance given by the photon mean free path at large temperature. These widely different time scales suggests that heat transport does not play a significant role in the very early cooling era.  Let us finally comment on a few aspects of our work that have not been adequately treated and should be addressed in future analyses:  (i) In our calculation of the emission rates and mean free paths we have restricted ourselves to the $SU(2)$-flavor case, \\ie,  did not include contributions from strangeness-carrying mesons in the pseudoscalar and vector octet. The analysis of Ref.~\\cite{sch:2001} showed that, while the $K^-$ is slightly heavier than the pion, the $K^+$ is considerably lighter. Since the sum of the kaon masses is less than the sum of the pion masses, one can expect that the emission rates from $K^+ K^-$ annihilation into photons will exceed the emission rates resulting from $\\pi^+ \\pi^-$ annihilation. Likewise, the kaon induced mean free path is expected to be smaller than the one induced by pions. Following our arguments of Sect.~\\ref{sec:neutrinos} we can assume, however, that the overall result for the flux will remain unchanged, being limited by black-body emission. At moderate densities, where the pion decay constant is of the order of the gap, one can also expect solitonic excitations in the baryon sector~\\cite{Jackson:2003dk} to become relevant.  (ii) Our discussion of the cooling behavior of a compact star featuring CFL matter was limited to the very early stages of the thermal evolution, and we have furthermore made the simplifying assumption that the star accommodates a CFL phase which extends all the way up to the surface.  We will postpone the investigation of the long-term cooling, including heat transport and more realistic scenarios in which the CFL phase is covered with a nuclear crust, to a future publication~\\cite{ouy:2003}.  \\vskip\\baselineskip  {\\bf Acknowledgements.} We would like to thank Emma Olsson, Chris Pethick, Robert Pi\\-sarski, Igor Shovkovy, Kim Splittorff and Dmitri Voskresensky for interesting discussions. We are especially indebted to Francesco Sannino for his continuous encouragement throughout the work. C.V. wishes to thank Deirdre Black and Sanjay Reddy for many enlightening conversations, and the Theory Group at Jefferson Lab for their kind hospitality during his visit, where parts of this work have been done. He also thanks Igor Musatov for providing his program FeynmanGraph. This work is partially funded by the European Commission IHP program under contract HPRN-CT-2000-00130. The research of R.~O. is supported by grants from the Natural Science and Engineering Council of Canada (NSERC).  \\appendix"
    },
    "0311/astro-ph0311507_arXiv.txt": {
        "abstract": "Studies of the broad emission-line regions (BLRs) in quasars have revealed solar or higher enrichment levels up to the highest redshifts. In combination with the presence of large amounts of dust in QSOs at $z \\sim$ 6, this implies that substantial amounts of star formation and nucleosynthesis took place at significantly earlier epochs. Here, we examine whether a top-heavy stellar initial mass function (IMF) is indicated by current data, by modelling the contributions from different regions of the IMF, including Type Ia/II and pair instability supernovae, to the metal synthesis in BLRs. We find that, in order to reproduce the observations of roughly solar values of N/C and Fe/Mg in these objects, (i) stars with a present-day IMF are sufficient, regardless of their metallicity, (ii) zero-metallicity stars with a top-heavy IMF severely underproduce N/C, and (iii) the contribution of Type Ia SNe is not strongly required by the data.  Therefore, stars of mass $\\sim$ 1-40 $M_\\odot$ must have existed at $z \\sim$ 10--20, possibly coeval with any hypothesized stars of masses $\\ga$ 100 $M_\\odot$ at these epochs. This is in agreement with the nucleosynthetic abundance pattern detected in extremely metal-poor stars in the galactic halo. ",
        "introduction": "Multiwavelength observations of increasing quality over the last decade have begun to strongly constrain the link between the earliest epochs of star formation in the universe and the metal enrichment of high-redshift ($z$) protogalaxies and the intergalactic medium (IGM). Constraints on abundances at the highest redshifts currently comes from data on QSOs, where the chemical composition of the gas can be estimated from spectral features from the broad emission-line regions (BLRs) as well as broad and narrow absorption lines, probing regions ranging from sub-pc scales to those exceeding several tens of parsecs.  Considerable progress has been made in recent years in observing and modelling such QSO environments, and detailed chemical evolution models have been developed to better assess reliable abundance diagnostics (Hamann \\& Ferland 1999).  The NV/CIV line ratio has received particular attention because of its high sensitivity to the overall metallicity of the gas. For typical values of BLR parameters, the models predict NV/CIV $\\sim$ 0.1 for solar N/C. The current data indicates that most QSOs have supersolar N/C and BLR gas metallicities, including the highest-$z$ ones at $z \\sim$ 6.28 and 5.99, which have NV/CIV = 0.35 and 0.67 respectively (Pentericci et al. 2002).  In addition, a roughly constant value of FeII/MgII ($\\sim$ 2--15) is detected in QSOs up to the highest redshifts at $z \\sim$ 6.4 (Freudling, Corbin \\& Korista 2003; Maiolino et al. 2003; Dietrich et al. 2003a and references therein). Assuming that FeII/MgII traces Fe/Mg and using FeII/MgII $\\sim$ 3 for solar Fe/Mg, the data imply BLR gas metallicities of near- or exceeding solar values (Dietrich et al. 2003b). Finally, recent submm/mm/IR observations of high-z QSOs have revealed the presence of large amounts of dust (thermal dust masses $\\sim 10^8 M_\\odot$) up to redshifts as high as 6.4 (Bertoldi et al. 2003), implying associated star formation rates of $\\sim 10^3 M_\\odot$ yr$^{-1}$.  Although the sizes and masses of the physical regions probed by BLR and dust observations are orders of magnitude apart, together they yield a consistent scenario: that significant stellar activity and nucleosynthesis occurred at $z \\ga$ 6. At these epochs, the age of the universe is approximately 1 Gyr and approaches the enrichment timescale of low-mass stars and Type Ia supernovae (SNe). These data may imply that star formation in such environments occurred rapidly and involved an initial mass function (IMF) favoring massive stars (Hamann \\& Ferland 1993).  Our aim here is to test whether the observed abundance ratios in combination with the constraints from stellar evolutionary timescales can place interesting limits on the nature of the stellar IMF at early epochs.  In particular, the presence of significant quantities of iron has led some authors to propose a Type Ia SN origin. Alternatively, Type II SNe and pair instability SNe (PISNe; Woosley \\& Weaver 1995; Heger \\& Woosley 2002) from stars of respective mass ranges 10--40 M$_\\odot$ and 140--160$M_\\odot$ could also be substantial sources of iron in QSO BLRs. These massive stars have the added advantages of providing the same nucleosynthesis site for alpha-elements such as magnesium, and of being sources that can create dust and generate prompt metal enrichment.  In this work, we examine the evolution of N/C and Fe/Mg for cases involving a top-heavy and present-day IMF in a burst scenario at early epochs. Our approach here is simple and chosen to reveal whether a preferred stellar mass range is indicated by the available data. For more detailed modelling of various star formation scenarios, we refer the reader to, e.g., Romano et al. (2002), Hamann \\& Ferland (1999) and references therein. ",
        "conclusions": "The study of element evolution in the deep potential wells associated with the BLRs near the most massive detected high-$z$ objects, rather than in the IGM (e.g., Ellison et al. 2000), has the potential advantage of not being subject to the uncertainties of metal transport mechanisms. BLR metals may therefore reflect the original element synthesis more accurately, barring any enriched gas that falls into the central black hole. Both IGM and BLR studies share, however, sizeable modelling uncertainties in translating observed line strengths to the underlying metallicity, which depends on the gas temperature, density and ionization balance.  Within these uncertainties, our results indicate that stars in a present-day IMF, regardless of stellar metallicity, can reproduce the data on those QSO BLR abundance ratios that we have considered here. A new population of stars (top-heavy IMF or otherwise) is not required from such data, although VMSs at $z \\sim$ 10--20 may be found to be attractive candidates to explain other observations, e.g., in order to match the observed amplitude and anisotropy of the near-infrared background (Salvaterra \\& Ferrara 2003; Magliocchetti, Salvaterra \\& Ferrara 2003; Cooray et al. 2003) and/or if the high Thomson optical depth initially detected by $WMAP$ is not revised by future data.  Therefore, the epochs at $z \\ga 6$ may not have been exclusively dominated by a top-heavy IMF and in fact appears to require the presence of stars in a present-day IMF in addition to any other coeval stellar population. Indeed, a prompt transition from a top-heavy IMF to a Salpeter IMF may well have occurred in the dense environments probed by QSO BLRs, as the metals released by the first episodes of star formation could remain confined and efficiently pollute the gas up to $Z_{\\rm cr}$.  Not requiring an exclusive VMS contribution by $z \\sim$ 6 in QSO BLRs would also be consistent with, e.g., the mild evolution of Fe/Mg over $z \\sim$ 0.1--6.4, since metal-free star formation is not observed at these lower redshifts.  If Type Ia SNe are assumed to be the dominant sources of iron, their generation timescales in relation to the BLR Fe abundances at $z \\sim$ 6 places the first star formation epochs at $z \\sim$ 10--20 (Freudling et al. 2003), consistent with $WMAP$'s findings. However, it is unclear if Type Ia SNe can occur when $Z_\\star=0$ stars are involved, and consequently, barring extremely short metal enrichment timescales, whether such SNe play a significant role in early cosmic metal generation. Observations of large amounts of dust in the highest-$z$ QSOs (Bertoldi et al. 2003) provide another avenue to constrain early SNe.  Although the connection between the synthesis timescales of metals and dust is not yet completely understood, it is currently believed that dust could be created from Type II SNe (Todini \\& Ferrara 2001) and/or PISNe (Schneider, Ferrara \\& Salvaterra 2003; Nozawa et al. 2003), but not from Type Ia SNe. As Bertoldi et al. (2003) additionally point out, there may not be enough time to create dust from intermediate-mass stars by $z \\sim$ 6.  Finally, the ashes of the earliest generations of stars may be reflected in the composition of very metal-poor halo stars in the local universe. Observations of such low-mass stellar relics in our Galaxy reveal the presence of elements heavier than Mg which cannot be synthesized in their interiors, and which must therefore have been already present in the pre-enriched gas cloud from which such stars formed.  The recent detection of the most iron-poor star observed to date (Christlieb et al. 2003) lies tantalizingly close to the hypothesized $Z_{\\rm cr}$ for the formation of VMSs (see section 2), although this star is relatively enriched in other metals such as C. This star's elemental abundance pattern indicates that the original starforming cloud was enriched with metals ejected by SNe from $\\sim$ 20--130 $M_\\odot$ stars (Umeda \\& Nomoto 2003, Tumlinson, Venkatesan \\& Shull 2004), and it is only partially consistent with enrichment by PISNe (Schneider et al. 2003a).  This derived mass range for a first-generation stellar population is roughly similar to that which could be responsible for the enrichment of QSO BLRs as shown in this work. Thus, a top-heavy IMF may not be required to explain the metal abundance patterns in QSO BLRs, which is consistent with the current data on extremely metal-poor stars in the galactic halo. It remains to be seen whether future observations of such local stellar relics and of the highest-$z$ objects yield a consistent picture of the sources of the earliest metal synthesis in the universe."
    },
    "0311/astro-ph0311488_arXiv.txt": {
        "abstract": "The realization that the total energy of GRBs is correlated with their jet break angles motivates the search for a similar relation between the peak luminosity, $L$, and the jet break angles, $L\\propto\\theta^{-2}$. Such a relation implies that the GRB luminosity function determines the angular distribution. We re-derive the GRB luminosity function using the BATSE peak flux distribution and compare the predicted distribution with the observed redshift distribution. The luminosity function can be approximated by a broken power law with a break peak luminosity of $4.4 \\times 10^{51}$ erg/sec, a typical jet angle of 0.12 rad and a local GRB rate of $0.44\\,h_{65}^{3}$~Gpc$^{-3}$yr$^{-1}$. The angular distribution implied by $L\\propto\\theta^{-2}$ agrees well with the observed one, and implies a correction factor to the local rate due to beaming of $75\\pm25$ (instead of 500 as commonly used).  The inferred overall local GRB rate is $33\\pm11\\,h_{65}^{3}$~Gpc$^{-3}$yr$^{-1}$.  The luminosity function and angle distribution obtained within the universal structured jet model, where the angular distribution is essentially $\\propto \\theta$ and hence the luminosity function must be $\\propto L^{-2}$, deviate from the observations at low peak fluxes and, correspondingly, at large angles. The corresponding correction factor for the universal structure jet is $\\sim 20 \\pm 10$. ",
        "introduction": "In spite of great progress, one of the missing links in our understanding of GRBs is a knowledge of the luminosity function and the rate of GRBs. The realization that GRB are beamed raised the question of what is the angular distribution. The discovery of Energy - angle relation (Frail et al., 2001; Panaitescu \\& Kumar 2001) suggested that the luminosity function and the angular distributions should be also related.  The number of GRBs with an observed afterglow is still rather limited. Redshift has been measured only for a fraction of these bursts and the selection effects that arise in this sample are not clear. Therefore, at present we cannot derive directly the  GRB luminosity function. A possible method to derive the luminosity function makes use of a luminosity indicator. Several such indicators have been suggested (Norris, Marani \\& Bonnell 2000; Fenimore \\& Ramirez-Ruiz 2000) but the robustness of those indicators is not yet clear.  In a recent paper  Schmidt (2003) shows that a luminosity function obtained using luminosity indicators fits the observations very poorly. An alternative is to consider the observed peak flux distribution and fit possible luminosity functions and burst rate distributions (Piran 1992, Cohen \\& Piran 1995, Fenimore \\& Bloom 1995, Loredo \\& Wasserman 1995, Horack \\& Hakkila 1997, Loredo \\& Wasserman 1998, Piran 1999, Schmidt 1999, Schmidt 2001, Sethi \\& Bhargavi 2001) This can be done by obtaining a best fit for the whole distribution or simply by just fitting the lowest moment, namely, $\\langle V/V_{max} \\rangle$.   Evidence of jetted GRBs arises  from long term radio observations (Waxman, Kulkarni \\& Frail 1998) and from observations of achromatic breaks in the afterglow light curves (Rhoads 1997, Sari, Piran \\& Halpern 1999). However the structure of these jets is still an open question. The two leading models are (1) the uniform jet model, where the energy per solid angle, $\\epsilon$, is roughly constant within some finite opening angle, $\\theta$, and sharply drops outside of $\\theta$, and (2) the universal structured jet (USJ) model (Rossi, Lazzati \\& Rees 2002), where all GRB jets are intrinsically identical, and $\\epsilon$ drops as the inverse square of angle from the jet axis. Within the uniform jet model, the observed break corresponds to the jet opening angle. Frail et al. (2001) and Panaitescu and Kumar (2001) have estimated the opening angles $\\theta$ for several GRBs  with known redshifts. They find that the total gamma-ray energy release, when corrected for beaming as inferred from the afterglow light curves, is clustered. A recent analysis (Bloom, Frail \\& Kulkarni, 2003) on a larger sample confirmed this clustering around $\\sim 1.3\\times 10^{51}$ergs. Within the USJ model the jet break corresponds to the viewing angle and the energy-angle relation given above implies  that the luminosity distribution within the jet is proportional to $\\theta^{-2}$  (Rossi, Lazzati \\& Rees 2002) and therefore the luminosity function has the form $\\Phi(L)\\propto L^{-2}$  (Perna, Sari \\& Frail 2003). Frail et al. (2001) also derive the observed $\\theta$ distribution. Taking into account the fact that for every observed burst there are $f_b^{-1} \\equiv (\\theta^2/2)^{-1}$ that are not observed they derived the true $\\theta$ distribution and used it to estimate the average ``beaming factor'',  $\\langle f_b^{-1} \\rangle \\simeq 500$. Then they multiply this factor times the local (isotropic estimated) GRB rate derived by Schmidt (2001) and obtain the true local rate of GRB.    We aim here at obtaining the combined L-$\\theta$ distribution function, $\\Phi(L,\\theta)$  in the uniform jet model. The present data is not sufficient to accurately constrain $\\Phi(L,\\theta)$. We therefore must make some assumptions.  First we estimate the {\\it isotropic} luminosity function, $\\Phi(L)=\\int d\\theta f_b(\\theta)\\Phi(L,\\theta)$, by fitting the peak flux distribution using a simple parametrization of this function. When considering the angular distribution we also examine (following Perna, Sari \\& Frail 2003) the USJ model  and show that the implied peak flux distribution is somewhat inconsistent with the observed one, with the problem most severe at the low end of the peak luminosity distribution.  We find that the data support $L\\times \\theta^2\\sim {\\rm const}$ (see also Van Putten \\& Regimbau 2003). There is a clustering of the true peak luminosity around $\\sim 3.2 \\times 10^{49}$ erg/s, but note that the distribution is not a narrow delta-function. We notice then that once we have such a relationship the luminosity function implies an angular distribution function. This is true even when some spread in the value of $L\\times \\theta^2$ is taken into account, reflecting the spread observed in current data. The resulting angular and redshift distributions are consistent with the observations. Our analysis is well motivated but approximate. Due to width of the luminosity angle distribution and given the data quality and selection effects, we believe there is no point in rigorous maximal likelihood analysis and the like. We give approximate estimates of the model parameters and approximate uncertainties.    The paper is structured in two main parts. In the first part we rederive (following Schmidt, 1999) the GRB luminosity function. In the second part we consider the angular distribution that follows from the luminosity function and the peak luminosity - angle relation. We compare the predicted angular distribution with the observed one. Using this distribution we re-calculate the correction to the rate of GRBs due to beaming. Unlike Frail et al., (2001) we use a weighted average of the angular distribution. This yields a significantly smaller correction factor. We also estimate a the correction factor needed for the USJ model. Note that this revised correction factor does not depend on the relation between the peak luminosity and the angle.   The observed sample we use to explore the peak luminosity - angle relation is small and its  selection effects  are hard to quantify. Therefore it is difficult to give robust conclusions from our analysis at this stage.  However, future missions and in particular SWIFT  will allow us, hopefully in the near future,  to use this procedure and test our conclusions with larger and less biased samples. ",
        "conclusions": "In this work a combined L-$\\theta$ distribution function of GRBs is derived, $\\Phi(L,\\theta)$, for the uniform jet model. To this aim we have rederived (following Schmidt, 1999) the luminosity function,  $\\Phi(L)$ by fitting the peak flux distribution and used the relation $L\\times \\theta^2\\sim$ const implied by the observation on the sample considered by Bloom et al. (2003). We have compared our results with those obtained in the framework of USJ jet model (Perna et al. 2003) showing that  the luminosity function implied by this model leads to a peak flux distribution some what inconsistent with the observed one.  The luminosity function that best fits the observed peak flux distribution is characterized by two power laws with slopes $\\alpha=-0.1$ and $\\beta=-2$ and isotropic-equivalent break luminosity $L^*\\sim 7.1\\times 10^{51}$ erg/s for a SF2 or $L^*\\sim 4.4\\times 10^{51}$ erg/s for a RR-SFR. Repeating the Schmidt analysis we have found the observed local rate of long GRBs $\\rho_0 \\sim 0.10$ Gpc$^{-3}$yr$^{-1}$ and $\\rho_0 \\sim 0.44$ Gpc$^{-3}$yr$^{-1}$ for the two SFR respectively. We have also shown that with these luminosity functions we find a reasonable agreement with the observed redshift distribution for a RR-SFR, while a SF2 predicts too few bursts at low redshift. Nevertheless since the redshift sample is strongly affected by selection effects that are hard to estimate the possibility that GRBs follow the SF2 cannot be ruled out by this analysis.  Using the sample of Bloom et al. (2003) we have shown that the energy-angle relation (Frail et al., 2001; Panaitescu \\& Kumar 2001)  applies also to the peak luminosity and that there is a clustering of the true peak luminosity around $\\sim 3.2\\times 10^{49}$ erg/s. However the distribution is not a narrow delta-function.  This implies that a luminosity function determines a $\\theta$ distribution as $L$ is related to $\\theta$. This is true even when the spread in  $L_j$ observed in the current data is taken into account. One important result of our work is that the resulting angular distribution found under this assumption is consistent with the observation.  We have re-calculated the correction to the rate of GRBs due to beaming using a weighted average of the predicted angular distribution instead of a simple average over the ``true'' angular distribution as done by Frail et al. (2001) and by van Putten \\& Regimbau (2003). We find a correction factor $75\\pm25$ (see the end of \\S~3). This is significantly smaller than the commonly used correction factor, $\\sim 500$, estimated by Frail et al. (2001) or $\\sim 475$ estimated by van Putten \\& Regimbau (2003). . This correction should also influence the estimated rates of both optical (Nakar, Piran \\& Granot, 2002) and radio (Levinson et al., 2002) orphan afterglows (see \\S~4).   The research was supported by the RTN ``GRBs - Enigma and a Tool\" and by an ISF grant for a Israeli Center for High Energy Astrophysics. DG thanks the Weizmann institute for the pleasant hospitality acknowledges the NSF grant AST-0307502 for financial support. We thank Marteen Schmidt, Ehud Nakar and Cristiano Guidorzi for helpful discussions."
    },
    "0311/astro-ph0311441_arXiv.txt": {
        "abstract": "The integrated profile of \\psr{} can be decomposed into four Gaussian components. The sum of these Gaussian components can be made to replicate the total intensity, linear and circular polarization of the integrated profile, and the discontinuity in the swing of the position angle.  Surprisingly, the single pulses from \\psr{} can also be decomposed into these same Gaussian components with only the amplitude of the components as a free parameter. The distributions of the fitted amplitudes are log-normal. Under the assumption that emission consists of 100\\% polarized orthogonal modes which are emitted simultaneously, I show that the polarization properties of the single pulses can also be replicated in a simple way. This lends support to models involving the superposition of orthogonal modes rather than disjoint emission of the modes. ",
        "introduction": "The integrated pulse profile from pulsars is one of the key characteristics of the emission process. Although single pulses vary widely both in intensity and shape, integrating a few thousand pulses produces a highly stable integrated profile. The integrated profile therefore yields clues as to the permanent structures associated with the pulsar magnetosphere. Kramer et al. (1994)\\nocite{kwj+94} showed that the integrated profile can be decomposed into a small number of Gaussian components. These components represent a 2-D Gaussian region of emission from a fixed height in the patchy pulsar beam \\cite{lm88}. In single pulses, some fraction of the component is emitting when the beam sweeps the line of sight and therefore the single pulses tend to have narrower features (sub-pulses) than those of the integrated profile. The sub-pulse features may also rotate about the magnetic axis, producing the drifting sub-pulse phenomenology \\cite{dr99}. Integrated profiles show high levels of linear polarization and moderate circular polarization. Single pulses tend to be more highly polarized than the integrated profile. The smooth position angle swing in integrated profiles as a function of pulse longitude can occasionally be broken by 90\\degr\\ jumps which shows that the emission is present in two orthogonal modes of radiation. There has been considerable debate in the literature over whether these two orthogonal modes are emitted simultaneously or disjointly. Recently, McKinnon \\& Stinebring (1998)\\nocite{ms98} and Karastergiou et al. (2002)\\nocite{kkjl+02} have shown that simultaneously occuring modes nicely explain many aspects of single pulse data.  \\psr{} is one of the brightest pulsars in the sky at 1.4 GHz. It has a pulse period of 455~ms and its pulse profile consists of a single component preceeded by a low amplitude `precursor'. Previous observations of this pulsar at a frequency of 1.6 GHz by Manchester, Hamilton \\& McCulloch (1980)\\nocite{mhm80} showed that it had moderate linear and circular polarization and that the leading `precursor' had a high degree of linear polarization orthogonal to the polarization of the main pulse emission. Rankin (1990)\\nocite{ran90} and van Ommen et al. (1997)\\nocite{vdhm97} classify the pulsar as a core single even though it has a relatively shallow position angle swing across the pulse. ",
        "conclusions": "I have shown that Gaussian decomposition of the integrated profile not only provides a good fit to the total intensity but can also fit the polarization profile, including the location of the orthogonal jump. In this pulsar, at least, these same Gaussians can be made to fit the single pulse data with only their amplitudes as a free parameter. Interestingly the distribution of these amplitudes is log-normal. Log-normal distributions are seen both in the time-averaged flux density of pulsars and in phase resolved statistics \\cite{cjd01,cjd03}. The implications are then that pulsar emission may be interpreted in terms of stochastic growth theory and that the emission mechanism is linear, involving either a direct linear instability or linear mode conversion of non-escaping waves driven by a linear instability as discussed in detail in Cairns et al. (2003)\\nocite{cjd03}.  Under the assumptions of 100\\% polarized orthogonal modes, the polarization distributions of the single pulse data can also be reproduced from a relatively small number of random variables. This extends the work of McKinnon \\& Stinebring (2000)\\nocite{ms00} by implying that the sub-pulses can be characterised in such a way (not just a particular phase bin) and that log-normal statistics are clearly the dominant statistic for the observed flux density distributions. Any physical model must therefore have as its basis an emission beam which is 100\\% polarized and which is subsequently split through e.g.  bi-refringence into two orthogonally polarized modes. However, it is clear that propagation effects, not taken into account in the simulation process, strongly affect the position angle and the degree of circular polarization seen in single pulses. This situation has also been seen previously by a number of observers and discussed in the context of theoretical models by Petrova (2001)\\nocite{pet01}. In her model, the polarization modes recombine in the `polarization limiting region'. This mode coupling then allows for significant deviations in the position angle and in the circular polarization, with greater position angle deviations leading to increased circular polarization. The data here do not entirely support this picture. Comparison between Figs \\ref{b960} and \\ref{b976} show that the former has a much narrower spread of position angles than the latter but that the percentage V distributions look very similar. The effects of refraction in the magnetosphere are also likely to be important as the two modes are affected differently. This is likely to result in a mode mixture at a given observer's longitude and such an effect would also serve to broaden the position angle distribution.  The applicability of this method to other pulsars is not clear. For example, in the Vela pulsar, Krishnamohan \\& Downs (1983)\\nocite{kd83} showed that four Gaussian components fitted the total intensity data and could account for some of the single pulse data. More recent observations with much higher time resolution \\cite{kjv02} showed that microstructure was prevalent in the single pulse data. Clearly, no amount of Gaussian fitting will reproduce microstructure. Other pulsars, such as PSR B0950+08 show a bewildering variation in the amplitude and location of single pulse components, and there is almost no resemblance between the single pulses and the integrated profile. Many pulsars, such as PSR B1133+16, show orthogonal mode jumps near the peak of a pulse component and again such a feature is hard to reproduce using simple Gaussian decomposition.  However, it may be that there is a class of pulsars similar to \\psr{} for which this method will work very well. One possibility is that the scatter broadening in this pulsar blurs the features with time-scales less than the sub-pulse time. The same was also true of the low resolution observations of the Vela pulsar where the dispersion measure smearing also blurred out the microstructure. Processing techniques which involve smoothing over the microstructure are also possible. For example, van Leeuwen et al. (2002)\\nocite{vlk+02} fitted Gaussian components to PSR B0809+74 in order to study the drift bands and noted how similar the Gaussians were in amplitude and width.  It therefore appears to be possible to get a handle on processes on the scale of tens of milliseconds (the subpulses) without delving into the complexities which occur on timescales significantly shorter. The data seem to favour the idea of simultaneously occuring modes of emission, with propagation through the magnetosphere influencing the degree of circular polarization and the position angle variations."
    },
    "0311/astro-ph0311394_arXiv.txt": {
        "abstract": "We carried out 1.25 pc resolution MHD simulations of the ISM -- including the large scale galactic fountain -- by fully tracking the time-dependent evolution of the magnetic field and the formation of shock compressed regions in a supernova-driven ISM. The simulations show that large scale gas streams emerge, driven by SN explosions, which are responsible for the formation and destruction of shocked compressed layers with lifetimes up to some 15 Myr. The T$\\leq 10^{3}$ K gas is distributed in filaments which tend to show a preferred orientation due to the anisotropy of the flow induced by the galactic magnetic field. The simulations also show that the magnetic field has a high variability, it is largely uncorrelated with the density and it is driven by inertial motions.  The latter is consistent with the fact that ram pressure dominates the flow for $10^{2}<$T$\\leq 10^{6}$ K. For T$> 10^{6}$ K thermal pressure dominates, while for T$\\leq 10^{2}$ K (stable branch) magnetic pressure takes over. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311327_arXiv.txt": {
        "abstract": "We present the results of a preliminary H$\\alpha$ survey of M81 for novae conducted over a 5 month interval using the 5' field of view camera (WFCAM) on the Calypso Telescope at Kitt Peak, AZ.  We observed M81 nearly every clear night during this interval, covering the entire galaxy, and discovered 12 novae.  Our comprehensive time coverage allowed us to produce the most complete set of H$\\alpha$ light curves for novae in M81 to date.  A raw nova rate for M81 gives 23 yr$^{-1}$ which, because of the nature of our survey, is a hard lower limit. An analysis of the completeness in our survey gives a corrected nova rate of 30 yr$^{-1}$.  This agrees well with the rate of 33$^{+13}_{-8}$ yr$^{-1}$, derived from Monte Carlo simulations using nova light curves and survey frame limits.  The spatial distribution of the novae we discovered follows the bulge light much better than the disk or total light according to Kolmogorov - Smirnov tests of their radial distributions.  The asymmetry in the distribution of novae across the major axis line of M81 implies a bulge-to-disk nova ratio of $>$ 9 and supports the idea that novae originate primarily in older stellar populations. ",
        "introduction": "Extragalactic novae are particularly appealing objects for studies of close binary evolution.  They are {\\it the} tracer of interacting close binary stars, visible to much greater distances than any other well studied standard candle except supernovae (novae display M$_v$ up to $-10$).  By studying the novae in a nearby galaxy, one can gather a homogeneous sample of objects, all at the same distance, that is not plagued by the selection effects hampering an analysis of the cataclysmic variable (CV) population in our own galaxy.  Given the large fraction of all stars that exist in binaries, the effort to understand binary formation and evolution has wide reaching implications for understanding stellar populations.  There are currently many unsolved problems in the theory of close binary formation and evolution that are difficult to tackle using CVs in our own galaxy because of the selection effects mentioned above.  The prediction that the nova rate should correlate with the star formation history in the underlying stellar population \\citep{yun97} is one example. This prediction is based on our understanding of how close binaries form and evolve, combined with our understanding of how the mass of a nova progenitor influences the nova outburst.  Massive white dwarfs come from massive progenitors.  These massive white dwarfs erupt as novae frequently as they need only accrete a small amount of hydrogen from their companions to explode as novae.  This, in turn, implies that stellar populations with a low star formation rate should have a correspondingly low nova rate, i.e., early-type galaxies with older stellar populations should have a lower luminosity specific nova rate (LSNR) than late-type galaxies with ongoing star formation.  Efforts have been made to detect a trend in LSNR with galaxy type \\citep{del94,sha00}, but the random errors are too large to date for a meaningful comparison.  Typical extragalactic nova surveys are carried out using short runs, and significant assumptions must be made about the mean lifetimes of novae in order to derive nova rates \\citep{sha01}.  It is also rare that an extragalactic survey has been able to spatially cover an entire galaxy and avoid making some assumption about the distribution of novae with light to derive a galactic nova rate.  The resulting errors in the nova rates are large, but systematic biases are far more pernicious.  In order to accurately compare nova rates with the underlying stellar population, we are pursuing a research program that uses the comprehensive time and spatial coverage afforded by a dedicated observatory.  We have begun our program with M81 and used 5 continuous months of observing time to produce a survey that requires no assumptions about nova distribution or mean lifetime.  Not since \\citet{arp56} surveyed the central parts of M31 has such extensive, continuous coverage of a galaxy for novae been attempted.  Our survey updates the 5 year photographic survey of M81 for novae reported by \\citet{sha99}.  They found 23 novae, evenly divided over the disk and central bulge.  Significant incompleteness must be present in that survey due to photographic saturation in the bright central regions of the galaxy and due to large gaps in their time coverage.  The results we present below show that a comprehensive spatial and temporal survey is required to minimize the systematic effects of incompleteness. ",
        "conclusions": "1) The raw nova rate for M81 provides a hard lower limit at 23 yr$^{-1}$.  Using a set of representative nova H$\\alpha$ light curves in random outburst simulations we derive a completeness corrected nova rate of 30 yr$^{-1}$.  Monte Carlo simulations using the same set of light curves provide our most reliable nova rate for M81 of 33$^{+13}_{-8}$ yr$^{-1}$. The high nova rate found here for M81, with a survey technique that uses comprehensive time and spatial coverage, implies that the nova rates for other galaxies derived from surveys with incomplete spatial coverage and widely spaced epochs are, at best, rough values and may be serious underestimates.  2) The LSNR for M81 is $\\rho_k = 3.96^{+1.8}_{-1.1}$ yr$^{-1} [10^{10} L_{\\odot,K}]^{-1}$ using our best nova rate and the 2MASS K-band photometry for M81.  This raises the LSNR for M81 up by a factor of two from that published by \\citet{sha00}, and implies that the LSNR for other galaxies could be systematically low by similar or greater amounts. A definitive comparison between galaxies of different Hubble type must await the results of comprehensive nova surveys, such as presented here.  3) The cumulative radial distribution of the novae matches the bulge light distribution significantly better than either the total or the disk light distribution, which is ruled out at a 99\\% confidence level.  The BTR value for M81 of 0.25 derived from the asymmetry in the spatial distribution of the apparent bulge novae across the major axis line, implies a bulge-to-disk nova ratio for M81 of $>$ 9, according to the models of \\citet{hat97} for novae in M31.  Both these facts lead to an association of the novae in M81 with the older bulge stellar population.  4) The disk novae reported here, Nova 3 and 10, have decline rates that place them in the slow class of novae, but their maximum magnitudes were not determined.  Nova 6 is most likely a fast hybrid nova, is 1\\minpnt6 from the nucleus of M81, but cannot be unambiguously assigned to either the bulge or the disk.  More comprehensive light curves of both bulge and disk novae are required to test for differences in their average maximum magnitudes and decline rates."
    },
    "0311/astro-ph0311057_arXiv.txt": {
        "abstract": "Turbulent accretion disks are the most probable origin of jets from young stellar objects. Using an extended version of ZEUS-3D code in the axisymmetry option we investigate the evolution of a disk wind into a collimated jet. Resistive jets propagate slower with increasing magnetic diffusivity, and the degree of collimation decreases. The resistive disk included in the simulations ejects persistent, well collimated jet. ",
        "introduction": "Astrophysical jets are observed in objects with an accretion disk. The similarity of jets from the different sources implies that the basic mechanism for the jet formation should be the same.  The mechanism which turns the inflowing, accreting matter into an disk outflow  is still not understood.  However, we know that for lifting the matter into the corona magnetic fields play a major role. In order to model the time-dependent evolution of jet formation we solve resistive MHD equations: \\[ {\\frac{\\partial\\rho}{\\partial t}}\\!+\\!\\nabla\\cdot(\\rho\\vec{u})\\!=\\!0,\\!\\ e\\!=\\!p/(\\gamma-1),\\!\\ \\rho\\left[{\\frac{\\partial\\vec{u}}{\\partial t}}\\!+\\!\\left(\\vec{u}\\cdot\\nabla\\right)\\vec{u} \\right]\\!+\\!\\nabla p\\!+\\!\\rho\\nabla\\Phi\\!-\\!\\frac{\\vec{j}\\times\\vec{B}}{c}\\!=\\!0 \\] \\[ {\\frac{\\partial\\vec{B}}{\\partial t}}\\!-\\!\\nabla\\times\\left(\\vec{u}\\times\\vec{B}\\!-\\!{\\frac{4\\pi}{c}}\\eta\\vec{j}\\right)\\!=\\!0,\\!\\ \\nabla\\cdot\\vec{B}\\!=\\!0,\\!\\ \\frac{4\\pi}{c}\\vec{j}\\!=\\!\\nabla \\times \\vec{B} \\]  \\begin{figure} \\plotone{cemeljicfig1.ps}\\\\ \\vspace{.1cm} \\plotone{cemeljicfig2.ps}\\\\ \\vspace{.1cm} \\plotone{cemeljicfig3.ps}  \\caption{Global evolution of the jet on a grid of (Z$\\times$R)=(280$\\times$40)R$_{\\rm i}$ with a resolution of 900$\\times$200 elements for different magnetic diffusivity. $\\eta$=0,0.01,0.1 for top, middle and bottom panel, respectively. Colors (blue to yellow in decreasing manner) indicate density and lines denote linearly spaced poloidal field lines. The bow shock advances slower with increasing magnetic diffusivity.} \\end{figure}  We perform two types of simulations: {\\bf 1)} The disk is defined as time-independent boundary condition (Ouyed \\& Pudritz 1997; Fendt \\& Elstner 2000; Fendt \\& \\v{C}emelji\\'{c} 2002). Here we investigate the acceleration and collimation of a resistive jet. {\\bf 2)} The evolution of the disk structure is included in the numerical simulation. ",
        "conclusions": ""
    },
    "0311/astro-ph0311261_arXiv.txt": {
        "abstract": "During 2001 and 2002, observations of several X-ray bright clusters of galaxies were conducted using the ACBAR Bolometer Array on the South Pole Viper telescope.  A multi-frequency analysis of these clusters is currently underway. This multi-frequency analysis includes 150, 220 and 275 GHz data from ACBAR, X-ray imaging and spectroscopy from Chandra and XMM-Newton, and weak lensing data from the CTIO 4m Blanco telescope. We describe here how ACBAR can be used to create fully sampled cluster images and present such images for four of the clusters in our sample; Abell 3266, Abell 3827, Abell S1063 and 1E0657-56. In these images, the Sunyaev-Zel'dovich Effect is clearly detected. ",
        "introduction": "The Arcminute Cosmology Bolometer Array Reciever instrument (ACBAR; Runyan et al. 2003) has been operating on the 2.1m millimeter/submillimeter Viper telescope since January 2001. The Viper telescope (Peterson et al. 2000) is located at one of the best sites in the world for millimeter/submillimeter astronomy (Peterson et al. 2002), the South Pole Martin A. Pomeranz Observatory. The ACBAR focal plane is tiled with 16 microlithographed ``spider-web'' bolometers arranged in a $4\\times4$ array. Each bolometer has a Gaussian beam with a $\\simeq4.5'$ FWHM (see Figures 2 and 12 in Runyan et al. 2003). Adjacent beams are separated on the sky by $\\simeq 15'$. During the 2001 observing season, the array was arranged so that each column of four bolometers was sensitive to one of the following mean frequencies; 150, 220, 275 and 345 GHz. During the 2002 season, the 345 GHz bolometers were replaced by 150GHz bolometers. The array was also re-arranged so that now each row, rather than each column, was sensitive to the same frequency. Two rows of 150GHz bolometers, and one row each of 220 and 275 GHz bolometers, were used. The ACBAR observing frequencies straddle the peak in the CMB spectrum and the 220 GHz ``null'' point in the thermal Sunyaev-Zel'dovich effect (Sunyaev \\& Zel'dovich 1972; SZE hereafter). The primary calibrators for the 2001 and 2002 seasons were Mars and Venus respectively. During the periods when these planets were not accessible from the South Pole, various Galactic sources were used to monitor variations in the calibration. The total calibration uncertainty is estimated to be 10\\% (Runyan et al. 2003).  The Galactic source observations were also used to determine the telescope pointing model. The overall pointing error for the 2001 season was $\\simeq1'.4$ in right ascension and $\\simeq20''$ in declination. For the 2002 season, a mechanical chopper encoder was replaced with an optical encoder and the pointing error in the right ascension direction improved to $\\simeq 20''$.  To date, ACBAR data have been applied to three different projects. First, to determine the Cosmic Microwave Background (CMB) power spectrum over the range $150<l<3000$ with a resolution of $\\Delta l=150$ (Kuo et al. 2003).  This spectrum has yielded constraints on a variety of cosmological parameters (Goldstein et al. 2003, Spergel et al. 2003). Second, we are using the 150 GHz maps generated for the CMB anisotropy project to search for high redshift clusters of galaxies via the SZE (Runyan 2002). Several cluster candidates have been selected using a matched spatial filtering technique. These candidates are being followed up with optical and X-ray imaging to confirm their identification and to measure redshifts. Ultimately, the high redshift, or ``blind'', ACBAR cluster survey will be used to derive complimentary (to the CMB power spectrum) constraints on cosmological parameters. Finally, we are using ACBAR, in conjunction with data from X-ray and optical telescopes, to study a small sample of well known, X-ray bright, clusters. We report on the recent progress of this cluster study below. ",
        "conclusions": ""
    },
    "0311/astro-ph0311275_arXiv.txt": {
        "abstract": "{ It has recently been proposed a theoretical model which accounts for the origin of radio mini--halos observed in some cooling flow clusters as related to electron re--acceleration by MHD turbulence (Gitti, Brunetti \\& Setti 2002). The MHD turbulence is assumed to be frozen into the flow of the thermal ICM and thus amplified in the cooling flow region. Here we present the application of this model to a new mini--halo candidate, the cluster A2626, and compare the results with those obtained for the mini--halo in the Perseus cluster. We present VLA data at 330 MHz and 1.5 GHz of the diffuse radio emission observed in A2626, and we show that its main properties can be explained by the model. We find that the power necessary for the re--acceleration of the relic electron population is only a factor $\\sim 0.7$ \\% of the maximum power that can be extracted by the cooling flow (as estimated on the basis of the standard model). We also discuss the observational properties of known mini--halos in connection with those of host clusters, showing that the radio power of mini--halos increases with the maximum power of cooling flows. This trend is expected in the framework of the model. Possible effects of new {\\it Chandra} and {\\it XMM--Newton} estimates of $\\dot{M}$ on this trend are considered: we conclude that even if earlier derived cooling rates were overestimated, cooling flow powers are still well above the radio powers emitted by mini--halos. ",
        "introduction": "A number of clusters of galaxies show extended synchrotron emission not directly associated with the galaxies but rather diffused into the intra--cluster medium (ICM). These radio sources have been classified in three classes: cluster--wide halos, relics and mini--halos (Feretti \\& Giovannini 1996). Both cluster--wide halos and relics have low surface brightness, large size ($\\sim$ 1 Mpc) and steep spectrum, but the former are located at the cluster centers and show low or negligible polarized emission, while the latter are located at the cluster peripheries and are highly polarized. They have been found in clusters which show significant evidence for an ongoing merger (e.g., Edge, Stewart \\& Fabian 1992; Giovannini \\& Feretti 2002). It was proposed that recent cluster mergers may play an important role in the re--acceleration of the radio--emitting relativistic particles, thus providing the energy to these extended sources (e.g., Schlickeiser, Sievers \\& Thiemann, 1987; Tribble 1993; Brunetti et al. 2001). The merger picture is consistent with the occurrence of large--scale radio halos in clusters without a cooling flow, since the major merger event is expected to disrupt a cooling flow (e.g., Sarazin 2002 and references therein).  In spite of the observed anti--correlation between the presence of cooling flows and extended radio emission in clusters of galaxies, there are several cooling flow clusters where the relativistic particles can be traced out quite far from the central galaxy, forming what is called a ``\\textit{mini--halo}'' (e.g. Perseus: Burns et al. 1992; Abell 2390: Bacchi et al. 2003). Mini--halos are diffuse steep--spectrum radio sources, extended on a moderate scale (up to $\\sim 500$ kpc), surrounding a dominant radio galaxy at the center of cooling flow clusters. Until recently, because of the presence of the central radio galaxy, these sources have been considered of different nature from that of extended halos and relics, and the problem of their origin has never been investigated in detail.  Mini--halos do not appear as extended lobes maintained by an Active Galactic Nucleus (AGN), as in classical radio galaxies (Giovannini \\& Feretti 2002), therefore their radio emission is indicative of the presence of diffuse relativistic particles and magnetic fields in the ICM. Rizza et al.(2000) presented three-dimensional numerical simulations of perturbed jet propagating through a cooling flow atmosphere in order to study the interaction between the radio plasma and the hot ICM in cooling flow clusters containing steep--spectrum radio sources. The evolution and spectrum of relativistic particles, however, is not considered in these simulations. The point is that the radiative lifetime of the radio--emitting electrons injected at a given time in the strong magnetic fields present in cooling flow regions is of the order of $\\sim 10^7$ - $10^8$ yr, much shorter than the transport time necessary to cover $\\sim$ hundred kpc scales, so that the diffuse radio emission from mini--halos may suggest the presence of re--accelerated electrons.  More specifically, Gitti, Brunetti \\& Setti (2002, hereafter GBS) suggested that the diffuse synchrotron emission from radio mini--halos is due to a relic population of relativistic electrons re--accelerated by MHD turbulence \\textit{via} Fermi--like processes, with the necessary energetics supplied by the cooling flow.  Alternatively, Pfrommer \\& En{\\ss}lin (2003) in a very recent paper discussed the possibility that the radiating electrons in radio mini--halos are of secondary origin and thus injected during proton-proton collision in the ICM.  The main aim of the present work is the application of GBS's model to a new mini-halo candidate, A2626 ($z=0.0604$, Rizza et al. 2000), and the discussion of the observational properties of the population of radio mini--halos so far discovered.  In Sect. 2 we briefly review GBS's model and its application to the Perseus  cluster. In Sect. 3 we consider the radio source observed in A2626: first, we present VLA data analysis and discuss the possibility that this source belongs to the mini--halo class, then we apply GBS's model to this cluster and discuss the results. In Sect. 4 we present the observational properties of other radio mini--halo candidates in relation to those of host clusters and discuss them in the framework of GBS's model.  For consistency with GBS, a Hubble constant $\\mbox{H}_0 = 50 \\mbox{ km s}^{-1} \\mbox{ Mpc}^{-1}$ is assumed in this paper, therefore at the distance of A2626 $1'$ corresponds to $\\sim$ 95 kpc. The radio spectral index $\\alpha$ is defined such as $S_{\\nu} \\propto \\nu^{-\\alpha}$ and, where not specified, all the formulae are in cgs system. ",
        "conclusions": "We have reported a detailed study of the radio properties of a new mini--halo candidate in A2626. We have shown that a particle re--acceleration model (GBS's model) with a set of physically--meaningful values of the parameters is able to account for the observed brightness profile, the integrated synchrotron spectrum and the radial spectral steepening. We conclude that A2626 has physical properties in between the case of M87, for which there is only marginal evidence for electron re--acceleration, and the prototype of mini--halos in the Perseus cluster. Moreover, we obtain that the maximum power of the cooling flow is more than 2 orders of magnitude larger than the emitted radio power, thus indicating that the cooling flow power (even if considerably reduced by the recent observational claims) may play a leading role in powering the radio mini--halo.   We have selected a small sample of well developed radio mini--halos and shown that the radio power of these objects correlates to the cooling flow power. GBS's model for particle re--acceleration in cooling flow is consistent with the observed trend. If confirmed, in the re--acceleration scenario this trend would indicate that the conversion of the cooling flow power into magneto--plasma turbulence and particle acceleration is similar in well developed radio mini--halos."
    },
    "0311/astro-ph0311043_arXiv.txt": {
        "abstract": "We have obtained low-resolution near-infrared spectra of five SDSS QSOs at 5.3 $<$ {\\it z} $<$ 6.3 using the NICMOS instrument of the {\\it Hubble Space Telescope}.  We find evidence of emission in the Fe II complex centered near $2500${\\AA} (rest) in all five objects.  We estimate Fe II / Mg II $\\lambda$2800 flux ratios comparable to those measured in QSOs at lower redshifts, which indicate metallicities near or above the solar level. We discuss the possible implications of this result assuming the iron enrichment to have been produced mainly by Type Ia supernovae. ",
        "introduction": "The relative abundances of iron and $\\alpha$-process elements in QSOs could provide an important constraint on the star-formation histories of their host galaxies (e.g. Yoshii, Tsujimoto \\& Kawara 1998; Korista et al., these proceedings).  We report here the results of a Cycle 11 {\\it Hubble Space Telescope} program to search for Fe II emission in five {\\it z} $\\approx$ 6 QSOs discovered in the Sloan Digital Sky Survey using the Near Infrared Camera and Multi-Object Spectrometer (NICMOS), and discuss the possible implications of our findings.  The results for the first three objects observed in this program are also presented in Freudling, Corbin \\& Korista (2003). ",
        "conclusions": "If the iron detected in {\\it z} $\\approx$ 6 QSOs were produced mainly by SNe Ia, the evidence that it has reached solar and super-solar abundance levels has several important implications.  First, the galaxy mass / metallicity relation indicates that the host galaxies of the QSOs are very massive, possibly the progenitors of giant ellipticals (see Hamann \\& Ferland 1999 and Fan et al., these proceedings). Such massive objects thus formed within $\\sim$ 1 Gyr of the Big Bang.  Second, recent calculations of SNe Ia formation and ISM enrichment timescales (Friaca \\& Terlevich 1998) indicate that solar Fe levels can be reached within 0.5 to 0.8 Gyr after the formation of the progenitor stars. Under currently favored cosmologies ({\\it h} $\\approx$ 0.7, $\\Omega_{\\Lambda}$  $\\approx$ 0.73, $\\Omega_{M}$ $\\approx$ 0.27) our results imply that the progenitor stars formed at {\\it z} $\\approx$ 20 $\\pm$ 10 (see Freudling et al. 2003). This is consistent with the estimate of the first reionization epoch made from the polarization of the cosmic microwave background (Bennett et al. 2003).  Finally, indirect and direct (Haiman \\& Cen 2002) constraints on the ages of z $\\approx$ 6 QSOs indicate that they reached their present luminosities near their observed redshifts.  The evidence that the stars of their host galaxies formed at higher redshifts thus supports models of an evolutionary connection between nuclear star formation and the AGN phenomenon (e.g. Norman \\& Scoville 1988; Kauffmann \\& Haenelt 2000; Kawakatu \\& Umemura 2003)."
    },
    "0311/astro-ph0311569_arXiv.txt": {
        "abstract": "The microlensing interpretation of the optical afterglow of \\GRB\\ seems naively surprising, since a simple estimate of the stellar microlensing rate gives less than one in four hundred for a flat $\\Omega_\\Lambda=0.7$ cosmology, whereas one event was seen in about thirty afterglows.  Considering baryonic MACHOs making up half of the baryons in the universe, the microlensing probability per burst can be roughly $5\\%$ for a GRB at redshift $z=2$.  We explore two effects that may enhance the probability of observing microlensed gamma-ray burst afterglows: binary lenses and {\\it double} magnification bias.  We find that the consideration of binary lenses can increase the rate only at the $\\sim 15 \\%$ level.  On the other hand, because gamma-ray bursts for which afterglow observations exist are typically selected based on fluxes at widely separated wavebands which are not necessarily well correlated (e.g.\\ localization in X-ray, afterglow in optical/infrared), magnification bias can operate at an enhanced level compared to the usual single--bias case.  Using a simple model for the selection process in two bands, we compute the enhancement to microlensing rate due to magnification bias in two cases: perfect correlation and complete independence of the flux in the two bands.  We find that existing estimates of the slope of the luminosity function of gamma-ray bursts, while as yet quite uncertain, point to enhancement factors of more than three above the simple estimates of the microlensing rate.  We find that the probability to observe at least one microlensing event in the sample of 27 measured afterglows can be $3-4\\%$ for stellar lenses, or as much as $25\\,\\Omega_{\\rm lens}$ for baryonic MACHOs.  We note that the probability to observe at least one event over the available sample of afterglows is significant only if a large fraction of the baryons in the universe are condensed in stellar--mass objects. ",
        "introduction": "The optical afterglow of \\GRB, at a redshift of $z=2.034$, exhibited variability prior to the break in the power law decline in its lightcurve. This variability has been interpreted as gravitational microlensing (Garnavich, Loeb \\& Stanek 2000), as first proposed by Loeb \\& Perna (1998).  While there is no obvious intervening galaxy which might provide lenses in the form of stars or MACHOs, we will take this possibility seriously.  Koopmans and Wambsganss (2001) considered lensing by MACHOs (see also Wyithe \\& Turner (2002b)), but as we will explain in \\S\\ref{discussion} we find a rate of MACHO lensing five times larger than theirs.  The rate of microlensing by stars is much lower, and is naively of order one in three hundred GRBs at redshift two. We note that microlensing is not the only possible explanation of this lightcurve (Panaitescu 2001).  Several authors have considered using microlensing to study the properties of the underlying afterglow (Granot \\& Loeb 2001; Gaudi \\& Loeb 2001; Gaudi, Granot \\& Loeb 2001; Ioka \\& Nakamura 2001).  In this work we will only make statements about afterglows in the aggregate.  Assuming that there is a fraction $f\\sim0.1-0.2$ of dark matter in MACHOs, it is not surprising that a lensing event was seen.  However, if the only cosmological microlenses are stars, the observation of an event is somewhat problematic.  Roughly thirty afterglows have been well--characterized (starting with Metzger et al. (1997) and Kulkarni et al. (1998) and summarized in Frail et al.\\ (2001) and Bloom, Frail \\& Kulkarni (2003)).  We investigate whether magnification bias can make a significant difference, considering the two extreme possibilities that the optical and gamma ray fluxes are perfectly correlated or perfectly uncorrelated.  In Appendix~\\ref{appdx:binary} we consider binary lenses, as discussed by Mao \\& Loeb (2001).  Neither of these two effects seems to be able to increase the probability by a factor of ten. We thus conclude that the observation of a second microlensing event by a star in a sample twice as large is quite unlikely, though with a large MACHO fraction the probability can be significant.  In \\S\\ref{physics} we summarize the basic physics involved in gamma ray burst fireballs, and give an overview of the multiwavelength observations.  In \\S\\ref{microlensing} we derive the cosmological microlensing rate (relegating the details for binary lenses to Appendix~\\ref{appdx:binary}).  In the remaining sections (\\S\\ref{bias} -- \\S\\ref{sample}) we discuss the enhancement of the lensing probability due to magnification bias (Turner, Gott \\& Ostriker 1984).  Magnification bias is the effect that a given {\\em detected} source is more likely to be lensed than a source chosen randomly, simply because magnified sources are easier to detect.  We will consider both the usual single magnification bias, as well as double magnification bias which might be relevant here because of the way gamma-ray burst afterglows are selected.  The selection is complicated, involving gamma ray, X-ray and optical wavelengths. We can only hope to capture the basic features.  In \\S\\ref{bias} \\& \\S\\ref{doublebias} we consider how the magnification bias affects bursts at a single redshift, while in \\S\\ref{sample} we consider the magnification bias on the sample of well--measured afterglows in Bloom et al. (2003).  We primarily consider the gravitational lensing caused by point mass lenses. We define the distance to the source $D_s$, lens $D_\\ell$, and deflector--source distance $D_{\\ell s}$.  These distances are all angular diameter distances in physical units, and we take a flat cosmology where $\\Omega_m+\\Omega_\\Lambda=1$.  We define the Einstein radius and the Einstein angle for the system as \\begin{equation} \\label{eq:R_E} R_E=\\sqrt{\\frac{4GM}{c^2}\\frac{D_\\ell D_{\\ell s}}{D_s}},\\hspace{0.2in} \\theta_E=\\frac{R_E}{D_\\ell}, \\end{equation} where $M$ is the mass of the point lens. ",
        "conclusions": "\\label{discussion}  Our estimate of the microlensing rate for gamma ray burst afterglows is in significant disagreement with previous literature, which we outline here.  If $f=1$, then the optical depth to $z=2$ is $\\tau_0\\approx0.2$, in significant disagreement with Koopmans and Wambsganss (2001).  Computing the mean value of $\\kappa$ from their Eq.~4 (this is identified with our $\\tau_0$), we find $\\overline{\\kappa}\\approx0.04$, a factor of five smaller.  They have cut their isothermal halos at radii that are too small, and have thus neglected 80\\% of their mass.  They have done this to enforce that the halos do not overlap along the line of sight, but the fact is that they do, and significantly.  They have thus underestimated the lensing probability by a factor of five.  Let us explain this in more detail.  In \\S\\ref{microlensing} we assumed an optical depth for microlensing consistent with a uniform distribution of lenses.  This is clearly not the case, as mass in clustered in halos.  We assume that all of the matter is in singular isothermal sphere halos with cutoffs ($\\rho(r)=\\sigma^2/(2\\pi Gr^2)\\Theta(r_0-r)$).  These will obey a luminosity--velocity dispersion relation, have luminosities taken from a Press-Schechter distribution, and have relation between cutoff radii and velocity dispersion, \\begin{eqnarray} L^\\star\\,\\frac{dn}{dL}&=&n^\\star\\left(\\frac{L}{L^\\star}\\right)^\\alpha\\, e^{-L/L^\\star},\\\\ \\frac{L}{L^\\star}&=&\\left(\\frac{\\sigma}{\\sigma^\\star}\\right)^\\gamma,\\\\ \\frac{r_0}{r_0^\\star}&=&\\left(\\frac{\\sigma}{\\sigma^\\star}\\right)^\\beta. \\end{eqnarray} Following Koopmans \\& Wambsganss (2001), we will take $\\alpha=-1$, $\\gamma=4$, as the fiducial case, with $n^\\star=0.0061\\,h^3$ Mpc$^{-3}$, $\\sigma^\\star=225$ km s$^{-1}$.  To prescribe the cutoff radii, we take $\\beta=1$, but the results are insensitive to the precise value of beta as long as it is not close to zero.  We will compute $r_0^\\star$ by normalizing the total mass density to account for $\\Omega_m=0.3$, namely all mass lies in halos.  We find \\begin{equation} r_0^\\star=\\left(\\frac{3}{16\\pi}\\right) \\frac{\\Omega_m H_0^2}{n^\\star\\sigma^{\\star2}}\\; \\Gamma\\left(1+\\alpha+\\frac{2+\\beta}{\\gamma}\\right)^{-1}=0.47 \\;h^{-1}\\;{\\rm Mpc}. \\end{equation}  We now can easily compute the geometric cross section of halos, and the mean number of halos intersected on a line of sight to a particular redshift, \\begin{eqnarray} \\tau_{\\rm GEOM}&=&\\left(\\frac{2\\pi}{3}\\right) \\left(\\frac{n^\\star r_0^{\\star2}c}{\\Omega_mH_0}\\right) \\Gamma\\left(1+\\alpha+2\\beta/\\gamma\\right)\\; \\left(\\sqrt{\\Omega_m(1+z)^3+\\Omega_\\Lambda}-1\\right),\\\\ &=&50.7\\left(\\sqrt{\\Omega_m(1+z)^3+\\Omega_\\Lambda}-1\\right). \\end{eqnarray} Even to moderate redshift, a significant number of halos are intersected.  For \\GRB\\ at $z=2.034$, $\\tau_{\\rm GEOM}\\approx 100$, and the approximation that halos do not overlap is therefore not accurate.  This is the origin of our disagreement.  We have shown that a large number of halos are typically intersected.  However, there will be variations in the optical depth (or equivalently, convergence) between lines of sight.  We have done a simple test, simulating the distribution of convergences and modifying the magnification distribution $dP/dA$ accordingly.  The magnification bias we find does not vary by more than 15\\% in this case.  We have shown that the consideration of magnification bias (both single and double) is important, in that the expected probability that a given observed gamma ray burst is microlensed can be roughly three times the naive optical depth estimate.  MACHOs making up half of the baryonic matter have a significant probability to cause such an event, but we find that stellar lenses alone have a very small probability (less than 3\\% for any event in the afterglow sample) to do so.  We find that binaries can affect the gamma ray burst microlensing rate at the 25\\% level, which though surprisingly large, is not very significant.  A single observed event with a low probability is not in itself problematic. However, should another microlensed gamma ray burst be observed in another thirty afterglows, we believe it would be a significant indication that there are significantly more lenses in the universe than can be inferred from stars. In that case, the effects we have calculated could be quite important in determining the exact density of lenses."
    },
    "0311/astro-ph0311619_arXiv.txt": {
        "abstract": "We quantify the importance of mass accretion during AGN phases in the growth of supermassive black holes (\\BH) by comparing the mass function of black holes in the local universe with that expected from AGN relics, which are black holes grown entirely with mass accretion during AGN phases. The local \\BH\\ mass function (\\BHMF) is estimated by applying the well-known correlations between \\BH\\ mass, bulge luminosity and stellar velocity dispersion to galaxy luminosity and velocity functions.  We find that different correlations provide the same \\BHMF\\ only if they have the same intrinsic dispersion. The density of supermassive black holes in the local universe which we estimate is $\\rhoBH = 4.6_{-1.4}^{+1.9}\\, h_{0.7}^2 \\rhoBHunits$. The relic \\BHMF\\ is derived from the continuity equation with the only assumption that AGN activity is due to accretion onto massive \\BH's and that merging is not important.  We find that the relic \\BHMF\\ at $z=0$ is generated mainly at $z<3$ where the major part of \\BH's growth takes place. Moreover, the \\BH\\ growth is anti-hierarchical in the sense that smaller \\BH's ($M_{BH}<10^{7}M_\\odot$) grow at lower redshifts ($z<1$) with respect to more massive one's ($z\\sim 1-3$). Unlike previous work, we find that the \\BHMF\\ of AGN relics is perfectly consistent with the local \\BHMF\\ indicating the local black holes were mainly grown during AGN activity.  This agreement is obtained while satisfying, at the same time, the constraints imposed from the X-ray background.  The comparison between the local and relic \\BHMF's also suggests that the merging process is not important in shaping the relic \\BHMF, at least at low redshifts ($z<3$), and allows us to estimate the average radiative efficiency ($\\varepsilon$), the ratio between emitted and Eddington luminosity ($\\lambda$) and the average lifetime of active \\BH's. Our analysis thus suggests the following scenario: local black holes grew during AGN phases in which accreting matter was converted into radiation with efficiencies $\\varepsilon = 0.04-0.16$ and emitted at a fraction $\\lambda = 0.1-1.7$ of the Eddington luminosity. The average total lifetime of these active phases ranges from $\\simeq 4.5\\times 10^{8}$ yr for $M_{BH}<10^{8}M_\\odot$ to $\\simeq 1.5\\xten{8}$ yr for $M_{BH}>10^{9}M_\\odot$ but can become as large as $\\sim \\ten{9}\\YR$ for the lowest acceptable $\\eff$ and $\\lambda$ values. ",
        "introduction": "Since the discovery of quasars (\\citealt{schmidt}) it has been suggested that Active Galactic Nuclei (AGN) are powered by mass accretion onto a supermassive Black Hole (\\BH) with mass in the range \\ten{6}-\\ten{10}\\Msun\\ (\\citealt{salpeter, zeldovich, lyndenbell}).  This paradigm combined with the observed evolution of AGNs implies that a significant number of galaxies in the local universe should host a supermassive \\BH\\ (e.g.~\\citealt{soltan, cavaliere89, chokshi}).  Supermassive \\BH's are now detected in $\\sim 40$ galaxies through gas and stellar dynamical methods (\\citealt{kormendy, merritt}).  Some galaxies are quiescent (e.g.~M32, \\citealt{vandermarel}; NGC 3250, \\citealt{barth}) and some are mildly or strongly active (e.g.~M87, \\citealt{marconi97,macchetto}; Centaurus A, \\citealt{marconi01};  Cygnus A, \\citealt{tadhunter}). The mass of the \\BH\\ correlates with some properties of the host galaxy such as spheroid (bulge) luminosity and mass (\\citealt{kr95,magorrian}), light concentration (Graham et al. 2001) and with the central stellar velocity dispersion (\\citealt{fm00, gebhardt}).  The latter correlation was thought to be the tightest but \\cite{MH03} have recently shown that all the correlations have similar intrinsic dispersion when considering only galaxies with secure BH detections (see also \\citealt{mclure, erwin}). Overall, the dispersion is of the order of 0.3 in \\lMBH\\ at a given value of \\Lsph, \\Msph\\ or \\sigstar. The existence of any correlations of \\BH\\ and host galaxy bulge properties has important implications for theories of galaxy formation in general and bulge formation in particular. Indeed, several attempts at explaining the origin of these correlations and the difficulties/constraints that they pose to galaxy formation models can be found in the literature (e.g.~\\citealt{silk98,cattaneo99,haehnelt00,ciotti01,cavaliere02,adams03}, and references therein).  To date, all newly found \\BH's have masses (or upper limits) in agreement with those expected from the above correlations suggesting that all galaxies host a massive \\BH\\ in their nuclei. By applying the correlations between \\MBH\\ and host galaxy properties it is then possible to estimate the mass function of local \\BH's or, more simply, their total mass density (\\rhoBH) in the local universe (e.g.~\\citealt{salucci,MS02, yu02, ferrarese, aller}).  It is important to verify if local BH's are exclusively relics of AGN activity or if other mechanisms, such as merging, play a role. Hereafter we will call 'AGN relics', or simply relics, those black holes which grew up from small seeds (1-\\ten{3}\\Msun) following mass accretion during AGN phases. For instance, an AGN relic of \\ten{9}\\Msun\\ is different from a \\BH\\ of the same mass but which was formed by the merging of many smaller \\BH's.  A simple comparison between local and relic BH's was performed by \\cite{salucci} and \\cite{fi99} who determined \\rhoBH\\ from the observed X-ray background emission.  A revised estimate was obtained by \\cite{erz02} who found $\\rhoBH =  (7.5 - 16.8)\\xten{5}\\, (\\eff/0.1)^{-1}\\rhoBHunits$ in disagreement with the estimate from local black holes ($\\rhoBH=(2.5\\pm 0.4)\\, h^2_{65}\\rhoBHunits$, \\citealt{yu02,aller}).  \\cite{erz02} thus suggested that, in order to reconcile this discrepancy, massive \\BH's should have large accretion efficiencies (i.e.~larger than the canonically adopted value of $\\eff =0.1$), hence they should be rapidly rotating.  A more detailed comparison was performed by \\cite{MS02} who found an agreement between the black hole mass functions (hereafter \\BHMF) of local and relic black holes.  Recently, however, \\cite{yu02} and \\cite{ferrarese} found a disagreement at large masses ($\\MBH\\ge \\ten{8}\\Msun$) where more AGN relics are expected relative to local \\BH's. As previously stated, a possibility to reconcile this discrepancy is to assume high accretion efficiencies but, clearly, this issue is still much debated.  The relation between AGN relics and local \\BH's is also being studied in the framework of coeval evolution of \\BH\\ and host galaxy.  Several physical models have been proposed in which the fueling of the \\BH, hence the AGN activity, is triggered by merging events (in the context of the hierarchical structure formation paradigm, see for instance \\citealt{kauffmann00,volonteri03,menci03,wyithe03,hatziminaoglou03,haehnelt03}) or is simply directly related to the star formation history of the host galaxy (e.g.~\\citealt{dimatteo03,granato04,haiman03}).  The \\BH\\ has then a feedback on the host galaxy through the energy released in the AGN phase (e.g.~\\citealt{silk98,blandford99,begelman}).  As a result of this double interaction (galaxy feeding the BH -- AGN feedback on the galaxy), these models can in general reproduce both the observed BH-host galaxy correlations and the AGN luminosity functions (e.g. \\citealt{haehnelt98,monaco00,nulsen00} and previous references).  However, this big effort in modeling cannot uniquely answer the question if local \\BH's are relics of AGN's, since a wide range of models with many different underlying assumptions cannot be ruled out with the available observational constraints.  The aim of this paper is to investigate the assumption that massive black holes in nearby galaxies are relics of AGN activity by comparing the local \\BHMF\\ with that of AGN relics.  We remark that in this paper we do not build a physical model of the coevolution between central \\BH\\ and host galaxy but we compare differential and integrated mass densities (local \\BH's) with differential and integrated energy densities (AGN's), with the only assumption that AGN activity is caused by mass accretion onto the central \\BH.  We refine the analysis by \\cite{MS02} and evaluate the discrepancies found by other authors between local and relic \\BHMF's.  In Section \\ref{sec:lbhmf} we estimate the local \\BHMF\\ by applying the known correlations between \\MBH\\ and host galaxy properties to the galaxy luminosity and stellar velocity dispersion function.  We also check the self-consistency of the results, and show that different \\MBH--host-galaxy-properties relations provide the same \\BHMF\\ within the uncertainties.  In Section \\ref{sec:rbhmf} we use the continuity equation to estimate the \\BHMF\\ of AGN relics and in Section \\ref{sec:localrelics}, we compare local and relic \\BHMF's, and find that local BH's are consistent with AGN relics.  We then show (Sec. \\ref{sec:xrb}) that the energetic constraints inferred from the X-ray background (XRB) are also satisfied and that there is no discrepancy between \\rhoBH\\ of local BH's and that expected from the XRB. In Sections \\ref{sec:efflambda}, \\ref{sec:hist}, and \\ref{sec:time} we discuss constraints on the accretion efficiency (\\eff) and on the Eddington ratio ($\\lambda= L/\\LEDD$, where $L$ is the AGN luminosity and \\LEDD\\ is the Eddington luminosity of the active BH), and we estimate the accretion history and the average lifetime of massive BH's.  We summarize our results and we draw our conclusions in Section \\ref{sec:conclusions}.  In this paper we adopt the current ``standard'' cosmological model, with $H_0=100 h\\,\\KM\\SEC\\1\\MPC\\1$ and $h=0.7$, $\\Omega_M=0.3$ and $\\Omega_\\Lambda=0.7$.   \\begin{table*} \\centering \\begin{minipage}{0.75\\linewidth} \\caption{\\label{tab:galtypes}Adopted galaxy properties per morphological types.} \\begin{tabular}{@{}lllllr@{}} \\hline & E & S0 & Sab & Scd & Notes \\\\ \\hline Morph. Fraction\t& 0.11$\\pm$0.03 & 0.21$\\pm$0.05 & 0.43$\\pm$0.07 & 0.19$\\pm$0.07\t& 1 \\\\ m(Bulge)-m(Total)\t& 0\t\t& $0.64\\pm0.30$ & $1.46\\pm0.56$ & $2.86\\pm0.59$ & 2 \\\\ m(Bulge)-m(Total)\t& 0\t\t& $0.60\\pm0.50$ & $1.78\\pm1.01$ & $2.78\\pm1.21$ & 3 \\\\ \\hline \\end{tabular}\\\\ 1. From \\protect\\citealt{fukugita98} (see text for details).\\\\ 2. B band. Estimate by \\protect\\cite{aller} based on data from \\protect\\citealt{simien}.\\\\ 3. H band. Based on data from \\citealt{hunt04}.\\\\ \\end{minipage} \\end{table*} ",
        "conclusions": "\\label{sec:conclusions}  We have quantified the importance of mass accretion during AGN phases in the growth of supermassive black holes (\\BH) by comparing the mass function of black holes in the local universe with that expected from AGN relics, which are black holes grown entirely during AGN phases.  The local \\BH\\ mass function (\\BHMF) has been estimated by applying the well-known correlations between \\BH\\ mass, bulge luminosity and stellar velocity dispersion to galaxy luminosity and velocity functions.  We have found that different \\BH-galaxy correlations provide the same \\BHMF\\ only if they have the same intrinsic dispersion, confirming the findings of \\cite{MH03}. The density of supermassive black holes in the local universe is $\\rhoBH = 4.6_{-1.4}^{+1.9}\\, h_{0.7}^2 \\rhoBHunits$.  The relic \\BHMF\\ is derived from the continuity equation with the only assumption that AGN activity is due to accretion onto massive \\BH's and that merging is not important.  We find that the relic \\BHMF\\ at $z=0$ is generated mainly at $z<3$ where the major part of \\BH's growth takes place.  The relic \\BHMF\\ at $z=0$ is very little dependent on its value at $\\zo=3$ since the main growth of \\BH's took place at $z<3$. Moreover, the \\BH\\ growth is anti-hierarchical in the sense that smaller \\BH's ($M_{BH}<10^{8}M_\\odot$) grow at lower redshifts ($z<1$) with respect to more massive one's ($z\\sim 1-3$). If the correlations between \\BH\\ mass and host-galaxy-properties hold at higher redshifts this would represent a potential problem for hierarchical models of galaxy formation.  Unlike previous work, we find that the \\BHMF\\ of AGN relics is perfectly consistent with the local \\BHMF\\ indicating the local black holes were mainly grown during AGN activity.  This agreement is obtained while satisfying, at the same time, the constraints imposed from the X-ray background both in terms of BH mass density and fraction of obscured AGN's. The reasons for the solution of the discrepancy at high masses found by other authors are  the following: \\begin{itemize} \\item we have taken into account the intrinsic dispersion of the \\MBHsig\\ and \\MBHlum\\ correlations in the determination of the local \\BHMF; \\item we have adopted the coefficients of the \\MBHsig\\ and \\MBHlum\\ ($K$ band) correlations derived by \\cite{MH03} after considering only `secure' BH masses; \\item we have derived improved bolometric corrections which do not take into account reprocessed IR emission in the estimate of the bolometric luminosity. \\end{itemize}  The comparison between the local and relic \\BHMF's also suggests that the merging process at low redshifts ($z<3$) is not important in shaping the relic \\BHMF, and allows us to estimate the average radiative efficiency ($\\varepsilon$), the ratio between emitted and Eddington luminosity ($\\lambda$) and the average lifetime of active \\BH's.  Our analysis thus suggests the following scenario: local black holes grew during AGN phases in which accreting matter was converted into radiation with efficiencies $\\varepsilon = 0.04-0.16$ and emitted at a fraction $\\lambda = 0.1-1.7$ of the Eddington luminosity. The average total lifetime of these active phases ranges from $\\simeq 4.5\\times 10^{8}$ yr for $M_{BH}<10^{8}M_\\odot$ to $\\simeq 1.5\\xten{8}$ yr for $M_{BH}>10^{9}M_\\odot$ but can become as large as $\\sim \\ten{9}\\YR$ for the lowest acceptable $\\eff$ and $\\lambda$ values."
    },
    "0311/astro-ph0311105_arXiv.txt": {
        "abstract": "There has been extensive recent progress in X-ray observations of clusters of galaxies with the analysis of the entire {\\it ASCA} database and recent new results from {\\it Beppo-SAX}, {\\it Chandra}, and {\\it XMM-Newton}.  The temperature profiles of most clusters are isothermal from 0.05--0.6 $R_{\\rm viral}$, contrary to theoretical expectations and early results from {\\it ASCA}. Similarly, the abundance profiles of Fe are roughly constant outside the central regions. The luminosity-temperature relation for a very large sample of clusters show that $L_{\\rm X}\\propto T^3$ over the whole observable luminosity range at low redshift, but the variance increases at low luminosity, explaining the previously claimed steepening at low luminosity. Recent accurate cluster photometry in red and infrared passbands have resulted in much better correlations of optical and X-ray properties, but there is still larger scatter than one might expect between total light and X-ray temperature and luminosity. The velocity dispersion and the X-ray temperature are strongly correlated, but the slope of the relation is somewhat steeper than expected. The surface brightness profiles of clusters are very well fit by the isothermal $\\beta$ model out to large radii and show scaling relations, outside the central regions, consistent with a $\\Lambda$-dominated Universe.   At high masses the gas mass fraction of clusters is quite uniform and is consistent with the low WMAP value of $\\Omega_{\\rm m}$.  The recent analysis of cluster mass-to-light ratio and the mass-to-light ratio of stars indicates that the ratio of gas to stellar mass is $\\sim$10:1 in massive clusters. There is an apparent decrease in gas mass fraction and increase in stellar mass fraction at lower mass scales, but the very flat surface brightness of the X-ray emission makes extension of this result to large scale lengths uncertain. The normalization of the scaling of mass with temperature, derived from measurements of density and temperature profiles and assuming hydrostatic equilibrium, is lower than predicted from simulations that do not include gas cooling or heating and has a slightly steeper slope.  Detailed {\\it Chandra} and {\\it XMM-Newton} imaging spectroscopy of several clusters show that the form of the potential is consistent with the parameterization of Navarro, Frenk, \\& White (1997) over a factor of 100 in length scale and that there is no evidence for a dark matter core. {\\it Chandra} X-ray images have revealed rather complex internal structures in the central regions of some clusters, which are probably due to the effects of mergers; however, their nature is still not completely clear.  There are now more than 100 clusters with well-determined Fe abundance, several with accurate values at redshifts $z \\approx 0.8$, with little or no evidence for evolution in the Fe abundance with redshift. There is real variance in the Fe abundance from cluster to cluster, with a trend for clusters with higher gas densities to have higher Fe abundances. The Si, S, and Ni abundances do not follow patterns consistent with simple sums of standard Type Ia and Type II supernova, indicating that the origin of the elements in clusters is different from that in the Milky Way. The Si/Fe abundance rises with cluster mass, but the S/Fe ratio does not. The high Ni/Fe ratio indicates the importance of Type Ia supernovae. {\\it XMM-Newton} grating spectra of the central regions of clusters have derived precise O, Ne, Mg, and Fe abundances.  {\\it XMM-Newton} CCD data are allowing O abundances to be measured for a large number of clusters. ",
        "introduction": "Clusters of galaxies are the largest and most massive collapsed objects in the Universe, and as such they are sensitive probes of the history of structure formation. While first discovered in the optical band in the 1930's (for a review, see Bahcall 1977a), in same ways the name is a misnomer since most of the baryons and metals are in the hot X-ray emitting intracluster medium and thus they are basically ``X-ray objects.''  Studies of their evolution can place strong constraints on all theories of large-scale structure and determine precise values for many of the cosmological parameters. As opposed to galaxies, clusters probably retain all the enriched material created in them and being essentially closed boxes they provide an unbiased record of nucleosynthesis in the Universe. Thus, measurement of the elemental abundances and their evolution provide fundamental data for the origin of the elements. The distribution of the elements in clusters reveals how the metals were removed from stellar systems into the intergalactic medium (IGM). Clusters should be fair samples of the Universe, and studies of their mass and their baryon fraction reveal the gross properties of the Universe as a whole. Since most of the baryons are in the gaseous phase and clusters are dark matter dominated, the detailed physics of cooling and star formation are much less important than in galaxies. This makes clusters much more amenable to detailed simulations than galaxies or other systems in which star formation has been an overriding process. Detailed measurements of their density and temperature profiles allow an accurate determination of the dark matter profile and total mass. While gravity is clearly dominant in massive systems, much of the entropy of the gas in low-mass systems maybe produced by nongravitational processes.  Clusters are luminous, extended X-ray sources and are visible out to high redshifts with present-day technology. The virial temperature of most groups and clusters corresponds to $kT \\approx (2-100) \\times 10^6$ K (velocity dispersions of 180--1200 km s$^{-1}$), and while lower mass systems certainly exist, we usually call them galaxies. Most of the baryons in groups and clusters of galaxies lie in the hot X-ray emitting gas, which is in virial equilibrium with the dark matter potential well [the ratio of gas to stellar mass is $\\sim$(2--10):1; Ettori \\& Fabian 1999]. This gas is enriched in heavy elements (Mushotzky et al. 1978) and is thus the reservoir of stellar evolution in these systems. The presence of heavy elements is revealed by line emission from H and He-like transitions, as well as L-shell transitions of the abundant elements. Most clusters and groups are too hot to have significant line emission from C or N, but all abundant elements with $Z > 8$ (O) have strong lines from H and He-like states in the X-ray band, and their abundances can be well determined.  Clusters of galaxies were discovered as X-ray sources in the late 1960's (see Mushotzky 2002 for a historical review), and large samples were first obtained with the {\\it Uhuru} satellite in the early 1970's (Jones \\& Forman 1978). Large samples of X-ray spectra and images were first obtained in the late 1970's with the {\\it HEAO} satellites (see Forman \\& Jones 1982 for an early review). The early 1990's brought large samples of high-quality images with the {\\it ROSAT} satellite and good quality spectra with {\\it ASCA} and {\\it Beppo-SAX}.  In the last three years there has been an enormous increase in the capabilities of X-ray instrumentation due to the launch and operation of {\\it Chandra} and {\\it XMM-Newton}. Both {\\it Chandra} and {\\it XMM-Newton} can find and identify clusters out to $z > 1.2$, and their morphologies can be clearly discerned to $z > 0.8$ (Fig. 1.1).  The cluster temperatures can be measured to $z \\approx 1.2$, and {\\it XMM-Newton} can determine their overall chemical abundances to $z \\approx 1$ with sufficiently long exposures (very recently the temperature and abundance of a cluster at $z$ = 1.15 was measured accurately in a 1~Ms {\\it XMM-Newton} exposure; Hasinger et al. 2004). Temperature and abundance profiles to $z \\approx 0.8$ can be well measured and large samples of X-ray selected clusters can be derived. {\\it Chandra} can observe correlated radio/X-ray structure out to $z > 0.1$ and has discovered internal structure in clusters. The {\\it XMM-Newton} grating spectra can determine accurate abundances for the central regions of clusters, in a model independent fashion, for O, Ne, Mg, Fe, and Si.  \\begin{figure*}[t] \\centering \\includegraphics[width=1.00\\columnwidth,angle=0,clip]{fig1.ps} \\vskip 0pt \\caption{ {\\it Chandra} (left panel) and {\\it XMM-Newton} (right panel) images of the Lynx region (Stanford et al. 2001), which contains three high-redshift clusters. The {\\it Chandra} image has $\\sim$170 ks and the {\\it XMM-Newton} image $\\sim$20 ks exposure.} \\end{figure*}   It is virtually impossible to give a balanced review of the present observational state of X-ray cluster research, with more than 100 papers published each year. I will not say much about those issues for which we have had detailed talks at this meeting: cooling flows, high-redshift clusters and evolution, X-ray data and the Sunyaev-Zel'dovich effect, radio source interaction, X-ray selected active galaxies in clusters, X-ray emission from groups, and detailed comparison of masses derived from lensing and X-ray observations. Other areas, such as the presence of nonthermal emission and the existence of very soft components, were not discussed. Even limiting the talk this much results in an abundance of material. However, for the purposes of continuity, I have included some material that overlaps with the reviews on chemical abundance given by Renzini (2004) and on groups by Mulchaey (2004). This review does not consider work published after February 2003. ",
        "conclusions": ""
    },
    "0311/astro-ph0311349_arXiv.txt": {
        "abstract": "The magnetic field strength in molecular clouds is a fundamental quantity for theories of star formation. It is estimated by Zeeman splitting measurements in a few dense molecular cores, but its volume--averaged value within large molecular clouds (over several parsecs) is still uncertain. In this work we provide a new method to constrain the average magnetic field strength in molecular clouds. We compare the power spectrum of gas density of molecular clouds with that of two $350^3$ numerical simulations of supersonic MHD turbulence. The numerical simulation with approximate equipartition of kinetic and magnetic energies (model A) yields the column density power spectrum $P(k)\\propto k^{-2.25\\pm 0.01}$, the super--Alfv\\'{e}nic simulation (model B) $P(k)\\propto k^{-2.71\\pm 0.01}$. The column density power spectrum of the Perseus, Taurus and Rosetta molecular cloud complexes is found to be well approximated by a power law, $P_{\\rm o}(k)\\propto k^{-a}$, with $a=2.74\\pm 0.07$, $2.74\\pm 0.08$ and $2.76\\pm 0.08$ respectively. We conclude that the observations are consistent with the presence of super--Alfv\\'{e}nic turbulence in molecular clouds (model B) while model A is inconsistent (more than 99\\% confidence) with the observations. ",
        "introduction": "The volume--averaged magnetic field strength in molecular clouds has never been measured directly. Zeeman splitting has been detected only from a few dense molecular cloud cores, where emission lines from molecules such as OH and CN are observed \\citep{Crutcher99,Bourke+2001}. Dense cores fill only a small fraction of the volume of molecular clouds. Therefore, the magnetic field strength averaged over the molecular cloud volume cannot be directly inferred from its value in the cores. This is true especially if the magnetic field strength has a very intermittent distribution and is correlated with the gas density, as suggested by \\cite{Padoan+Nordlund99MHD}.  Estimates of magnetic field strength in molecular clouds have been inferred from the dispersion in the polarization angle (e.g. Myers \\& Goodman, 1991; Chrysostomou et al., 1994; Lai et al., 2001; Matthews \\& Wilson, 2002; Lai et al., 2002) \\nocite{Myers+Goodman91,Chrysostomou+94,Lai+2001,Matthews+Wilson2002,Lai+2002} as originally suggested by \\cite{Davis51} and \\cite{Chandrasekhar+Fermi53}. This method was tested in numerical simulations of MHD turbulence by \\cite{Ostriker+2001}, \\cite{Padoan+2001pol} and \\cite{Heitsch+2001}. The relative motion of ions and neutral molecules, as manifested by a comparison of their spectral lines, has also been used to estimate the magnetic field strength in molecular clouds \\citep{Houde+2000,Houde+2002}.  In this work we present a new way to constrain the average magnetic field strength in molecular clouds, based on the density power spectrum. In \\S~2 we show numerical simulations of supersonic MHD turbulence with different magnetic field strength yield different power spectra of gas density. The power spectrum is then computed from maps of molecular clouds in \\S~3. We find that only a rather low value for the average magnetic field strength, leading to super--Alfv\\'{e}nic turbulence, is consistent with the observations. Conclusions are summarized in \\S~4.  The density power spectrum depends on the rms sonic Mach number of the turbulence, $M_{\\rm S}$ (the ratio of the rms flow velocity and the sound speed). In this work we use only simulations with $M_{\\rm S}\\approx 10$ because that is the approximate value in the molecular cloud complexes we have studied. We do not study the dependence of the power spectrum on the value of $M_{\\rm S}$. In general, smaller values of $M_{\\rm S}$ yield steeper density power spectra than the $M_{\\rm S}\\approx 10$ models (this was verified with a set of simulations that will be presented elsewhere). Turbulent flows with $M_{\\rm S} \\ll 1$, for example, are expected to generate a Kolmogorov density power spectrum, proportional to $k^{-11/3}$. This slope is comparable with the electron density power spectrum on very small scale estimated from scintillation studies \\citep{Armstrong+95}. HI surveys of our galaxy and the Magellanic Clouds \\citep{Crovisier+Dickey83,Elmegreen+2001LMC,Stanimirovic+Lazarian2001,Dickey+2001} produced power spectra with slope intermediate between the present results in molecular clouds and the scintillation studies. ",
        "conclusions": "The average magnetic field strength in molecular clouds cannot be measured directly. However, it can be inferred from observational data due to its effect on the gas dynamics. In this work we have found the power spectrum of the density field is a sensitive diagnostic of the magnetic field strength. Numerical simulations of supersonic MHD turbulence with rms sonic Mach number $M_{\\rm S}\\approx 10$ develop a power law power spectrum of gas density, $P_{\\rm MHD}(k)\\propto k^{-a_{\\rm MHD}}$. The value of the power law exponent is $a_{\\rm MHD}=2.25\\pm 0.01$ when the flow rms velocity is comparable to the Alfv\\'{e}n velocity and $a_{\\rm MHD}=2.71\\pm 0.01$ in the super--Alfv\\'{e}nic simulation.  The density power spectrum can be measured also in molecular cloud complexes, for example using \\jco\\ maps. However, the column density (and its power spectrum) estimated using only the \\jco\\ line and the LTE method are biased by a number of uncertainties. The most significant uncertainties are the 3D distribution of the gas kinetic temperature in the molecular cloud complexes and the saturation of the \\jco\\ line in very dense regions. We have investigated the effect of these uncertainties in the density power spectrum using synthetic \\jco\\ spectral maps computed with a non--LTE Monte Carlo radiative transfer code. The 3D equilibrium temperature distribution is computed self--consistently as part of the radiative transfer solution by balancing cosmic ray heating with molecular and atomic cooling. The correct power spectrum slope, $a_{\\rm MHD}$, is found to be smaller than the slope estimated with the LTE method, $a_{\\rm LTE}$, with $a_{\\rm MHD}=a_{\\rm LTE}-0.13\\pm 0.06$.  With this correction, the power spectrum slope is $a=2.74\\pm 0.07$ for Perseus,$a=2.74\\pm 0.08$  for Taurus and $a=2.76\\pm 0.08$ for Rosetta. The super--Alfv\\'{e}nic model is consistent with this result, while the model with rms flow velocity comparable to the Alfv\\'{e}n velocity is ruled out by the observations. This is yet another indication that super--Alfv\\'{e}nic turbulence provides a good description of molecular cloud dynamics. As first proposed by \\cite{Padoan+Nordlund97MHD,Padoan+Nordlund99MHD}, the average magnetic field strength in molecular clouds may be much smaller than required to support them against the gravitational collapse. Evidence of super--Alfv\\'{e}nic turbulence was also found recently by Troland and Heiles (2001) in HI clouds.  \\nocite{Troland+Heiles2001}"
    },
    "0311/astro-ph0311380_arXiv.txt": {
        "abstract": "Q-LET is a FORTRAN 77 code that enables a quick estimate of the gravitational lensing effects on a point- or an extended source. The user provided input consists of the redshifts, angular positions relative to the source, mass or velocity dispersion estimate and halo type for the lens galaxies. The considered halo types are the Navarro-Frenk-White and the Singular Isothermal Sphere. The code uses the so-called multiple lens-plane method to find the magnification and intrinsic shape of the source. This method takes into account the multiple deflections that may arise when several mass accumulations are situated at different redshifts close to the line-of-sight.  The Q-LET code is applied to the recently disco\\-vered super\\-nova, SN2003es, which is likely to be of Type Ia as its host galaxy is classified as an elliptical. We find that SN2003es is likely to have been significantly magnified by gravitational lensing and that this should be considered in high-$z$ studies if this SN is to be used to determine the cosmological parameters.  Q-LET was motivated by the supernova searches, where lensing can be a problem, but it can also be applied to any simple lens system where a quick estimate is wanted, e.g.~the single lens case. ",
        "introduction": "The recent advance in modern cosmology due to new and improved telescopes and techniques seems to overthrow the long-lasting paradigm or hope of an Einstein-de Sitter (EdS) universe. From the Cosmic Microwave Background (CMB) \\cite{wmap,acbar,cbi}, along with studies of Large Scale Structure (LSS), \\cite{2df} and the Lyman $\\alpha$ forest \\cite{lyalpha1,lyalpha2} (Ly$\\alpha$) etc.~it is inferred that the universe consists of roughly 30 \\% matter (Cold Dark Matter, CDM and ordinary matter) and 70 \\% Dark Energy. Furthermore, measurements of the apparent magnitude of Type Ia supernovae (SNe) have been used by two collaborations \\cite{scp,highz} to estimate the values of the total matter energy density $\\Omega_{\\rm M}$, which includes CDM, and the density in Dark Energy, e.g.~$\\Omega_{\\Lambda}$. Type Ia SNe are believed to be so-called standard candles, i.e.~they show a very small scatter in intrinsic luminosity after empirical corrections \\cite{snscattcorr1,snscattcorr2}, and by measuring the apparent magnitudes, the cosmology-dependent luminosity distance can be inferred. The supernovae look fainter than they would in an EdS universe favouring both a non-zero $\\Omega_{\\Lambda}$ and $\\Omega_{\\rm M}\\neq 1$. Combined with CMB data from Boomerang and MAXIMA \\cite{boomer,maxima}, they end up at approximately the same values as the above-mentioned CMB, LSS and Ly$\\alpha$ measurements. However, there is a number of potential systematic errors that can contaminate the measurements. First, inter- and/or intra-galactic dust may obscure the light from the SNe and make them look fainter. This has been investigated in e.g.~\\cite{edvariel1,aguirre1,aguirre2}. It is also still unclear whether there is any evolution of the SN brightness with redshift. Many potential sources of evolution have been proposed, such as differing progenitor composition \\cite{umeda,hoflich,domingez} or host galaxy morphology (and thereby redshift) dependence \\cite{snevol} just to name a few. Furthermore, if photons can oscillate into axions, the SNe will also look fainter due to the reduced number of received photons. This has been addressed e.g.~in reference \\cite{edvard1}. A fourth possible contaminant is magnification or de-magnification by gravitational lensing. Having a large sample of SNe this can be corrected for statistically \\cite{edvard2}, but when studying single SNe, the lensing effects need a careful treatment. This was the case with the farthest known SN so far, SN1997ff, at a redshift of 1.7. The vicinity of the line-of-sight to this SN was unusually dense in galaxies and modelling of the individual lens galaxies was required. This was done by \\cite{lewiba,cgedvariel,riess1,benitez} who conclude a large possible magnification depending on the lens masses and concentrations. Ben\\'{\\i}tez et al.~\\cite{benitez} conclude a magnification of 0.34$\\pm$0.12 mag.  In this paper we present the well-known multiple lens-plane method applied to the calculation of the magnification of any point- or extended source. Section \\ref{sec:lensing} describes this method in a general way. Furthermore, we model the lenses as spherically symmetric Navarro-Frenk-White (NFW) or Singular Isothermal Sphere (SIS) halos and the lensing properties of these are described in sections \\ref{subs:nfw} and \\ref{subs:sis}. For more information on under which circumstances each halo type is appropriate, see e.g.~reference \\cite{liost}. The method and models are manifested in a \\texttt{FORTRAN 77} program, Q-LET, obtainable from the author upon request or it can be downloaded at \\texttt{http://www.physto.se/\\~{}cg/qlet/qlet.htm}. The program is described in the text in section \\ref{sec:program}. In reference \\cite{cgedvariel}, an early version of Q-LET was used to estimate the magnification of the abovementioned SN1997ff. Section \\ref{sec:ramone} describes the application of the code to the recently discovered supernova SN2003es. ",
        "conclusions": "\\label{subs:ramdisc} One of the approximations that may be criticised is the fact that we extend the halos out to infinity. However, the lensing effect at large impact radii is very small. For SN2003es, the lenses with large impact radii contribute very little to the total magnification which is dominated by galaxies no.~631 and 662.  Putting ${\\rm M}_{200}$ of the NFW equal to that of the SIS for a given velocity dispersion is not obviously a good approximation. The most likely measured quantities are the luminosity or velocity dispersion of the lens and the luminosity may be used to estimate a velocity dispersion either via the Tully-Fisher or Faber-Jackson relation. When a velocity dispersion is obtained, the best way would be to use it directly on the NFW profile. However, as opposed to the SIS, the NFW velocity dispersion depends on the radius. By following {\\L}okas and Mamon in reference \\cite{lokas} it is possible to obtain the line-of-sight velocity dispersion, $\\sigma_{\\rm los}$, for well resolved sources, or, for small or distant sources, the ``aperture velocity dispersion'', $\\sigma_{\\rm ap}$, which is the average over an aperture centered on the object. We calculated both these quantities as a function of radius and found that $\\sigma_{\\rm ap}$ did not fall off dramatically. As an example, in a case where we put the SIS velocity dispersion ($\\sigma_{\\rm v}^{\\rm SIS}$) to 170 km/s to get the parameters for the NFW and then compute the corresponding NFW velocity dispersions, $\\sigma_{\\rm ap}$ varies roughly by $+5$\\% and $-15$\\% around 170 km/s between radii $0.1r_{200}$ and $2r_{200}$. The line-of-sight velocity dispersion varies more, $+10$\\% to $-40$\\% in the same range. However, at a radius of \\emph{roughly} $0.3r_{200}$, there is a quite small scatter around the SIS value of less than $\\sim$15\\%. This seems to be valid for a large range of redshifts and velocity dispersions. It also seems valid in low $H_0$- and low density universes. Thus in almost all cases it is possible to find a radius at which $\\sigma_{\\rm los}\\sim\\sigma_{\\rm ap}\\sim\\sigma_{\\rm v}^{\\rm SIS}$, with this radius being roughly equal to $0.3r_{200}$.  Inverting the velocity dispersion-mass relation for the NFW is quite time-consuming, and this, along with the ambiguity due to the dependence of $\\sigma_{\\rm v}$ on radius and the fact that the difference seems not to be significant, motivates using the simpler conversion via the SIS. However, if this conversion is unsatisfactory, it is always possible for the user to supply ${\\rm M}_{200}$ instead, in which case of course no conversion is needed.  For the irregular galaxies we have used the Faber-Jackson relation which is valid for ellipticals. An approximation which is difficult to say to what extent it affects the results, and in which direction. Using the Tully-Fisher relation on the irregulars would increase the magnification even further, but at most by $\\sim$15-20\\%.  Whether the halo models we use are good descriptions of real halos is still under debate. We are using spherically symmetric models, which might be an oversimplification but the ellipticity of the dark matter in the observed halos is in practice very seldom known (see e.g.~reference \\cite{elllens} for the effects of ellipticity in lensing).  So, unless some new model is very different from the SIS and NFW, the results for SN2003es is likely to remain.    \\subsection{Conclusions} \\label{subs:ramconcl} We have found that SN2003es is likely to have been significantly magnified by gravitational lensing, a fact that must be taken into account when trying to use this (likely Type Ia) SN to infer cosmological parameters. Either a more thorough analysis of the lensing galaxies\\footnote{For instance one could try to measure the velocity dispersions.} is needed to take the effect into account or the SN should be used only to set a lower limit on the distance or maybe even be excluded from the sample if not having large statistics. In that case some SNe are likely to be de-magnified and the effect can be assumed to average to $\\sim 0$. The fact that the intrinsic shape of the host galaxy looks perfectly normal makes it practically impossible to use this shape to put constraints on the magnification of SN2003es. If an extremely unusual shape would have been obtained, it could perhaps have affected the reliability of the model.   We have presented the multiple lens-plane method of gravitational lensing along with its implementation in Q-LET, a \\texttt{FORTRAN 77} code that can be used to quickly estimate the lensing effects on either a point source or to obtain the intrinsic shape of an elliptical image. A square grid in the image plane is also possible for user-defined image shapes.  The code was applied to SN2003es, a SN (most likely Type Ia) discovered in the HDF-N. Using both NFW and SIS type halos and estimating the velocity dispersions from the observed luminosities via the Faber-Jackson and Tully-Fisher relations we found that SN2003es is likely to have been significantly magnified by gravitational lensing and that this problem should be addressed in detail before using SN2003es in estimates of the cosmological parameters.  \\ack The author would like to thank Julio Navarro for permission to use his code for calculation of the NFW parameters, Phillip Helbig for permission to use his code for calculations of cosmological distances, Ariel Goobar, Joakim Edsj\\\"o and Edvard M\\\"ortsell for useful discussions, suggestions and proofreading. Finally, the author thanks Tomas Dahl\\'en. This project was funded by the Swedish Research Council."
    },
    "0311/astro-ph0311455_arXiv.txt": {
        "abstract": "We explain the linear polarization of the Ca {\\sc ii} infrared triplet observed close to the edge of the solar disk. In particular, we demonstrate that the physical origin of the enigmatic polarizations of the 866.2 nm and 854.2 nm lines lies in the existence of atomic polarization in their metastable $^2{\\rm D}_{3/2\\,,\\,5/2}$ lower levels, which produces differential absorption of polarization components (dichroism). To this end, we have solved the problem of the generation and transfer of polarized radiation by taking fully into account all the relevant optical pumping mechanisms in multilevel atomic models. We argue that `zero-field' dichroism may be of great diagnostic value in astrophysics.  \\vskip 2truecm  {\\bf Published in Physical Review Letters: Volume 91, Number 11, 111102, (2003)}  \\vskip 5truecm  . ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311039_arXiv.txt": {
        "abstract": "{We study the long term evolution of the relationship between the root mean square (rms) variability and flux (the ``rms-flux relation'') for the black hole Cygnus~X-1 as monitored from 1996 to 2003 with the Rossi X-ray Timing Explorer (\\textsl{RXTE}). We confirm earlier results by \\citet{uttley:01} of a linear relationship between rms and flux in the hard state on time scales $>5$\\,s reflecting in its slope the fractional rms variability. We demonstrate the perpetuation of the linear rms-flux relation in the soft and the intermediate state. The existence of a non-zero intercept in the linear rms-flux relation argues for two lightcurve components, for example, one variable and one non-variable component, or a possible constant rms component. The relationship between these two hypothesized components can be described by a fundamental dependence of slope and intercept at time scales $\\lesssim 10$\\,ksec with \\emph{long term} averages of the flux. ",
        "introduction": "\\label{sec:intro}  Galactic Black holes (BHs) are generally found in two major states which are mainly characterized by their X-ray and radio spectral and timing properties \\citep[see, e.g.,][and references therein]{nowak:02a}. The ``hard state'', occuring at luminosities $L\\lesssim 0.05\\,L_{\\text{Edd}}$, where $L_{\\text{Edd}}$ is the Eddington luminosity, is characterized by $\\sim$30\\% root mean square (rms) variability, a hard X-ray spectrum that can be roughly described by an exponentially cutoff power law, and radio emission. The hard state energy spectrum can physically be modelled by thermal Comptonization of cool seed photons in a hot electron plasma \\citep{thorne:75,shapiro:76,sunyaev:79a,dove:97c}. At higher luminosities, BHs are seen in the ``soft state'' with a thermal X-ray spectrum with a characteristic temperature of at most a few keV, X-ray variability of only a few percent rms, and no observed radio emission \\citep{fender:02a}. The thermal X-ray spectrum in the soft state is generally associated with a geometrically thin, optically thick accretion disk \\citep{novikov:73a,shakura:73,gierlinski:98a}. In addition to these two canonical states, further states have been identified which are characterized either by an even larger luminosity than in the soft state (the ``very high state''), or by variability and X-ray spectral properties which are mostly intermediate between the hard and the soft state \\citep[the ``intermediate state'';][]{belloni:96}.  The analysis of the short term X-ray variability of BHs is especially well suited for gaining better insight into the physical processes at work near the compact object. In order to analyze this behavior in a systematic way, in 1997 we initiated a long term monitoring campaign of the canonical hard state BH Cygnus~X-1. During the campaign, which is still ongoing, Cyg~X-1 is observed at $\\sim$7\\,d to $\\sim$14\\,d intervals simultaneously with the Rossi X-ray Timing Explorer (\\textsl{RXTE}), the Ryle radio telescope, and several optical instruments.  These observations enable us to study the evolution of the observational properties of Cyg~X-1 (e.g., X-ray spectrum and timing parameters, radio flux) and thus the source history on time scales longer than the orbital time scale of $\\sim 5.6$\\,days.  Earlier results from our campaign have been reported by \\citet{pottschmidt:00a} and by \\citet[][hereafter paper~I]{pottschmidt:03}, where we studied the evolution of the Fourier-frequency dependent X-ray time lags and of the power spectral density (PSD), respectively. Following \\citet{nowak:00a}, in paper~I we were able to show that for a very wide range of source parameters, the PSD of Cyg~X-1 is well described by the superposition of broad Lorentzian functions. Such a description is an alternative to the more common description of the PSD in terms of broken power laws \\citep{klis:95} or shot-noise profiles \\citep[][and references therein]{lochner:91}.  During changes of the source into the ``intermediate state'' or the ``soft state'', the characteristic frequency of these components shifts towards higher frequencies in a systematic manner, some of the Lorentzians change their strength or vanish, and the X-ray time lag increases in the frequency band dominated by those Lorentzians which become stronger \\citep[][paper~I]{pottschmidt:00a}. Once the ``soft state'' is reached, the X-ray time lags are again similar to those seen during the hard state.  In paper~I we interpreted these results in terms of models invoking damped oscillations in the accretion disk and/or the Comptonizing plasma \\citep{dimatteo:99a,psaltis:00a,churazov:01a,nowak:02b}. Because of the complicated physics of the system ``accretion disk -- accretion disk corona -- radio outflow'', none of these models are fully self-consistent and thus none of the models describes all observed properties of Cyg~X-1.  It is our hope that additional empirical data for the hard state behavior might help to further constrain these different emission models of BHs.  Recently, \\citet{uttley:01} found a remarkable linear relationship between the rms amplitude of broadband noise variability and flux in the X-ray lightcurves of neutron star and BH X-ray binaries and active galactic nuclei (AGN). This linear relation between flux and rms variability has an offset on the flux axis, leading to the suggestion that the lightcurves of X-ray binaries are made from at least two components: one component with a linear dependence of rms variability on flux, and one component which contributes a constant rms to the lightcurve or does not vary at all. The fact that the same rms-flux relation was found in the lightcurves of active galactic nuclei suggests that this behavior is intrinsic to compact accreting systems.  In the case of X-ray binaries, the rms-flux relation has until now only been investigated for single observations of SAX~J1808.4$-$3658 and for a few hard state observations of Cyg~X-1.  It is important to discover whether this relation applies generally for many observations of a given object and across a range of states.  Therefore in this paper we investigate the evolution and behavior of the rms-flux relation in Cyg X-1 using the \\textsl{RXTE} monitoring data. Preliminary results have already been presented elsewhere \\citep{gleissner:02a}.  The remainder of this paper is structured as follows. In Sect.~\\ref{sec:data} we describe the data and the computation of the rms variability. We then present our results on the rms-flux relation (Sect.~\\ref{sec:results}). We show that the rms-flux relation is valid throughout all spectral states, for all energies and on all time scales. It is demonstrated how mean flux, $\\langle F\\rangle$, and the relation fit parameters slope, $k$, and intercept on the flux axis, $C$, map out a fundamental plane in the hard state. After analyzing spectral dependences, we examine the rms-flux relation on short time scales by calculating the variability of the rms variability. We end the paper with a discussion of the physical interpretations, and explore the meaning of the long term rms-flux relation as well as of the hard state $k$-$C$-$\\langle F\\rangle$ fundamental plane (Sect.~\\ref{sec:disc}). ",
        "conclusions": "\\label{sec:disc}  \\subsection{The origin of the rms-flux relation} \\citet{uttley:01} have interpreted the rms-flux relation in terms of the disk fluctuation model of \\citet{lyubarskii:97a}, where variations in the accretion rate occur at various radii and propagate inwards to modulate the X-ray emission.  The fact that variations at larger radii should have longer characteristic time scales naturally explains the fact that the short term rms variations are modulated by longer time scale flux changes.  As the slower variations in the accretion flow propagate inwards they are able to modulate the variations on shorter time scales. \\citet{kotov:01a} have extended this model by introducing an extended X-ray emitting region with a temperature gradient, to explain the frequency dependent time lags between hard and soft bands observed in Cyg~X-1 \\citep{pottschmidt:00a}.  \\citet{churazov:01a} have pointed out that in a standard thin disk, viscous damping will smooth accretion rate fluctuations on a viscous time scale, i.e., before they reach the inner disk, implying that for the \\citet{lyubarskii:97a} model to work, the fluctuations in accretion rate must happen in an optically thin or geometrically thick flow, such as an advection dominated accretion flow (ADAF) or another type of thick disk.  We have shown that the linear rms-flux relation is seen in both high and low flux states, so that it very well might be associated with some component of the system that is common to both states.  Since the X-ray variability during the soft state appears to be associated with the power law emission rather than thermal emission from the disk \\citep[which may be constant,][]{churazov:01a}, it seems likely that the rms-flux relation is associated directly with a corona which is present in both hard and soft states.  We therefore speculate that the rms-flux relation in Cyg~X-1 (and probably also other objects showing linear rms-flux relations) may present evidence for accretion rate variations in a \\emph{fluctuating coronal accretion flow} \\citep[see, e.g.,][and references therein]{smith:01b, klis:01a}. The suggestion that the corona itself accretes has been used to explain a number of observational features of both AGN and galactic BH systems \\citep{witt:97}, and since an optically thin coronal accretion flow is similar to an ADAF, such a model would not encounter the problems faced by a thin disk origin for the accretion variations \\citep{churazov:01a}.  \\subsection{The long term rms-flux relation}\\label{sec:longterm} In the previous sections we have shown that the linear rms-flux relation seems to apply on very long time scales, since the average 1--32\\,Hz rms of an observation scales linearly with the average flux of the observation.  The hard and soft state data form two separate long term rms-flux relations, with different slopes $k_\\text{long}$ (somewhat lower in the soft state, $k_\\text{long}=0.19$ versus $k_\\text{long}=0.22$ in the hard state).  Intriguingly, the flux-offsets $C_\\text{long}$ in the long term rms-flux relations of both hard and soft states are very close to zero, and certainly of much smaller amplitude than the typical offsets observed in the short term rms-flux relations of either state.  The scatter in the hard state long term rms-flux relation is particularly small.  Combined with the almost-zero offset of the long term relation, this tight correlation implies that the fractional rms of the 1--32\\,Hz variability in the hard state is remarkably constant (similar to $k_\\text{long}$, i.e., $\\sim$22\\%).  This result is rather surprising because the hard state PSD is clearly not stationary, i.e., the Lorentzians which successfully describe the PSD change significantly in both peak frequency and fractional rms (paper~I).  The most obvious example of such a change in the PSD can be seen in 1998~May (see paper~I, Fig.~3) when the peak frequencies of the Lorentzians simultaneously increase, and the fractional rms of the $L_{3}$ Lorentzian (which lies between 3--10\\,Hz) and the fractional rms of the total PSD both decrease significantly.  However, despite such a marked change in the PSD, the data before and after 1998~May lie along the same long term hard state rms-flux relation.  The fact that the fractional rms in the 1--32\\,Hz band seems undisturbed by significant changes in the PSD suggests some sort of conservation of fractional rms at high frequencies, even though the low-frequency (and hence the total) fractional rms changes significantly.  Although the specific 1--32\\,Hz band chosen to measure the rms-flux relation is somewhat arbitrary, we point out that the bulk of the power in this band is associated with the $L_2$ and $L_3$ Lorentzians, and the band fully encompasses the range of frequencies between 3--10\\,Hz in which a strong and well-defined anti-correlation between Lorentzian peak frequency and fractional rms is clearly seen (see paper~I, Fig.~6). We suggest, therefore, that some process which conserves fractional rms above a few~Hz may be at work, with fractional rms in the $L_{3}$ Lorentzian decreasing as the $L_{2}$ Lorentzian moves further into the 1--32\\,Hz band, resulting in a constant fractional rms in that band. We now examine how such a phenomenological model relates to the observed fundamental plane of short term rms-flux relation parameters $k$, $C$, and flux.  \\subsection{The hard state $k$-$C$-$\\langle F\\rangle$ fundamental plane} We have shown that the ``fan shape'' of the correlation between the rms-flux relation slope and its flux-offset in the hard state can be directly related to the long term rms-flux relation.  The fact that $k$ and $C$ vary but the average 1--32\\,Hz rms \\emph{must} track the average flux explains the degeneracy in the correlation, which is the result of the various fluxes observed for a given pair of $k$ and $C$ values.  Thus $k$, $C$, and $\\langle F\\rangle$ track out a fundamental plane which describes the forms the rms-flux relation can take in the hard state.  As noted by \\citet{uttley:01}, $C$ can be interpreted either simply as a constant-flux component to the lightcurve, or as a constant rms component. We concentrate on these two interpretations since they are the simplest possibilities, but it should be emphasized that there are numerous other forms the variability could take, e.g., a component of constant fractional rms.  We have shown that, in the soft state at least, $C$ can be negative, which is only possible if the offset in this state represents a constant rms component.  We argue that the existence of the $k$-$C$-$\\langle F\\rangle$ fundamental plane also suggests that $C$ represents a constant rms component in the hard state.  For example, consider an increase in the slope $k$ between two observations, while keeping the average flux of the observations constant.  The increase in $k$ corresponds to an increase in the fractional rms of the linear rms-flux component of the lightcurve.  In order to maintain the constant total fractional rms as given by the observed long term rms-flux relation, $C$ must also increase.  Therefore we suggest that both a constant rms component and a linear rms-flux component exist in the lightcurve and that the relative contribution of each varies in such a way that the sum of both components, the total fractional rms, is constant.  As a constraint to possible models it should be mentioned, as already noted in Sect.~\\ref{subsec:spectral}, that the spectral shapes of $k$ and $C$ are generally similar, with $C$ tending to a flatter/harder spectrum, and there is no obvious correlation between spectral index $\\Gamma$ and $C$, $C_\\text{long}$, and $k$, respectively.  \\subsection{Summary} To summarize, in this paper we analyzed Cyg~X-1 lightcurves from 1996 to 2003 in an attempt to describe the long term evolution of X-ray rms variability.  Our main results are \\begin{enumerate} \\item The linear rms-flux relation is valid throughout all states and energy ranges observed. \\item The slope of the linear rms-flux relation is steeper in the hard state than in the soft and intermediate states. \\item The rms-flux relation is valid not only on short time scales (seconds) but also on longer time scales (weeks--months). \\item In the hard state, the slope of the flux-rms-relationship is correlated with its flux intercept, but the correlation is increasingly degenerate at higher values of $k$ and $C$, forming a fan-like distribution in the $k$-$C$ plane.  The hidden parameter governing the degeneracy is the long term average flux, $\\langle F\\rangle$, and the $k$-$C$-$\\langle F\\rangle$ relation maps out a ``fundamental plane'' in the low state. \\item The intercept of the linear rms-flux relation is relatively constant and positive in the hard state, but is diverse in the soft state: in the 1996 soft state $C$ is positive and slightly higher than during the hard state, in the atypical 2001/2002 ``soft state'' $C$ is decreasing relative to the hard state level, with several observations showing a negative $C$. During ``failed state transitions'', $C$ is smaller than in the hard state, similar to the soft state behavior. \\item The $F$-$\\sigma$-$\\Gamma$ correlation can be confirmed for the hard state. \\item The clear dichotomy of $k$ in a hard and a soft state level in energy band~2 is weaker at higher energies. \\end{enumerate}"
    },
    "0311/astro-ph0311513_arXiv.txt": {
        "abstract": "The free-electron population during the reionized epoch rescatters CMB temperature quadrupole and generates a now well-known polarization signal at large angular scales. While this contribution has been detected in the temperature-polarization cross power spectrum measured with WMAP data, due to the large cosmic variance associated with anisotropy measurements at tens of degree angular scales only limited information related to reionization, such as the optical depth to electron scattering, can be extracted. The inhomogeneities in the free-electron population lead to an additional secondary polarization anisotropy contribution at arcminute scales. While the fluctuation amplitude, relative to dominant primordial fluctuations, is small, we suggest that a cross-correlation between arcminute scale CMB polarization data and a tracer field of the high redshift universe, such as through fluctuations captured by the 21 cm neutral Hydrogen background or those in the infrared background related to first proto-galaxies, may allow one to study additional details related to reionization. For this purpose, we discuss an optimized higher order correlation measurement, in the form of a three-point function, including information from large angular scale CMB temperature anisotropies in addition to arcminute scale polarization signal related to inhomogeneous reionization. We suggest that the proposed bispectrum can be measured with a substantial signal-to-noise ratio and does not require  all-sky maps of CMB polarization or that of the tracer field. A measurement such as the one proposed may allow one to establish the epoch when CMB polarization related to reionization is generated and to address if the universe was reionized once or twice. ",
        "introduction": "The increase in sensitivity of upcoming cosmic microwave background (CMB) polarization experiments, both from ground and space, raises the possibility for detailed studies related to reionization. The main reason for this is the existence of a large angular scale polarization contribution due to rescattering of the temperature quadrupole by free-electrons in the reionized epoch \\cite{Zaldarriaga:1996ke}. This contribution peaks at angular scales corresponding to the horizon at the rescattering surface. While such a signal has now been detected with the temperature-polarization cross-correlation power spectrum measured with first-year data of the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{Bennett:2003bz}, information related to reionization from it, unfortunately, is limited.  While the amplitude of the signal depends on the total optical depth to electron scattering, with an estimated value of 0.17 $\\pm$ 0.04 \\cite{Kogetal03}, the reionization history related to this optical depth is still unknown. While  slight modifications to the large angular scale polarization power spectra exist with complex reionization histories, such as a two stage reionization process advocated by Ref.~\\cite{Cen:2002zc}, due to large cosmic variance associated with anisotropy measurements at few tens degree scales, one cannot fully reconstruct the reionization history as a function of redshift \\cite{Kaplinghat:2002vt}.  Given the importance of understanding reionization for both cosmological and astrophysical reasons, and the increase in sensitivity and angular resolution of upcoming CMB polarization data, it is useful to consider alternative possibilities for additional studies beyond the large angular scale signal. Under standard expectations for reionization, mainly due to UV light emitted by first luminous objects, the process of reionization is expected to be patchy and inhomogeneous \\cite{Barkana:2000fd}. This leads to fluctuations in the electron scattering optical depth and a modulation to the polarization contribution such that new anisotropy fluctuations, at arcminute scales corresponding to inhomogeneities in the visibility function, is generated \\cite{Hu:1999vq,Santos:2003jb}. Even if the reionization were to be more uniform, as expected from alternative models involving X-ray backgrounds \\cite{Oh:2000zx} and reionization via particle decays \\cite{Hansen:2003yj}, density fluctuations in the electron field can modulate the polarization contribution and generate secondary anisotropies \\cite{Baumann:2002es}. The new secondary polarization fluctuations, whether due to patchiness or density inhomogeneities, however, is hardly detectable in the polarization power spectrum given the dominant primordial polarization contribution, with a peak at arcminute scales, related to the velocity field at the last scattering surface. Thus, information related to reionization is limited from measurements that involve simply the two point correlation function of polarization or the cross-correlation between polarization and the temperature. To extract additional information one must move to a higher order and consider measurements related to, for example, the three-point correlation function or the bispectrum \\cite{Cooray:1999kg}.  In this paper, we consider a potentially interesting study that has the capability to allow detailed measurements related to the reionization history. The proposed measurement involves the cross-correlation between CMB polarization maps and images of the high redshift universe around the era of reionization or its end. The existence of such a correlation arises from the fact that the polarization contribution related to reionization is generated at the same epoch as that related to the tracer field selected from high redshifts. Since the reionization contribution to polarization is generated by the scattering of the temperature quadrupole, direct cross-correlation between polarization maps and a tracer field is not useful. However, since the rescattering quadrupole correlates with CMB temperature maps, one must consider higher order statistics that also use information from CMB temperature data. We consider a possibility in the form of a three point correlation function, or a bispectrum in Fourier space, involving CMB polarization maps at arcminute-scale resolution, CMB temperature maps with large angular scale anisotropy information, and maps of the high redshift universe that trace fluctuations during the reionization and prior to reionization. Examples of useful fields include maps of the infrared background anisotropies related to the redshifted UV emission by first proto galaxies \\cite{Cooray:2003yf} and fluctuations in the 21 cm background due to the neutral Hydrogen content \\cite{Tozzi:1999zh}. Additional possibilities also include maps of the $z \\sim 3$ universe since one can then establish the low-redshift part of the electron scattering visibility function. Cross-correlation studies with the 21 cm background may be the most interesting possibility given that one is allowed to a priori select redshift ranges from which fluctuations arise to the redshifted 21 cm emission, based on frequency information of the observations.  The paper is organized as follows.  In \\S~\\ref{sec:deriv}, we derive the existence of a higher order correlation between CMB temperature, polarization and a tracer field of the large scale structure that overlaps with fluctuations in the visibility function. In \\S~\\ref{sec:results}, we discuss our results and suggest that there is adequate signal-to-noise to perform such a measurements and consider several applications. We conclude with a summary in \\S~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  The free-electron population during the reionized epoch rescatters CMB temperature quadrupole and generates a now well-known polarization signal at large angular scales. While this contribution has been detected in the temperature-polarization cross power spectrum from WMAP data, due to the large cosmic variance associated with anisotropy measurements at relevant scales, only limited information related to reionization, such as the optical depth to electron scattering, can be extracted. The inhomogeneities in the free-electron population lead to an additional secondary polarization anisotropy contribution at arcminute scales.  While the fluctuation amplitude, relative to dominant primordial fluctuations, is small, we suggest that a cross-correlation between arcminute scale CMB polarization data and a tracer field of the high redshift universe, such as through fluctuations captured by the 21 cm neutral Hydrogen background or those in the infrared background related to first proto-galaxies, may allow one to study additional details related to reionization. This includes a possibility to recover the visibility function related to electron scattering. For this purpose, we discuss an optimized higher-order correlation measurement, in the form of a three-point function, involving information from large angular scale CMB temperature anisotropies in addition to arcminute scale polarization data.  The proposed bispectrum can be measured with a substantial signal-to-noise ratio with lens-cleaned B-mode maps where the confusion related to primordial anisotropies and the arcminute scale reionization-related polarization is the least; this can be compared to other sources of non-Gaussianity in CMB data, such as due to gravitational lensing, which with the same combination of temperature, polarization, and the tracer field is at the level of at most unity in terms of the signal-to-noise ratio.  Given that the arcminute scale polarization fluctuations are correlated with fluctuations at high redshift at same angular scales, for a reliable measurement of the three point correlation function suggested here, one does not require all-sky maps of CMB polarization or that of the tracer field. This is helpful since one can obtain high signal-to-noise maps of the polarization from ground-based experiments, and in certain cases fluctuation data on certain tracers, by concentrating on limited sky-area instead of the whole sky. The required information from temperature anisotropies is at large angular scales and this information is already available to the limit allowed by the cosmic variance with WMAP data and is expected to improve with Planck in terms of the foreground confusion and separation. A study such as the one proposed may allow one to establish the epoch when CMB polarization related to reionization is generated and to address if the universe was reionized once or twice."
    },
    "0311/astro-ph0311220.txt": {
        "abstract": "We present a comprehensive examination of the X-ray variability of the narrow line Seyfert 1 (NLS1) galaxy NGC~4051, one of the most variable AGN in the sky. We combine over 6.5 years of frequent monitoring observations by {\\it RXTE} with a $>100$~ks continuous observation by {\\em XMM-Newton} and so present an overall 2-10 keV powerspectral density (PSD) covering an unprecedent frequency range of over 6.5 decades from $<10^{-8}$ to $>10^{-2}$ Hz.  The combined {\\it RXTE} and {\\em XMM-Newton} PSD is a very good match to the PSD of the galactic black hole binary system (GBH) Cyg~X-1 when in a `high', rather than `low', state providing the first definite confirmation of an AGN in a `high' state. We also find that a bending powerlaw, rather than a sharply broken powerlaw, besides being more physical, is a much better description of the high state PSD of Cyg~X-1 and is also a better description of the PSD of NGC~4051.  At low frequencies the PSD of NGC~4051 has a slope of -1.1 bending, at a frequency $\\nu_{B}=8^{+4}_{-3}\\times10^{-4}$Hz, to a slope of $\\alpha_{H} \\sim -2$. Although $\\nu_{B}$ does not depend on photon energy, $\\alpha_{H}$ is steeper at lower energies.  If $\\nu_{B}$ scales with mass, we imply a black hole mass of $3^{+2}_{-1} \\times 10^{5} M_{\\odot}$ in NGC~4051, which is consistent with the recently reported reverberation value of $5^{+6}_{-3} \\times 10^{5} M_{\\odot}$. Hence NGC~4051 is emitting at $\\sim30\\% \\rm~ L_{Edd}$.  NGC~4051 follows the same rms-flux relationship as GBHs, consistent with higher Fourier frequencies being associated with smaller radii.  From the cross-powerspectra and cross-correlation functions between {\\em XMM-Newton} lightcurves in different energy bands, we note that the higher energy photons lag the lower energy ones. We also note that the lag is greater for variations of longer Fourier period and increases with the energy separation of the bands.  Variations in different wavebands are very coherent at long Fourier periods but the coherence decreases at shorter periods and as the energy separation between bands increases. This behaviour is again similar to that of GBHs, and of MCG-6-30-15, and suggests a radial distribution of frequencies and photon energies with higher energies and higher frequencies being associated with smaller radii.  Combining our observations with observations from the literature we find it is not possible to fit all AGN to the same linear scaling of break timescale with black hole mass. However broad line AGN are consistent with a linear scaling of break timescale with mass from Cyg~X-1 in its low state and NLS1 galaxies scale better with Cyg~X-1 in its high state. We suggest that the relationship between black hole mass and break timescale is a function of at least one other underlying parameter which may be accretion rate or black hole spin or both. ",
        "introduction": "Early observations with EXOSAT showed that, on short timescales (day to minutes) AGN X-ray variability was quasi-scale invariant, with power spectral densities (PSDs) being reasonably described by simple powerlaws or perhaps slightly curving functions (Lawrence \\etal 1987; \\mch and Czerny 1987).  However examination of archival data (\\mch 1988) showed that, on longer timescales ($\\gtsim$weeks) the scale invariance of at least one AGN, NGC~5506, broke down and its PSD flattened below a bend, or `knee' frequency of $\\sim 10^{-6}$Hz.  The combined long and short timescale PSD of NGC~5506 resembled that of galactic black hole X-ray binary systems (GBHs) such as Cyg~X-1 (eg Nowak \\etal 1999) which, when in a `low' state, has a `knee' frequency of $\\sim 0.2$Hz, below which the PSD flattens to a slope of zero. If the knee frequencies scale with mass, then a reasonable black hole mass in NGC~5506 of $\\sim$few $\\times 10^{6}M_{\\circ}$ was implied. Thus the possibility arose of determining AGN black hole masses from observations of long timescale X-ray variability.  Again using a variety of archival data, Papadakis and \\mch (1995) demonstrated that the PSD of the brighter AGN, NGC~4151, also flattened at low frequencies ($<10^{-7}$Hz) but long term X-ray lightcurves remained of poor quality until the launch of the Rossi X-ray Timing Explorer ({\\it RXTE}; Swank 1998) in December 1995.  {\\it RXTE} can sample on sub-daily timescales and so, since shortly after launch we, and other groups, have been monitoring a small sample of AGN in order to produce high quality long timescale PSDs. One of our main targets is NGC~4051 and results from the first 1.5 years of observation (\\mch \\etal 1998) indicated that the slope at low frequencies was probably flatter than that at high frequencies but neither the break frequency, or the slopes, were particularly well defined. Edelson and Nandra (1999), Pounds \\etal (2001) and Markowitz \\etal (2003) have all now demonstrated PSD flattening in other AGN at low frequencies, using {\\it RXTE} data.  Although {\\it RXTE} provides much better long term X-ray lightcurves than any previous observatory, the lightcurves are still not continously sampled, or even completely uniformly sampled. Thus the `raw' or `dirty' PSDs will be contaminated by artefacts which depend on the window sampling pattern. In order to overcome such problems we have developed a simulation-based modelling process (Uttley, M$\\rm^{c}$Hardy and Papadakis 2002) which can determine not only the value of any break frequencies, or PSD slopes, but can also derive errors for these parameters, and also quantify the goodness of fit of any model. Uttley \\etal (2002) applied this method to the first 3 years of observations of four AGN. They found evidence for bends, or breaks, in the PSDs in three AGN, NGC~5506, NGC~3516 and MCG-6-30-15, although not in NGC~5548.   GBHs are found in a variety of `states' (eg McClintock and Remillard 2003; Cui \\etal 1997a), the most common of which are the so called low flux, hard spectrum (or `low') state, and the high flux, soft spectrum (or `high') state. In the standard Comptonisation models for the production of X-ray emission, different physical regimes apply to low and high states, with a higher seed photon flux in the high state, leading to a cooler Comptonising corona (eg Zdziarski \\etal 2002).  If we are to attempt physical modelling of AGN we therefore need to know which state they are in. For GBHs, the PSDs of the high (eg Cui \\etal 1997b, Churazov \\etal 2001) and low (eg Nowak \\etal 1999; Revnivtsev, Gilfanov and Churazov 2000) states are different, thereby providing us with a means of discrimination.  In Cyg~X-1 in the low state, the PSD slope above the `knee' (refered to above, at $\\sim 0.2$Hz) is $\\sim -1$. However at $\\sim 3$Hz there is a second, higher frequency, break above which the PSD slope is $\\sim-2$. In the high state only one break has been detected so far in the PSD of Cyg~X-1.  The break frequency is $\\sim10-20$Hz and, above the break, the PSD slope is steeper than -2. Below the break the PSD has a slope of -1 for at least 4 decades of frequency (Reig \\etal 2002).  For AGN, there are two cases (Akn~564, Papadakis \\etal 2002 and NGC~3783, Markowitz \\etal 2003) where two breaks have been found, indicating a possible similarity to low state systems. However in most other AGN only one break has been found and, although the slopes below the break are not that well defined (eg Uttley \\etal 2002; Markowitz \\etal 2003), in general they are closer to -1 than to zero. However it is not clear whether these breaks correspond to the high frequency break in low state systems or the single break in high state systems.   If we assume that timescales scale approximately with mass, then to properly compare AGN with GBHs we require AGN PSDs which cover timescales from $\\sim$few years to tens of seconds.  {\\it RXTE} monitoring observations principally sample timescales from $\\sim$few years to $\\sim$day but {\\em XMM-Newton} can provide continous observations of up to almost 2 days with sampling on $\\sim$second, or shorter, timescales. NGC~4051 is the AGN which is best observed on long timescales by {\\it RXTE} and is also one of the best observed on short timescale with {\\em XMM-Newton}.  Here we present a combined PSD analysis of these two datasets.   In Section~\\ref{sec:obs} we describe the observations and the data analysis procedures. In Section~\\ref{sec:psd} we discuss the combined {\\it RXTE} and {\\em XMM-Newton} PSD, covering more than 6 decades in frequency, and compare it to that of GBHs, particularly Cyg~X-1, in different states.  As further diagnostics of the Comptonisation process we discuss the high frequency {\\em XMM-Newton} data in more detail, and note that a similar analysis of a long {\\em XMM-Newton} observation of MCG-6-30-15 has been carried out by Vaughan, Fabian and Nandra (2003), producing broadly similar results.  In Section~\\ref{sec:psd} we note a variation of PSD slope with photon energy at high frequencies.  In Section~\\ref{sec:rms} we determine the rms-flux relation (Uttley and \\mch 2001) from the {\\em XMM-Newton} data.  We discuss the relationship between variations in different wavebands both in terms of the cross-correlation function (Section~\\ref{sec:crosscor}) and using cross-spectrum phase spectrum analysis (Section~\\ref{sec:crosspectrum}). In Section~\\ref{sec:coherence} we discuss the coherence of the variations between different wavebands, ie whether the same variations occur in all bands or whether there are independent contributions in different wavebands.  In Section~\\ref{sec:summary} we summarise the main observational results from this paper. Finally, in Section~\\ref{sec:implications}, we compare NGC~4051 with high and low state GBHs and discuss possible physical models for the X-ray emission region.  We also examine the relationship between PSD break timescale and black hole mass in AGN, noting that the relationship is probably different for broad and narrow line Seyfert 1 galaxies, perhaps because of a difference in mass accretion rate or spin. ",
        "conclusions": "\\label{sec:summary}  In this Section we summarise the main observational results of this paper. In the next Section (\\S~\\ref{sec:implications}) we discuss the implications of these results.  \\begin{enumerate}  \\item We present $>6.5$ years of well sampled X-ray observations of NGC~4051 by {\\it RXTE}, together with $\\sim100$~ks of continuous observations with {\\em XMM-Newton}. The resulting PSD spans over 6.5 decades of frequency and is the best determined AGN PSD yet published.  \\item A gently bending powerlaw model is a better fit to the combined {\\it RXTE} and {\\em XMM-Newton} PSD of NGC~4051, although a sharply broken powerlaw is still an acceptable fit.  The gently bending model is also a much better fit to the high S/N high state PSD of Cyg~X-1 than a sharply broken powerlaw. As a gentle bend is also more plausible physically than a sharp break, we use it as our standard description of the PSDs of NGC~4051.  \\item The combined {\\it RXTE} and {\\em XMM-Newton} (4-10 keV) PSD has a slope, $\\alpha_{L}$, of $\\sim-1.1$ at low frequencies and steepens, above a break frequency, $\\nu_{B}$, of $8 \\times 10^{-4}$Hz, to a slope, $\\alpha_{H}$, of $\\sim -2$. There is no evidence of a second break at lower frequencies.  \\item In NGC~4051 $\\nu_{B}$ is almost independent of energy, as in Cyg~X-1 in the high state.  \\item In NGC~4051 $\\alpha_{H}$ is steeper at lower energies ($\\sim -3$, 0.1-2 keV; $\\sim -2$, 2-10 keV) whereas in Cyg~X-1 in the high state $\\alpha_{H} \\sim -3$ over all of the presently observable spectral range (2-13 keV).  \\item NGC~4051 shows the same rms-flux relationship (Uttley and \\mch 2001) which is displayed by GBHs, including the presence of small component of constant rms.  \\item As shown by the cross-correlation analysis (Fig.~\\ref{fig:ccf}), variations in different energy bands are very well correlated in NGC~4051. However as the phase-spectral analysis shows (Fig.~\\ref{fig:lags}), there are complex lags between bands, with the hard photons always lagging the soft photons. The lag is greatest for variations of large Fourier period, and also increases as the energy separation between bands increases.  \\item The coherence between wavebands is very high at long Fourier periods but decreases at shorter periods, and with increasing separation between wavebands. We note that the decrease in coherence becomes noticeable at the same timescale as the break timescale in the PSD. \\end{enumerate}"
    },
    "0311/cond-mat0311304_arXiv.txt": {
        "abstract": "The sun provides an explosive, heavenly example of self-organized criticality.  Sudden bursts of intense radiation emanate from rapid rearrangements of the magnetic field network in the corona. Avalanches are triggered by loops of  flux that reconnect or snap into lower energy configurations when they are overly stressed. Our recent analysis of observational data reveals that the loops (links) and footpoints (nodes), where they attach on the photosphere, embody a scale free network.  The statistics of the avalanches and of the  network structure are unified through a simple dynamical model where the avalanches and network co-generate each other into a complex, critical state.  This particular example points toward a general dynamical mechanism for self-generation of complex networks. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311548_arXiv.txt": {
        "abstract": "By extending the constant-acceleration analysis of Smith, Mao, \\& Paczy\\'nski to include jerk, I show that microlens parallax measurements are subject to a four-fold discrete degeneracy.  The new degeneracy is characterized by a projected velocity $\\tilde v_j = -(3/4)\\csc\\beta_\\ec(\\cos^2\\psi\\sin^2\\beta_\\ec+ \\sin^2\\psi)^{3/2}v_\\oplus$, where $\\beta_\\ec$ is the ecliptic latitude, $\\psi$ is the phase of the Earth's orbit relative to opposition at the time of the event maximum, and $v_\\oplus=30\\,\\kms$ is the speed of the Earth.  The degeneracy becomes important when the lens projected velocity $\\tilde v$ is of order $\\tilde v_j$.  For events toward the Large Magellanic Cloud, $\\tilde v_j\\simeq (3/4)v_\\oplus$, so this degeneracy is important primarily for lenses in the Milky Way disk.  In particular, it solves the puzzle of MACHO-LMC-5, whose microlens parallax measurement had yielded mass and distance determinations for the lens that were inconsistent with photometric estimates.  Toward the Galactic bulge, $\\tilde v_j$ ranges from $\\sim 0.2\\,\\kms$ at the solstice to $\\sim 200\\,\\kms$ at the equinoxes, so the effect of the degeneracy depends strongly on the peak time of the event.  The degeneracy applies mainly to events with Einstein timescales, $t_\\e\\la {\\rm yr}/2\\pi$. ",
        "introduction": "}  The spectacular high-magnification microlensing event MACHO-LMC-5 has been a puzzle since its discovery was reported by the MACHO collaboration \\citep{alcock97}. The ``source star'' for the event is quite red at baseline but lies well below the Large Magellanic Cloud (LMC) giant branch.  On the other hand, the event itself is quite blue, indicating that the true lensed source is blue and that the red star at baseline is not being microlensed. \\citet{gbf} proposed that the red star was a foreground M dwarf in the Milky Way disk and that this M dwarf was in fact the microlens.  While the optical depth of the Milky Way disk is extremely low, $\\tau\\sim 10^{-9}$, the MACHO survey was sufficiently big that one should expect of order one such event.  {\\it Hubble Space Telescope (HST)} observations carried out by \\citet{alcockprelim} and further analyzed by \\citet{alcock01} virtually proved this conjecture.  The {\\it HST} Planetary Camera images taken 6.3 years after the event clearly show two stars, one blue and one red, separated by 134 mas.  The chance that an unrelated foreground M dwarf would lie so close to a given microlensed source is only $\\sim 10^{-4}$.  Even the conditional probability of such an alignment, given that a red star was known to lie within the $\\sim 2''$ MACHO PSF, is only a few percent.  \\citet{alcock01} therefore concluded that the red dwarf was in fact the lens.  However, this identification turned out to compound rather than resolve the puzzle of MACHO-LMC-5.  The lightcurve of this event shows a clear asymmetry about its peak (see Fig.~1).  Such asymmetries can be induced by parallax effects due to the Earth's orbital motion \\citep{gould92}, as was first observed by \\citet{al1} and subsequently in about a dozen other events.  Measuring the parallax effect yields the (two-dimensional) velocity projected onto the observer plane, $\\tilde \\bv$, of the lens relative to the source.  The lens mass and lens-source relative parallax are then given by, \\begin{equation} M = {\\tilde v\\mu_\\rel t_\\e^2\\over 4 G}, \\qquad \\pi_\\rel = {\\mu_\\rel\\over\\tilde v}, \\label{eqn:mpi} \\end{equation} where $t_\\e$ is the Einstein crossing time (measured during the event), and $\\mu=134\\,{\\rm mas}/6.3\\,{\\rm yr} = 21.4\\pm 0.7\\,\\rm mas\\,yr^{-1}$ (measured from the {\\it HST} image).  Combining these various measurements, \\citet{alcock01} found, \\begin{equation} M = 0.036^{+0.009}_{-0.004}\\,M_\\odot,\\qquad \\pi_\\rel = 5\\pm 1\\,{\\rm mas}, \\qquad (\\rm microlensing), \\label{eqn:mpieval} \\end{equation} the latter value being equivalent to a lens distance of $D_l\\sim 200\\,$pc.  Of course, the very fact that the lens could be seen would seem to argue against such a small, substellar mass.  Indeed, from the observed color and magnitude of the lens, \\citet{alcock01} estimated its mass and distance to be, \\begin{equation} M = 0.11\\pm 0.02\\,M_\\odot,\\qquad D_l = 650\\,\\pm 190\\,{\\rm pc}\\qquad (\\rm photometric). \\label{eqn:mpieval2} \\end{equation}  One way out of this conflict would be simply to assume that the parallax measurement was wrong.  With a measured timescale of $t_\\e\\sim 30\\,$days, MACHO-LMC-5 was about a factor three shorter than any other event for which there is a reliable parallax.  Moreover, the parallax detection was only at about the $5\\,\\sigma$ level, small enough that it could conceivably be the product of small and unrecognized systematic errors.  However, \\citet{alcock01} argued that the parallax measurement could not be so easily dismissed because the direction of $\\tilde \\bv$ derived from this measurement agreed with the measured direction of $\\bmu_\\rel$ to within about $6^\\circ$.  The chance of such an agreement, if the parallax measurement were indeed spurious, is only about 3\\%.  Thus, a decade after its discovery, the event remains truly a puzzle.  Here I show that the parallax solution for MACHO-LMC-5 is subject to a four-fold degeneracy.  While two of these solutions are virtually identical to their counterparts, the other two lead to quite different estimates of the mass and distance.  One pair of solutions is equivalent to the solution reported by \\citet{alcock01}. The other pair yields both a larger mass and larger distance. I show that these are consistent with the photometric determinations. ",
        "conclusions": "}  \\subsection{Levels of Parallax Degeneracy \\label{sec:levels}}  The Taylor expansion of the (squared) lens-source separation, equations~(\\ref{eqn:usquared})-(\\ref{eqn:C4}), allows one to understand analytically the various levels of microlensing parallax degeneracy. If there is only enough information in the lightcurve to measure the first two terms (out to $C_2$), then there is no information at all about $\\pi_\\parallel$, while $\\pi_\\perp$ is completely degenerate with the timescale $t_\\e$.  Hence, nothing can be learned about the parallax, and it is usually assumed to vanish.  If the next ($C_3$) term can be measured (and if the second term in eq.~[\\ref{eqn:C3}] can be ignored) then this equation allows a determination of $\\pi_\\parallel$, but $\\pi_\\perp$ remains indeterminate. This is the degeneracy identified by \\citet{gmb}.  Taking account of the second term in equation~(\\ref{eqn:C3}) leaves a linear continuous (i.e., line-like) degeneracy in the $\\bpi_\\e$ plane, but rotates this line by an angle, \\begin{equation} \\theta_\\pi = \\tan^{-1}{\\eta\\pi_{j,\\parallel}\\over 1+\\eta\\pi_{j,\\perp}}, \\qquad\\eta\\equiv {1\\over 4}{u_0\\alpha t_\\e^2}. \\label{eqn:thetapi} \\end{equation} Since typically $\\eta\\pi_j\\ll 1$, this angle is also usually very small, but it can be significant in some cases (see \\S~\\ref{sec:implications}).  If $C_4$ can be measured, then the continuous degeneracy is broken, but it is replaced by a four-fold discrete degeneracy.  In this case, all four solutions lie along the line of the continuous degeneracy described in the previous paragraph.  The transition from the continuous to the discrete degeneracy is itself continuous. If $C_4$ is measured, but with only modest significance, then the error ellipses for each of the four solution will be extremely elongated along the line of continuous degeneracy.  In this case, the two solutions with $u_0>0$ may merge.  Similarly for the two solutions with $u_0<0$.  However, the $u_0>0$ and $u_0<0$ solutions cannot merge with each other unless $u_0^2$ is consistent with zero. To break the discrete degeneracy requires sensitivity to the higher order terms, which is typically obtained only in relatively long events.  \\subsection{Broad Implications \\label{sec:implications}}  For which microlensing events is the jerk-parallax degeneracy likely to be important?  To address this question, it is best to think in terms of the ``inverse projected velocity'', $\\bLambda$, \\begin{equation} \\bLambda \\equiv {\\tilde \\bv\\over \\tilde v^2} ={t_\\e\\over {\\rm AU}}\\bpi_\\e. \\label{eqn:bLambdadef} \\end{equation} (I choose ``$\\bLambda$'' for this quantity because it looks like an inverted ``$\\bv$''.) Since $t_\\e$ is virtually the same for all solutions, $\\bLambda$ is basically just a linear rescaling of $\\bpi_\\e$.  Its usefulness derives from the fact that the projected velocity (and so its inverse) is related solely to the kinematics of the event and thus is independent of the mass.  Expressed in terms of this quantity, equation~(\\ref{eqn:pisols}) becomes \\begin{equation} \\Lambda_\\parallel'=\\Lambda_\\parallel, \\qquad \\Lambda_\\perp'=-(\\Lambda_\\perp + \\tilde v_j^{-1}), \\label{eqn:bLambdasols} \\end{equation} where, \\begin{equation} \\tilde v_j \\equiv {3\\over 4}\\,{\\alpha^3\\over\\balpha\\times\\bj}{\\rm AU}. \\label{eqn:bLambdajdef} \\end{equation} (Note that since $\\balpha$ and $\\bj$ are 2-dimensional vectors, the denominator is a signed scalar.  See also eq.~[\\ref{eqn:C3}].)\\ \\ If we approximate the Earth's orbit as circular, then $\\tilde v_j$ may be evaluated, \\begin{equation} \\tilde v_j  = -{3\\over 4}\\, {(\\cos^2\\psi \\sin^2\\beta_\\ec + \\sin^2\\psi)^{3/2}\\over \\sin\\beta_\\ec}\\,v_\\oplus, \\label{eqn:Lambdajeval} \\end{equation} where $\\beta_\\ec$ is the ecliptic latitude of the event, $\\psi$ is the phase of the Earth's orbit at $t_0$ relative to opposition, and $v_\\oplus=30\\,\\kms$ is the speed of the Earth.  Note first that for LMC events, $\\tilde v_j \\simeq (3/4)v_\\oplus$. For the great majority of LMC events, those with lenses in the Galactic halo or the LMC itself, $\\tilde v\\gg v_\\oplus$.  Hence, $\\tilde v_j^{-1}\\gg \\Lambda$, so it would seem at first sight that the degeneracy would be very important.  In fact, for events with $\\tilde v\\gg v_\\oplus$, parallax effects will not generally be detectable unless the event is quite long, $\\Omega_\\oplus t_\\e>1$.  Such long events would most likely not in fact be degenerate because the corresponding high-parallax solution would have strong signatures at Taylor-expansion orders beyond $t^4$. It is only the disk-lens events, such as MACHO-LMC-5, for which this degeneracy is likely to cause confusion.  For events seen toward the bulge, the situation depends very strongly on the time of year.  For definiteness, consider events toward Baade's Window, for which $\\sin\\beta_\\ec\\sim -0.1$.  At opposition (roughly the summer solstice), $\\tilde v_j = 0.0075\\,v_\\oplus\\sim 0.2\\,\\kms$, which is much slower than the $\\tilde v$ of any plausible event. In this case, the reasoning just given for the majority of LMC events applies, so the degeneracy is not likely to be important.  At quadrature (roughly the equinoxes) $\\tilde v_j = 7.5\\,v_\\oplus\\sim 225\\,\\kms$, which is comparable to typical $\\tilde v$ for lensing of bulge sources by disk lenses and is smaller by a factor of a few than typical $\\tilde v$ for bulge-bulge lensing.  Hence, this degeneracy could have significant impact on the interpretation of events peaking near the equinoxes.  At intermediate times, $v_j$ scales roughly as $v_j\\sim 7.5\\,v_\\oplus\\sin^3\\psi$, so that, for example, at $\\psi=\\pm 60^\\circ$, $v_j\\sim 145\\,\\kms$, still large enough to cause significant confusion."
    },
    "0311/astro-ph0311581_arXiv.txt": {
        "abstract": "We analyze an X-ray flux-limited, complete sample of 93 AGN at z$< $0.1, selected from the ROSAT Bright Survey. Two thirds of the sample are Seyfert 1 galaxies (Sy1) and one third are Narrow-Line Seyfert 1 galaxies (NLSy1). We have obtained optical images of all objects. By modeling the host galaxy and the AGN central component we decompose the optical emission into nuclear, bulge and disk components, respectively.  We find that the nuclear optical luminosity, thought to be associated with the accretion disk surrounding the active black hole, correlates with the X-ray luminosity, the radio luminosity and the black hole mass. ",
        "introduction": "The last few years have brought fundamental progress in understanding the suspected connection between active galactic nuclei (AGN), characterized by non-thermal emission from a massive ($10^6 - 10^9$ \\msun) black hole (BH), and normal galaxies, whose extended emission is thermal.  In particular, now 1) there is robust evidence that quasars reside in galaxies (Bahcall et al. 1997); 2) it is established that massive ``quiescent'' BHs are also present in the central regions of normal galaxies (Magorrian et al. 1998); 3) a statistical correlation between BH mass and mass or velocity dispersion of the host galaxy's bulge exists (Gebhardt et al. 2000; Ferrarese \\& Merritt 2000), which suggests a common evolution for galaxies and their central, massive BHs; 4) the previous correlations hold for normal as well as AGN-host galaxies (McLure \\& Dunlop 2002).  For AGN with typical masses of $10^6 - 10^9$ \\msun, the bulk of the accretion radiation is emitted in the soft X-ray and EUV/UV range.  A fraction of this optically thick radiation is Compton-upscattered by a corona of electrons surrounding the BH. This process is thought to produce the typical power law X-ray spectra observed in all AGN. The primary soft X-ray/EUV radiation of the accretion disk also ionizes the surrounding gas, producing the well-known AGN emission lines. At radio frequencies, the core luminosity of the AGN is due to synchrotron emission by relativistic electrons immersed in a seemingly global magnetic field.  It is conceivable that these electrons are the same which Compton-upscatter the disk photons. It is therefore interesting to ask whether there is a direct correlation between radio, optical and X-ray emission in the central part of the AGN where the accretion disk is located.  While information in the radio and X-ray bands is already available from all-sky surveys like the NRAO VLA Sky Survey (NVSS, Condon \\etal 1998), FIRST (Faint Images of the Radio Sky at Twenty centimeters; Becker \\etal 1995) and the ROSAT Bright Survey (RBS, Schwope \\etal 2000), optical data for the central kiloparsec of AGN are more difficult to obtain.  Indeed, in modeling AGN and their host galaxies, the information contained in the central pixels has been usually rejected (see e.g. Virani \\etal 2000).  With a new well-tested, well-performing and robust fitting algorithm (Kuhlbrodt \\etal 2001), it has recently been possible to model both the host galaxy {\\em and} the nucleus from optical images.  Here we apply this decomposition algorithm to a new sample of objects, and use the results of the decomposition to study the relations between the properties of the central point source, in particular the X-ray ($0.5-2.0$ keV), optical (R band) and radio (1.49 GHz) emission. After presenting the sample selection in \\S 1 and the catalog data in \\S 2, the relations between the different parameters are described in \\S 3, and discussed in \\S 4. Unless specified differently, a value of H$_o$=50 km Mpc$^{-1}$s$^{-1}$ has been used throughout this paper. ",
        "conclusions": ""
    },
    "0311/astro-ph0311062_arXiv.txt": {
        "abstract": "We analyze the mass accretion histories (MAHs) and density profiles of cluster-size halos with virial masses of $0.6-2.5\\times 10^{14}h^{-1}\\rm\\ M_{\\odot}$ in a flat $\\Lambda$CDM cosmology. We find that most MAHs have a similar shape: an early, merger-dominated mass increase followed by a more gradual, accretion-dominated growth.  For some clusters the intense merger activity and rapid mass growth continue until the present-day epoch.  In agreement with previous studies, we find that the concentration of the density distribution is tightly correlated with the halo's MAH and with its formation redshift.  During the period of fast mass growth the concentration remains approximately constant and low $c_{\\rm v}\\approx 3-4$, while during the slow accretion stages the concentration increases with decreasing redshift as $c_{\\rm v}\\propto (1+z)^{-1}$.  We consider fits of three widely discussed analytic density profiles to the simulated clusters focusing on the most relaxed inner regions.  We find that there is no unique best fit analytic profile for all the systems. At the same time, if a cluster is best fit by a particular analytic profile at $z=0$, the same is usually true at earlier epochs out to $z\\sim 1-2$. The local logarithmic slope of the density profiles at $3 \\%$ of the virial radius ranges from $-1.2$ to $-2.0$, a remarkable diversity for the relatively narrow mass range of our cluster sample.  Interestingly, for all the studied clusters the logarithmic slope becomes shallower with decreasing radius without reaching an asymptotic value down to the smallest resolved scale ($\\lesssim 1\\%$ of the virial radius). We do not find a clear correlation of the inner slope with the formation redshift or the shape of the halo's MAH. We do find, however, that during the period of rapid mass growth the density profiles can be well described by a single power law $\\rho(r)\\propto r^{-\\gamma}$ with $\\gamma\\sim 1.5-2$.  The relatively shallow power law slopes result in low concentrations at these stages of evolution, as the scale radius where the density profiles reaches the slope of $-2$ is at large radii.  This indicates that the inner power law like density distribution of halos is built up during the periods of rapid mass accretion and active merging, while outer steeper profile is formed when the mass accretion slows down. To check the convergence and robustness of our conclusions, we resimulate one of our clusters using eight times more particles and twice better force resolution.  We find good agreement between the two simulations in all of the results discussed in our study. ",
        "introduction": "\\label{sec:intro}  During the last decade, there has been an increasingly growing interest in testing the predictions of variants of the cold dark matter models (CDM) on small scales. The interest was spurred by indications that the density distribution in the inner regions of dark matter halos predicted by CDM is at odds with the observed galactic rotation curves \\citep{flores_primack94,moore94}. This discrepancy is yet to be convincingly resolved and is still a subject of active debate \\citep[e.g.,][]{cote_etal00,vandenbosch_swaters01, blais_ouellette_etal01,deblok_etal01,deblok_etal03,swaters_etal03}. In addition, the CDM models face other apparent discrepancies with observations on galactic scales such as the amount of substructure in galactic halos \\citep{klypin_etal99,moore_etal99}, the incorrect normalization of the Tully-Fisher relation, the angular momentum of disk galaxies \\citep{navarro_steinmetz97,navarro_steinmetz00}, the ellipticity of dark matter halos \\citep{ibata_etal01}, and others.  In the past several years, the density distribution in the cores of galaxy clusters has also become a subject of a related debate. CDM models predict cuspy density profiles without flat cores \\citep{frenk_etal85,quinn_etal86,dubinski_carlberg91}. \\citet[NFW]{nfw_96,nfw_97} argued that the CDM halo profiles can be described by the following simple formula in all cosmologies and at all epochs, \\begin{equation} \\label{eq:nfw} \\rho(r)=\\frac{\\rho_{0}}{(r/r_{s})(1+r/r_{s})^{2}}. \\end{equation} This analytic formula describes the density profile of a halo using two parameters: a characteristic density, $\\rho_{0}$, and a scale radius, $r_{s}$. These parameters are determined by the halo virial mass, $M_{v}$, and concentration index, $c \\equiv r_{v} /r_{s}$, where $r_{v}$ is the virial radius of the halo. In addition, NFW argued that there is a tight correlation between $c$ and $M_{v}$, which means that the halo profiles of different mass objects form a one parameter family.  \\citet{moore_etal98} \\citep[see also][]{ghigna_etal00} carried out a convergence study of the dark matter profiles and concluded that high mass resolution is required to resolve the inner density distribution robustly. They advocated the analytic density profile of the form $\\rho(r) \\propto (r/r_{s})^{-1.5}[1+(r/r_{s})^{1.5}]^{-1}$, as a better description of the density distribution of their simulated halos.  This profile behaves similarly to the NFW profile at large radii ($\\propto r^{-3}$), but is steeper at small radii ($\\propto r^{-1.5}$). \\citet{fukushige_makino97,fukushige_makino01,fukushige_makino03} reached similar conclusions using a set of independent simulations. \\citet{jing_suto00} presented a systematic study of the density profiles of halos with masses in the range $2 \\times 10^{12}-5\\times 10^{14} \\ h^{-1} \\ M_{\\odot}$.  They found that the inner slope at a radius of 1\\% of the virial radius is shallower ($\\approx -1.1$) for cluster mass halos than for galactic halos ($\\approx -1.5$). Recently, \\citet{hayashi_etal03, navarro_etal03} found that often the logarithmic slope of the density distribution at the convergence radius is steeper than -1 as expected from the NFW profile, but significantly shallower than the -1.5  inner slope found by \\citet{moore_etal98}. Several other studies \\citep{kravtsov_etal97, kravtsov_etal98, avila_reese_etal99, jing00, bullock_etal01, klypin_etal01,fukushige_etal03} found a significant scatter in both the shape of the density profiles and halo concentrations, likely related to the details of the mass accretion histories of individual objects \\citep{wechsler_etal02,zhao_etal03a}. \\citet{mucket_hoeft03} and \\citet{hoeft_etal03} study the radial dependence of the gravitational potential and the velocity dispersion and come to the conclusion that there does not exist a slope asymptote of the density profile over a wide range but the slope increases with decreasing radius and reaches the value -0.58 as $r \\rightarrow 0$.  Observational constraints on the inner slope of the dark matter density distribution in galactic halos are difficult because the distribution is affected by the cooling and dynamics of baryons. The dark matter profiles in clusters, on the other hand, should be less affected by cooling as a much smaller fraction of cluster baryons is observed to be in the cold condensed phase.  Observational studies of the mass distribution in clusters using weak lensing and hydrostatic analysis of the X-ray emitting gas show that the overall mass distribution is in general agreement with CDM predictions \\citep{allen_98,clowe_etal00,willick_padmanabhan00, clowe_schneider01,sheldon_etal01,arabadjis_etal02, athreya_etal02, bautz_arabadjis03}.  Strong lensing studies can probe the mass distribution in the inner region of clusters and thus test the ``cuspiness'' of cluster halos. However, the results of strong lensing analyses have so far been contradictory, even in the case when the same system was studied. \\citet{tyson_etal98}, for example, argue that the density profile of cluster CL0024$+$1654 has a constant density core, while \\citet{broadhurst_etal00} find that the mass distribution in this cluster is cuspy. \\citet{czoske_etal02} argue that CL0024$+$1654 is undergoing a major merger and its density profile may not be representative.  \\citet{sand_etal02} find that the inner slope of the density profile in cluster MS$2137-23$ is flatter than expected in CDM models, a conclusion they recently confirmed for six more clusters \\citep{sand_etal03}. \\citet{gavazzi_etal03}, reanalyzing the same observations, argue that if the fifth demagnified image near the center of the lensing potential is not taken into account then the inner slope may be consistent with CDM predictions.   Given the disagreement among the different analytical fits proposed for the density profiles of dark matter halos found in simulations and a possible discrepancy with strong lensing observations, it is interesting to conduct a systematic study of the density profiles of clusters in the concordance $\\Lambda$CDM model. The study of cluster mass halos is also interesting because the typical concentrations of their matter distribution are lower than those of galactic halos. Thus, if an asymptotic inner slope, suggested by the analytic profiles, does exist it should be reached at a larger fraction of the virial radius in cluster halos and should be easier to detect.  The paper is organized as follows.  In the following two sections, we describe the numerical simulations and halo finding algorithm used in our analysis. In $\\S$~\\ref{sec:mah} we discuss the mass accretion histories of the analyzed clusters. In $\\S$~\\ref{sec:density_prof} we present the convergence test, discuss the fitting procedure, and our results on the shapes and inner slopes of the density profiles. We summarize our results and conclusions in $\\S$~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have studied the MAHs and density profiles of 14 cluster-size halos simulated using the ART code in a flat $\\Lambda$CDM cosmology. In agreement with previous studies, we find that most MAHs have a similar shape: an early, merger-dominated mass increase followed by a more gradual, accretion-dominated growth.  To obtain a formation redshift that characterizes the overall shape of the MAH we perform analytic functional fits. The typical MAHs are well described by the one parameter exponential function proposed by W02 (Eq.~[\\ref{eq:w02mah}]).  Two of the clusters in our sample experience intense merger activity and rapid mass growth until the present-day epoch. The MAHs of these systems are better described by a one parameter power-law function in the scale factor. We thus generalize the form proposed by W02 into a two parameter form (Eq.~[\\ref{eq:our_function}]) to encompass both exponential and power-law MAHs.  For each class, however, the fit reduces to a one parameter fit.  We check the convergence of halo density profiles using a re-simulation of one of the clusters with eight times more particles and better force resolution.  We show that both the halo profiles and their local logarithmic slopes converge at scales larger than about four times the formal resolution of the low resolution run, in agreement with a previous convergence study of the ART code by \\citet{klypin_etal01}. We fit the density distribution of the clusters with the NFW, M, and JS analytic profiles. Experiments show that the choice of merit function, weighting, and binning affect the absolute quality of a fit and may bias conclusions about how well a particular analytic profile fits simulation results.  We find, however, that the relative goodness of fit for the three analytic profiles and our conclusions about the best fit profile are robust to the changes in binning and merit function.  The main general result of our study is a remarkable diversity of the mass accretion histories, profile shapes, concentrations, and inner slopes for cluster-size halos in a relatively narrow mass range.  The concentrations of cluster-size halos at the present-day epoch exhibit a scatter of $\\sigma_{\\log{c_v}}\\approx 0.20$. This scatter is related to the diversity of halo MAHs and formation redshifts.  We find a statistically significant correlation between the formation redshift and the concentration of a halo, in agreement with results of W02. There is a more detailed connection between the MAH and concentration. The concentration of a halo is approximately constant at $c_{\\rm v}\\approx 3-4$ during the period of rapid mass growth and frequent major mergers ($z>z_{\\rm f}$) and increases with decreasing redshift when the mass accretion rates slows down at $z<z_{\\rm f}$. This behavior was recently pointed out by \\citet{zhao_etal03a,zhao_etal03b}. The implied ``floor'' in the concentration is not accounted for in the currently used models for $c_{\\rm v}(M)$ \\citep{bullock_etal01,eke_etal01}, which predict a monotonic decrease of concentration with increasing mass and may thus underestimate concentrations of the most massive, $\\gtrsim 5\\times 10^{14}h^{-1}\\ \\rm M_{\\odot}$, halos.  This may have important implications for estimates of expected number of wide-separation quasar lenses \\cite[e.g.,][]{kuhlen_etal03} and other results sensitive to the concentrations of very massive clusters.  The inner logarithmic slope of cluster profiles at $3\\%$ of the virial radius (or $10-50$~kpc) ranges from $-1.2$ to $-2$. In the best resolved clusters the logarithmic slope does not seem to reach a specific asymptotic value down to the smallest resolved scales in our simulations ($r/r_{180} \\simeq 0.007$). A similar conclusion was reached by \\citet{klypin_etal01} for galaxy-size halos and several recent studies \\citep{power_etal03,ascasibar_etal03,fukushige_etal03,hoeft_etal03,hayashi_etal03}. It is still not clear whether the density profiles in our simulations are consistent with density distribution of observed clusters. We note, however, that at the scales probed in observations the slope is not expected to be shallower than $-1$.  The asymptotic value of the slope has been a subject of much numerical effort in the last several years. Our results indicate that a universal asymptotic slope may not exist. We should note that the resolution of current dissipationless simulations is sufficiently high to converge on the density profile at scales smaller than the size of a typical central galaxy in clusters and groups ($\\sim 30-50$~kpc). Further improvement in profile modeling should therefore include realistic dynamics and cooling of the baryonic component as contraction of gas is expected to significantly modify dark matter distribution at these scales.  One of the most interesting results of our study is existence of systems with density profiles that can be well described by a power law $\\rho\\propto r^{-\\gamma}$ with $\\gamma$ ranging from $\\approx -1.5$ to $\\approx -2$. All of these systems are still in their rapid mass growth stage and experienced a recent major or minor merger. Remarkably, these halos maintain the power law density profiles at earlier epochs out to at least $z\\sim 1.5$.  The relatively shallow $\\gamma>-2$ power law slopes result in low concentrations as the scale radius where the density profiles reaches the slope of $-2$ is at large radii.  There are also indications that the profiles of all clusters are power law like during their rapid mass growth stages. We find, for example, that the power law provides an increasingly better fit with increasing redshifts for all of our clusters. At early epochs ($z\\gtrsim 1.5-2$), the power law fit is always either comparable or better than the NFW, M, and JS analytic profiles. We did not find any correlation of the power slope with the details of the cluster MAH. It would be interesting to look for such correlations using a larger sample of objects.  When the mass growth slows down at $z>z_{\\rm f}$, an outer steeper density profile is built up. As pointed out by \\citet{zhao_etal03a}, the difference in density profiles during the two mass accretion regimes may be due to more violent and thorough relaxation during the period of rapid mass growth. Although \\citet{zhao_etal03a} focused on halo concentrations and did not consider density profile shapes, their results are consistent with our conclusions. In particular, they find that during the rapid mass growth stage the circular velocity is nearly constant from the scale radius to the virial radius. This behavior is consistent with a power law density distribution with a slope close to $-2$.  In a recent study, \\citet{ascasibar_etal03} find that objects which experienced a recent merger event\\footnote{Note that \\citet{ascasibar_etal03} identify a recent merger by the presence of massive substructures within virial radii of their systems. This is different from our definition, which identifies major mergers directly from mass accretion tracks.} have lower concentrations and steeper inner profiles than more relaxed systems. This is consistent with our findings described above.  At the same time, \\citet{ascasibar_etal03} find that relaxed systems are better fit by the NFW, while systems with a recent major or minor merger by the M profile \\citep[see also][]{ascasibar_03}. They thus associate a particular shape of the profile with a recent merger history. In contrast, our results show the shape of density profiles is set early in the halo evolution and is usually stable over the past ten billion years.  Clusters with density profiles best described by the NFW rather than a JS at $z=0$, tend to have NFW-like profiles at earlier epochs as well.  The reverse is also true. We tested this conclusion using the high-resolution re-simulation of one of the clusters in our sample.  Also, we do not find any correlation between the redshift of last major merger (in our definition) or the formation redshift and the best fit analytic profile.  The origin of the distinctive density profile shape of the CDM halos remains poorly understood. Our results and results of other recent studies indicate that the shape is tightly linked to the halo mass accretion history. During the period of rapid mass accretion the violent relaxation is significant and results in a power-law like density distribution. This stage of evolution usually occurs early when the universe is dense and builds up the inner dense regions of halo. At this point, the logarithmic slope of the density distribution is shallower than $-2$ over a large fraction of the halo volume and its concentration is small. At later epochs, as the mass accretion rate slows down, the outer regions of the halo are built while its central regions remains nearly intact. This can be seen in Figure~\\ref{fig:bestfit_CL7and5} and Figures~10-13 of \\citet{fukushige_makino01}. This picture can explain why the best fit analytic profiles tend to be the same at various redshifts. The fits are sensitive to the density distribution at small and intermediate ($\\sim r_{\\rm s}$) radii which are set early. However, it is still unclear which process(es) determines a particular shape of the profiles. The key to understanding these processes appears to be in the details of early evolutionary stages of CDM halos, which will be the subject of a future study."
    },
    "0311/astro-ph0311254_arXiv.txt": {
        "abstract": " ",
        "introduction": "Galactic cirrus clouds are typically complexes of diffuse gas and dust most easily seen at high galactic latitudes. On large scales they are dominated by turbulent motions, however, on small scales dense cores have been found and there is evidence that these cores are gravitationally bound (Heithausen et al. 2002). In one core, even the spectral signature of infall motion was discovered (Heithausen 1999).\\\\ Core formation constitutes the very first step which leads to star-formation and is still rather poorly understood. The knowledge of how cores form and how they evolve is important for understanding the minimal physical conditions needed for stars to form. Studying the evolution of dense cores in cirrus clouds, which live in a rather quiescent environment and are mainly heated by the interstellar radiation field, may prove to be crucial to interpret observations of star-forming regions correctly.\\\\ One of the best studied cirrus cloud cores so far is the one in MCLD123.5+24.9, located in the Polaris Flare, a huge cirrus cloud complex in the direction of the north celestial pole. Its distance is between 130 and 240\\,pc (cf. Heithausen et al. 1993), and here we adopt a value of 150\\,pc. The temperature of the dust was determined to 13\\,K (Bernard et al. 1999) and the kinetic temperatures of the gas between 6 and 15\\,K  (e.g. Gro\\ss mann \\& Heithausen 1992).\\\\ Here we present new high-angular resolution HC$_3$N and CS observations of this core which allow a much closer insight into the fragmentation and core formation processes. ",
        "conclusions": "The most likely explanation for this appearance is that it is caused by strong abundance variations of the observed molecules due to time dependent chemical evolution. There are several chemical models in the literature which can account for the observed behavior. One that matches our derived properties of the core quite good is the one of Markwick et al. (2000), constructed for TMC-1, a low-mass star forming region. The initial conditions in this model are a temperature of 10\\,K, a hydrogen number density $\\rm{n}(\\rm{H}_2)=2\\times 10^{4}\\rm{cm}^{-3}$, a cosmic ray ionization rate of $1 \\times 10^{-17} s^{-1}$, and the UMIST Database for Astrochemistry 1995 reaction set (Millar et al. 1995). This model predicts a time delay of HC$_3$N compared to CS in reaching a measurable high abundance of $(3-10) \\times 10^5$\\,years (see Fig. 2). By this time the CS, which attained its highest abundance very fast after only $10^3$\\,years, may already start to deplete due to freezing out on dust grains.  \\begin{figure} \\centering {\\includegraphics[height=4cm]{fig_cbo2.eps}} \\caption{Abundances of CS and HC$_3$N relative to H$_2$ with time in the chemical model by Markwick et al. (2000). HC$_3$N needs about $3\\times 10^5$\\,years more than CS to reach a measurable high abundance. (from www.astrochemistry.net)} \\label{fig:2}       % \\end{figure}  This result is further supported by a time-dependent chemical model of core formation by Taylor et al. (1998), which basically comes to the same results with somewhat different initial conditions. Their main conclusion is, that CS can be seen in more or less the whole core, whereas the HC$_3$N on the other hand can only be observed in longer-lasting clumps. This behavior is essentially a result of the cold ($\\sim 10$\\,K) gas chemistry where also neutral-neutral reactions play a surprisingly important role."
    },
    "0311/astro-ph0311124_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{s:intro}  Most young stars are surrounded by accretion disks \\cite{beckwith:nature}, which are likely progenitors of planetary systems. The study of the physical structure, chemical composition, and evolution of these disks is therefore clearly a topic of great interest. The thermal emission of dust grains in disks has provided a wealth of information on the disks' properties. At near-infrared wavelengths, the thermal emission is dominated by a small amount ($\\sim 10^{-5}$~M$_\\odot$) of warm material close to the star, and the emission is often optically thick. These wavelengths serve as excellent probes of the presence or absence of a disk. At the longer (sub) millimeter wavelengths, the emission probes the cold material at larger radii ($>1$ AU), where the emission is usually mostly optically thin. These wavelengths are good probes of the total mass of a disk.  Dust, however, only makes up 1\\% of the total mass of the disk, if we adopt the standard interstellar gas-to-dust mass ratio. Ignoring the difficult to observe H$_2$ molecule (but see, e.g., \\cite{thi:h2_co_disks, richter:h2disk}), circumstellar disks are now commonly observed through lines of molecules such as CO, HCO$^+$, HCN, CN, {\\it etcetera\\/}. Although it is found that the molecular abundances are often strongly depleted through freezing out onto dust grains (\\S \\ref{s:models}) and the lines correspondingly weaker, molecular-line observations offer a powerful array of diagnostics to study disks: from the relative intensities of transitions, the molecular excitation can be derived and the density and temperature of the gas constrained; from the line strengths, the molecular abundances are found and the chemical composition inferred; and from the resolved line profiles, the velocity field in the disk is directly probed. This contribution discusses the methods with which this wealth of information can be extracted from molecular lines, and presents examples of recent research that illustrates this. ",
        "conclusions": ""
    },
    "0311/astro-ph0311312_arXiv.txt": {
        "abstract": "We present a flexible framework for constructing physical models of quasar evolution that can incorporate a wide variety of observational constraints, such as multi-wavelength luminosity functions, estimated masses and accretion rates of active black holes, space densities of quasar host galaxies, clustering measurements, and the mass function of black holes in the local universe.  The central actor in this formulation is the accretion rate distribution $p(\\mdot|M,z)$, the probability that a black hole of mass $M$ at redshift $z$ accretes at a rate $\\mdot$ in Eddington units.  Given a model of accretion physics that specifies the radiative efficiency and SED shape as a function of $\\mdot$, the quasar luminosity function (QLF) is determined by a convolution of $p(\\mdot|M,z)$ with the black hole mass function $n(M,z)$. In the absence of mergers, $p(\\mdot|M,z)$ also determines the full evolution of $n(M,z)$, given a ``boundary value'' of $n(M)$ at some redshift.  If $p(\\mdot|z)$ is independent of mass, then the asymptotic slopes of the QLF match the asymptotic slopes of $n(M)$, and $n(M)$ evolves in a self-similar fashion, retaining its shape while shifting to higher masses.  Matching the observed decline of the QLF ``break'' luminosity at $z<2$ requires either a shift in $p(\\mdot|z)$ that increases the relative probability of low accretion rates or an evolving mass dependence of $p(\\mdot|M,z)$ that preferentially shuts off accretion onto high mass black holes at low $z$. These two scenarios make different predictions for the masses and accretion rates of active black holes.  If the first mechanism dominates, then the QLF changes character between $z=2$ and $z=0$, shifting from a sequence of black hole mass towards a sequence of $L/\\Ledd$.  We use our framework to compare the predictions of five models that illustrate different assumptions about the quasar population: two dominated by unobscured thin-disk accretion with short and long quasar lifetimes, respectively, one with a 4:1 ratio of obscured to unobscured systems, one with substantial black hole merger activity at low redshift, and one with substantial low redshift growth in radiatively inefficient flows.  We discuss the observational advances that would be most valuable for distinguishing such models and for pinning down the physics that drives black hole and quasar evolution. ",
        "introduction": "The study of the quasar and AGN populations has been transformed in recent years by ambitious new optical (e.g., \\citealt{boyle00,schneider02,wolf03}), X-ray (e.g., \\citealt{brandt01,giacconi02,anderson03,ueda03}), and radio (e.g., \\citealt{white00}) surveys, by the recognition that low efficiency accretion modes may become important when the accretion rate itself is low (e.g., \\citealt{narayan98} and references therein), by detailed studies of low luminosity AGN in the local universe (e.g., \\citealt{ho01}), and, perhaps most of all, by the accumulating evidence that supermassive black holes are ubiquitous in the bulges of present-day galaxies (e.g., \\citealt{richstone98,merritt00,gebhardt00}). The dynamical studies of nearby galaxies strengthen the long-standing hypothesis that quasars are powered by black hole accretion \\citep[e.g.,][]{lyndenbell69,rees78}, and the ``demography'' of the local black hole population provides a powerful constraint on models of quasar evolution and its connection to galaxy evolution. These developments have inspired increasingly sophisticated theoretical models that place quasar evolution in the context of hierarchical clustering models for the formation of dark matter halos and galaxies \\citep[e.g.,][]{kauffmann00,cavaliere02,haimanloeb,wyithe03}.  This paper presents a physically motivated calculational framework that is intermediate in complexity between such {\\it ab initio} models of the quasar population and older descriptions in terms of ``luminosity evolution'' or ``density evolution.'' The central actor in our formulation is the accretion probability distribution $p(\\mdot|M,z)$, the probability that a black hole of mass $M$ at redshift $z$ is accreting mass at a rate $\\mdot$ in Eddington units (discussed in \\S\\ref{sec:framework} below). The key supporting players are the black hole mass function $n(M,z)$ and a physical model of accretion that predicts the radiative efficiency for a given $\\mdot$.\\footnote{Henceforth, we will usually drop the explicit dependence on $z$ and refer only to $p(\\mdot|M)$ or $n(M)$, but these should always be understood to refer to the distribution at some particular redshift.} An example of an accretion model would be thin-disk accretion with efficiency $\\epsilon \\equiv L/\\Mdot c^2 \\sim 0.1$ when $\\mdot \\sim 1$, changing to low efficiency advection-dominated (ADAF) accretion when $\\mdot$ is below some critical value. At a given redshift, $p(\\mdot|M)$ and $n(M)$ together determine the quasar luminosity function, and $p(\\mdot|M)$ also determines the accretion driven growth of the black hole population, and hence the evolution of $n(M)$.  Thus, given physical assumptions about radiative efficiencies and a ``boundary condition'' specifying $n(M)$ at one redshift, the history of $p(\\mdot|M)$ determines the complete evolution of the black hole population and the quasar luminosity function. An essential caveat is that mergers of black holes following mergers of their host galaxies could alter $n(M)$ independently of the accretion characterized by $p(\\mdot|M)$.  The simplest scenario connecting black hole and quasar evolution is that black holes ``shine as they grow'': a luminous quasar is powered by a black hole radiating at near-Eddington luminosity with efficiency $\\epsilon \\sim 0.1$, and no significant growth occurs in a non-luminous phase. In this case, the bolometric luminosity function at a given redshift is just $\\Phi(L) = \\fon n(L/l) l^{-1}$ , where $l$ is the (universal) ratio of Eddington luminosity to black hole mass and $\\fon$ is the fraction of black holes that are accreting. In a time interval $\\Delta t$, black holes on average increase their mass by a factor $\\exp(\\fon \\Delta t/t_g)$, where $t_g = M/(\\Ledd/\\epsilon c^2) = 4.5\\times 10^7(\\epsilon/0.1)\\;\\yr$ is the $e$-folding time for growth at the Eddington luminosity (\\citealt{salpeter64}; see discussion in \\S\\ref{sec:framework}). The mass density in black holes at the present day is simply related to the emissivity of the quasar population integrated over luminosity and redshift, $\\rho_{\\rm BH} = \\int_0^{t_0} U(t) dt / \\epsilon c^2$ (\\citealt{soltan82}, updated by, e.g., \\citealt{chokshi92,richstone98,yu02,fabian03}). In our language, this is a model in which all active black holes have the same radiative efficiency and $p(\\mdot|M)=\\fon\\delta_D(\\mdot-1)$, where $\\delta_D$ is the Dirac-delta function.  The evolution of quasars and black holes is determined by a boundary condition on $n(M)$ and the redshift history of the active fraction $\\fon(z)$. The ``quasar era'' $z\\sim 2-4$ when the emissivity of the population peaks is also the era in which today's black holes grew to their current mass.  This simple scenario may not be too far from the truth, but the possible complications raise a number of questions. Did today's black holes gain a significant fraction of their mass through low-$\\mdot$, low efficiency, ADAF-type accretion, thus growing at low luminosity?  Have black hole mergers substantially altered $n(M)$, leaving the integrated density $\\rho_{\\rm BH}$ fixed but changing the relative numbers of high and low mass black holes? Are some quasars accreting mass at super-Eddington rates, radiating at high luminosity but low efficiency in ``smothered,'' optically thick ADAF modes \\citep{katz77,begelman78,abramowicz88}? Do some quasars radiate substantially above the Eddington limit \\citep{begelman02}? Is a significant fraction of quasar activity obscured by gas and dust, as hypothesized in synthesis models of the X-ray background (e.g., \\citealt{setti89,comastri95,fiore99,fabian99,gsh01}), thus redistributing bolometric luminosity from the optical-UV-soft X-ray to the far-IR? Are low luminosity AGNs powered mainly by low mass black holes radiating at Eddington luminosity, by more massive black holes with thin-disk efficiencies but sub-Eddington accretion rates, or by still more massive black holes with sub-Eddington accretion rates {\\it and} low efficiency?  The methods developed here provide useful tools for addressing these questions, allowing us to construct concrete, quantitative models that answer them in different ways, then examine how observational data might distinguish among such models.  Our framework complements, but by no means replaces, the {\\it ab initio} approach that connects the evolution of quasars and black holes to that of the underlying dark halo and galaxy populations.  This approach has yielded many valuable insights, including the recognition that the rise of the quasar population probably traces the formation of the first dark halos large enough to host massive black holes, that the rapid decline of the population at low redshift probably reflects the combined impact of declining galaxy interaction rates and decreasing gas supplies in quasar hosts, and that the clustering of quasars with themselves or with galaxies can provide a valuable diagnostic of typical quasar lifetimes \\citep[e.g.,][]{efstathiou88,haehnelt93,hnr98,salucci99,kauffmann00, cavaliere00,haiman01,martini01,mhn01,wyithe03}. In combination with semi-analytic models of galaxy formation, these quasar evolution models can also predict the properties and environments of quasar hosts and the relation between the properties of present day galaxies and the masses of their central black holes. However, the models necessarily rely on specific assumptions about the mechanisms that trigger quasar activity and the accretion rates that these mechanisms produce.  To put things in our terms, the {\\it ab initio} models adopt a particular set of hypotheses about quasar activity in order to predict $p(\\mdot|M,z)$ from first principles.  Our framework is designed to model observational data in a flexible way with relatively few assumptions, while retaining the basic physical picture of black hole accretion that underlies nearly all modern interpretations of the quasar population.  One hope is that measurements of the luminosity function and the local black hole mass function will eventually allow us to integrate backwards in time and determine $p(\\mdot|M,z)$ empirically, drawing on a variety of observations to test the assumptions that enter such a reconstruction. We may find that the data are not powerful enough to tie down $p(\\mdot|M,z)$ without some {\\it a priori} constraints on its expected form, but that finding in itself would be a valuable, if disappointing, lesson. More generally, we hope to illuminate the connections between black hole evolution and quasar activity and learn what observations can and cannot tell us about these connections.  The \\cite{hnr98} paper has had the strongest impact on our thinking about these issues, but the most direct antecedent that we know of to the approach taken here is the lucid paper of \\cite{small92}.  They adopted a similar description of black hole evolution and its connection to quasar activity, and they applied this description to the observational data available at the time.  Advances in the observational data and the theoretical models of accretion make this an opportune time to revive and extend this approach.  \\cite{yu02} have recently used a similar method in assessing constraints on black hole accretion and mergers, though their assumptions and goals are more restrictive than ours --- in particular, they assume that quasars radiate at Eddington luminosity and thus that $p(\\mdot)$ consists of $\\delta$-functions at $\\mdot=1$ and $\\mdot=0$.  The difference in the form of $p(\\mdot)$ is one of the most significant differences between the models presented in this paper and most models of the quasar population in the literature.  These typically assume that $p(\\mdot)$ for $\\mdot>0$ is sharply peaked at some value close to Eddington, such as a $\\delta$-function (e.g., \\citealt{small92}) or a spike followed by an exponential decline (e.g., \\citealt{hnr98,kauffmann00}). A sharply peaked $p(\\mdot)$ could arise physically if the central black hole typically plays a large role in controlling its own fuel supply, through feedback or influence on stellar dynamics. However, we think it is more likely that fueling is driven by galactic scale events --- galaxy mergers, interactions, and bar formation, for example --- that are minimally influenced by the central black hole, and therefore do not ``know'' that they should feed it at any particular rate.  In particular, it seems reasonable that for every major event that leads to Eddington-like fueling of a central black hole there are many minor events that fuel it at a sub-Eddington rate.  The particular functional forms that we adopt here to represent this scenario, a power-law or broken power-law $p(\\mdot)$ between some $\\mdotmin$ and some $\\mdotmax$, are arbitrary, and chosen largely for mathematical convenience, but they reflect this general thinking about the process of quasar fueling. In the long run, one goal of our effort is to test observationally whether $p(\\mdot)$ is in fact a broad function or a peaked function, which would in turn have implications for the mechanisms of quasar fueling. At a qualitative level, the wide $L/\\Ledd$ distribution of AGN activity in the local universe (e.g, \\citealt{ho01}) seems to support the idea of a broad distribution of accretion rates.  In this paper we will keep our contact with observations relatively loose, focusing instead on presenting our framework in a clear way and illustrating how observations might distinguish among different scenarios for the quasar population. We will carry out a detailed analysis of multi-wavelength measurements of the quasar luminosity function and constraints on active black hole masses and the local black hole mass function a subsequent paper. We adopt a cosmological model with $\\Omega_m=1$ and $h \\equiv H_0/100\\;\\hubunits=0.5$ because all of the observational papers include results for this model, and not always for the low density, $\\Lambda$-dominated, $h\\sim 0.7$ model favored by recent cosmological observations.  Cosmology affects our models indirectly through its influence on observationally inferred luminosity functions and directly through the time-redshift relation.  We would not expect a change of cosmological model to have any qualitative impact on our results, and we expect that the quantitative impact could be compensated by modest changes to mass scales and accretion rates, especially since the age of the universe is similar for $\\Omega_m=1$, $h=0.5$, and for $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $h=0.7$.  We present the definitions and key equations of our framework in the following section.  We then present mathematical results for luminosity functions in \\S\\ref{sec:lumfunction} and for luminosity and black hole mass evolution in \\S\\ref{sec:evolution}, focusing on analytically solvable cases that illustrate general points. In \\S\\ref{sec:illus} we construct five illustrative models of the quasar population, each designed to match the observed evolution of the optical luminosity function but differing from one another in the typical quasar lifetime or in the importance of mergers, ADAF growth, or obscuration. We show how measurements of the black hole mass function, luminosity functions at other wavelengths, the masses and accretion rates of active black holes, and the space density of host galaxies might distinguish among these scenarios. We summarize our results in \\S\\ref{sec:summary}.  This is a long paper, and a reader who wants to get quickly to the main points can read \\S\\ref{sec:framework}, look through the figures and captions, and read \\S\\ref{sec:summary}.  The definitions of the models illustrated in Figures~\\ref{fig:bbandevol}--\\ref{fig:spacedens} are summarized in Table~\\ref{tbl:parameters} and its accompanying caption. ",
        "conclusions": "\\label{sec:summary}  In the framework developed here, the central actor in black hole and quasar evolution is the accretion rate distribution $p(\\mdot|M,z)$, the probability that a black hole of mass $M$ accretes at a rate $\\mdot$ (in Eddington units) at redshift $z$. Given a model for the accretion efficiency as a function of $\\mdot$, which can be inferred from observations and theoretical considerations that are largely independent of QLF evolution {\\it per se}, the combination of $p(\\mdot|M)$ and the black hole mass function $n(M)$ determines the bolometric luminosity function via equation~(\\ref{eqn:lumfun}).  Furthermore, in the absence of mergers, $p(\\mdot|M,z)$ determines the evolution of $n(M,z)$ given a ``boundary value'' of $n(M)$ at some redshift.  Mergers can complicate the picture by changing $n(M)$ independently of $p(\\mdot|M,z)$. We have generally made the plausible but not incontrovertible assumption that black holes accreting in the range $0.01 < \\mdot < 1$ have ``thin-disk'' efficiencies $\\epsd\\approx 1$ and that efficiencies decrease at higher (super-Eddington) and lower (ADAF) accretion rates (eq.~\\ref{eqn:effofmdot}). We have derived many of our results under the mathematically simplifying assumption that $p(\\mdot)$ is independent of mass.  While this assumption is unlikely to hold to high accuracy, it may be a reasonable approximation at redshifts near the peak of quasar activity, since black holes that grow faster than their peers tend to reduce their accretion rates in Eddington units and {\\it vice versa}. Most of our specific examples assume that $n(M)$ and $p(\\mdot)$ are double power-laws with breaks at $M_*$ and $\\mdotstar$, respectively, with $p(\\mdot)$ truncated at $\\mdotmin=10^{-4}$ and $\\mdotmax=10$.  \\subsection{Basic Results} \\label{sec:summary1}  Our framework yields a number of mathematical results that give insight into the relations among the black hole mass function, the accretion rate distribution, and the QLF. When $p(\\mdot)$ is independent of mass, the convolution integral~(\\ref{eqn:lumfun}) for $\\Phi(L)$ can be understood as follows: for each range $\\mdot \\rightarrow \\mdot+d\\mdot$, the mass function $n(M)$ is mapped to a luminosity $L=\\epsd \\mdot l M$, multiplied by $p(\\mdot)d\\mdot$, and added to a running total. If $n(M)$ is a double power-law and the range of $\\mdot$ is bounded, then the asymptotic slopes of $\\Phi(L)$ must equal the low and high mass slopes of $n(M)$, since high luminosity objects must come from black holes with $M>M_*$ and low luminosity objects must come from black holes with $M<M_*$. In the intermediate luminosity regime, where black holes above and below $M_*$ can both contribute, the QLF turns over in a way that depends on the slopes of $n(M)$ and the shape of $p(\\mdot)$. For our usual double power-law $p(\\mdot)$, the QLF break occurs at a luminosity $\\Lbrk \\sim \\mdotstar l M_*$. If $\\mdotstar$ is close to one, then the break luminosity corresponds roughly to the Eddington luminosity of $M_*$ objects, and the slope above the break corresponds to the high mass slope of $n(M)$.  However, if $\\mdotstar$ is low, then the turnover may be associated largely with the change in slope of $p(\\mdot)$, and the asymptotic regime where the high-$L$ slope of $\\Phi(L)$ matches the high-$M$ slope of $n(M)$ may only be reached beyond the observed range of luminosities. For most of the models that we have presented here, which are designed to match the \\cite{boyle00} optical luminosity function, the first case applies at high redshift and the second at low redshift.  A key feature of our model of accretion physics is that objects do not radiate at super-Eddington luminosities --- instead, we assume that the efficiency is $\\epsd = \\mdot^{-1}$ for $\\mdot > 1$, so that super-Eddington accretors radiate at $L=\\Ledd$. Although we have generally adopted $\\mdotstar = 0.5-1$ for fitting data at $z \\geq 2$, our results would not be very different if we took $\\mdotstar > 1$, or if we changed the maximum accretion rate or the high-$\\mdot$ slope, because the change in efficiency would cut off the luminosity distribution at $\\Ledd$ anyway. Our results would also not be very different if we assumed that black holes fed at a super-Eddington rate by their host galaxy regulate their accretion by outflows or convection \\citep{blandford99,quataert00} so that they gain mass at $\\mdot\\approx 1$ and radiate at $L \\approx \\Ledd$. If such flows drove gas out into the galaxy halo, they would effectively truncate $p(\\mdot)$ at $\\mdot=1$, while if they returned gas to a reservoir from which it would eventually be accreted, they would transform the $\\mdot>1$ tail of $p(\\mdot)$ into a spike at $\\mdot=1$.  With our standard efficiency assumptions, the latter scenario would increase the average radiative efficiency of high-$\\mdot$ accretion, by moving it from the super-Eddington regime to the thin-disk regime, thus yielding more luminosity for a given amount of black hole growth. Changing the accretion physics to allow $\\epsd \\approx 1$ with $\\mdot > 1$, and thus to allow substantially super-Eddington luminosities when the accretion rate is high \\citep{begelman02}, would have a more drastic effect on our results.  In this case, only a truncation of $p(\\mdot)$ would cut off the luminosity distribution at a given black hole mass, and the predicted luminosity function would therefore be sensitive to the values of $\\mdotstar$ and $\\mdotmax$ and to the shape of $p(\\mdot)$ at $\\mdot > 1$.  The fraction of black holes of mass $M$ that are active at a luminosity $L$ is proportional to $n(M) p\\left(\\mdot=\\frac{L}{\\epsd \\mdot l M}\\right)$, the underlying black hole space density times the probability of having the accretion rate required to shine at $L$. The same product appears in the integrand for the luminosity function itself (eq.~\\ref{eqn:lumfun}), but by constraining the integrand rather than the integral, measurements of active black hole masses can discriminate among models that produce similar $\\Phi(L)$ with different $p(\\mdot)$ and $n(M)$. In terms of our double power-law models, the active mass distribution is particularly useful for distinguishing cases with high $M_*$ and low $\\mdotstar$ from models with low $M_*$ and high $\\mdotstar$.  A general mass-dependent $p(\\mdot)$ can be written in the form $p(\\mdot|M) = p_0 (\\mdot) D(M | \\mdot)$, with $p_0(\\mdot) \\equiv p(\\mdot|M_0)$ for some fiducial mass $M_0$ and $D(M_0|\\mdot) \\equiv 1$. With such a mass dependence, we can still understand the convolution integral for $\\Phi(L)$ by considering the contribution from each range $\\mdot \\rightarrow \\mdot + d\\mdot$, but where the sum before involved only horizontally and vertically shifted versions of the mass function $n(M)$, now the mass function can be tilted or distorted by $D(M | \\mdot)$ before being added to the running total. Mass-dependence of $p(\\mdot)$ thus breaks the tight link between $n(M)$ and $\\Phi(L)$ and adds freedom to models of the QLF. Specifically, for a model with a given $n(M)$ and a mass-independent $p(\\mdot)$, there is a family of models with different $n(M)$ and mass-dependent $p(\\mdot)$ that yield identical predictions for $\\Phi(L)$ and the distribution of {\\it active} black hole masses.  However, this degeneracy applies only at a single redshift; the evolution of models with mass-dependent $p(\\mdot)$ is different because the shape of $n(M)$ changes with time. Furthermore, while the active black hole mass distributions are the same, the underlying black hole mass functions are different, and a measurement of the full $n(M)$ at $z=0$ may be sufficient to diagnose the mass-dependence of $p(\\mdot)$ at higher redshifts. If two models have similar underlying $n(M)$, then the masses of active black holes or space densities of their host galaxies are powerful diagnostics for mass dependence of $p(\\mdot)$, as illustrated in Figures~\\ref{fig:shortvsmassdep} and~\\ref{fig:spacedens}.  For evolutionary calculations, a crucial simplification (eq.~\\ref{eqn:acc}) is that the accretion driven growth of $n(M)$ depends only on the mean accretion rate $\\langle \\mdot(M,z) \\rangle \\equiv \\int_0^\\infty \\mdot p(\\mdot|M,z) d\\mdot$, not on the full form of the distribution function. Between times $t_1$ and $t_2$, a black hole with accretion rate $\\mdot(t)$ grows in mass by a factor $\\exp (\\tacc/t_s)$, where $\\tacc = \\int_{t_1}^{t_2} \\mdot dt$ is the accretion weighted lifetime (eq.~\\ref{eqn:tacc}) and $t_s$ is the Salpeter time (eq.~\\ref{eqn:tsdef}). In any particular mass bin, some of the black holes grow faster than average and some grow slower, but $n(M)$ evolves as if all of them grew at the average rate for that bin.  If $\\avmdot$ is independent of mass, then the growth factor is the same for all mass bins, and the mass function simply shifts in a self-similar fashion, preserving its shape (eq.~\\ref{eqn:accgrowth}).  One might expect this self-similar behavior to emerge, approximately, as a ``fixed point'' solution during the epoch of rapid black hole growth.  If $\\avmdot$ depends on mass, then the shape of $n(M)$ changes with time, becoming steeper if low mass black holes grow more rapidly and shallower if high mass black holes grow more rapidly.  Our results also imply a number of critical values for the logarithmic slopes of $p(\\mdot)$ and $n(M)$, where the character of the solutions changes.  We summarize these critical slopes in Appendix~\\ref{sec:appx2}. The most important result of this sort is that growth by accretion drives $n(M)$ at fixed mass up if the mass function is steeper than $M^{-1}$ and down if it is shallower, and that the corresponding critical slope for merger driven growth is $-2$ rather than $-1$.  Generically, one expects mergers to increase $n(M)$ at high masses and decrease it at low masses.  \\subsection{Implications of the Observed Luminosity Evolution} \\label{sec:summary2}  We have kept the comparison to observations in this paper rather loose, saving detailed tests of models against multi-wavelength luminosity functions and estimates of black hole mass distributions for future work. However, we are able to draw some interesting general conclusions from our efforts to reproduce the observed evolution of the optical luminosity function.  At $z\\leq 2$, we have focused on the \\cite{boyle00} evolution results, and particularly on their finding that the QLF has a break that shifts to lower luminosities at lower redshifts.  In any model where $p(\\mdot)$ is independent of mass, reproducing this result requires a shift towards lower typical accretion rates over time (Fig.~\\ref{fig:dblpowevol}). Decreasing the duty cycle uniformly by reducing the amplitude of $p(\\mdot)$ can lower the normalization of $\\Phi(L)$, but on its own it cannot shift the break to lower luminosities, since black hole masses, and thus the location of the break in $n(M)$, can only increase with time. Lowering the break luminosity requires decreasing the probability of high accretion rates {\\it relative} to low accretion rates.  In the context of our double power-law models, this change is achieved by reducing $\\mdotstar$, from slightly below unity at $z\\sim 2$ to far below unity at $z \\sim 0-0.5$. Physically, such a change could arise because of decreasing gas supplies and increasing dynamical times in galaxies at lower redshifts \\citep{kauffmann00}.  A consequence of this reduction in $\\mdotstar$ is a change in the nature of the QLF between high and low redshift. At high redshift, the sequence of quasar luminosities is primarily a sequence of black hole masses.  Because $\\mdotstar$ is close to unity and $p(\\mdot)$ is shallow below the break, most high luminosity quasars are produced by the more numerous low mass black holes accreting at $\\mdot \\approx 1$ rather than extremely rare high mass black holes with low accretion rates. The typical mass at a given luminosity is $M \\approx L/l$, and a factor of ten increase in luminosity roughly corresponds to a factor of ten increase in black hole mass. However, once $\\mdotstar$ falls well below unity, there is a large range $\\mdotstar < \\mdot < 1$ over which $p(\\mdot)$ is steep ($\\propto \\mdot^{-3}$ for most of our models), thus increasing the probability that a given luminosity is generated by a high mass black hole accreting at a low $\\mdot$ close to $\\mdotstar$. The typical active black hole mass still increases with luminosity, but at low redshift a factor of ten increase in $L$ corresponds to only a factor $\\sim 3$ increase in median black hole mass, and the range in black hole mass at fixed luminosity is much larger than at high redshift.  At low redshift, the sequence of quasar luminosities remains partly a sequence of black hole mass, but it is also in large part a sequence of $\\mdot$ (Figs.~\\ref{fig:massbh} and~\\ref{fig:mdot}).  This change in character arises only because we reduce the radiative efficiency in the super-Eddington regime, forcing systems with $\\mdot > 1$ to radiate at the Eddington luminosity. Without this transition in the accretion physics, there would be no preferred scale to change the relative importance of mass and $\\mdot$ between high and low redshift, since we assume the same double power law {\\it form} of $p(\\mdot)$ at all redshifts and it is only the value of $\\mdotstar$ relative to unity that changes. If the efficiency does not decrease in the super-Eddington regime, then the shape of the luminosity function depends in detail on the form and cutoff of $p(\\mdot)$ at $\\mdot > 1$. However, the limited data on masses of high redshift black holes generally shows systems with $L/\\Ledd \\approx 1$ but not substantially larger \\citep{mclure03,vestergaard03}. This result suggests that there is indeed a break in the radiative efficiency at $\\mdot \\approx 1$, or else that black hole accretion self-regulates to enforce $\\mdot \\la 1$, since otherwise it would require $p(\\mdot)$ to cut off coincidentally at $\\mdot > 1$ and be relatively high at $\\mdot$ just below one.  If $p(\\mdot)$ is mass-dependent, then there is a fundamentally different alternative for explaining the shift of the QLF break to lower luminosities: instead of a decline in characteristic accretion rates, the mass-dependence can itself evolve so that activity is preferentially suppressed in high mass black holes at low redshift (Fig.~\\ref{fig:massdepevol}). Physically, such behavior could arise because high mass black holes reside in early type galaxies with large bulges, which tend to exhaust their gas supplies earlier. In this scenario, quasars at a given position on the luminosity function (relative to $\\Lbrk$) are always associated with the same distribution of accretion rates, but the associated black hole mass declines with redshift as the high mass black holes turn off. In comparison with the mass-independent $p(\\mdot)$ models, the median black hole mass at fixed luminosity is lower, and the range of black hole masses is smaller, with near-Eddington accretors making a large contribution to the high end of the luminosity function at all redshifts (see Figure~\\ref{fig:shortvsmassdep}). The low black hole mass implies a high space density of hosts, so this model predicts that a larger fraction of low redshift quasars reside in late type or low luminosity galaxies.  Based on anecdotal evidence, it is hard to say whether a decline in characteristic accretion rates or a decline in the activity of high mass black holes is more important in producing the observed decline of $\\Lbrk$ at $z<2$.  The first scenario's prediction of a wider range of $L/\\Ledd$ at lower redshifts seems in qualitative agreement with studies of black hole masses \\citep{woo02,vestergaard03}. The second scenario seems in qualitative agreement with the relative quiescence of the black holes in most massive ellipticals (such as M87), though \\cite{kauffmann03} find that the most luminous AGN in the local universe do reside in massive, early type systems. The two scenarios make quantitatively different predictions, and careful comparison to studies of active black hole masses and host galaxy properties should be able to show whether one mechanism dominates or both are comparably important. The results of \\cite{vestergaard03} suggest that active black hole masses at low redshift are intermediate between the values predicted by the two models shown in Figure~\\ref{fig:shortvsmassdep}.  At high redshifts, constraints on the form of the QLF are weaker; in particular, the SDSS measurements of \\cite{fan01} probe only the high luminosity end of $\\Phi(L)$.  For matching the observed evolution over the range $z\\sim 5$ to $z\\sim 2$, we find one acceptable solution in which $p(\\mdot)$ is roughly constant in Eddington units and the growth of the black holes themselves drives the growth in amplitude of the QLF (Fig.~\\ref{fig:growth}). With our adopted parameters, the black holes grow by a factor $\\sim 10$ between $z=5$ and $z=2$, and the space density at $z=2$ is $n_* M_* = 2.012 \\times 10^{-5} \\dunits$ at $M_*=10^9 M_\\odot$. Significantly lower normalizations of $n(M)$ at $z=2$ are not allowed because they would require a negative black hole mass density at $z=5$. Higher normalizations are allowed, in which case the quasar duty cycle is shorter, the rate of black hole growth is smaller, and the growth of the QLF is driven by a steady increase in $p(\\mdot)$. While a solution with a high black hole space density can give an acceptable match to the QLF, it seems physically unattractive because it requires that most of the black hole mass density was already in place at $z=5$, before the main epoch of quasar activity.  In this latter scenario, the luminous phases of quasars represent the addition of a small amount of mass to already formed black holes.  If we assume instead that the observed optical QLF {\\it does} trace the growth of black holes, and that the former model is therefore more realistic, then our analysis predicts a black hole space density $Mn(M) \\sim 2-3\\times 10^{-5}\\dunits$ at $M=10^9 M_\\odot$ at $z=2$, and continuation to $z=0$ implies a similar space density at $M \\sim 2\\times 10^9 M_\\odot$.  However, we have not investigated the sensitivity of this prediction to our specific choices of parameters, such as the double power-law form and adopted slopes of $p(\\mdot)$ and $n(M)$.  Black hole mergers and obscured accretion could also alter the prediction significantly, especially at low redshift.  \\subsection{Distinguishing Scenarios} \\label{sec:summary3}  Many of the qualitative results mentioned above could have been anticipated without detailed calculations.  Our framework, however, allows one to compute quantitative predictions of concrete models that illustrate distinct ideas about the nature of black hole and quasar evolution.  We did this in \\S\\ref{sec:illus} for the low redshift $(z \\leq 2)$ regime, adopting as our baseline a model with quasar emission and black hole growth dominated by unobscured thin-disk accretion and a normalization of $n(M,z=2)$ implying a short quasar lifetime, and, consequently, little growth of black hole masses from $z=2$ to $z=0$.  We compared this model to four variants: one with a lower $n(M,z=2)$ and correspondingly longer quasar lifetime, one with a 4:1 ratio of obscured to unobscured systems, one with a large amount of merger driven growth of black hole masses, and one with a boosted probability of low-$\\mdot$ accretion leading to substantial ADAF growth of black holes.  For each scenario, we are able to find parameters that acceptably reproduce the \\cite{boyle00} optical QLF at $z=2$, 1, and 0.5.  Relative to the \\stq\\ model, the \\ltq\\ model starts at $z=2$ with a factor $\\sim 6$ lower $n(M)$ at every mass.  Because of the larger amount of accretion per black hole, however, the \\ltq\\ $n(M)$ overtakes the \\stq\\ $n(M)$ at high masses by $z=0$, while remaining below it at low masses.  The two models have similar QLFs at every wavelength, since they match in the optical by construction and are dominated by systems with the \\cite{elvis94} SED. However, the growth of $M_*$ in the \\ltq\\ model requires a more rapid decline of $\\mdotstar$ to compensate, so at lower redshifts it predicts lower median $\\mdot$ and higher median black hole mass at a given luminosity. The two models can thus be distinguished observationally by the $z=0$ black hole mass function, by the mass distributions of active black holes at $z\\sim 0.5-1$, and by the space densities of host systems, which are lower in the \\ltq\\ model at every luminosity and redshift.  In the obscured model, the large amount of obscured accretion produces more black hole growth than in the \\stq\\ model, leading to a higher $n(M)$ and $\\rhobh$ at low redshift.  By $z=0$, $n(M)$ is higher by a factor $\\sim 6$ at high masses.  However, the median black hole mass and host space density at a given optical luminosity are only slightly higher.  The clearest distinguishing feature of this model is the relative amplitude of luminosity functions in different wavelength bands, a consequence of the different SED shapes of obscured and unobscured accretors.  At $z=2$, the 2-10 keV luminosity function is elevated by nearly a factor of five, the ratio of all systems to unobscured systems.  This boost decreases towards low redshift because of the increasing importance of obscuration in the (observed-frame) 2-10 keV band.  The 0.5-2 keV band is heavily obscured at $z<2$, so the soft X-ray luminosity function of the obscured model is similar to that of the \\stq\\ and \\ltq\\ models, except at $z=2$ where it has a higher amplitude at low luminosity.  The strongest departure of all is in the FIR, where the re-radiated emission of obscured accretors boosts $\\Phi(L)$ by factors of ten (low luminosity) to one hundred (high luminosity), relative to the \\stq\\ and \\ltq\\ models. This crucial prediction should soon be testable by {\\it SIRTF} and by other sub-mm and mm-wavelength observations.  In the merger model, low redshift mergers strongly distort the initial black hole mass function, depleting it at low masses and boosting it at high masses, with a factor of 16 increase in the high mass end at $z=0$ relative to the \\stq\\ model. The quasar population of this model is still dominated by systems with a thin-disk SED, and since it matches the observed optical luminosity function by construction, its predictions at other wavelengths are close to those of the \\stq\\ and \\ltq\\ models. However, the high space density of high mass black holes leads to a high median mass of active black holes at fixed luminosity, similar to that of \\ltq\\ at low luminosities and higher still at high luminosities. Merger driven distortions of the mass function also lead to distinctive redshift and luminosity dependence of the host space density.  Finding parameters that yield significant ADAF growth and an acceptable match to the optical luminosity function proves quite difficult, requiring an artificial boost to the probability of accretion rates below $\\mdotcrit$. With our adopted parameters, the ADAF model predicts a large amount of black hole growth between $z=2$ and $z=0$, and thus a high $n(M)$ and $\\rhobh$ at $z=0$.  With the combination of high $M_*$ and high probability of low $\\mdot$, the ADAF model predicts the largest median black hole masses and lowest median accretion rates at fixed luminosity, with the median accretion rate approaching $\\mdotcrit$ even for high luminosity AGN at $z \\leq 0.5$.  The high X-ray fraction of the ADAF SED boosts the soft and hard X-ray luminosity functions relative to the optical, especially at low redshift, and the model predicts that a majority of X-ray selected systems at $z\\sim 0.5$ should be ADAF accretors, even at high luminosities.  This prediction appears observationally untenable, and the model requires rather implausible parameter choices in the first place, so our results suggest that ADAFs are unlikely to make an important contribution to black hole growth in the real universe, even at $z<2$ (\\citeauthor{hnr98} [\\citeyear{hnr98}] reach a similar conclusion). Low radiative efficiency at low $\\mdot$ may nonetheless help explain the remarkable quiescence of most black holes in the local universe \\citep{narayan98}.  \\subsection{Prospects}  Our results illustrate how a variety of observational constraints can be brought to bear on the key questions of quasar and black hole evolution. In particular, we have extended the ideas of \\cite{soltan82} and \\cite{small92} to show that incorporating the link between luminous accretion and black hole growth allows one to construct concrete physical models that are simultaneously constrained by multi-wavelength luminosity function measurements and estimates of black hole masses and accretion rates. For our models in \\S\\ref{sec:illus}, we chose a plausible but not unique set of initial conditions at $z=2$ and evolved them forward in time under varying assumptions, always matching the observed optical QLF. The models then make distinguishable predictions for other observables.  With our current parameter choices, all of our models face some difficulty when confronted with recent estimates of X-ray luminosity functions and the local black hole mass function, as illustrated in Figures~\\ref{fig:comparenofm}, \\ref{fig:comparesxdata}, and \\ref{fig:comparehxdata}. Models that fit the \\cite{ueda03} hard X-ray QLF generally do not fit the \\cite{miyaji01} soft X-ray QLF at the same redshift, and vice versa.  A model incorporating luminosity and redshift dependence of the obscured quasar fraction might fare better, though some of the problem may still lie with the observational estimates themselves.  Combining Tremaine et al.'s (\\citeyear{tremaine02}) estimate of the $M-\\sigma$ relation with Sheth et al.'s (\\citeyear{sheth03}) estimate of the distribution of galaxy velocity dispersions yields a mass function that lies well below our model predictions for $M > 10^9 M_\\odot$, unless the intrinsic scatter of the $M-\\sigma$ relation is $\\sim 0.5$ dex, compared to Tremaine et al.'s estimate of $\\leq 0.3$ dex.  Repairing this discrepancy would require either reducing the break mass at $z=2$ substantially below $M_*=10^9 M_\\odot$ or dropping our assumed double power-law form of $n(M)$ and adopting a mass-dependent $p(\\mdot)$ to reproduce the \\cite{boyle00} QLF. For the present, we do not want to draw strong conclusions from these discrepancies, since we have not thoroughly assessed the observational uncertainties, and we have not investigated the extent to which a failing model can be ``fixed up'' by adjusting its parameters (e.g., the initial shape of the black hole mass function), while retaining its essential features (e.g., a large fraction of obscured systems).  The $z=0$ black hole mass function can be reasonably well estimated from current data, at least at masses $M \\leq 10^9 M_\\odot$ where the form and scatter of the $M-\\sigma$ relation are well constrained, and it is a fundamental boundary constraint on any evolution model.  For a comprehensive attempt to match observations, therefore, it probably makes sense to impose this constraint {\\it a priori} on all models, and integrate the evolutionary equations {\\it backward} in time.  In our framework, this approach is just as easy as integrating forward, even if it is less intuitive.  In effect, one takes the known black hole mass function today, infers the accretion rate distribution by matching the luminosity function, steps backward by removing the implied amount of mass from each bin of the mass function, and repeats. We will apply this approach to available observations in future work. Given the inevitable uncertainties in the observational data, radiative efficiencies, and bolometric corrections, there are likely to be some degeneracies in the solutions, but we can hope that models that differ in fundamental rather than incidental features will remain observationally distinguishable.  Based on the results found here, we suspect that the primary source of uncertainty will be the mass dependence of $p(\\mdot|M)$, which requires accurate measurements of both the QLF and the masses of active black holes to pin down empirically. Mergers look like the other most difficult problem, though in this case there are good theoretical ideas about what the merger rates of dark halos and galaxies should be (e.g., \\citealt{taylor01}), and these can be incorporated into model calculations.  The traditional picture of quasars as a population of supermassive black holes growing by accretion seems more secure than ever. Many open questions remain about the roles of black hole mass, accretion rate, radiative efficiency, and SED shape in determining quasar luminosities, about the properties of accretion flows at low and high accretion rates, about the importance of black hole mergers and obscured accretion as drivers of black hole growth, and about the relations among populations observed at different wavelengths. Our work highlights a number of areas where observational advances will be crucial to answering these questions. These include improved determination of the local black hole mass function, better understanding of the dependence of radiative efficiency and SED shape on accretion rate, measurements of the luminosity function at different wavelengths over the widest achievable range in luminosity and redshift, estimates of masses and accretion rates of active black holes as a function of redshift and luminosity, and indirect estimates of black hole space densities from host galaxy and quasar clustering studies. Fortunately, the observational situation is advancing rapidly, and many of these areas have seen substantial progress in the last few months alone, as discussed in \\S\\ref{sec:illus}. It is worth emphasizing the value of luminosity function determinations and black hole mass estimates that traverse the break in the QLF and extend as far below as possible. Accurate characterization of this regime is crucial for separating the roles of $n(M)$ and $p(\\mdot)$ in shaping the luminosity function, which in turn is necessary for understanding the contribution of sub-Eddington accretion rates to black hole mass evolution. These lower luminosities are also where optical and X-ray evolution appear to be radically different, and better measurements of the joint X-ray, optical, and IR luminosity functions are needed to pin down the origin of these differences. The emerging data on black hole and quasar evolution are complex, complementary, and rich.  We hope that the physical modeling approach described in this paper will prove useful in exploiting their power."
    },
    "0311/astro-ph0311197_arXiv.txt": {
        "abstract": "{\\citet{Mcs} discovered intense UV (1100 \\AA) emission from the direction of the Coalsack Nebula. We have used their results in conjunction with  a Monte Carlo model for the scattering in the region to show that the scattering is from dust in the foreground of the Coalsack molecular cloud. We have constrained the albedo of the grains to 0.4 $\\pm$ 0.1. This is the first determination of the albedo of dust in the diffuse ISM in the FUV. ",
        "introduction": "Although a potentially useful test of models of interstellar dust grains, measurements of the optical parameters --- the albedo ($\\it{a}$) and the phase function asymmetry factor ($\\it{g}$) --- of the grains have been too uncertain to have been of much utility \\citep[for a recent review see][]{Dr03}. Both methods used to investigate the scattering properties have been problematic: reflection nebulae because an uncertain geometry can heavily influence the derived parameters \\citep{Ma02} ; and the diffuse background because of its faintness and because there is often a tradeoff between $a$ and $g$ which allows neither to be tightly constrained \\citep{Dr03}.  The Coalsack Nebula, one of the most prominent dark nebulae in the Southern Milky Way, offers an excellent location for the determination of the scattering function of the diffuse grains, particularly in the UV where \\citet{Mcs} found it to be one of the brightest sources of diffuse emission in the sky. Without detailed modeling, they were unable to provide useful contraints on the optical constants of the grains but did suggest that most of the observed emission was due to forward scattering of photons from three of the brightest UV stars in the sky (Table \\ref{star_inf}) by foreground dust, rather than back-scattering from dust in the molecular cloud.  We note that \\citet{Mat70} observed scattering from the Coalsack in the visible which he interpreted as being due to scattering from the Coalsack.  In this paper, we have reinterpreted the $\\it{Voyager}$ observations of \\citeauthor{Mcs} using improved distances for the three stars, a detailed model for the interstellar dust distribution and a Monte Carlo model for the grain scattering. In agreement with \\citeauthor{Mcs}, the observed radiation is  dominated by scattering from dust in the foreground cloud, rather than the Coalsack molecular cloud. The albedo is tightly constrained to a value of 0.4 $\\pm$ 0.1 at 1100 \\AA. \\begin{table} \\caption[]{Properties of stars in our model} \\label{star_inf} \\begin{tabular}{r c c c c} \\hline \\hline Name&$\\it{l}$&$\\it{b}$&distance&flux(1100 \\AA )\\\\ & (deg) & (deg) & (pc) & (photons s$^{-1}$ \\AA$^{-1}$ ) \\\\ \\hline $\\alpha$ Cru &300.13&-0.36&98.3&8.86$\\times$10$^{8}$\\\\ $\\beta$ Cru &302.46&3.18&108.1&9.70$\\times$10$^{8}$\\\\ $\\beta$ Cen &311.77&1.25&161.3&2.53$\\times$10$^{9}$\\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": ""
    },
    "0311/astro-ph0311474_arXiv.txt": {
        "abstract": "We construct the first ``face-on'' Tully-Fisher (TF) relation for 24 galaxies with inclinations between 16$^\\circ$ and 41$^\\circ$. The enabling measurements are integral-field, echelle spectroscopy from the WIYN 3.5m telescope, which yield accurate kinematic estimates of disk inclination to $\\sim15^\\circ$. Kinematic inclinations are of sufficient accuracy that our measured TF scatter of 0.42 mag is comparable to other surveys even without internal-absorption corrections. Three of four galaxies with significant kinematic and photometric asymmetries also have the largest deviations from our TF relation, suggesting that asymmetries make an important contribution to TF scatter. By measuring inclinations below 40$^\\circ$, we establish a direct path to linking this scatter to the unprojected structure of disks and making non-degenerate dynamical mass-decompositions of spiral galaxies. ",
        "introduction": "One common characteristic of Tully-Fisher (TF; Tully \\& Fisher 1977) studies of the relation between luminosity and rotation speed is the selection of spiral galaxies with inclinations ($i$) above $45^\\circ$, and typically with $i\\sim60^\\circ$.  The inability to determine accurate and precise inclinations for galaxies with $i< 40^\\circ$ has limited studies attempting to link structural parameters, such as lopsidedness and ellipticity, to TF scatter (Franx \\& de Zeeuw 1992, Zaritsky \\& Rix 1997). Likewise, combining vertical stellar velocity dispersions and disk rotation speeds has been thwarted for lack of precise inclinations for nearly face-on galaxies (e.g., Bottema 1997). Hence, non-degenerate dynamical mass-decompositions of galaxies have not yet been made.  One solution to the above problems is to use nearly face-on galaxies to study the TF relation. Here the detailed planar structure of the disk can be studied unprojected, while the vertical component of the disk stellar velocity dispersion is favorably projected. The problem is now observational: How can accurate and precise inclinations be measured for small inclination angles?  For most TF studies, inclinations are calculated from photometric axis-ratios, assuming circular disks, and correcting for a constant, intrinsic flattening (e.g. Tully \\& Fouqu\\'e 1985; Courteau 1996). Bars, spiral structure, and intrinsic ellipticity cause systematic errors in photometric inclination and position angles (PA), dominate the apparent ellipticity at low inclination, and skew apparent inclinations to larger values (Andersen \\etal 2001; Courteau \\etal 2003). Hence there is a triple penalty at low inclinations for TF studies: (1) Erroneously large inclinations lead to under-correcting the velocity projection. (2) Random errors diverge since photometric disk axis ratio measurements have constant errors as a function of axis ratio. (3) For long-slit studies, PA-mismatch will always yield systematic underestimates of the projected rotation speed.  These pitfalls can be mitigated if kinematic inclinations are derived from two-dimensional velocity fields (e.g., Schommer \\etal 1993; Verheijen 2001) because disk scale heights are irrelevant, and velocity fields between 1 and 3 scale lengths are typically less aberrated than the light distribution by spiral structure, bars or warps (Courteau \\etal 2003; Briggs 1990). A limit of $\\sim40^\\circ$ had been set on the minimum derived kinematic inclinations for nearby galaxies from HI maps (Begeman 1989), but this minimum value is not fundamental; it is notionally adopted as such.  We have discovered that kinematic inclinations, derived from high-quality, optical kinematic data collected with integral field units, are sufficiently accurate and precise to construct a TF relation at inclinations as low as 15$^\\circ$. In this letter we present an analysis of the first TF relation for a sample of nearly face-on galaxies. ",
        "conclusions": "We have presented a TF relationship for 24 galaxies with a mean inclination of 26$^\\circ$ -- the first time the TF relation has been measured for a sample of galaxies with inclinations less than $40^\\circ$. Accurate and precise kinematic inclinations are the key to studying TF relation at low inclination, and velocity fields are the essential ingredient for their measurement.  We are able to construct this FOTF relation from integral-field H$\\alpha$ velocity fields for four primary reasons: (1) DensePak and SparsePak are efficient for gathering H$\\alpha$ velocity fields with moderately high spatial resolution and medium spectral resolution at high signal to noise. (2) H$\\alpha$ velocity fields generally extend out to $\\sim3$ disk scale-lengths so that warping, which dominates the outer radii of some galaxy disks, is unimportant. (3) Instead of using multiple tilted rings, we successfully fit galaxy disks with a single velocity-field model, simultaneously using data at all radii and phase angles. (4) We equally-weight all data, making our fits sensitive to the signal constraining inclination, which peaks in the velocity field at $\\sim45^\\circ$ from the major axis.  The FOTF relation measured from our data has a zeropoint and scatter comparable to conventional TF relations but possibly a shallower slope. The bulk of the FOTF scatter is consistent with measurement errors, as found by Verheijen (2001). Of the four galaxies which show large kinematic and photometric asymmetries, three fall well off the TF relation. While our statistics are small, this is consistent with the claim by Barton \\etal (2000) that kinematic asymmetries are an important source of TF scatter.  The FOTF can play an important role in understanding the nature and interpretation of the TF slope, zeropoint and scatter in three ways. First, internal extinction is less important for face-on galaxies minimizing inclination-dependent effects.  For inclined TF samples, if correlated errors exist between de-reddened luminosities and de-projected rotation-speeds, TF scatter will be increased. If extinction depends on luminosity or velocity (mass), the corrected TF slopes may differ from their true values.  Second, measurements of asymmetry, ellipticity, and lopsidedness ideally are made at low inclinations. Several studies have suggested non-axisymmetric dark halo mass distributions such as lopsidedness (Zaritsky \\& Rix 1997; Swaters \\etal 1999) or ellipticity (Franx \\& de Zeeuw 1992; Andersen \\etal 2001) could be sources of TF scatter. However, without access to inclination measurements for face-on systems, these claims could not be directly evaluated. Our small FOTF sample suggests such links are possible, and show that such studies can now be pursued wholesale.  Finally, the mass-to-light ratio of spiral disks is an essential ingredient for interpreting the TF zeropoint, slope, and scatter. Only limits can be placed on mass decompositions (Bell \\& de Jong 2001) without direct dynamical estimates of total {\\it and} disk mass. This can be accomplished via rotation curves {\\it and} vertical stellar velocity dispersions of disks. The latter are best measured in face-on systems, and hence the FOTF unlocks the door for determining the fundamental mass-budget of spiral galaxies."
    },
    "0311/astro-ph0311018_arXiv.txt": {
        "abstract": "Circuit models involving bulk currents and inductances are often used to estimate the energies of coronal magnetic field configurations, in particular configurations associated with solar flares. The accuracy of circuit models is tested by comparing calculated energies of linear force-free fields with specified boundary conditions with corresponding circuit estimates. The circuit models are found to provide reasonable (order of magnitude) estimates for the energies of the non-potential components of the fields, and to reproduce observed functional dependences of the energies. However, substantial departure from the circuit estimates is observed for large values of the force-free parameter, and this is attributed to the influence of the non-potential component of the field on the path taken by the current. ",
        "introduction": "\\noindent Solar flares are understood to involve the conversion of magnetic energy in the solar corona into the energy of particles and radiation. In particular flare energy is believed to derive from non-potential coronal magnetic fields, i.e.\\ that part of the magnetic field produced by currents flowing in the solar corona (e.g.\\ Tandberg-Hanssen and Emslie, 1988).  The origin of flare energy in coronal currents has suggested to many authors the use of a circuit analogy (e.g.\\ Alfv\\'{e}n and Carlqvist 1967; Spicer, 1982; Somov 1992; Melrose, 1995; Melrose 1997; Longcope and Noonan 2000; Khodachenko and Zaitsev, 1998; Khodachenko, Haerendel and Rucker, 2003). It is well known that the energy of a system of $N$ distinct current-carrying circuits may be written (e.g.\\ Jackson, 1998) \\begin{equation}\\label{eq:general_E} E=\\frac{1}{2}\\sum_{i=1}^{N}L_iI_i^2 +\\sum_{i=1}^{N}\\sum_{j>i}M_{ij}I_iI_j, \\end{equation} where \\begin{equation}\\label{eq:general_L} L_i=\\frac{\\mu_0}{4\\pi I_i^2}\\int_{C_i}d^3x_i\\int_{C_i}d^3x_i^{\\prime} \\frac{{\\bf J}({\\bf x}_i)\\cdot {\\bf J}({\\bf x}_i^{\\prime})} {|{\\bf x}_i-{\\bf x}_i^{\\prime}|} \\end{equation} is the self-inductance and \\begin{equation}\\label{eq:general_M} M_{ij}=\\frac{\\mu_0}{4\\pi I_iI_j}\\int_{C_i}d^3x_i\\int_{C_j}d^3x_j^{\\prime} \\frac{{\\bf J}({\\bf x}_i)\\cdot {\\bf J}({\\bf x}_j^{\\prime})} {|{\\bf x}_i-{\\bf x}_j^{\\prime}|} \\end{equation} is the mutual inductance, and where ${\\bf J}({\\bf x})$ is the current density at a point ${\\bf x}$ in space, and $C_i$ is the volume of the $i^{\\rm th}$ circuit. Although Equations~(\\ref{eq:general_E})---(\\ref{eq:general_M}) are quite general, their utility relies on having a small number of well-defined circuits with simple geometries for which analytic expressions for $L_i$ and $M_{ij}$ may be evaluated, and in particular for which $L_i$ and $M_{ij}$ are independent of $I_i$ and $I_j$. This is the situation for currents flowing in a small number of ``wire'' loops, i.e.\\ loops with a fixed geometry.  Circuit models for coronal magnetic fields adopt wire loop configurations. For example, in the flare model of Melrose (1997) based on reconnection between current-carrying magnetic flux tubes, each flux tube is represented by fixed toroidal structures carrying a uniform axial current. The self-inductance of such a structure is (e.g.\\ Landau and Lifshitz, 1960) \\begin{equation}\\label{eq:L} L_i=\\mu_0 a_i C_i, \\quad \\quad C_i= \\ln \\frac{8a_i}{r_i}-\\frac{7}{4}, \\end{equation} where $a_i$ is the major radius of the toroid and $r_i$ is the minor radius. Melrose (1997) also adopts the following approximate expression for the mutual inductance between two loops whose centres are separated by a distance $d_{ij}$ and whose axes subtend an angle $\\theta_{ij}$: \\begin{equation}\\label{eq:M_don} M_{ij}=\\mu_0 \\frac{8a_i^2a_j^2 (C_iC_j)^{1/2}\\cos\\theta_{ij}} {\\left[(a_i+a_j)^2+d_{ij}^2\\right]^{3/2}}. \\end{equation} This formula is not exact but is based on interpolation between known results. An exact formula which holds for two axially-aligned line-current loops separated by a distance $d_{ij}$ is (Stratton, 1941) \\begin{equation}\\label{eq:stratton} M_{ij}=\\mu_0(a_ia_j)^{1/2}\\left[\\left(\\frac{2}{k_{ij}}-k_{ij}\\right) K(k_{ij}) -\\frac{2}{k_{ij}}E(k_{ij}) \\right], \\end{equation} where \\begin{equation} \\quad k_{ij}^2=\\frac{4a_ia_j}{(a_i+a_j)^2+d_{ij}^2}, \\end{equation} and where $K(x)$ and $E(x)$ are complete elliptic integrals.  Real coronal currents do not flow along simple wire loops --- they are distributed in the coronal volume, and describe complicated paths. A common model for current-carrying coronal fields is the force-free model, in which the magnetic field ${\\bf B}$ satisfies ${\\bf J}\\times {\\bf B}=0$, (as well as $\\nabla\\cdot {\\bf B}=0$) where ${\\bf J}=\\frac{1}{\\mu_0}\\nabla\\times {\\bf B}$. In other words the current density is everywhere parallel to the magnetic field. Unfortunately the force-free field equation is in general nonlinear and hence difficult to solve. A restricted set of force-free fields --- linear force free fields --- are easy to calculate (e.g.\\ Nakagawa and Raadu 1972; Chiu and Hilton 1977; Alissandrakis, 1981; Gary, 1989). For these fields $\\nabla\\times {\\bf B}=\\alpha {\\bf B}$ where $\\alpha$ is a constant, the force-free parameter. Linear force-free models for coronal magnetic fields permit detailed descriptions of current paths in the corona, although they have limitations as large scale models for solar fields (e.g.\\ Sturrock, 1994).  In this paper the accuracies of simple circuit models for coronal magnetic field energies are tested by application to linear force-free field configurations. To our knowledge this is the first quantitative test of the accuracy of the circuit model in the solar context. For simplicity, the comparison is limited to field configurations representing one and two magnetic loops. The order of presentation is as follows. In \\S\\,2 the method is outlined, including the adopted boundary conditions (\\S\\,2.1), and the details of the linear force-free (\\S\\,2.2) and circuit (\\S\\,2.3) models. In \\S\\,3 the results of the calculations for one loop (\\S\\,3.1) and two loops (\\S\\,3.2) are presented. Finally in \\S\\,4 the results are discussed. ",
        "conclusions": "\\noindent In this paper calculated energies of non-potential components of linear force-free fields are compared with simple circuit model estimates, for boundary conditions suggestive of one and two magnetic loops. The motivation is to test the circuit models, which are often invoked to estimate coronal magnetic field energies. In particular circuit models have been used to describe flare-related energy changes. In the Melrose (1997) model, preferred magnetic configurations for flaring are identified on the basis of geometrical dependences of circuit energy estimates.  In \\S\\,3.1 results for a single loop configuration are presented. An important result is the $\\overline{\\alpha}$-dependence of the energy of the non-potential component of the field. For a range of small values of the force-free parameter $\\overline{\\alpha}$ the free energy is found to be proportional to $\\overline{\\alpha}^2$, or equivalently to be proportional to the square of the current in the loop, consistent with the circuit model. This result is explained mathematically in terms of the $\\overline{\\alpha}^2$ dependence of the free energy appearing in Equation~(\\ref{eq:binomial}) for small $\\overline{\\alpha}$. For larger values of current (or $\\overline{\\alpha}$) the free energy increases more rapidly than the circuit prediction. This result is also expected for the two-loop configuration, since Equation~(\\ref{eq:binomial}) applies for any boundary conditions. A simple physical interpretation of this result is as follows. For small values of $\\overline{\\alpha}$ the non-potential component of the field is small compared to the potential component, and the current essentially follows the fieldlines of the potential field. Hence as $\\overline{\\alpha}$ is varied the geometry of the current does not change and the inductances [which are geometrical quantities --- see Equations~(\\ref{eq:general_L}) and~(\\ref{eq:general_M})] are constant. This situation corresponds to the free energy scaling with the square of the current (or $\\overline{\\alpha}^2$). For larger values of $\\overline{\\alpha}$ the non-potential component of the field becomes comparable to the potential component, and influences the path of the current. In this case the geometry [and hence the inductances as defined by Equations~(\\ref{eq:general_L}) and~(\\ref{eq:general_M})] depend on $\\overline{\\alpha}$, which corresponds to the free energy scaling with a higher power of $\\overline{\\alpha}$. Although demonstrated here for a class of linear force-free fields, this effect is expected to be quite general and to limit the accuracy of simple circuit models in application to any force-free field for large values of current.  In \\S\\,3.1 the dependence of the free energy of the field on the geometry of the loop is also investigated (for a modest value of $\\overline{\\alpha}$). The functional dependences on minor radius and major radius are found to be qualitatively consistent with the circuit model.  Section \\S\\,3.2 presents the results for the two-loop configuration. The interaction energy of the two loops is obtained by calculating the free energy of the loops together and then subtracting off the individual free energies of each loop. This procedure is used to calculate an effective mutual inductance $\\overline{M}_{12}^{\\rm eff}$ for the field. This was then compared with simple circuit model formulae, namely an exact expression for the mutual inductance of two axially-aligned line-current loops due to Stratton (1941), and an approximate expression for the mutual inductance of oblique loops with finite cross section due to Melrose (1997). The Stratton formula is found to qualitatively reproduce the observed variation of $\\overline{M}_{12}^{\\rm eff}$ with the major radius of one loop. The Melrose formula is compared with $\\overline{M}_{12}^{\\rm eff}$ as one loop is rotated with respect to the other. The effective mutual inductance of the field reproduces the cosine dependence in the Melrose formula, and in particular the inferred mutual inductance is close to zero for the case of perpendicular loops. However, the calculated values of $\\overline{M}_{12}^{\\rm eff}$ are almost an order of magnitude larger than the Melrose formula values. Comparison with the Stratton formula for the case of parallel loops confirms that, for the chosen parameters, the Melrose formula underestimates the actual interaction energy. The Melrose expression is an interpolation between known results, and so the discrepancy can be attributed to the approximate nature of the formula.  In this paper linear force-free fields have been used because they are straightforward to calculate. It is likely that real coronal magnetic fields involve currents that are spatially highly concentrated, and hence require a nonlinear force-free model (or a description involving non-zero pressure forces). In this case circuit models may provide quite accurate energy estimates because the configuration closely resembles a set of isolated circuits. However, it is also possible that in the nonlinear case the current path sensitively depends on the value of the current, in particular for large currents, and this may limit the accuracy of simple circuit models, as discussed above.  Another point to note is that circuit models are really only useful if there are a small number of circuits which are easily identified. In this paper we have considered widely separated magnetic structures, which as a result have simple connectivity. In general coronal fields have a complicated topology, and it may be difficult to identify coronal current paths.  Circuit models appear to provide reasonably accurate free energy estimates for coronal magnetic field configurations for a range of values of the current, and also qualitatively describe various functional dependences of the free energy (the dependence on loop major radii, the dependence on the angle between two loops, etc.). Simple circuit models are less accurate for very large values of the current, when the current path depends on the value of the current. Despite this reservation, the results of this paper confirm the general utility of circuit models, in particular in their application to the flare phenomenon."
    },
    "0311/gr-qc0311082_arXiv.txt": {
        "abstract": "{This article is meant as a summary and introduction to the ideas of effective field theory as applied to gravitational systems, ideas which provide the theoretical foundations for the modern use of General Relativity as a theory from which precise predictions are possible.} ",
        "introduction": "\\label{section:introduction} Quantum mechanics and General Relativity were discovered within a decade of one another close to a century ago, and almost immediately the search for a quantum theory of gravity had begun. Ever since it has been a puzzle as to what theoretical framework might ultimately reconcile these two theories with one another. This reconciliation has proven to be difficult to achieve, and although several promising proposals now exist none can yet claim complete theoretical and experimental vindication.  \\subsection{Against the Split Brain} The long-standing nature of this difficulty has driven some physicists to a state of intellectual despair, wherein they conclude that a crisis exists in physics which might be called the crisis of the \\emph{split brain}. On one hand, quantum mechanics (and its offspring quantum field theory) provides an incredibly successful description of all known non-gravitational phenomena, with agreement between predictions and experiment sometimes taking place at the part-per-billion level \\cite{QED,EWdata}. On the other hand, classical general relativity is also extremely successful, with its predictions being well tested within the solar system and for some binary pulsar systems \\cite{GRdata}. (The cosmological evidence for dark matter and dark energy is sometimes proposed as indicating the failure of gravity over long distances \\cite{MOND}, but at present the evidence for new gravitational physics at large distances does not seem compelling \\cite{DMvsNP}.) The perceived crisis is the absence of an over-arching theoretical framework within which both successes can be accommodated. Our brains are effectively split into two incommunicative hemispheres, with quantum physics living in one and classical general relativity in the other.  The absence of such a framework would indeed be a crisis for theoretical physics, since real theoretical predictions are necessarily approximate. Controllable results always require some understanding of the size of the contributions being neglected in any given calculation. If quantum effects in General Relativity cannot be quantified, this must undermine our satisfaction with the experimental success of its classical predictions.  It is the purpose of this article to present the modern point of view on these issues, which has emerged since the early 1980's. According to this point of view there is no such crisis, because the problems of quantizing gravity within the experimentally accessible situations are similar to those which arise in a host of other non-gravitational applications throughout physics. As such, the size of quantum corrections can be safely estimated and are extremely small. The theoretical framework which allows this quantification is the formalism of \\emph{effective field theories}, whose explanation makes up the better part of this article. In so doing we shall see that although there can be little doubt of the final outcome, the explicit determination of the size of sub-leading quantum effects in gravity has in many cases come only relatively recently, and a complete quantitative analysis of the size of quantum corrections remains a work in progress.  \\subsection{Identifying Where the Problems Lie} This is not to say that there are no challenging problems remaining in reconciling quantum mechanics with gravity. On the contrary, many of the most interesting issues remain to be solved, including the identification of what the right observables should be, and understanding how space and time might emerge from more microscopic considerations. For the rest of the discussion it is useful to separate these deep, unsolved issues of principle from the more prosaic, technical problem of General Relativity's \\emph{non-renormalizability}.  There have been a number of heroic attempts to quantize gravity along the lines of other field theories \\cite{EarlyQG}, \\cite{CanonQG}, \\cite{LivRev-QG} and it was recognized early on that General Relativity is not renormalizable. It is this technical problem of non-renormalizability which in practice has been the obstruction to performing quantum calculations with General Relativity. As usually stated, the difficulty with non-renormalizable theories is that they are not predictive, since the obtention of well-defined predictions potentially requires an infinite number of divergent renormalizations.  It is not the main point of the present review to recap the techniques used when quantizing the gravitational field, nor to describe in detail its renormalizability. Rather, this review is intended to describe the modern picture of what renormalization means, and why non-renormalizable theories need not preclude making meaningful predictions. This point of view is now well-established in many areas --- such as particle, nuclear and condensed-matter physics --- where non-renormalizable theories arise. In these other areas of physics predictions can be made with non-renormalizable theories (including quantum corrections) and the resulting predictions are well-verified experimentally. The key to making these predictions is to recognize that they must be made within the context of a low-energy expansion, in powers of $E/M$ (energy divided by some heavy scale intrinsic to the problem). Within the validity of this expansion theoretical predictions are under complete control.  The lesson for quantum gravity is clear: non-renormalizability is not in itself an obstruction to performing predictive quantum calculations provided the low-energy nature of these predictions in powers of $E/M$, for some $M$, is borne in mind. What plays the role of the heavy scale $M$ in the case of quantum gravity? It is tempting to identify this scale with the Planck mass, $M_p^{-2} = 8 \\pi G$ (with $G$ denoting Newton's constant), and in some circumstances this is the right choice. But as we shall see $M$ need not be $M_p$, and for some applications might instead be the electron mass, $m_e$, or some other scale. One of the points of quantifying the size of quantum corrections is to identify more precisely what the important scales are for a given quantum-gravity application.  Once it is understood how to use non-renormalizable theories, the size of quantum effects can be quantified and it becomes clear where the real problems of quantum gravity are pressing and where they are not. In particular, the low-energy expansion proves to be an extremely good approximation for all of the present experimental tests of gravity, making quantum corrections negligible for these tests. By contrast, the low-energy nature of quantum-gravity predictions implies that quantum effects are important where gravitational fields become very strong, such as inside black holes or near cosmological singularities. This is what makes the study of these situations so interesting: it is through their study that progress on the more fundamental issues of quantum gravity is likely to come.  \\subsection{A Road Map} The remainder of this article is organized in the following way.  \\begin{itemize} \\item {\\it Section 2: Effective Field Theories} does not involve gravity at all, but instead first describes why effective field theories are useful in other branches of physics. The discussion is kept concrete by considering a simple toy model, for which it is argued how some applications make it useful to keep track of how small ratios of energy scales appear in physical observables. In particular, considerable simplification can be achieved if an expansion in small energy ratios is performed as early as possible in the calculation of low-energy observables. The theoretical tool for achieving this simplification is the effective lagrangian, and its definition and use is briefly summarized using the toy model as an explicit example.  \\item {\\it Section 3: Quantum Gravity as an Effective Theory} describes how the tools of the previous section may be applied to calculating quantum effects including the gravitational field. In particular, it is shown how to make predictions despite General Relativity's non-renormalizability, since effective lagrangians are generically not renormalizable. As we shall see, however, some of the main results one would like to have regarding the size of quantum corrections to arbitrary loop orders remain incomplete.  \\item {\\it Section 4: Explicit Applications} of these ideas are described in this section, which use the above results to compute quantum corrections to several gravitational results for two kinds of sources. These calculations compute the leading quantum corrections to Newton's Law between two slowly-moving point particles, and to the gravitational force between two cosmic strings (both in 3+1 spacetime dimensions).  \\item {\\it Section 5: Conclusions} are briefly summarized in the final section.  \\end{itemize} ",
        "conclusions": "\\label{section:suggestions}  The goal of this review has been to summarize the modern picture of quantum gravity, within which the perturbative non-renormalizability of General Relativity is recognized as being a particular instance of a more general phenomena: the widespread application of non-renormalizable quantum field theories throughout many branches of physics. Regarding quantum gravity in this way shows how quantitative predictions can be made: one must simply apply the rules of effective field theories, which are known to give an accurate description of experiments in low-energy nuclear, particle, atomic and condensed matter physics.  Thinking of General Relativity as an effective theory in this way is not a new development, and underlies most approaches to quantum gravity either explicitly or implicitly. Neither is it new to calculate explicitly the behaviour of quantum fields in curved space (sometimes including the graviton). What \\emph{is} new (over the last few years) is proceeding beyond the qualitative statement that General Relativity is an effective theory to obtain the quantitative next-to-leading predictions within a controlled semiclassical approximation. Although much of the mechanics of such calculations leans on experience obtained when calculating with quantum fields in curved space, the crucial new difference is the quantitative power-counting arguments which identify precisely which quantum effects contribute to any given order in small quantities.  What emerges from this summary is a snapshot of a work which is very much still in progress. The following loom large among the missing results. \\begin{itemize} \\item Although a general statement of the power-counting result for very light, relativistic particles near flat space has been known for some time \\cite{DonoghuePC}, \\cite{Ode}, the central general power-counting results are not yet demonstrated to all orders for the most interesting case for practical calculations: the gravitational interactions of very massive, non-relativistic sources which are weakly interacting gravitationally ({\\it i.e.} in spacetimes which are perturbatively close to Minkowski space). \\item Although the theoretical tools exist, in the form of the operator product expansion, similarly general power-counting arguments are not yet given for quantum fluctuations about more general curved spaces. \\item Some explicit calculations which are cast within the power-counting framework have been done, but many more can surely be done. \\end{itemize}  There can be little doubt that quantum effects are extremely small in the classical systems for which gravitational measurements are possible (like the solar system), but this need not undermine the motivation for their computation. The point of such calculations is not their relevance for practical experiments (we wish!). Rather, their point is conceptual. It is only through the careful calculation of quantum effects that the theory of their size can be solidly established. In particular, any precise comparison between observations and the predictions of classical gravity is ultimately incomplete unless the quantitative size of the quantum corrections is explicitly established, as a systematic, all-orders power-counting argument would do.  Furthermore, we can always hope to get lucky, even if only theoretically. A clean understanding of how the size of quantum corrections depends on the variables (mass, size, separation, {\\it etc.})in a given system, one might hope to find larger-than-generic quantum phenomena in special systems. Even if these lie beyond the reach of present-day experimenters, they may furnish instructive theoretical laboratories within which differing approaches to quantum gravity might be more starkly compared.  In the last analysis, I hope the reader has become convinced of the utility of effective field theory techniques, and that the effective field theory point of view lifts the experimental triumphs of classical General Relativity to precision tests of the leading-order implications of the Quantum Theory of Gravity."
    },
    "0311/astro-ph0311526_arXiv.txt": {
        "abstract": "There are point-like sources in central regions of several supernova remnants which have not been detected outside the X-ray range. The X-ray spectra of these Central Compact Objects (CCOs) have thermal components with blackbody temperatures of 0.2--0.5 keV and characteristic sizes of 0.3--3 km. Most likely, the CCOs are neutron stars born in supernova explosions.  We overview their observational properties, emphasizing the \\cxo\\ data, and compare them with magnetars. ",
        "introduction": "X-ray observations have shown a new population of radio-quiet compact objects, presumably isolated neutron stars (INSs), apparently different from both isolated rotation-powered pulsars and accretion-powered X-ray pulsars in close binary systems. About 10 of these objects (AXPs and SGRs), which show pulsations with periods in a 5--12 s range, large period derivatives, $\\dot{P}\\sim 10^{-12}$--$10^{-10}$, and/or strong bursts are believed to be {\\em magnetars} (Thompson \\& Duncan 1996), i.e., INSs with superstrong magnetic fields, $B\\sim 10^{14}$--$10^{15}$ G. The magnetars in quiescence have typical X-ray luminosities $10^{34}$--$10^{35}$ erg s$^{-1}$ and show two-component spectra, a thermal component with a blackbody (BB) temperature of 0.4--0.7 keV and a non-thermal component described by a power-law (PL) with a photon index around 3. At least two magnetars have associated supernova remnants (SNRs). The second class of radio-quiet INSs includes colder (apparently older) nearby objects, showing soft thermal spectra with temperatures in a 50--150 eV range and luminosities $\\sim 10^{32}$--$10^{33}$ erg s$^{-1}$. At least some of these objects, sometimes dubbed {\\em Dim Isolated Neutron Stars}, are likely ordinary rotation-powered pulsars whose radio beams are not observable from Earth. Third is a class of point-like X-ray sources found near the centers of SNRs that cannot be identified as active radio pulsars or magnetars.  These {\\em Central Compact Objects} (CCOs), presumably INSs, have thermal spectra with BB temperatures  0.2--0.5 keV and X-ray luminosities $L_{\\rm x} \\sim 10^{33}$--$10^{34}$ erg s$^{-1}$.  We define a CCO as an X-ray point source which (1) is found near the center of a SNR, (2) shows no radio/$\\gamma$-ray counterpart, (3) shows no pulsar wind nebula (PWN), (4) has a soft thermal-like spectrum. Current observational status of CCOs is discussed below.  \\vspace{-0.1in} ",
        "conclusions": ""
    },
    "0311/astro-ph0311460_arXiv.txt": {
        "abstract": "Over the last three years, the Galactic center region has been monitored with the {\\it Chandra X-Ray Observatory}. Besides the X-ray emission from the target object, Sgr A*, diffuse X-ray emission was detected throughout most of the 17$'$ $\\times$ 17$'$ field of view. With 11 {\\it Chandra} observations through 2002 June, the total effective exposure reaches $\\sim$590 ks, providing significant photon statistics on much of the detailed structure within the faint, often filamentary, diffuse X-ray emission. The {\\it true-color} X-ray image and the {\\it equivalent width} images for the detected elemental species of the Galactic center region demonstrate that the diffuse X-ray features have a broad range of spatio-spectral properties. Enhancements of the low-ionization-state Fe line emission ($E$ $\\sim$ 6.4 keV) to the northeast of Sgr A* can be interpreted as fluorescence within the dense interstellar medium resulting from irradiation by hard, external X-ray sources. They may also be explained by emission induced by the bombardments by high energy particles on the ambient medium, such as may accompany unresolved supernova ejecta intruding into dense interstellar medium. The detection of molecular cloud counterparts to the 6.4 keV Fe line emitting features indicates that these Fe line features are associated with dense Galactic center clouds and/or active star-forming regions, which strongly supports the proposed origins of the X-ray reflection and/or supernova ejecta for the Fe line emission. We detect highly ionized S and Si lines which are generally coincident with the low-ionization-state Fe line emission and the dense molecular clouds in the northeast of Sgr A*. These hot plasmas are then most likely produced by massive star-forming activities and/or supernova remnants in the Galactic center. In contrast, we find that highly ionized He-like Fe line emission ($E$ $\\sim$ 6.7 keV) is primarily distributed along the plane instead of being concentrated in the northeast of Sgr A*. The implied high temperature and the relatively uniform, but, strong alignment along the plane are consistent with the magnetic confinement model, suggesting this hot gas component has been re-heated by the strong interstellar magnetic fields in the Galactic center to produce the observed He-like Fe line emission. ",
        "introduction": "OBSERVATIONS \\& DATA REDUCTION}  For the last three years, the Galactic supermassive black hole candidate Sgr A* has been monitored, as part of Guaranteed Time Observation (GTO) and Guest Observer (GO) programs, with the Advanced CCD Imaging Spectrometer (ACIS) \\citep{garmire03} on board the {\\it Chandra X-Ray Observatory} (Table \\ref{tbl:tab1}). As of 2002 June, combining 11 {\\it Chandra}/ACIS observations (except for ObsID 1561a, which was severely contaminated by a bright transient source within the FOV), the total exposure has reached $\\sim$590 ks, which is the deepest ever observation of the Galactic center region in X-rays. We first applied the algorithm developed by Townsley et al. (2002a) to correct the spatial and spectral degradation of the ACIS data caused by radiation damage, known as charge transfer inefficiency (CTI; Townsley et al. [2000]). The expected effects of the CTI correction include an increase in the number of detected events and improved event energies and energy resolution \\citep{townsley00,townsley02a}. We then screened the data by status, grade, and the flight timeline filter. We have also removed observation time intervals of strong flaring in the background. All individual event files were then reprojected to the tangent plane at the radio position of Sgr A* (RA[J2000] = 17$^h$ 45$^m$ 40$^s$.0409, Dec[J2000] = $-$29$^{\\circ}$ 00$'$ 28$\\farcs$118) in order to generate the composite data. The detailed descriptions of these data reduction processes and the resulting broadband raw image from the composite data may be found in Muno et al. (2003a). ",
        "conclusions": "SUMMARY AND CONCLUSIONS}  Deep {\\it Chandra}/ACIS observations of the Galactic center reveal a complex spatial and spectral structure of diffuse X-ray emission with significant photon statistics.  The thermal (e.g., hot Si and S) and nonthermal (e.g., neutral Fe) X-ray line emission is generally enhanced in a region lying to the northeast of Sgr A*, a region in which massive star-formation has been active over the past 5 million years (e.g., the Arches cluster), and in which dense Galactic center molecular clouds are currently present.  The highly ionized line emission from Si and S in these regions is therefore most likely associated with the hot ISM heated by the shocks from SNe and/or massive stellar winds, while shocks interacting with dense, cold ISM may produce emission in the neutral Fe line.  Curiously, we detect little or no enhancement of the emission lines to the south and southwest of Sgr A*.  This weak line emission to the south is apparently the result of an absence of activity there, in contrast to the northeast regions, and it is owed in part to the substantial absorption by the foreground 20 km s$^{-1}$ cloud.  In contrast, the highly ionized He-like Fe line emission shows a more uniform distribution along the plane, with no apparent enhancements to the northeast of Sgr A*. This feature suggests that the young star-forming and/or SNe activities have less effect on the He-like Fe distribution than they do on the neutral Fe line emission. The hot Fe line emission might rather be caused by the hot ISM trapped and then heated by the strong Galactic center magnetic fields.  The spectral characteristics of the strong neutral Fe line emitting features to the northeast of Sgr A* are plausibly described by the XRN model. We identify dense molecular counterparts for these Fe knots, which is strong support for the hypothesized XRN origin.  If past activity of Sgr A* is the primary external illuminating source, the required luminosities are $L_X$ $\\sim$ 10$^{35}$$-$10$^{36}$ ergs s$^{-1}$ with durations of at least a few years. Alternatively, nonthermally-accelerated electrons by SN shocks penetrating back into dense, cold supernova ejecta may also explain the observed neutral Fe line features. The presence of active star-forming regions and the dense molecular clouds near the neutral Fe features are supportive of this scenario. Unambiguous discriminations between these models are however difficult with the current data.  Follow-up observations with the current and upcoming observatories may be helpful for a detailed understanding of the origins for these features."
    },
    "0311/astro-ph0311183_arXiv.txt": {
        "abstract": "We report high-resolution spectroscopic observations of the young star HD~34700, which confirm it to be a double-lined spectroscopic binary. We derive an accurate orbital solution with a period of $23.4877 \\pm 0.0013$~days and an eccentricity of $e = 0.2501 \\pm 0.0068$. The stars are found to be of similar mass ($M_2/M_1 = 0.987 \\pm 0.014$) and luminosity. We derive also the effective temperatures (5900~K and 5800~K) and projected rotational velocities (28~\\kms\\ and 22~\\kms) of the components. These values of $v \\sin i$ are much higher than expected for main-sequence stars of similar spectral type (G0), and are not due to tidal synchronization. We discuss also the indicators of youth available for the object. Although there is considerable evidence that the system is young ---strong infrared excess, X-ray emission, \\ion{Li}{1}~$\\lambda$6708 absorption (0.17~\\AA\\ equivalent width), H$\\alpha$ emission ($\\sim$0.6~\\AA), rapid rotation--- the precise age cannot yet be established because the distance is unknown. ",
        "introduction": "\\label{sec:intro}  Over the past decade numerous surveys of young stars in star-forming regions have found evidence that the overall frequency of binary and multiple systems is consistently higher than in the general field \\citep[e.g.,][]{Ghez:93,Leinert:93,Ghez:97a,Ghez:97b,Kohler:98,Simon:99a}, although the enhancement may vary somewhat from region to region \\citep[see, e.g.,][]{Simon:99b}. In general those surveys have indicated roughly twice as many binaries as in the classical study by \\cite{Duquennoy:91}, which focussed on solar-type stars in the solar neighborhood that are typically much older. The results for the younger binaries refer mostly to relatively wide systems ($> 5$~AU), spatially resolved by high-resolution imaging techniques such as infrared speckle interferometry or adaptive optics \\citep[see, e.g.,][]{Takami:03}.  Studies of binaries with smaller separations (spectroscopic binaries) in star-forming regions have not been as systematic, but may indicate a similar trend \\citep[see, e.g.,][]{Mathieu:94,Corporon:99,Zinnecker:01,Melo:03}.  Partial lists of known spectroscopic binaries among pre--main sequence (PMS) stars have been presented by \\cite{Melo:01}, \\cite{Prato:02}, and others, and new binaries are slowly but continuously being added. More than two dozen of the systems with known orbits are double-lined spectroscopic binaries, which provide the most information on the individual masses of the components (compared to single-lined spectroscopic binaries).  The young star HD~34700 has recently been reported to be a double-lined spectroscopic binary by \\cite{Arellano:03}, who also provided evidence of velocity variations. HD~34700 was originally identified as a young object from its strong infrared excess as observed by IRAS \\citep{Odenwald:86,Walker:88,Oudmaijer:92}, a characteristic usually interpreted in terms of a circumstellar disk in other similar ``Vega-excess\" stars. It is located in Orion, although it is not clear whether it is associated with any of the star-forming complexes in the general vicinity.  The optical, infrared, and millimeter-wave properties of HD~34700 have been modeled extensively by \\cite{Sylvester:96a} and \\cite{Sylvester:96b}.  They inferred that the disk has inner and outer radii of roughly 25--50~AU and 550~AU, respectively, depending on the model, and that the circumstellar material is composed mostly of relatively small dust grains (size $\\leq 10~\\mu$m).  Sub-millimeter observations have led to the detection of $^{12}$CO and $^{13}$CO emission, from which a total mass of $\\sim$1~M$_{\\earth}$ in dust particles has been derived \\citep{Zuckerman:95,Coulson:98}.  Other estimates have varied between 0.2~M$_{\\earth}$ and 3.7~M$_{\\earth}$ \\citep{Sylvester:96b,Walker:00,Sylvester:01}, depending also on the assumed distance.  Molecular hydrogen in the disk has been searched for, but not seen \\citep{Bary:03}. Low-level optical linear polarization has been detected by \\cite{Bhatt:00} and \\cite{Oudmaijer:01}, consistent with a non-spherically symmetric distribution of dust (a disk).  The infrared luminosity of the star is a considerable fraction of its bolometric luminosity \\citep[$\\sim$8--35\\%;][]{Zuckerman:95,Sylvester:96a,Coulson:98,Bhatt:00,Sylvester:01}. As in other young objects, the \\ion{Li}{1}~$\\lambda$6708 absorption line is quite prominent in HD~34700, and the H$\\alpha$ line is seen in emission with highly variable profiles \\citep{Zuckerman:94,Arellano:03}.  It is also a strong X-ray source. The spectral type of the star has been listed as G0 \\citep{Odenwald:86}, or more recently as G\\ion{0}{4}e \\citep{Mora:01}.  Given the interest in the object and the evidence that it may be a spectroscopic binary, the main motivation for this paper is to present high-resolution spectroscopic observations that indeed confirm its double-lined nature. From these observations we derive an accurate double-lined orbital solution. We discuss also the available indicators of youth. ",
        "conclusions": "We have confirmed the young star HD~34700 to be a double-lined spectroscopic binary, and presented an accurate orbital solution with a period of 23.4877~days and a significant eccentricity ($e = 0.2501 \\pm 0.0068$). The components are of very nearly equal mass, temperature, and luminosity. The measured projected rotational velocities indicate super-synchronous rotation in both stars.  We have examined the available indicators of youth, and find that when considered together, the following present fairly compelling evidence in favor of PMS status:  \\noindent (a) Significant infrared excess suggesting the presence of a sizeable circumstellar disk. Polarization and CO measurements consistent with this idea. Large fractional luminosity in the infrared.  \\noindent (b) H$\\alpha$ in emission, with a strength ($\\sim$0.6~\\AA\\ equivalent width) that places the object in the weak-line T~Tauri class.  \\noindent (c) Strong \\ion{Li}{1}~$\\lambda$6708 absorption (equivalent width 0.17~\\AA).  \\noindent (d) Variable line profiles for H$\\alpha$, and possibly other spectral lines. Possibility of veiling.  \\noindent (e) Strong X-ray emission as seen by the ROSAT satellite, consistent with that detected in other weak-line T~Tauri stars.  \\noindent (f) Rapid rotation of the components (28~\\kms\\ and 22~\\kms, $\\sim$10 times faster than main-sequence stars of similar spectral type) in an orbit that is wide enough that the stars are not tidally synchronized.  \\noindent (g) Location of the star in the general vicinity of the Orion star-forming region, and a radial velocity ($+21$~\\kms) fairly close to typical values for other young stars in that area.  It is difficult to estimate a precise age for the system due to the lack of an accurate parallax. If it is at the distance of Orion, theoretical isochrones indicate it is only a few Myr old. However, it could be closer and therefore be approaching the ZAMS, in which case the age could be several tens of Myr. The \\ion{Li}{1}~$\\lambda$6708 absorption is not quite as strong as seen in other very young stars, but the Li measurements for HD~34700 may perhaps be affected by veiling.  HD~34700 thus joins the relatively reduced group of a few dozen young spectroscopic binaries with known orbits, with its double-lined nature making it a particularly interesting case."
    },
    "0311/astro-ph0311130_arXiv.txt": {
        "abstract": "We present the preliminary results of a wide-field photometric survey of individual red giant branch (RGB) and asymptotic giant branch (AGB) stars in the M81 group, performed with the CFH12K mosaic camera of the CFHT.  We use deep VI images of 0.65 sq.~deg.~of sky to map out the two-dimensional distribution of intragroup stars and to search for stars associated with the many HI tidal tails in the group.  We place an upper limit on the presence of metal-poor RGB stars in a field located 50-80 kpc from M81, and derive an `intragroup' fraction of $<2$\\%.  In a field sampling the M81-NGC3077 HI tidal tail, we find blue stars associated with some of the tidal features, including 2 clumps which we tentatively describe as tidal dwarf candidates.  These objects are $\\sim 1$ kpc in size, and, based on their color-magnitude diagrams, have formed stars as recently as $\\sim30-70$ Myr ago, long after the group's most recent interactions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311289_arXiv.txt": {
        "abstract": "He I $\\lambda$10830 profiles acquired with Keck's NIRSPEC for 6 young low mass stars with high disk accretion rates (AS 353A, DG Tau, DL Tau, DR Tau, HL Tau and SVS 13) provide new insight into accretion-driven winds. In 4 stars the profiles have the signature of resonance scattering, and possess a deep and broad blueshifted absorption that penetrates more than 50\\% into the 1 $\\micron$ continuum over a continuous range of velocities from near the stellar rest velocity to the terminal velocity of the wind, unlike inner wind signatures seen in other spectral features. This deep and broad absorption provides the first observational tracer of the acceleration region of the inner wind and suggests that this acceleration region is situated such that it occults a significant portion of the stellar disk. The remaining 2 stars also have blue absorption extending below the continuum although here the profiles are dominated by emission, requiring an additional source of helium excitation beyond resonant scattering. This is likely the same process that produces the emission profiles seen at He I $\\lambda$5876. ",
        "introduction": "Outflows from young stars undergoing disk accretion are suspected to play a vital role in the process of star formation and in the evolution of protoplanetary disks. Although their ubiquity among accreting YSOs is well documented, and there is some evidence for a rough proportion of ${\\dot M}_{wind}/{\\dot M}_{acc} \\sim 0.1$ \\citep{kon00}, the origin of the outflows remains a mystery. Most outflow models tap magneto-centrifugal ejection as the heart of the launching mechanism, where rotating fields in the disk fling material along inclined field lines. A distinguishing characteristic among disk wind models is the location in the system where mass loading onto field lines occurs -- either from the inner disk over a range of radii \\citep{kon00} or from the radius at which the stellar magnetosphere truncates the disk, lifting most of the accreting material toward the star into magnetic funnel flows and ejecting the rest along opened stellar field lines emerging from this point \\citep{shu94}. Others have explored the possibility that accretion-driven winds emerge via magnetohydrodynamic acceleration from the star rather than, or in addition to, the disk \\citep{hir97,kwa88}. Each approach has different implications for angular momentum evolution in accretion disk systems, both for the disk, where the dominant mode of angular momentum transport remains unknown, and for the accreting star, where the means of stellar spin-down in the face of accumulation of high angular momentum material is uncertain.  Observational probes of winds in accreting systems have not yet provided decisive evidence for where the wind originates. Either outflowing gas is observed far from the launch site, i.e. in collimated jets and molecular outflows, or else inner winds close to the star produce kinematic signatures that are entangled with emission deriving from other phenomena such as funnel flows and accretion shocks associated with magnetospheric accretion. Among the classical T Tauri stars (cTTS), the traditional indicator of an inner wind is blueshifted absorption superposed on strong and broad emission lines, as seen in H$\\alpha$, Na D, Ca II H\\&K, and MgII h\\&k \\citep{naj00}. The line emission is attributed primarily to infall in magnetic funnel flows \\citep{lee94, muz01} while the accompanying blueshifted absorption signifies the presence of a high velocity wind close to the star \\citep{cal97, ard02}. However little progress has been made in understanding the properties of this inner wind or in establishing whether it originates from a range of radii in the inner disk, the disk truncation radius, or the star.  A recent analysis of profiles of He I $\\lambda$5876 and $\\lambda$6678 in a sample of 31 accreting cTTS revealed a new means of diagnosing the inner wind \\citep{ber01}. The helium lines were selected for study because of their high excitation potential, restricting line formation to a region either of high temperature or close proximity to a source of ionizing radiation. In spite of these restrictions, the helium emission lines were found to have a composite origin, including contributions from a wind, from the funnel flow, and from an accretion shock.  The wind component, characterized by broad blueshifted emission (centroid velocities exceeding $-$30 km~s$^{-1}$) and extended blue wings (maximum velocities exceeding $-$200 km~s$^{-1}$), was most prominent among stars with the highest disk accretion rates. Curiously, the high accretion rate stars with a strong hot helium wind component were also found to have anomalously weak or absent signatures of the narrow emission component formed in the accretion shock where the magnetosphere directs accreting material from the disk. We interpreted these findings as indicating that a hot accretion driven wind originates in the vicinity of the stellar corona, and that when the hot wind is present, the conditions of magnetospheric accretion and its accompanying accretion shock may be modified.  This intriguing result led us to pursue an investigation of an additional line in the He I triplet series, $\\lambda$10830, immediately following $\\lambda$5876 in a recombination-cascade sequence. The lower $2s^3S^{\\circ}$ level of He I $\\lambda$10830 is metastable, resulting in conditions ripe for resonance scattering, which provides an opportunity to search for outflowing gas in absorption. The sensitivity of $\\lambda$10830 absorption to the dynamics of winds in cool stars is illustrated in Dupree et al. (1992). This feature is in a spectral region that was until recently not readily accessible to either CCD or infrared imaging detectors. An early study by Ulrich and Wood in 1981, using pre-ccd technology, detected this feature in a few T Tauri stars, but the data were not of sufficient quality to interpret the line profile. The next published work featuring $\\lambda$10830 in a T Tauri star is for DG Tau {\\citep{tak02}, showing broad emission and a blueshifted absorption feature, confirming the prediction of BEK that the inner wind is traced by helium.  Using Keck II's NIRSPEC, we are undertaking a census of He I $\\lambda$10830 in accreting young stars to evaluate the constraints that helium places on inner winds. In this Letter, we present He I $\\lambda$10830 profiles for 6 low mass YSOs with higher than average disk accretion and mass outflow rates to demonstrate that this feature offers new insight into mass loss from young stars with accretion disks. ",
        "conclusions": "The He I $\\lambda$10830 profiles for 6 high accretion rate cTTS are shown in Figures 1 and 2. The $\\lambda$10830 profiles all show unequivocal evidence for mass loss with blueshifted absorption penetrating below the continuum level, with depths ranging from 20\\% (DG Tau) to 90\\% (DR Tau) and with maximum velocites ranging from $-$200 km~s$^{-1}$ (HL Tau) to $-$450 km~s$^{-1}$(DG Tau). In Figure 1 we compare them to non-simultaneous profiles for He I $\\lambda$5876 from BEK for the 4 stars in common with that study and in Figure 2 we compare them to non-simultaneous H$\\alpha$ profiles from BEK or Hillenbrand and White (in preparation).  While blueshifted absorption is not seen in the $\\lambda$5876 transition (discussed below), it is present in traditional indicators of the inner wind, such as the H$\\alpha$ profiles shown here. However what distinguishes He I $\\lambda$10830 is both the prominence and depth of its blueshifted absorption, offering the promise of a more direct probe of the inner wind. At H$\\alpha$ all profiles are overwhelmingly dominated by broad and strong emission, while for $\\lambda$10830 4 of the 6 stars are either predominantly in absorption (DR Tau and SVS 13) or else resemble classic P Cygni profiles, with comparable levels of emission and absorption (HL Tau and AS353A). Moreover at $\\lambda$10830 the blue absorption is far deeper and broader than at H$\\alpha$, penetrating the continuum in all 6 stars. The most significant aspect of the blueshifted absorption at helium is that in 4 stars (DR Tau, HL Tau, SVS 13 and AS 353A) it penetrates to depths of more than 50\\% into the continuum over a continuous and broad velocity range extending from near the stellar rest velocity to the terminal wind velocity, as inferred from velocities characterizing their spatially extended jets (see Table 1).  The uniqueness of He I $\\lambda$10830 as a probe of the inner wind derives from the metastability of its lower level ($2s^3S^{\\circ}$), which, although energetically far above (20 eV) the singlet ground state, is radiatively isolated from it. Whether this metastable level is populated by recombination and cascade or by collisional excitation from the ground state, it will become significantly populated relative to other excited levels owing to its weak de-excitation rate via collisions to singlet states, making it an ideal candidate to form an absorption line. This absorption is essentially a resonant-scattering process since the $\\lambda$10830 transition is the only permitted radiative transition from its upper state to a lower state and the electron density is unlikely to be so high as to cause collisional excitation or de-excitation. The absence of any absorption or emission at $\\lambda$10830 in the non-accreting wTTS V827 Tau (not shown), which like all wTTS has strong coronal x-rays \\citep{fei02}, clarifies that the presence of an inner T Tauri wind traced via resonant scattering from the helium metastable level requires the additional presence of disk accretion and its associated energetic phenomena.  New insight regarding the inner wind in cTTS can be made from visual inspection of the helium profiles for those 4 stars where resonance scattering is the dominant formation mechanism. Their deep and broad absorption troughs reveal new information on the kinematics of the inner wind, probing its acceleration region in an accreting star for the first time, and on the geometry of the inner wind, suggesting that the acceleration region is likely close to the star and occults a significant portion of it.  The insight into the kinematics of the wind is provided by the large and continuous span of velocities over which the blueshifted absorption is observed. This suggests that we are probing the acceleration region of the inner wind, as the absorption extends from velocities at rest relative to the star continuously through to velocities comparable to or often larger than the terminal wind velocities inferred from spatially extended jets. There is also a suggestion of turbulence in the inner wind, since in 2 of the stars the blueshifted absorption trough extends to velocities as red as 50 km~s$^{-1}$.  The insight into the geometry of the inner wind comes from the depth of the absorption into the $1~\\mu m$ continuum over such a wide range of velocities, from rest velocity to the terminal velocity of the wind. In fact, the actual depletion of blueshifted continuum photons will exceed what is observed, as some photons scattered into the line of sight will fall within the velocity range of the blue absorption, making it appear shallower and less extended. At $1~\\mu m$ the continuum in a heavily veiled star will have a larger fraction contributed by the stellar photosphere than at other wavelengths, because the spectral energy distributions of the excess emission arising from both the accretion shock on the stellar surface \\citep{gul98} and the inner disk \\citep{muz03} are rapidly decreasing toward this wavelength while the stellar spectral energy distribution is peaking. For example, at $1~\\mu m$ veiling from an accretion shock at T$\\sim$8000K will be reduced by a factor of 5 compared to the V band and veiling from the inner disk with T$\\sim$1400K will be reduced by a factor of 4 compared to the K band.  We have determined the veiling at $1~\\mu m$ for our high accretion rate stars by comparing the depth of photospheric features in the echelle order centered at 1.075$\\mu m$ with the wTTS V827 Tau. Our veiling estimates, listed in Table 1, indicate that the photosphere contributes between 30\\% to 50\\% of the $1~\\mu m$ continuum in stars known for much higher veilings in the optical and at 2$\\mu m$ \\citep {har95,muz03}. (Photospheric lines could not be detected in SVS 13 because of obscuration by emission lines, however we estimate the veiling is $\\sim 2$ based on the equivalent width of these lines compared to emission features in the other stars.) With the clarification that a significant fraction of the continuum arises from the star, including both the stellar photosphere and the accretion shock on the stellar surface, then the implication of the observed depth and velocity range of the absorption in the 4 stars where $\\lambda$10830 is predominantly formed by resonance scattering is that the full span of the wind acceleration region lies along the line of sight to the stellar disk.  An unusual feature of the $\\lambda$10830 profiles shown here is that 2 of the 4 stars with resonance scattering profiles have significant net negative equivalent widths (see Table 1). This is unexpected in the standard formation scenario for a P-Cygni line by resonant scattering, which yields a net zero equivalent width in the absence of stellar occultation, indicating that in these two stars far fewer photons are being scattered into the line of sight than out of it. In the most extreme case, DR Tau, the net equivalent width has a value of -10.2 \\AA~ or, equivalently in velocity units, of 283 km~s$^{-1}$. We can demonstrate analytically that this unusual behavior can be accounted for in a resonant scattering scenario if occultation of scattered photons is provided by both the star and the accretion disk. For example, if the disk were to extend into the stellar surface, thereby screening photons scattered from receding material, then in a pole-on system with a spherical isotropically scattering shell at R = 2 R$_\\star$, the photons scattered into the line of sight amount to only 40\\% of the continuum photons scattered out of the line of sight. This fraction increases to 76\\% as the inclination changes from 0$^{\\circ}$ to 90$^{\\circ}$. The corresponding net negative equivalent widths would thus be characterized by 60\\% to 24\\% removal of the continuum, values consistent with the large negative values found here for DR Tau and SVS 13. Moving the scattering surface to different radii does not significantly alter this result, although if the disks of these 2 stars are truncated at a radius outside the scattering surface then accounting for their large negative equivalent widths will be a challenge.  Additional perspective on the inner wind comes from comparing optical He I profiles to those of $\\lambda$10830. In Figure 1 we have compared (non-simultaneous) $\\lambda$10830 and $\\lambda$5876 profiles for the 4 stars also in the BEK study. For these 4 high accretion rate stars BEK concluded that their $\\lambda$5876 emission primarily originates in a wind rather than in funnel flows or accretion shocks. This conclusion was based on the large blueshifted centroids and/or the large blue wing velocities found in the broad component of the $\\lambda$5876 emission coupled with no conspicuous narrow component coming from the accretion shock. This interpretation is strengthened by the unequivocal wind signatures found at $\\lambda$10830 in the cTTS studied here, where the blue wings are now in absorption rather than emission. The absence of blue absorption at $\\lambda$5876 is readily understood as due to the optical thinness of this line in the wind, given the strong radiative decay rate of the lower level of the $\\lambda$5876 transition.  The strength of the emission at $\\lambda$5876 also suggests that the helium wind can have a significant optical depth at $\\lambda$10830. As this optical depth increases, helium emission will be enhanced over the resonant scattering process, either via recombination/cascade and/or collisional excitation. This effect could then account for the 2 stars investigated here, DG Tau and DL Tau, which have large positive equivalent widths at $\\lambda$10830. They could still have an underlying trough-like absorption from resonance scattering, extending from the stellar rest velocity to the wind terminal velocity as seen in the other 4 stars, which would be filled in by the additional emission processes in the wind. Since both resonant scattering and line emission contribute to $\\lambda$10830 formation, while the primary contributor to $\\lambda$5876 is line emission, we anticipate that simultaneous observations of both He I lines, together with theoretical modelling, may yield the velocity, density, and temperature structures of the inner wind.   In summary, we conclude that the He I $\\lambda$10830 line offers a unique new probe of inner winds in accretion disk systems, revealing the wind acceleration region and providing constraints on the geometry of this region. Most importantly, the deep and broad absorption seen in 4 stars seems to require that the full span of the wind acceleration region is seen in projection against the stellar disk. We find it more straightforward to interpret this finding in a scenario where the wind arises from stellar coronal regions, since it would be difficult for wind material arising from a disk to absorb the radiation from the stellar disk at all velocities, from rest to the terminal wind speed. If this is the case, it would not be a normal \"stellar wind\", since this phenomenon is strongest in stars with the highest disk accretion rates. We anticipate a fuller understanding of the physical properties of the inner wind, its geometry and its launch distance from the star when the behavior of this line has been examined for a large sample of accreting young stars covering a wide range of disk accretion rates."
    },
    "0311/astro-ph0311240_arXiv.txt": {
        "abstract": " ",
        "introduction": "Pulsar velocities range from 100 to 1600~km~s$^{-1}$~\\cite{astro, astro_1}. They present a long-standing puzzle.  According to numerical simulations of supernova explosions~\\cite{explosion}, the asymmetries in the core collapse could not account for a kick velocity of more than 300-600~km~s$^{-1}$.  Although the average pulsar velocity is in this range~\\cite{astro}, there is a substantial population of pulsars with velocities in excess of 700~km/s, while as many as 15\\% of all pulsars have speeds over 1000~km~s$^{-1}$~\\cite{astro_1}.  It has been suggested that an asymmetric emission of electroweak-singlet neutrinos from a cooling neutron star could explain the pulsar velocities~\\cite{ks97,fkmp}. Since most of the supernova energy is carried away by neutrinos, only a few per cent asymmetry is sufficient.  The range of parameters (mass and mixing angle) consistent with this explanation overlaps with the allowed region for the singlet neutrinos to make up the cosmological dark matter~\\cite{fkmp,Fuller,dh}.  Let us consider a singlet (or sterile) neutrino $\\nu_s$ that has a small mixing with the electron neutrino $\\nu_e$.  The mixing angle $\\theta $ we will consider is $10^{-5} - 10^{-4}$, much smaller than those accessible to laboratory experiments.  However, neutrino oscillations could result in a production of such a neutrino in the early universe, as well as in a supernova.  There are actually two possible mechanisms for the pulsar kick due to the sterile neutrinos.  One is based on off-resonance production and requires the singlet neutrino to have a mass of a few keV~\\cite{fkmp}.  The other possibility is a resonant $\\nu_e \\rightarrow \\nu_s$ conversion, which can occur for the mass of several to 20 keV. Let us briefly discuss both of these possibilities. ",
        "conclusions": ""
    },
    "0311/astro-ph0311076_arXiv.txt": {
        "abstract": "{In the course of the Hamburg/ESO survey we have discovered a white dwarf whose spectrum exhibits  many similarities with the prototype of magnetic white dwarfs Grw+70$^{\\circ}$8247. In particular several  stationary line components indicative for magnetic fields between about 150 and 400 MG are found in both objects. However, the features between 5000 and 5500\\,\\AA\\ in the spectrum of HE0241-0155 cannot be explained by stationary line components and demand a relatively homogeneous magnetic field with clustering around 200\\,MG. For this reason a pure dipole model failed to reproduce this spectral region. An offset-dipole configuration led to some improvement in the fit but a good agreement was only possible for a geometry -- described by an expansion into spherical harmonics -- where most of the surface is covered with magnetic field strengths strongly clustered around 200\\,MG. This may indicate the presence of a large magnetic spot whose presence  could be tested  with time resolved spectro-polarimetry. ",
        "introduction": "A few percent (5\\%) of all known white dwarfs have detectable magnetic fields, and the true fraction may be as high as \\mbox{10 \\%} if selection effects are taken into account (Harris et al. 2003). However, the number of white dwarfs with field strengths above 100 MG, where the spectral line pattern is no longer similar to normal H or He I spectra, is still small, and each star is unique. Most of these stars have been detected serendipitously during follow-up spectroscopy by surveys designed for other purposes, e.g. quasar surveys like the Palomar-Green Survey (Schmidt \\& Green 1983) or the Hamburg/ESO survey (Wisotzki et al. 1996, Reimers \\& Wisotzki, 1997). The first star of this type was Grw+70$^{\\circ}$ 8247, whose unique spectrum was first observed by by Minkowski (1938). Its shallow spectral features resisted a physical identification for decades. Kemp (1970), following an idea by J. Landstreet,  proposed that a magnetic field would produce  circular dichroism, and discovered  circular dichroism in Grw+70$^{\\circ}$ 8247. Only progress in quantum mechanical calculations of the hydrogen energy levels in strong magnetic fields made possible an identification (Praddande, 1972, R\\\"osner et al. 1984). Details of the Grw+70$^{\\circ}$ 8247 spectrum and its interpretation are given by Greenstein (1984), Angel et al. (1985) and Wickramasinghe \\& Ferrario (1988). In this letter we report on the serendipitous discovery of a near ``twin'' to Grw+70$^{\\circ}$8247 in the course of the Hamburg/ESO survey. We give a physical interpretation of its spectra using Jordan's code (Jordan, 1988, 1992, Euchner et al. 2002) in terms of de-centered magnetic dipole fields and we briefly discuss the remaining discrepancies between observations and theory. ",
        "conclusions": "The efficiency in modeling the spectra and polarization data of magnetic white dwarfs has considerably improved due to two factors: Firstly, instead of solving the radiative transfer equations on thousands of surface elements on the star for each model, databases for flux and polarization spectra computed for a wide range of field strengths, viewing angles, and effective temperatures are used. Secondly, either evolutionary strategies or genetic algorithms, and a subsequent downhill simplex algorithm (Press et al., 1992) are used in order to find the best fit  to the observations. The approach used in this paper is very similar to the one described by Euchner et al. (2002) with the exceptions that a genetic algorithm (Caroll 2001) has been used, more viewing angles (12 instead of 9), and an interpolation (12 $\\mu$ angles) instead of a linear approximation are used.  If a pure magnetic dipole is assumed, the best fit (Fig.\\,\\ref{bestfit}) to the observed data was obtained for a polar field strength of 398 MG, with the dipole axis oriented almost ($i=-86^\\circ$) perpendicular to the observer, and an  effective temperature of 12000\\,K.  However, this model fails to account for the features at 5093, 5218, and 5287\\,\\AA. These features do not correspond to any stationary line component in any reasonable range of magnetic field strengths; the closest agreement would have the 4f$_{-1}\\rightarrow$ 2s$0$ components with a maximum at 5119\\,\\AA\\ and 61\\,MG. The wavelengths of stationary line components  go through maxima or minima as functions of the magnetic field strength; they   are visible in the spectra of magnetic white dwarfs despite a considerable variation of the field strengths as in the case of a pure magnetic dipole (where the field strength is a factor of two higher at the pole than at the equator, and the distribution of the magnetic field strengths over the visible stellar surface is relatively flat (see Fig.\\,\\ref{histo}).  As indicated by the circles in Fig\\,\\ref{observation}, the features missing in the fit could be due to a relatively homogeneous field of about 200\\,MG. One possibility to describe a field distribution where the majority of the surface is covered by a magnetic field close to 200\\,MG is a dipole field with a large offset of the dipole relative to the center of the star. Although the total variation of field strengths is larger in this case, more magnetic field  strengths cluster around one relatively homogeneous magnetic field on one pole.  The best fit for an offset dipole (only offsets along the magnetic axis are considered) is obtained for the same polar field strength (399 MG) and the same orientation ($i=-86^\\circ$), but for an offset of 0.16 stellar radii away from the observer. The histogram Fig.\\,\\ref{histo} shows the effect: The magnetic field strengths are now even more strongly peaked at 200\\,MG and the vast majority of surface elements has a magnetic field between 200 and 250\\,MG. The resulting fit (upper curve of Fig.\\,\\ref{bestfit}) has improved considerably, and particularly the missing three line features are now visible -- although still somewhat too shallow -- in the theoretical model.  Finally, a more general approach, an expansion of the magnetic field into spherical harmonics, assuming that the source of the magnetic field lie entirely within the star (Jacobs 1987):  \\begin{displaymath} \\begin{array}{lll} B_r=&-\\sum_{l=1}^\\infty \\sum_{m=0}^l (l+1) (g_l^m \\cos m\\phi +h_l^m \\sin m \\phi)\\\\ &\\hfill  P_l^m(\\cos\\theta)\\\\ B_\\theta=+ &\\sum_{l=1}^\\infty \\sum_{m=0}^l  (g_l^m \\cos m\\phi + h_l^m \\sin m \\phi)\\\\ & \\hfill dP_l^m(\\cos\\theta)/d\\theta\\\\ B_\\phi=&-\\sum_{l=1}^\\infty \\sum_{m=0}^l  m (g_l^m \\cos m\\phi + h_l^m \\sin m \\phi)\\\\&\\hfill dP_l^m(\\cos\\theta)/d\\sin \\theta \\end{array} \\end{displaymath} with the associated Legendre polynomials $P_l^m$. The components are given in spherical coordinates $r, \\theta,$ and $\\phi$. For this paper we limited ourselves to $l\\le 4$ and $m=0$, leading to 4 free parameters beside the orientation of the magnetic axis relative to the observer. Any attempt to account for the 20 $m \\ne 0$ parameters, did only slightly improve the solution. Since no polarization data additionally constrain the distribution of the longitudinal component of the magnetic field, we present only the best fit for $m=0$: $g_1^0=172$\\,MG, $g_2^0=74$\\,MG, $g_3^0=0.3$\\,MG, $g_4^0=-0.1$\\,MG, and $i=-97^\\circ$.  Fig.\\,\\ref{bestfit} shows  now that the overall agreement between observation and theoretical fit is  almost perfect. The corresponding histogram is now even stronger peaked around 200\\,MG.  It is clear that the limitation to spectra without polarization data cannot lead to a unique solution. All distributions that cluster around 200\\,MG should lead to similar results. Such a strong clustering could be  a hint for a magnetic spot with a relative homogeneous magnetic field covering a large portion of the visible hemisphere, as discussed for the magnetic white dwarfs PG\\,1031+234 (Latter et al. 1987) and RE\\,0317-853 (Vennes et al. 2003). The origin of such a spot-like feature, is however unclear although star spots have been found on the possible progenitors of magnetic white dwarfs, the Ap stars, where the spots are connected to abundance anomalies and can be detected via Doppler imaging (Khokhlova et al. 1986, Rice \\&\\ Wehlau 1994).  An extremely homogeneous field was also found in phase-resolved low-state spectra of the polar MR Ser. Schwope et al. (1993) revealed a strongly decentered dipolar ($z_{\\rm off}\\approx 0.3$ stellar radii).  However, it must be pointed out that in Ap stars, which also have fossil magnetic fields, low order multipole expansions describe statistical aspects of overall field geometry, but do not provide a realistic representation of the actual distribution of field vectors over the surface; the latter would be more constrained by observations of circular polarization. Therefore, our solution only means that the field modulus is quite uniform and large over the visible hemisphere of the star. The relatively small region with a higher field modulus does not strongly contribute to the strength of the line features.  In the framework of multipole exansions Muslimov et al. (1995) have shown that a weak quadrupole (or octupole, etc.) component on the surface magnetic field of a white dwarf may survive the dipole component und specific initial conditions: Particularly the evolution of the quadrupole mode is very sensitive (via Hall effect) to the presence of internal toroidal field. For a 0.6\\,$M_{\\odot}$ white dwarf with a toroidal fossil magnetic field of strength $<10^9$\\,G the dipole component declines by a factor of three in $10^9$\\,yr, while the quadrupole component is practically unaffected. Without an internal toroidal field the dipole component still declines by a factor of three but the quadrupole component is a factor of six smaller after 10\\,Gyr. This is much longer than the cooling age of a 12000\\,K white dwarf, which well under 1\\,Gyr for  0.6\\,$M_{\\odot}$ and larger than 3\\,Gyr only for masses larger than 1.2\\,$M_{\\odot}$.  Whether spot-like structures could also survive for the long cooling time of white dwarfs has to be investigated theoretically.  If a magnetic spot is indeed present  a strong temporal variation of the spectrum and particularly the polarization is expected, if the star is rotating. However, no significant differences could be found between the spectra taken in September, 1998 and November, 1998, respectively. This may be an indication for only slow rotation as known in the case of Grw+70$^{\\circ}$8247 where a rotational period of decades or centuries is possible (Schmidt \\&\\ Norsworthy 1991, Friedrich \\&\\ Jordan 2001). Therefore, a more detailed analysis and a decision whether a spot is present has to wait until  time resolved spectro-polarimetry of this object becomes available."
    },
    "0311/astro-ph0311399_arXiv.txt": {
        "abstract": "We report observations of the hot post-asymptotic giant branch star ZNG~1 in the globular cluster M5 (NGC~5904) with the {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)}.\\/ From the resulting spectrum, we derive an effective temperature \\teff\\ = $44300 \\pm 300$ K, a surface gravity \\logg\\ = $4.3 \\pm 0.1$, a rotational velocity \\vrot\\ = $170 \\pm 20$ \\kms, and a luminosity $\\log L/L_\\sun = 3.52 \\pm 0.04$.  The atmosphere is helium-rich (Y = 0.93), with enhanced carbon (2.6\\% by mass), nitrogen (0.51\\%) and oxygen (0.37\\%) abundances.   The spectrum shows evidence for a wind with terminal velocity near 1000 \\kms\\ and an expanding shell of carbon- and nitrogen-rich material around the star. The abundance pattern of ZNG~1 is suggestive of the ``born-again'' scenario, whereby a star on the white-dwarf cooling curve undergoes a very late shell flash and returns to the AGB, but the star's rapid rotation is more easily explained by a previous interaction with a binary companion. ",
        "introduction": "Post-asymptotic giant branch (PAGB) stars represent a brief ($10^4$--$10^5$ yr) phase of stellar evolution during which stars move rapidly across the HR diagram at constant luminosity from the cool tip of the AGB to the beginning of the hot white-dwarf cooling curve. As such, they are keys to both the past and the future:  their atmospheric abundances provide important information about the mixing and mass-loss processes at work in AGB stars, while their masses yield constraints on the masses of white dwarfs currently forming in globular clusters. To investigate these issues, we observed the PAGB star ZNG~1 (Table~\\ref{obs_parms}) in the globular cluster M5 (NGC~5904) with the {\\it Far Ultraviolet Spectroscopic Explorer (FUSE).}  First identified as a UV-bright star by \\citet*{ZNG:72} on the basis of its extreme $U-V$ color, ZNG~1 was observed with \\iue\\/ by \\citet{Bohlin:83}, who suggested that an apparent \\ion{N}{4} $\\lambda 1487$ emission feature is the signature of a planetary nebula (PN). Combining this spectrum with other archival \\iue\\/ data, \\citet{deBoer:85} argued that the apparent \\ion{N}{4} feature is spurious and pointed out the P-Cygni shape of the strong \\ion{N}{5} $\\lambda \\lambda 1238, 1242$ doublet. To test the PN hypothesis, \\citet{Napiwotzki:97} used the {\\it Hubble Space Telescope} to obtain a WFPC2 H$\\alpha$ image and GHRS ultraviolet spectra of the star.  The H$\\alpha$ image shows a bright G-type star only 0\\farcs5 from ZNG~1, but no evidence for extended emission around the star.  The GHRS spectra clearly show P-Cygni profiles in the \\nfive\\ doublet and discrete absorption components (DACs), blueshifted by $\\sim$ 900 \\kms, in both \\nfive\\ and \\cfour\\ $\\lambda \\lambda 1548, 1550$.  No nebular emission is present. ",
        "conclusions": "Hot PAGB stars in globular clusters differ from their field counterparts in that they rarely show evidence for third dredge-up (by which the products of helium burning are brought to the surface) and are seldom accompanied by PNe \\citep{Jacoby:97}.  These characteristics can be understood qualitatively as a result of the expected low core mass ($< 0.55  M_\\sun$) of globular cluster AGB stars. While the minimum core mass for third dredge-up is still a subject of active research \\citep*{Karakas:02}, nearly all current AGB models agree that third dredge-up should not occur in stars with initial mass $\\lesssim 0.8 \\; M_\\sun$.     The rarity of PNe can be explained by the lower luminosities and longer crossing times of PAGB stars with low core masses; the nebula disperses before the central star becomes hot enough to ionize it \\citep{Jacoby:97}.   \\citet*{Alves:00} therefore suggest that binary evolution is necessary for the formation of PNe in globular clusters.      Interestingly, the best-studied PN in a globular cluster, K648 in M15, shows clear evidence for third dredge-up in its central star \\citep*{Rauch:02}, but not the strong helium and nitrogen enhancements seen in ZNG~1.  One way to produce a hydrogen-deficient star with enhanced CNO abundances is for the star to experience a very late shell flash while descending the white dwarf cooling curve and subsequently return to the AGB \\citep{Herwig:99}.     This ``born-again'' scenario might explain the absence of a PN surrounding ZNG~1, because the star is so old (since it is on a second crossing), that any nebula has long since dispersed.     For the born-again scenario to work for ZNG~1, the hydrogen must not be completely burned, and the carbon enhancement must not be too high.  \\citet{Lawlor:03} provide a born-again model of an initial 1 $M_{\\sun}$ star with  Z = 0.001 that predicts stellar parameters similar to those of ZNG~1 (see their Table 3), although the CNO abundances are about a factor of three larger.     The difficulty with the born-again scenario (or any single-star evolution scenario) is providing an explanation for the observed rapid rotation. One possibility: to explain the observed rotational velocity of up to 40 \\kms\\ in some globular horizontal-branch stars \\citep{behr:03},  \\citet{sills:00} postulate that red giants may have a rapidly rotating core.  If so, then perhaps its angular momentum can be transferred to the surface during the convective mixing in the final shell flash.  Might the properties of ZNG~1 be explained through interaction with another star? First, we note that its very blue \\bv\\ color indicates that any current companion must be quite faint. Given the proximity of the star to the projected cluster center (4\\farcs2) and the moderately high stellar interaction rate in M5 \\citep{Pooley:03}, one possiblity is a stellar collision, which is thought to be the primary production mechanism for the blue stragglers found at the centers of globular clusters.  Collisionally-created blue stragglers are expected to be born with a high angular momentum.  \\citet*{Shara:97} measure a rotation rate of $155 \\pm 55$ \\kms\\ and estimate a mass of $1.7 \\pm 0.4 \\; M_{\\sun}$ for a blue straggler in 47 Tuc.  A star of this mass would likely experience third dredge-up on the AGB, but would also be expected to eject a PN, which is not seen for ZNG~1.  Furthermore, it is unclear how much of a blue straggler's angular momentum would survive its post-main sequence evolution.  A second scenario, in which a low-mass companion spirals into the AGB progenitor, has been proposed to explain K648 in M15 by \\citet{Alves:00} and discussed in general by \\citet{DeMarco:02}. The increased AGB mass due to the ingestion of the companion might be sufficient to allow third dredge-up to occur, but additional mechanisms such as shear or rotationally-induced mixing are required to create a hydrogen-deficient star \\citep{DeMarco:02}.  A merger process is thought to be the origin of the high rotational velocity of FK Comae stars ($v \\sin i = 162.5 \\pm 3.5$ \\kms\\ for the prototype FK Comae; \\citealt{Huenemoerder:93}). More detailed analysis is required to assess the viability of a merger scenario for ZNG~1; whether the merger remnant should eject a PN is of particular interest.  \\citet{Pauldrach:88} compute terminal velocities and mass-loss rates for the winds of (Pop.\\ I) central stars of PNe.    Assuming an evolutionary mass of 0.565 $M_{\\sun}$ (\\S \\ref{stellarparms}) and  temperature of \\teff\\ = 45 K, their models predict mass-loss rates of a few times $10^{-9}$ $M_{\\sun}$ yr$^{-1}$. Given this modest wind,  it appears unlikely that ZNG~1 will dissipate significant angular momentum in the brief time ($\\sim 15,000$ years, according to the 0.565 \\msun\\ track of \\citealt{Schoenberner:83}) prior to its arrival on the white-dwarf cooling curve. Previous studies have yet to find white dwarfs with such rapid rotation (see review by \\citealt{Kawaler:03}) but the example of ZNG~1 suggests that further searches may be fruitful."
    },
    "0311/astro-ph0311485_arXiv.txt": {
        "abstract": "We present multicolor photometry of bright star cluster candidates in the nearby starburst galaxies NGC~3077 and NGC~5253, observed with the {\\it Hubble Space Telescope} Wide Field and Planetary Camera~2 in both broadband (F300W, F547M and F814W) and narrowband (F487N and F656N) filters.  By comparing the photometry with theoretical population synthesis models, we estimate the age and mass of each star cluster, which provides constraints on the recent star formation histories of the host galaxies.  We compare the star cluster populations in these dwarf starburst galaxies to those of the nuclear starburst in the barred spiral M~83, and discuss implications for our understanding of the nature and evolution of starburst events. ",
        "introduction": "\\label{sec:intro}  Throughout the Universe's history, a significant fraction of star formation has apparently occurred in high-intensity bursts.  These rapid star-formation events play important roles in the ionization, chemical enrichment, and overall evolution of their host galaxy. Such starburst events are often seen in interacting galaxies, and it seems reasonable to infer a causal connection between galaxy interactions and starburst activity.  However, interactions are probably not the only causal agent for starburst activity; internal dynamics can also play an important role (\\eg, M~83). Combined with the dynamic merger histories predicted by the standard hierarchical collapse model of galaxy formation, it is likely that much of a typical galaxy's stellar content was formed during starburst events.  Despite this ubiquity, we currently have only a rudimentary understanding of the cause, nature and evolution of starburst events, and in particular, of the physical processes which connect the external environment to the internal star formation rate.  Nearby starburst galaxies provide important laboratories for investigating the processes involved in galactic-scale star-formation events in detail.  Local starburst galaxies appear to be similar to the active galaxies seen at high-redshift ($z\\simgreat2.5$) in UV colors \\citep{meu97, mhc99, as00} and spectral morphology \\citep{ste96, low97, pet00}, although the high-redshift systems typically have star-formation rates that are several times higher than local starbursts.  Because of their similarity, we can reasonably expect that detailed photometric and spectroscopic studies of local starbursts will lead to a general understanding of the starburst phenomenon at all redshifts.  In particular, we wish to address the following questions regarding the physical processes governing the starburst phenomenon: (1) Are starburst properties determined by the global properties of the host galaxies? (2) How long do starbursts last? (3) What are the mechanisms that cause and sustain starbursts? (4) Does star-formation activity propagate across galaxies?  If so, what is the propagation mechanism? (5) On what timescale do star clusters dissolve in starbursts? (6) Is the field stellar population composed primarily of evaporated cluster stars, or is there a significant population that is native to the diffuse field?  To investigate these questions, we are undertaking a multi-band photometric survey of nearby starburst galaxies using the Hubble Space Telescope ({\\it HST}) Wide Field and Planetary Camera 2 ({\\it WFPC2}).  The program's goal is to survey the stellar populations and gas content of nearby starburst galaxies spanning a range of internal properties (mass, metallicity, internal dynamics) and external environments, in order to understand better what role these factors play in determining the nature and evolution of the galaxies' starbursts.  The focus of the present work is to constrain the recent star-cluster formation histories of two nearby dwarf galaxies (NGC~3077 and NGC~5253) which have similar mass ($M\\sim10^9M_\\odot$) and luminosity ($L_B\\sim10^9L_{B\\odot}$), but large contrasts in metallicity and external environment.  Until recently, it was thought that NGC~3077 and NGC~5253 were both interacting satellites of large spiral galaxies.  NGC~3077's membership in the M~81 group is not in doubt, but recent Cepheid distance measurements place NGC~5253 about 600~kpc from its nearest neighbor, M~83 \\citep{kar02b, sch94}, making it unlikely that M~83 has influenced NGC~5253's recent history.  We also compare their recent star-cluster formation histories to that of the large spiral galaxy M~83 (NGC~5236), which was studied in Paper~1 of this series \\citep{har01}.  The global properties of NGC~3077, NGC~5253 and M~83 are compared in Table~\\ref{tab:galprop}.  The paper is organized as follows.  We present the observations and data reduction in Section~\\ref{sec:data}.  In Section~\\ref{sec:analysis}, we present the photometric catalogs, the extinction corrections, the star/cluster separation analysis, and our estimates of the star cluster ages and masses, based on both broad-band photometry and the equivalent width of the $H\\alpha$ emission line.  We interpret and discuss the results in Section~\\ref{sec:discuss} and summarize the work in Section~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discuss}  \\subsection{Cluster Formation Histories}\\label{cfh}  The spatial distributions of our cluster-age estimates for NGC~3077 and NGC~5253 are shown in Figures~\\ref{fig:n3077.agemap} and \\ref{fig:n5253.agemap}, respectively.  Each cluster position in the F547M image is labeled with its best age estimate, in Myr.  The starburst region in NGC~3077 is dominated by a heavily-extincted central dust cloud (see Figure~\\ref{fig:color}).  There is a large group of clusters younger than 10~Myr in a diffuse fan-shaped clump to the north of the central dust cloud.  Many of these clusters appear to be nearly coeval, with ages between 5 and 7~Myr.  The most massive of these is strikingly apparent as a blue knot to the northeast of the dusty core in Figure~\\ref{fig:color} (object~\\#12 in Table~\\ref{tab:n3077clusters}).  The youngest cluster in our NGC~3077 sample (aged 1~Myr) lies inside the central dust cloud (object~\\#43 in Table~\\ref{tab:n3077clusters}); there are also two objects nearby that are so heavily extincted that they were undetected in the F300W image; these objects are likely also very young clusters that have yet to emerge from the surrounding dust cloud.  There are several clusters in NGC~3077 aged between 10 and 20~Myr, and these are scattered throughout the central starburst region.  The visually brightest cluster in our NGC~3077 sample falls in this age range (object~\\#3 in Table~\\ref{tab:n3077clusters}, aged 13~Myr), and is located adjacent to the central dust cloud.  This cluster's heavy extinction ($E(B-V)=0.7$~mag) and close proximity to the central dust cloud suggest that its age may be overestimated.  Indeed, if we adopted an even larger extinction, it would be possible to place the cluster on the youngest portion of the Starburst99 model track in Figure~\\ref{fig:2cd}.  However, the cluster has a very low $EW(H\\alpha)$ (see Table~\\ref{tab:n3077clusters}), making it more likely that its present age estimate is the correct one.  NGC~3077 contains several clusters older than 20~Myr, up to 300~Myr old.  Most of these are located to the west and southwest of the central dust cloud, and are relatively far from the galaxy's center.  In NGC~5253, we see a tight group of very young (age $<5$~Myr) clusters near the center of the galaxy, and there are similarly young clusters scattered throughout the central starburst region, including the \"super star cluster candidate\" of \\cite{cal97b}; object \\#1 in Table~\\ref{tab:n5253clusters}.  In contrast to NGC~3077, the clusters with ages between 5 and 10~Myr in NGC~5253 are scattered throughout the starburst region, with no obvious concentrations.  Looking at slightly older objects, there are two very bright clusters or groups of clusters in NGC~5253 that have ages around 10~Myr.  The cluster populations in NGC~3077 and NGC~5253 suggest that star formation has proceeded in discrete, highly-correlated regions of activity producing tight groups of clusters which dissipate their parent molecular cloud on a timescale of a few~Myr.  However, there appears to be a more diffuse component as well, since some of the very young clusters are isolated.  The clusters themselves are probably not bound, and will likely dissipate on a short timescale (see Section~\\ref{sec:dissolve}).  In Figure~\\ref{fig:mass-age}, we plot our best estimate of each cluster's age against its best mass estimate.  The curves in Figure~\\ref{fig:mass-age} represent the mass corresponding to the 90\\% completeness limit as a function of cluster age.  To construct this curve, we first determine the 90\\% completeness magnitude, $m_{90}$, in each band using artificial object tests.  We then scale the mass of each model cluster along the Starburst99 track until its magnitude is $m_{90}$.  The artificial objects are generated using the best-fit cluster profiles which we used to subtract objects for our series of neighbor-cleaned images (see Section~\\ref{sec:analysis}).  Artificial clusters are distributed randomly in each image, but only within a central subregion where most of the real clusters are found.  The completeness curves indicate that evolutionary fading biases our sample against old, low-mass clusters: roughly half of the young (age $<20$~Myr) clusters would be unobservable if they were 100-200~Myr old.  When the possiblility of mass-loss by evaporation is considered, the bias with age becomes even more significant.  Figure~\\ref{fig:mass-age} also indicates that while NGC~3077 contains many more clusters than NGC~5253, the distribution of cluster masses appears to be similar in both galaxies.  We plot the distribution of cluster ages in Figure~\\ref{fig:agehist}. The shaded histograms are the ages of clusters brighter than M$_V=-9$~mag.  Both galaxies show a strong peak of clusters with ages between 1 and 20~Myr, but the peak in NGC~3077 is much sharper.  It has around 20 clusters aged 5--7~Myr, while NGC~5253 has only a few clusters in this age range.  NGC~3077 has a number of clusters older than 20~Myr, but there are no clusters in this age range in NGC~5253.   \\subsection{The Interstellar Medium}\\label{ism}  NGC~3077 and NGC~5253 exhibit some interesting contrasts in the relationship between their cluster populations and their interstellar gas and dust content.  NGC~3077 appears to be a dustier galaxy overall, as can be seen in Figures~\\ref{fig:color}, \\ref{fig:ebvhist} and \\ref{fig:ebvmap} (the two regions of apparently-zero extinction at the center of NGC~3077 in Figure~\\ref{fig:ebvmap} are actually extremely extincted; the $H\\beta$ flux is too low in these regions to obtain an extinction estimate).  Furthermore, Figure~\\ref{fig:ebvmap} suggests that the clusters in NGC~3077 are more likely to be found in substantially extincted regions, compared to the clusters in NGC~5253.  The distribution of $H\\alpha$-emitting gas in these galaxies (Figure~\\ref{fig:halpha}) reveals more contrasts between these galaxies.  The $H\\alpha$ gas in NGC~3077 exhibits a shell/bubble morphology \\citep{mar98}, and the largest peak in its $H\\alpha$ is adjacent to (but not coincident with) the location of a grouping of several young clusters.  NGC~5253 has much more $H\\alpha$ emission overall (note the different greyscales used in the panels of Figure~\\ref{fig:halpha}), and it is strongly concentrated into a single, dominant clump of $H\\alpha$ gas.  The super star cluster candidate of \\cite{cal97b} lies at the very center of this $H\\alpha$ concentration (object \\#1 in Table~\\ref{tab:n5253clusters}).  The larger $H\\alpha$ flux from NGC~5253 suggests a larger population of very young stars.  Yet we actually observe a {\\it smaller} population of young clusters in NGC~5253.  Since NGC~5253 contains many more objects which may be individual stars in a diffuse field component, one might expect that a substantial portion of the $H\\alpha$ may have been ionized by this diffuse population.  However, \\cite{tre01} concluded from long-slit UV STIS spectra that the intercluster field in NGC~5253 is {\\it devoid} of the massive O stars that are largely responible for photoionization.  NGC~5253 contains several $H\\alpha$ ``hotspots'', peaks that have no optical counterpart whatsoever.  These hotspots are likely extremely young clusters that have yet to emerge from the molecular cloud complexes from which they formed.  These embedded young clusters may explain the apparent overabundance of $H\\alpha$ in NGC~5253.   \\subsection{Cluster Dissolution Timescales and the Diffuse Field Population}\\label{sec:dissolve}  Determining cluster dissolution timescales and studying the field stellar populations of starburst galaxies can provide constraints on important questions regarding the origin of field populations.  Are the field populations of galaxies composed of dissolved clusters, or did they form {\\it in situ}, from a diffuse field mode of star formation?  Is the origin of the field population different for different galaxies?  In NGC~5253, we observe a complete lack of clusters older than $\\sim20$~Myr, while NGC~3077 contains many such clusters, and several which are older than 100~Myr.  This contrast cannot be entirely understood as a difference in the completeness rates in the two galaxies.  Applying NGC~5253's completeness curve to NGC~3077's cluster population in Figure~\\ref{fig:mass-age} results in seven clusters older than 10~Myr becoming unobservable, which is less than half the total number of its clusters in this age range.  Does the lack of older clusters in NGC~5253 indicate that cluster dissolution is more efficent in NGC~5253, or simply that the star-formation episode in NGC~5253 began only $\\sim20$~Myr ago?  We examine the dynamic environments of the two galaxies to address this question.  As \\cite{tre01} already determined, evaporation timescales for clusters in the center of NGC~5253 are around 20~Myr, using the figures of \\cite{kml99} for the Milky Way rescaled to the case of NGC~5253. Although NGC~5253 is about two orders of magnitudes less massive than the Milky Way, the evaporation timescales for stellar clusters within $\\sim$100-150~pc from its center are only $\\sim$4~times longer than for our Galaxy under the same conditions. The main reason resides in the large velocity dispersion measured by \\cite{cp89}, $\\sigma\\sim$50~km~s$^{-1}$, in the center of the dwarf galaxy, much larger than for instance, its central rotation velocity (a few~km~s$^{-1}$).  By applying the virial theorem to this central dispersion velocity, we obtain an estimate for the mass contained within the central 100~pc of NGC~5253, which is only $\\sim$15~times smaller than the mass contained within the same radius in the Milky Way \\citep{kml99}.  This suprisingly large central mass concentration leads to small predicted evaporation timescales for the star clusters in the center of NGC~5253. For the range of cluster masses we detect, the evaporation timescales are in the range 16--50~Myr. These values may be uncertain by as much as a factor of two, because of the uncertainty in the techniques used to model the evolution of clusters \\citep[\\eg\\ ][]{por02}. In addition, mass segregation can contribute to accelerate tidal evaporation of clusters \\citep{kml99}.  Combining the cluster dissolution information with the lack of massive stellar populations in the intercluster field of NGC~5253 \\citep{tre01} suggests that most stars in NGC~5253 form in star clusters, which then dissolve into the field population on a timescale of $\\sim20$~Myr.  The two large clusters in NGC~5253 which are $\\sim10$~Myr old are interesting in this context; perhaps these objects are massive enough to survive the harsh dynamical environment of NGC~5253's core.  NGC~3077 appears to have a comparable rotation velocity to that of NGC~5253, but a much smaller central dispersion velocity; direct measurements from broadening of stellar features are not available for this galaxy; the small dispersion is inferred from molecular clouds velocities \\citep{wal02}. If these velocities are representative, the central dispersion in this galaxy is $<$20~km~s$^{-1}$, or about a factor 2.5 smaller than in NGC~5253. Because of the strong tidal fields in the environments of the cluster populations, we associate the evaporation timescale of star clusters with the dynamical time of the star cluster's orbit in the host galaxy \\citep[see Eq. 2 in\\ ][]{kml99}.  Under this assumption, the difference in central velocity dispersions implies that cluster evaporation timescales are 2.5 times longer in NGC~3077, compared to NGC~5253, which can perhaps account for the older clusters present in NGC~3077.   \\subsection{Comparing to the Starburst in M~83}\\label{sec:m83}  It is instructive to compare the nature of the starbursts in the dwarf galaxies NGC~3077 and NGC~5253 to the nuclear starburst of the grand-design spiral galaxy, M~83.  While the three starburst regions contain similar total numbers of clusters, the M~83 starburst has many more bright clusters (33 clusters with $M_V < -9$~mag); than NGC~3077 or NGC~5253, which have 12 and 11, respectively.  Since M~83 is 5--10 times more massive than these dwarf galaxies, it is perhaps not suprising that its starburst contains a proportionately larger number of massive star clusters.  However, while the mass distribution is more top-heavy in M~83, we note that the mass of the largest star clusters in all three galaxies is about the same: $\\sim10^5 M_\\odot$ (see Figure~\\ref{fig:mass-age} and Figure~11 in \\citet{har01}).  Furthermore, we find no correlation between the physical size of the starburst region and the size of the galaxy.  In all three galaxies, the starburst regions (as defined by the presence of young clusters) span $\\sim300$~pc.  In M~83, this amounts to only the innermost few percent of the galaxy's total radial extent; while in NGC~3077 and NGC~5253, the starburst regions encompass a much larger fraction of each galaxy's optical extent.  This contrast in the starburst's significance to the whole galaxy may point to a fundamental difference in the nature of the starbursts.  In M~83 it is likely that gas and dust are funneled into the nuclear region by the galaxy's bar instability, creating a localized region overdense with prestellar material, and ripe for intense star formation. However, these dynamical conditions do not exist in the dwarf galaxies NGC~3077 and NGC~5253, so the conditions leading them to starburst may be fundamentally different.  Another interesting contrast between M~83 and these dwarf galaxies is in the spatial distribution of their youngest clusters.  In \\citet{har01}, we noted the presence of an ``active ringlet''; a semicircular annulus centered on the optical nucleus of the galaxy. This ringlet contains the bulk of the young clusters in M~83's starburst, and also appears to be a hot cavity of ionized interstellar material, surrounded by denser gas.  The young clusters span a very narrow range in age (5--10~Myr), and appear to be more-or-less uniformly distributed throughout the ringlet.  There is some evidence that a handful of even younger clusters are distributed along the rim of the cavity, perhaps partially embedded in the surrounding denser gas.  The history implied by this evidence is of a highly-correlated star formation event that formed most of the clusters in the ringlet at about the same time.  In NGC~3077 and NGC~5253, on the other hand, we see starburst regions that are not so well-defined, and cluster populations spanning a slightly wider range of ages (1--20~Myr).  In addition, we see discrete groups of coeval star clusters, distinctly younger than the ages of the more diffusely distributed clusters.  The history implied is one where cluster formation is more stochastic in these dwarf galaxies; small groups of clusters form together in the densest regions of the interstellar material, which is then dispersed by the energy output of the clusters after only a few~Myr.     We present {\\it HST} photometry of stars and star clusters in the nearby dwarf starburst galaxies NGC~3077 and NGC~5253.  We use standard aperture photometry in our analysis, but the photometry is performed on images in which nearby neighbors have been subtracted using a best-fit profile model.  We separate star clusters from individual stars using the profile-shape information provided by the WFPC2 images.  We have measured photometry in the F300W, F547M, nd F814W bands for 55 star clusters in NGC~3077 and 33 star clusters in NGC~5253.  The photometry is extinction-corrected using the ratio of $H\\alpha$ to $H\\beta$.  By comparing the extinction-corrected cluster photometry to Starburst99 population synthesis models, we derive age and mass estimates for each cluster.  We also measure $EW(H\\alpha)$ toward each cluster, which provides an additional age constraint for young clusters.  Both galaxies have abundant populations of star clusters with ages less than 20~Myr, and masses between a few thousand and $10^5 M_\\odot$.  We discuss some interesting contrasts between these two cluster populations.  NGC~3077 has more clusters overall, and about half of them lie in the very narrow age range between 5 and 7~Myr.  These coeval clusters are distributed in a loose, fan-shaped region to the northeast of the central dust cloud in NGC~3077.  NGC~5253 has a large number of very young (age $<5$~Myr) star clusters, and about half of these are in one tight clump near the center of the galaxy. NGC~3077 has only five clusters younger than 5~Myr, and they don't appear to be spatially correlated.  However, the youngest cluster in NGC~3077 lies within the central dust cloud, and there is some evidence for other young objects nearby that have yet to fully emerge from the dusty core.  NGC~3077 contains several intermediate-age clusters, up to several hundred~Myr old (although all clusters older than $\\sim20$~Myr are relatively far from the galaxy's center).  In contrast, NGC~5253 contains no clusters older than $\\sim20$~Myr. This difference is possibly due to the different cluster dissolution timescales in these two galaxies.  We estimate that typical dissolution timescales are $\\sim2.5$ times longer in NGC~3077 than in NGC~5253.  The star cluster populations suggest that star formation has proceeded in discrete, highly-correlated clumps in these galaxies, and that this initial structure is dissipated on a short timescale (of order 20~Myr in NGC~5253) by the harsh dynamic environment of the galaxies' centers.  We compare the star clusters in these dwarf starburst galaxies to the clusters in the nuclear starburst of M~83, a giant spiral galaxy. The three cluster populations have similar age distributions, and they cover similar ranges in mass.  However, the clusters in M~83 have a mass distribution that is more top-heavy.  Furthermore, the spatial distribution of clusters in M~83 is much more structured than in either dwarf.  Most of its clusters lie within a distinct annular ring structure centered on M~83's nucleus.  This contrast may result from a fundamental difference in the conditions that lead these galaxies to starburst.  M~83's strong bar feature is likely funneling material into the nuclear region, creating a localized region rich in gas and dust; an environment ripe for vigorous star formation.  No such internal dynamics mechanism exists in the dwarf galaxies, so something else must have caused them to starburst.  External interactions seem like a likely candidate for NGC~3077, given its well-known interactions within the M~81 group.  However, NGC~5253 may be too isolated to have been recently triggered by an interaction, so the cause for its current starburst activity is less clear.   \\vskip 1in \\noindent Acknowledgements: This work has been supported by the NASA LTSA grant NAG5-9173 and by the NASA HST grant GO-9144.01-A."
    },
    "0311/astro-ph0311166_arXiv.txt": {
        "abstract": "The single-pulse properties of the millisecond radio pulsar PSR B1937+21 are studied in the 1410 MHz radio band. Aside from occasional ``giant pulses'' occurring in restricted regions of pulse phase, the emission appears to be remarkably stable, showing no pulse-to-pulse fluctuations other then those induced by propagation through the interstellar medium. This type of behavior has not been seen in any other pulsar although it was seen in previous 430 MHz observations of this source.  The stability of PSR B1937+21 can be understood in the context of the sparking gap model of radio pulsar emission. Given the emission properties of this source at 430 MHz, this model predicts that the emission at all higher frequencies will be just as stable. Since the stability depends on the outflow velocity of the emitting plasma, an upper bound may be placed on its Lorentz factor. ",
        "introduction": "This letter reports on observations of the radio pulsar PSR B1937+21 taken with a center frequency of 1410 MHz at the 305m Arecibo radio telescope in Puerto Rico. Previous observations of this source at 430 MHz revealed a behavior that was completely different from other known sources \\citep[]{jenal01}. At 430 MHz this source exhibits almost no detectable pulse-to-pulse fluctuations. Occasional bursts of radio radiation, or ``giant pulses'', are observed but they are restricted to small regions in pulse phase \\citep{kt00, cst+96}. The observations presented here show that this remarkable pulse-to-pulse stability of the non-giant pulse emission also occurs at 1410 MHz. \\citet{es03} attempted to measure the modulation properties of PSR B1937+21 at 1300 MHz with the Westerbork synthesis radio telescope. Unfortunately, the signal-to-noise ratio in their observations was not high enough to measure the pulse-to-pulse properties.  Observations of bright pulsars have shown that the shapes and intensities of individual pulses are unique, although they average together to form a stable mean profile. The characteristic widths of individual pulses, typically referred to as sub-pulses, are usually smaller than the average profile width. This typical behavior has not been detected in the pulses from PSR B1937+21. Instead, this source appears to be highly stable in the sense that each pulse is identical to the next.          It will be shown that the stability of this source at both frequencies can be understood within the context of the sparking gap model of pulsar emission originally proposed by \\citet{rs75} and later reworked in more detail by \\citet[][ hereafter GS00] {gs00}. In this model, the stability is a result of a large number of emission regions, or ``sparks'', occurring on the polar cap. The GS00 model predicts that if the emission is stable at a given frequency, then it will be stable at all higher frequencies. Since this stability will also depend on the velocity of the emitting plasma, an upper bound may be placed on its Lorentz factor.             The sparking-gap model of \\citet{rs75} and its revised form (GS00) have received much attention in recent years.  They have been used to interpret the sub-pulse properties of slow pulsar conal emission \\citep{esv03,rr03,vsr+03,ad01,dr01} and of millisecond pulsars core emission \\citep{es03}. These models have also been used in pulsar population studies \\citep{acc02,fcm01}. Given their widespread use, it is important to develop observational tests which can determine the validity of these models. Since the GS00 model predicts that the radio emission from PSR B1937+21 should be stable at all frequencies higher than 430 MHz, high frequency observations of this source will test the current form of the GS00 model.   The next section describes the 1410 MHz observations performed at the Arecibo radio observatory. It also describes the techniques used to measure the intrinsic fluctuation properties of this source. Section 3 discusses the sparking gap model in more detail with emphasis on the frequency dependence of the modulation index. This letter is summarized in section 4. ",
        "conclusions": "Radio observations of the millisecond pulsar PSR B1937+21 were taken at a center frequency of 1410 MHz. Statistical techniques were used to determine if this source exhibits any of the standard single pulse phenomenology that has been seen in other radio pulsars. Remarkably, none of the usual phenomena were detected. The data are consistent with each pulse having the same shape and intensity. Previously, \\citet{jenal01} showed that this source exhibits the same single pulse stability at 430 MHz. This stability is currently unique to this source. Hence, PSR B1937+21 maybe the prototype for a new class of radio pulsars. On the other hand, the sparking gap model provides a unified framework that can explain this source's behaviour as an extreme case of the normal pulsar emission process. In the context of the GS00 model, the stability is a result of a large number of sparks occurring on the polar cap together with a low Lorentz factor of the emitting plasma. Since the 430 MHz observations show no pulse-to-pulse fluctuations, the GS00 model predicts that all higher frequencies will be just as stable provided higher frequencies are emitted at lower altitudes. This result is independent of the exact physical model used to describe the vacuum gap. Future observations of this source at several frequencies above 430 MHz will further test the GS00 model. If fluctuations are seen, then the current form of the sparking gap model is incorrect. Observations of this source at frequencies below 430 MHz will enable a better measurement of the Lorentz factor. If pulse-to-pulse fluctuations are seen at lower frequencies, a lower bound may be placed on $\\gamma$ for a given model of the vacuum gap."
    },
    "0311/astro-ph0311216_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{intro}  The past few years has seen dramatic improvements in our ability to extract chemical information from molecular line observations. These advances have been driven by new observational platforms and techniques. In the following I will outline the various methods and uncertainties in computing chemical information from observations across the spectrum of star formation activity (from diffuse clouds, to star-forming cores, to molecular outflows). I will show how improved chemical knowledge can be used to place stronger constraints on theoretical astrochemical models, but also open new avenues in the use of molecular emission to examine the physics of molecular clouds and star formation. ",
        "conclusions": ""
    },
    "0311/astro-ph0311020_arXiv.txt": {
        "abstract": "We address the issue of accuracy in recovering density profiles from observations of rotation curves of galaxies. We ``observe'' and analyze our models in much the same way as observers do the real galaxies.  Our models include stellar disks, disks with bars, and small bulges. We find that the tilted ring model analysis produces an underestimate of the central rotational velocity. In some cases the galaxy halo density profile seems to have a flat core, while in reality it does not.  We identify three effects, which explain the systematic biases: (1) inclination (2), small bulge, and (3) bar. Inclination effects are due to finite thickness of disk, bar, or bulge.  Admixture of a non-rotating bulge component reduces the rotation velocity. A small (200-500~pc) bulge may be overlooked leading to systematic bias even on relatively large $\\sim 1~\\kpc$ distances.  In the case of a disk with a bar, the underestimate of the circular velocity is larger due to a combination of non-circular motions and random velocities.  The effect of the bar depends on the angle that the bar makes with the line of sight. Signatures of bars can be difficult to detect in the surface brightness profiles of the model galaxies. The variation of inclination angle and isophote position angle with radius are more reliable indicators of bar presence than the surface brightness profiles. The systematic biases in the central $\\sim 1\\kpc$ of galaxies are not large. Each effect separately gives typically a few \\kms error, but the effects add up. In some cases the error in circular velocity was a factor of two, but typically we get about 20 percent effect.  The result is the false inference that the density profile of the halo flattens in the central parts.  Our observations of  real galaxies show that for a large fraction of galaxies the velocity of gas rotation (as measured by emission lines) is very close to the rotation of stellar component (as measured by absorption lines). This implies that the systematic effects discussed in this paper are also applicable both for the stars and emission-line gas. ",
        "introduction": "The cosmological cold dark matter (CDM) model and its present variant with cosmological constant ($\\Lambda$CDM) have been very successful in accounting for the properties of galaxies and their distribution in space. A current problem for CDM lies in observations of dwarf and low surface brightness galaxies.  The rotation curves and derived density profiles of these galaxies appear to be in conflict with the CDM model \\citep{M94, FP94, Burkert95, HG98}.  The problem arises because numerical simulations of the formation of halos in the CDM model reveal steeply cusped density profiles while observations favor shallow (core) profiles. The theory predicts that the mass density in the inner parts of the dark matter halos is best described as an $r^{-\\alpha}$, $\\alpha=1$ power law \\citep{NFW97,Power2003,Colin2004} or possibly by a power-law with a slightly steeper $\\alpha=1.2-1.5$ slope \\citep{Moore1999, Ghigna2000, Klypin2001, Reed03, Tasitsiomi2004}. The observed rotation curves of a significant fraction of dwarf galaxies, however, suggest that their cores do not have steep density cusps \\citep{dBM97, BAC01, dBMBR01, MRdB01, Bolato2002, WdBW02, Simon2003}. Note that this is not true for some galaxies and that some observational results are still consistent with cosmological predictions \\citep{vdB99, SMT2000, vdBS01}.  \\citet{MRdB01} have analyzed H$\\alpha$ observations of 50 low surface brightness galaxies of which 36 are of sufficient quality that rotation curves can be derived. The spatial resolution is improved by an order of magnitude over old 21~cm data.  In a companion paper \\ \\citet{dBMBR01}\\ conclude that at small (0.5-1~kpc) radii, the mass density distribution is dominated by a nearly constant density core.  The distribution of inner slopes is strongly peaked around $\\alpha= -0.2$ with clear contradiction with the CDM predictions. Thus, there appears a controversy concerning the density distribution of matter in the inner parts of dwarf galaxies.  We propose a somewhat different approach to this problem which we believe can help to clarify these debates.  We construct realistic models of galaxies and ``observe'' them as real galaxies using the same techniques or algorithms that observers apply to real data.  We study the details of the analysis of observations and we find discrepancies between recovered and true properties of the simulated galaxies.  We investigate projection effects and the effects of bars and bulges on the observed rotation curves. We show that these effects are often far from negligible.  Some of the effects which we study, have already been discussed. \\citet{Lia1984} discussed the biases in the observed rotation curves of high surface brightness galaxies created by strong bars.  More recently, \\citet{KW02} have investigated the effects of a weak bar on mass determination in the central parts of spiral galaxies. They examine rotation curves of edge-on model galaxies observed from different viewing angles. The models had a two-dimensional solid rotating barred potential. Motion of gas in this potential was followed using an SPH code.  \\citet{KW02} find that the mass estimated from rotation curves is larger than the true mass by a factor of 5 for 15\\% of the viewing angles. Overestimation in mass occurs more frequently than underestimation.  \\citet{KW02} focus on rotation curves defined as the highest velocity envelope of position-velocity diagrams.  When the rotation curve is defined as the intensity weighted mean of position-velocity diagrams, the rotation velocity is underestimated for most of the view angles.  \\citet{KW02} conclusions are only for edge-on galaxies; the situation with less inclined galaxies could be more complex. Another important aspect studied by \\citet{WMN02} and \\citet{WK01} is the gravity driven turbulence induced by a weak bar in a clumpy and filamentary gaseous media. If was found that the velocity dispersion of a cold neutral gas as well as of an ionized component is significantly lager than naively expected if the distribution is clumpy \\citep{WMN02} or filamentary \\citep{WK01}. These results challenge the common expectation that the gas component in a galaxy is a cold medium with low velocity dispersion.   The effect of bulges on the observed kinematics of galaxies has its own history \\citep[e.g.,][]{Fillmore86, Kormedy89, Bertola1995, Beltran2001,Funes2002,Pizzella2004}.  An interesting issue here is the difference between the rotation of stars and rotation of gas in the central 1-2~kpc region.  Because our simulations have only the stellar component (no gas), it is important to know how different are rotation of gas and of stars. Naively, one expects that gas as a ``colder'' component should rotate faster.  The reality is more complicated. A detailed discussion of this issue is given in \\S 2.  We also use our own observations to measure rotation velocities in stellar and gas components. To summarize, like other workers in this field, we find that often in the central region the rotation of gas (as measured by $H\\alpha$ and $NII$ emission) is not very different from the rotation of the stellar component.   Some other systematic effects were considered by \\citet{Swaters2003} who discuss effects of seeing, finite slit width, and slit misalignment for ideal thin exponential disks. It is found that all effects systematically underestimate the central slope. The magnitude of the bias depends on the distance to the galaxy and on slit misalignment.  \\citet{dB03} considered the effects of an offset between optical and kinematical centers, as well as a very simplistic attempt to quantify the effects of non-circular motions.  Both effects tend to lower the estimated density profile slope. \\citet{Hayashi2003} discuss effects of triaxial halos on the rotation of gas in CDM simulations, which lead to non-circular motions.    This paper is organized as follows.  In \\S 2 we present results of our observations of rotation velocities as measured by gas and stars. In \\S 3 we discuss different observations, focusing on systematic biases, which may affect rotation curves.  We also present results of our observations of rotation of gas and stars. In \\S 4 we present and describe our models.  Simulations are discussed in \\S 5.  We have constructed four detailed galaxy models and run them as N-body simulations.  In \\S 6 we present our methods of analysis of the models and in \\S 7 we show our results. In \\S 8 we present our conclusions. ",
        "conclusions": "The rotation curves of dwarf and LSB galaxies arguably present the most challenging problem for current cosmological models. The key to this problem lies in reliably estimating the accuracy with which one can recover the mass distribution and, thus, the dark matter density using observed rotation curves. This is the main question which we address in this paper by studying realistic numerical models. The models mimic real galaxies in many respects. For example, they have thin ($\\sim 100~$pc) and cold ($\\sim 10~\\kms$) exponential disks. The latter condition is very important because many effects are related to the magnitude of random motions.  The main lesson to be learned from our analysis is that there are many systematic biases.  Each bias can be very small (only few km/s), yet they all add up to produce significant effects.  We find that the tilted ring analysis of spectroscopic data results in large 30\\% to 50\\% underestimates of the rotation velocity in the central 1~kpc region of the galaxy models. These  underestimates in velocity in turn result in the inference of slopes in the central density distributions of our models that are shallower than the true slopes.  A quantitative understanding of these biases is important if one is to accurately interpret rotation curves. If one ignores these biases, the analysis of rotation curves can produce misleading results. The extent to which this has occurred for real galaxies can only be resolved by detailed analyses of individual objects. One cannot be sure that the analysis of real observations data was done properly when the same analysis produces false results when applied to realistic models.  One of the difficulties in interpreting observational data lies in the fact that galaxies are treated in an overly simplified manner (e.g., cold gas inside thin exponential disks). Observational results (including those presented in this paper) suggest that this is far too naive a picture. In the central region, the stellar component often rotates as fast as the emission-line gas - in contradiction with the prediction of a naive picture of a cold gas embedded in a hot stellar component.  The presence small (200-500~pc) bulges and bars, which can be easily overlooked, further complicates the interpretation of the data."
    },
    "0311/astro-ph0311343.txt": {
        "abstract": "This review describes telescopes designed to study neutrinos from astrophysical sources. These sources include the Sun and Supernovae emitting neutrino energies up to tens of MeV, atmospheric neutrino sources caused by cosmic ray interactions and other sources generating neutrino energies ranging up to $1\\times 10^{20}$ eV. Measurements with these telescopes also provide information on neutrino properties including clear evidence for neutrino flavor change. Telescopes in operation in the past and present are described, along with plans for future instruments to expand this rapidly growing field of particle astrophysics. ",
        "introduction": "A new generation of astronomical instruments has evolved, starting with the radiochemical solar neutrino detector of Ray Davis and his collaborators~\\cite{davis1} more than 30 years ago. This generation of instruments is part of the developing field of particle astrophysics, where particles from astrophysical objects are used to provide information about those objects and about the properties of the particles themselves. Since astrophysical objects can provide intensities and energies far outstripping any terrestrial source, this area of study has become a significant one for the field of particle physics. Neutrinos have a number of favorable properties that make them attractive as a means to study astrophysical objects, so neutrino telescopes have significant value for astrophysics as well.  Neutrinos have no electric charge and are not sensitive to the strong interaction. Their mass and magnetic moments are very small and so they undergo very few interactions even in the densest matter. This makes them ideal candidates for the study of otherwise inaccessible astrophysical regions, such as the center of the Sun, or distant astrophysical sources without distortion of their trajectories by intergalactic magnetic fields or material. Neutrinos are known to have three active species (electron, muon, tau) associated with these three lepton generations. Experiments on the decay properties of the Z Weak Boson \\cite{PartData} have shown that there are three (2.981 $\\pm$ 0.008) such active species.  The energies of neutrinos detected by the telescopes described in this article include: from a few 100 keV to 30 MeV for solar and supernova neutrinos; 0.1 GeV to a few TeV for neutrinos generated by cosmic rays interacting with the atmosphere; and much higher energies, ranging up to $1\\times 10^{20}$ eV for the highest energy neutrino sources. The detection techniques differ substantially, but all detectors are situated in locations where interfering backgrounds are shielded to a substantial degree because the experiments tend to have very low neutrino detection rates as a result of the small interaction probability. The experiments seeking neutrinos below about 1000 GeV are typically in underground locations to range out cosmic ray background and the highest energy experiments are typically under ice or under water in order to combine effective shielding with very large detection volumes.  The sections of the paper include: 2. Solar Neutrino Detectors; 3. Supernova Neutrino Detectors; 4. Atmospheric Neutrino Detectors; 5. Very High Energy Neutrino Detectors. This subdivision of detector groups is somewhat arbitrary since there is substantial overlap in the capabilities of the detectors, and all of them have some degree of sensitivity for supernova neutrino detection. However, this choice has some advantages for description, as the detectors are described in the section covering the primary scientific motivation for the detector and reference to additional capabilities may be mentioned further in another section. This paper is mainly intended to describe the principal features of the detectors themselves, but in each case the physics motivation for the detectors is presented, along with the detector scientific capability and a succinct summary of the results obtained to date or anticipated for the future.  There are many astrophysical motivations for the detection of neutrinos with neutrino telescopes, as well as opportunities to determine particle physics properties of neutrinos themselves through the use of astrophysical sources. The complementary nature of this information has made neutrino telescopes a central part of the rapidly developing field of Particle Astrophysics.  Some of the principal astrophysical motivations are as follows. The neutrinos produced in the Sun arise from the principal nuclear reactions powering the Sun. Detailed models of the Sun have been developed that match other measured solar properties very well and predict a spectrum of neutrinos from a number of the nuclear reactions. The study of neutrinos from the Sun over the full range of energies provides a very detailed test of solar models, including otherwise inaccessible information from the core region.  Atmospheric neutrinos arise primarily from the decay of muons, pions or kaons produced by energetic cosmic rays interacting with the atmosphere. Measurements of neutrino fluxes over a broad range of energies can be combined with other sources of information to obtain an understanding of cosmic rays from astrophysical sources. The highest energy neutrinos arise from high energy cosmic rays, protons or nuclei. One of the objectives of neutrino detection is to determine the origins and acceleration mechanisms of these very high energy charged particles. Neutrinos from Supernovae include electron neutrinos arising from the conversion of neutrons to protons during the initial infall as well as electron, mu and tau neutrinos and anti-neutrinos produced during the expansion phase following the point of maximum compression. Neutrinos carry away the majority of the energy from a supernova and can provide very detailed information on the processes involved in these spectacular astronomical events.  Neutrinos from astrophysical sources can also be used to study detailed properties of neutrinos and results from these measurements have provided very valuable information, including the first clear indications of neutrino flavor change (i.e. the change of a neutrino of one family into one of another family) and information on neutrino mass differences. Some of the particle physics information that has been obtained from astrophysical neutrinos is as follows.  In the case of neutrino mixing, the neutrino flavor fields $\\nu_{{\\ell}}$ are superpositions of the components $\\nu_{k}$ of the fields of neutrinos with definite masses $m_k$: \\begin{equation} \\nu_{{\\ell}} = \\sum_{k=1}^{3} U_{{\\ell}k} \\, \\nu_{k} \\qquad (\\ell=e,\\mu,\\tau) \\,, \\label{mixing} \\end{equation}  \\noindent where $U$ is the 3 x 3 unitary Maki-Nakagawa-Sakata-Pontecorvo (MNSP) mixing matrix \\cite{MNSP}.  When the neutrinos travel in a vacuum or low density region, the linear combinations are changed so that the detected neutrino types appear to \u0093oscillate\u0094 among the neutrino flavors. If the masses are not strongly degenerate, then the 3 x 3 MNSP matrix can be separated into 2 x 2 matrices affecting oscillation among two masses almost independently of the third. In this case there is one $\\Delta m^2$ such as $m_2^2-m_1^2$ for the masses involved and the mixing matrix \\begin{equation}\\label{vacangle} U = \\left( \\begin{array}{rr} \\cos \\theta & \\sin \\theta \\\\ -\\sin \\theta & \\cos \\theta \\end{array} \\right) \\end{equation}  For this simplest case where two mass eigenstates dominate the process, the following probability is predicted for subsequent detection of a given neutrino type after it has traveled for a distance L in vacuum:  \\begin{displaymath} P = 1 - sin^2(2\\theta) sin^2(1.27 \\Delta m^2 L/E) %P = 1 - 1/2 sin^2 (2 \\theta_{12}) (1 - cos (2.54 \\Delta m^2 L/E) \\end{displaymath} \\noindent where $\\Delta m^2$ is in eV$^{2}$, $L$ is the source-detector distance in meters, $E$ is the neutrino energy in MeV and $\\theta$ is a mixing angle. In order for a significant fraction of neutrinos to change their flavor, $sin^2 (2 \\theta)$ must be reasonably large and $\\Delta m^2 L/E$ must be appropriate for the source and detector involved. This process has been used to interpret atmospheric neutrino data observed by the Super-Kamiokande detector as described in Section 4 wherein muon neutrinos appear to oscillate predominantly to tau neutrinos as they traverse the Earth. These measurements can be used to define elements of the MNSP matrix, in particular, $\\Delta m_{23}^2$ and $sin^2 \\theta_{23}$.  Additional effects can occur when neutrinos pass through regions of dense matter. These effects are associated with the fact that electron neutrinos can have interactions with electrons through W exchange in addition to the Z exchange occurring for all flavors. These effects can result in \u0093matter enhancement\u0093 of the oscillation process \\cite{MSW}, the so-called MSW (Mikheyev-Smirnov-Wolfenstein) effect. This enhancement greatly broadens the region of $\\Delta m^2_{12}$ and $sin^2 \\theta_{12}$ where significant effects can occur and introduces an additional dependence on the neutrino energy.  The Sudbury Neutrino Observatory (SNO) has provided evidence \\cite{SNO} for neutrino flavor change by observing clear differences between one reaction sensitive to only electron neutrinos and another sensitive equally to all neutrino types. The results show that about 2/3 of the electron neutrinos produced in the Sun appear to have transformed to other types before detection on Earth. A similar conclusion is reached by comparing the SNO results with the measurements of solar neutrinos by Super-Kamiokande \\cite{SK}, using a reaction with a small sensitivity to other neutrino types in addition to electron neutrinos. These results, together with a number of other measurements of solar neutrinos with different energy thresholds, can be interpreted in terms of neutrino oscillations, providing additional information about elements of the MNSP matrix, particularly $\\Delta m_{12}^2$ and $sin^2\\theta_{12}$.  Many other particle physics properties of neutrinos can be defined by measurements with neutrino telescopes, including limits on neutrino decay lifetimes and magnetic moments. Future experiments will also improve significantly on present neutrino mass and mixing parameters, particularly when combined with measurements from terrestrial sources such as reactors and accelerators.  The improved definition of neutrino properties through these measurements from astrophysical sources can, in turn, be used to improve our astrophysical information. For example, with the oscillation properties of neutrinos defined, it is possible to understand the original fluxes of solar electron neutrinos as a function of energy. This provides a very accurate test of solar models and adds to our understanding of the astrophysics of the Sun. ",
        "conclusions": ""
    },
    "0311/astro-ph0311570_arXiv.txt": {
        "abstract": "Most of the proposed associations between magnetars and supernova remnant suffer from age problems. Usually, supernova remnants ages are determined from an approximation of the Sedov-Taylor phase relation between radius and age, for a fixed energy of the explosion $\\sim 10^{51}$ erg. Those ages do not generally agree with the characteristic ages of the (proposed) associated magnetars. We show quantitatively that, by taking into account the energy injected on the supernova remnant by magnetar spin-down, a faster expansion results, improving matches between characteristic ages and supernova remnants ages. However, the magnetar velocities inferred from observations would inviabilize some associations. Since characteristic ages may not be good age estimators, their influence on the likelihood of the association may not be as important.  In this work we present simple numerical simulations of supernova remnants expansion with internal magnetars, and apply it to the observed objects. A short initial spin period, thought to be important for the very generation of the magnetic field, is also relevant for the modified expansion of the remnant. We next analyze all proposed associations case-by-case, addressing the likelyhood of each one, according to this perspective. We consider a larger explosion energy and reasses the characteristic age issue, and conclude that about 50\\% of the associations can be true ones, provided SGRs and AXPs are magnetars. ",
        "introduction": "Soft Gamma Repeaters (SGRs) and Anomalous X-ray Pulsars (AXPs) are two classes of objects being candidates to magnetars, neutron stars with magnetic fields above the quantum critical threshold $B_c \\, = \\, 4.41 \\times 10^{13}$ G. In the model proposed by \\citet{mag}, which we will follow through this work, dynamo action in a fast rotating proto-neutron star amplifies the ``seed'' magnetic field ($10^{11} - 10^{13}$ G) by one to three orders of magnitude. The simplest estimate of the magnetic field intensity is derived from measurements of their rotational periods ($P$) and period derivatives ($\\dot P$), through the well known expression  \\begin{equation} B \\sin \\chi \\, = \\, {\\left( {{3 c^3 I P \\dot P} \\over {8 \\pi R_{psr}^6}}\\right)}^{1/2} \\; , \\end{equation}  \\noindent where $B$ is the magnetic field intensity, $R_{psr}$ is the pulsar radius ($\\sim 10^6$ cm), $\\chi$ is the angle between the rotational and magnetic field axis, $I$ is the moment of inertia ($\\sim 10^{45}$ g cm$^2$), and $c$ is the speed of light. Their ages can also be estimated through the characteristic age  \\begin{equation} \\tau \\, = \\, {P \\over {2 \\dot P}} \\; , \\end{equation}  \\noindent but both calculations assume continuous spin-down driven by magnetic dipole braking, an assumption which has been questioned by several authors \\citep{nn,venrel,PTH95,CHN00,kou99}, in different contexts, regarding AXPs and SGRs. The derived figures are shown in Table 1. As most of the objects are younger than 10$^4$ years, it seems reasonable to look for the  Supernova Remnant (SNR) that was originated at the same supernova explosion. In fact, almost all SGRs and AXPs were tentatively associated to some SNR. Though, the likelihood of those associations is put in doubt because the age estimated for those SNRs is generally larger than the characteristic age of the associated magnetar.  This situation leads to the ``age'' and ``velocity'' problems of SGRs and AXPs. Both problems are related to the relatively low characteristic ages found for these objects. The SNRs usually associated to them seem to be systematically older, thus the age problem. Some SGRs and AXPs have been found to lie beyond, at or near the border of their proposed SNRs, and, given their low ages, the implied velocities are much higher than those found for ordinary pulsars (see \\citealp{sgraxp} for a discussion about these objects). As seen on Tables 1 and 2, there is a large uncertainty in those estimates, mainly due to uncertainties in the distances. This implies uncertainties in the radii on which age estimates are based through models of the expansion.  Several authors have addressed magnetar-SNR associations. Some of them have discarded either the magnetar hypothesis or the association, based on those problems. This became known as the ``nature versus nurture'' debate, where one side proposes that the characteristics of SGRs and AXPs are derived from their special nature (see \\citealt{axp}), while the other side finds more plausible that the environment is the actual responsible, nurturing the otherwise ordinary neutron stars with a fossil accretion disk \\citep{nn,CHN00,PTH95}.  Associations between pulsars and SNRs are generally examined according to the following criteria: 1) positional coincidence in sky; 2) distance estimates; 3) age estimates; 4) evidences for interaction between the neutron star and the SNR. Additional criteria may be derived from these, like the estimate of the projected velocity on sky of the compact object. \\citet{gae01} did a comprehensive study of positional coincidences for magnetar-SNR associations. Distance can be determined by neutral hydrogen column density measurements, evidence of association with another object (e. g., HII region or molecular cloud) whose distance is already known, or through surface brightness-diameter ($\\Sigma -D$) relations. This last method is not considered reliable enough, because it ignores the effect of density and structure of the local interstellar medium, resulting in a large dispersion of values from the best fit, as can be seen in \\citet{sigd}. Some SNRs on the proposed associations show plerions, usually taken as positive signal of the presence of a neutron star, but most do not present that feature. Finally, measurements of HI column density are subject to systematic errors from models, and usually are not better than 50\\%, implying on similar uncertainties in distance determinations. Radio pulsars can have their distance estimated by use of dispersion measure, however SGRs and AXPs have not yet been detected in radio waves. There are reasons to believe they cannot emmit radio signals, either because they are magnetars, where the huge magnetic field induces photon splitting rather than pair creation (\\citealp{split}, but see \\citealt{Bgrande}), or in the case they are accretors, because accretion quenches radio emission \\citep{CHN00,PTH95}.  SNR age estimates are usually derived from the knowledge of the radius by means of an approximate version of the expansive blast-wave evolution equations. We shall see in next Section that those approximations which ignore the previous phase(s) can introduce an apreciable error in cases near phase transition, and such approximations imply a reduction of the number of variables to two by linking the explosion energy, the local interstellar medium density and the ejected mass. The present local density of the interstellar medium can be estimated within 50\\% uncertainty, though this may not be representative of the structure of the interstellar medium prior to the supernova explosion. for example, an explosion which occurred inside a HII region, (low density medium) and have run into a nearby molecular cloud (high density medium) would be misleading, since the observer would be able to estimate the molecular cloud density only. Even inside the HII region it would be likely that density is not constant. Situations like those can be modelled (see a good review on this in the work of \\citealt{TM}), but the point is that the unknown structure of the pre-supernova environment introduces one more possible systematic error on age estimates. The ejected mass can be estimated from the observation of how much Fe was produced, coupled to models of supernova yields, when the SNR is young. Explosion energy is usually considered to be within a factor of 2 of the canonical value of 10$^{51}$ erg, due to previous both observational and theoretical knowledge. The review of \\citet{SN} considers the energy range $0.5 - 5.5 \\; \\times 10^{51}$ erg for classical SNII, from a sample with a dozen well-observed events. However, these estimates also contribute with their own uncertainties, because the type of explosion event giving rise to a SGR-AXP is completely unknown. As the radius picks the same relative uncertainty as the distance, the combined uncertainty on the age estimate will not be smaller than a factor of 2, and probably much larger than that. We must recall that it is not uncommon to be unsure about the phase a given SNR is in, and that near the transition time between phases approximations do not strictly hold. More accurate relations can be used, reducing systematic errors, though usually most authors feel the rather small gain in precision is not worth the trouble.  Blast-wave velocity estimates could also be used to determinate age and radius, through relations similar to the ones for radius and age. Knowledge of both radius and velocity would determine the phase of the expansion, removing a source of uncertainties. Unfortunately, such estimates must come from X-rays measurements difficult to be made, and up to now they could not be obtained for any of the proposed associations.  Finally there are evidences of interaction between the neutron star and the SNR in some of the proposed associations, which show plerions (or filled-centre morphology). Additionally, one of the associations was reported as presenting a jet-like feature \\citep{sgr}.  The overall reliability of the described estimates is not good, although several papers have addressed the association issue, and drawn conclusions about AXP and SGR nature. In this work, we will show that at least an important factor for those estimates is still missing, namely the injection of energy by the internal magnetar, and we analyze the proposed associations bearing that factor in mind. The next Section is dedicated to SNR expansion considerations, Section 3 introduces the effects of the injected energy on the SNR, while Section 4 deals with an application of those considerations to the observed sample. In Section 5 we discuss our results. ",
        "conclusions": "We have analyzed in this work the proposed SGR-AXP/SNR associations. Previous analisis of the associations arrived at different results. According to \\citet{nn}, all associations can be considered likely. \\citet{gae01} contend that only AXP J1845-0258/G29.6+0.1, AXP 1841-045/G27.4+0.0 and AXP 2259+586/CTB 109 could be valid. \\citet{ank} considered SGR 1806-20/G10.0-0.3 and SGR 0526-66/N 49 as plausible, in addition to the 3 already mentioned by \\citet{gae01}.  It is expected that if magnetars exist and are born in supernova explosions, they can inject enough energy to enhance the expansion of their associated SNRs. If AXPs and SGRs are indeed magnetars, their associated SNRs should appear older than the ages derived from standard expansion models. We analyzed all the proposed associations, and have come to the following results:  1) If the characteristic age of the neutron star is regarded as the true age of the association, then SGRs and AXPs may not be magnetars at all (or the model of \\citet{mag} is not correct regarding magnetar origin), for only one case (SGR 1806-20/G10.0-0.3) has shown good agreement within the model, and even the true nature of G10.0-0.3 was put in doubt;  2) If characteristic ages are in turn ignored (see \\citealt{venrel,gae01,nn,kes75}, among others), then two SGRs (out of 5) and three AXPs (out of 5 with proposed associations) seem to be associated within our model, so about 50\\% of general agreement. It should be noted that two associations which agreed within our model (AXP J1845-0258/G29.6+0.1 and AXP 2259+586/CTB 109) were already believed to be true by previous works \\citep{gae01,nn,ank}, meaning that uncertainties in distance are large enough to allow for both possible scenarios (standard and with energy injection);  3) AXP 1841-045/G27.4+0.0 and SGR 0526-66/N 49 can only be considered as true associations if the magnetars were born with $P_0 \\geq 6$ ms, because in that way the energy injected would be insufficient to directly affect SNR expansion.  These results from our model are tied to the dynamical evolution of magnetars. While we have assumed for simplicity the standard magnetic dipole braking with braking index equal to 3, several proposals have been done that argue for different braking models: fossil or fallback accretion disks \\citep{nn,CHN00,PTH95}, episodes of relativistic wind emission \\citep{venrel}, a different constant braking index \\citep{china}, magnetic field decrease \\citep{mag,mon}. As the measurements of braking indices for 5 young pulsars revealed that all but one have significant departures from the canonical value, it should not come as a surprise that magnetars do not follow the standard model of spindown. The main effects of alternative models are to change the estimates of magnetic field strength and spindown age. Nonetheless, the influence in our model would be very small, since $\\tau_0$ can be increased up to one month without appreciable modifications on the model, and rotational energy is not dependent on spindown models. That is why we feel justified to leave item (1) above aside of the discussion. Spindown models that allow for ages significantly smaller or larger than the conventional characteristic age are required to explain certain associations, if they are true ones, either within our model, or considering the standard scenario.  Supposing that our results from item (2) above represent the actual situation, we could ask where are the SNR that are associated to the other SGRs and AXPs. The answer can be that they were not detected yet, since they lie generally at regions of coexistence of HII regions, other SNRs, and other types of objects (variable stars, young stars, molecular clouds, etc.). Including AXP 0142+614 and AXP J0720-312 in the sample, we would have slightly less than 50\\% of the true associations identified. On the other hand, few of the latter can be considered as firm, due to several other factors. For example, the angular size of SNRs are often quoted without uncertainties, which certainly is an understatement, given the dificulty to recognize a SNR and to assess its shape and size.  It is worthy to remark that a similar analysis, ignoring the effects of the injection of energy, and desconsidering conventional characteristic ages, can be done to the same associations, and will result in different age and velocity ranges than those obtained here. Indeed, even associations considered unlikely by us may turn out to be likely. This is because the standard energy SNRs would take more time to reach the observed sizes, allowing for higher ages for the associations than in our model, and thus producing lower velocities than those we found. Nevertheless, that scenario requires non-standard spindown and no injection of energy by the neutron star, be it a magnetar or not.  It is interesting to check that the most important factor to decide on the plausibility of associations is $\\beta$. As $\\beta$ increases, associations are considered increasingly unlikely. However, \\citet{gvar} points to the possibility that a SNR expanding in a region with anisotropic interstellar medium densities will be distorted and/or expands faster in one or more directions. This means that the geometrical center of the SNR can be displaced from the actual explosion site. Although it is difficult to take this effect into account quantitatively, one should be cautious to do not dismiss it entirely. It is possible that one or more of the proposed associations are affected by this effect, which can increase or decrease $\\beta$ randomly. That effect will be more accentuated for SNRs that have been expanding in low-density regions (``bubbles'') amidst high-density ``walls'', and it increases with age. The first consideration is the case for the most massive stars, which do not live long enough to move far away from the sites where they had been born, while their stellars winds create a low-density cavity. The second consideration means we must expect $\\beta$ estimates to be a little bit more scattered as age (and radius) increases.  Besides studying associations including AXPs and SGRs, we also examined the association PSR J1846-0258/G29.7-0.3 \\citep{gvk75,kes75}, which also suffers from the age problem, since it has the smallest characteristic age known among pulsars (723 years), and the SNR age was estimated as at least 1800 years. The magnetic field of this pulsar is $\\sim 5 \\times 10^{13}$ G, considering magnetic braking spindown, slightly above the quantum critical field, and by so marginally qualifying as a magnetar. It was not detected at radio wavelenghts, only in X-rays. The association is considered very likely, as the pulsar is located at the geometrical center of the SNR, coincident with a radio/X-ray nebula, probably powered by the pulsar. Proceeding as in the other cases investigated in the previous Sections, we see from Fig. \\ref{f-known} that the range of radius quoted on the literature is nearly coincident with the range allowed by our two extreme cases, considering the characteristic age. This way, there are no preferences for high or low densities or ejected mass. As before, we will ignore the characteristic age, to found that the age range allowed is 250-700 years, in case of high $M_{ej}$ or $n$, and 650-6000 years, in case of low $M_{ej}$ or $n$. Intermediate values of the parameters $M_{ej}$ and $n$ would provide intermediate ranges of ages. The placement of the neutron star at the geometrical center of the SNR imply in low velocities or alignment between the velocity vector and the line of sight. It is important to notice that \\citet{kes75} find a braking index $\\sim 1,9$ and age $\\sim 1700$ years for this pulsar, in which case the magnetic field can be under the quantum critical one. Nonetheless, this association represents one more evidence against the consideration of the characteristic age (without braking index information) as a good age estimate.  We have left for a future work (Allen \\& Horvath, in preparation) the study of an alternative origin scenario for magnetars, the collapse of a white dwarf star, induced by accretion or merging from a binary companion. Simulations for this scenario \\citep{aic} reveal that $\\sim 0.1$ $M_{\\sun}$ can be ejected, with an explosion energy of $10^{50}$ erg, implying on higher initial velocities.  Finally we would like to point that magnetars lose very little rotational energy through gravitational waves, when compared to typical pulsars, unless r-modes can play an important role, which is not clear to us. Statistical studies on pulsar gravitational wave detectability, like the one performed by \\citet{pach}, shall not be affected by this consideration, at least for the current detectors, like VIRGO, because of the scarcity of magnetars."
    },
    "0311/astro-ph0311436_arXiv.txt": {
        "abstract": "After 40 years we have no `low' theory of quasars to predict the atomic absorption and emission features so abundant in quasar spectra. Here I take the first step, using results selected for their diagnostic power to build an `observational paradigm' unifying all the atomic features in quasars. This paradigm bears a remarkable resemblance to the one I proposed in Elvis (2000)! Yet all the results used here were not used in that earlier study, with one exception. The Elvis (2000) model has been tested by the new data and has survived, and so has been strengthened. The structure is readily tested, is physically suggestive, and opens up many areas for detailed modelling.  From this I believe a truly predictive `low' theory of quasars will emerge. ",
        "introduction": "\\vspace{-.3cm}  Enlightenment does not scale with published mass.  The quasar/AGN literature has over 12,000 refereed published papers, and papers now appear at a rate of about 2/day. Yet few would contend that we have a much deeper understanding of quasars today than we did 30 years ago.  By 1974, just over 10 years after the discovery of quasars, all of the current model for a quasar were in place: a central massive black hole, an accretion disk and a relativistic jet. We have made great observational progress mapping out quasar properties at all wavelengths, and in the Unified Scheme (Padovani \\& Urry 1999) have resolved much of the confusing properties of quasars to be due to the effects of obscuration.  But most quasar observations are of atomic features.  The black hole/disk/jet theory is a `high' theory, dealing only with overall energetics, and describes a naked quasar, devoid of any of the veiling gas that creates these atomic features. We need a `low' theory with sufficient detail to predict the emission and absorption phenomenology of quasars. But to get to such a theory in one step from 10$^4$ papers is not realistic. First we need to build an `observational paradigm' to unify the phenomenology into a coherent, over-constrained, and so robust, structure of what can accurately be called the {\\bf quasar atmosphere}.  \\vspace{-0.2cm} ",
        "conclusions": "\\vspace{-0.2cm}  I have presented a structure for quasar atmospheres. This structure is simple, and increasingly well supported by the data. The structure explains a variety of, once unrelated, quasar observations and puts them into a self-consistent physical picture.  This picture is physically suggestive, and gives us the prospect not only of understanding quasar winds and atmospheres, but of working back to the physics of accretion, and out to the impact of quasar winds on their larger scale environment. If the view presented here continues to fit the data, then we have the beginnings of a low theory of quasars, starting with their atmospheres. Quasar studies can then enter a stage in which they connect with detailed physics and so link quasars to the rest of astrophysics.  \\vspace{-0.5cm}"
    },
    "0311/hep-ph0311103_arXiv.txt": {
        "abstract": "The nonlinear propagation of intense incoherent photons in a photon gas is considered. The photon-photon interactions are governed by a pair of equations comprising a wave-kinetic equation for the incoherent photons in the presence of the slowly varying energy density perturbations of sound-like waves, and an equation for the latter waves in a background where the photon coupling is caused by quantum electrodynamical effects. The coupled equations are used to derive a dispersion relation, which admits new classes of modulational instabilities of incoherent photons. The present instabilities can lead to fragmentation of broadband short photon pulses in astrophysical and laboratory settings. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311422_arXiv.txt": {
        "abstract": "We present the expressions of the three- and four-point correlation functions of a self interacting light scalar field in a de Sitter spacetime at tree order respectively for a cubic and a quartic potential. Exact expressions are derived and their limiting behaviour on super-horizon scales are presented. Their essential features are shown to be similar to those obtained in a classical approach. ",
        "introduction": "The computation of high order correlation functions of scalar fields in a de Sitter spacetime can be of interest for investigations of the physics of the early universe. There are indeed a growing number of indications that in its early phase the universe underwent an inflationary period~\\cite{WMAP} that can be accurately described by a de Sitter phase during which the energy density of the universe is thought to be dominated by the self energy of a scalar field~\\cite{guth81}. Its fluctuations, or the fluctuations of any other light scalar fields, are thought to be the progenitors of the large-scale structure of the universe~\\cite{inflation}. The statistical properties of the induced metric fluctuations then depend on the potential landscape in which the fields evolve. It is unlikely that the fluctuations produced along the inflaton direction can be significantly non-Gaussian as stressed in recent works~\\cite{maldacena,Creminelli}. It is possible however that if there exist other self interacting light scalar fields during that period primordial non Gaussian metric fluctuations can be generated~\\cite{bartolo2,bartolo1,BU1,BU2}. The details of the effects induced by such self-interacting fields is based on the statistical properties of the metric fluctuations that quantum generated field fluctuations can induced.  The purpose of this paper is to investigate the foundation of these calculations by calculating the correlation properties of a test scalar field in a de Sitter background.  The content of the paper is the following. In the second part we present the basis of such computations. In the third we explore the expression of the leading order term of high order correlators for different self-interacting potentials. In particular we give explicit results for the superhorizon limit that corresponds to modes that can be observed today. Finally we give insights on the physical interpretations of those results. ",
        "conclusions": ""
    },
    "0311/astro-ph0311108_arXiv.txt": {
        "abstract": "We discuss the crucial role of a dynamical picture in the analysis of AGN intrinsic absorbers data.  High quality FUSE data of Mrk 279 are used to demonstrate that the line of sight covering fraction is a strong function of velocity.  In Mrk 279, as well as in most cases where the data is of high enough quality, the shape of the absorption troughs is mainly determined by the velocity-dependent covering fraction. We argue that the traditional ``cloud'' picture of AGN outflows is hard pressed to account for the velocity-dependent covering fraction, as well as for the highly super-thermal width of the troughs and the detached trough phenomenon.  A disk outflow picture naturally explains these features and furthermore, is using the simplest reservoir for the outflowing material: The accretion disk around the black hole. Accounting for velocity dependent covering can drastically increase the inferred ionic column density of the analyzed trough, an increase which is amplified in the total column density and ionization parameter solution for the AGN outflow trough. ",
        "introduction": "AGN outflows are evident by resonance line absorption troughs, which are blueshifted with respect to the systemic redshift of their emission counterparts. In Seyfert galaxies, velocities of several hundred \\kms\\ (Crenshaw et~al.\\ 1999; Kriss et~al.\\ 2000) are observed in both UV resonance lines (e.g., \\civ~$\\lambda\\lambda$1548.20,1550.77, \\nv~$\\lambda\\lambda$1238.82,1242.80, \\ovi~$\\lambda\\lambda$1031.93,1037.62 and \\Ly), as well as in X-ray resonance lines (Kaastra et~al.\\ 2000; Kaspi et~al.\\ 2000). Similar outflows (often with significantly higher velocities) are seen in quasars (Weymann et~al.\\ 1991; Korista, Voit, Morris, \\& Weymann 1993; Arav et~al.\\ 2001a). Reliable measurement of the absorption column densities in the troughs are crucial for determining the ionization equilibrium and abundances of the outflows, and the relationship between the UV and the ionized X-ray absorbers.  In the last few years our group (Arav 1997; Arav et~al.\\ 1999a; Arav et~al.\\ 1999b; de~Kool et~al.\\ 2001; Arav et~al.\\ 2001a) and others (Barlow 1997, Telfer et~al.\\ 1998, Churchill et~al.\\ 1999, Ganguly et~al.\\ 1999) have shown that in quasar outflows most lines are saturated even when not black. We have also shown that in many cases the shapes of the troughs are almost entirely due to changes in the line of sight covering as a function of velocity, rather than to differences in optical depth (Arav et~al.\\ 1999b; de~Kool et~al.\\ 2001; Arav et~al.\\ 2001b). Gabel et al.\\ (2003) show the same effect in the outflow troughs of NGC~3783. As a consequence, the column densities inferred from the depths of the troughs are only lower limits.   \\clearpage   \\vspace{-0.2cm} \\begin{figure} \\plotfiddle{sdss_conf_mrk279_o6.ps}{3.1in}{0}{50}{35}{-150}{-20} \\caption{In the bottom panel we show the $C(v)$ solution for the \\ovi\\ troughs in Mrk 279. The straight lines below the trough are our attempt to attach a single covering factor to each of the troughs (covering fraction values are read from the right axis). It is clear that the shape of trough A is dominated by $C(v)$. The top panel shows three solutions for the optical depth: Apparent $\\tau$, which is derived from the residual intensity of the blue doublet component assuming full coverage (i.e., the ISM and IGM case); a solution (dashed line) assuming a constant covering (shown at the bottom of troughs A and B); and the real $\\tau$ using the $C(v)$ curve shown on the bottom panel. } \\vspace{-0.2cm} \\end{figure} ",
        "conclusions": ""
    },
    "0311/astro-ph0311614_arXiv.txt": {
        "abstract": "We present preliminary results from a series of high-resolution \\nbody\\ simulations that focus on 8 dark matter halos each of order a million particles within the virial radius. We follow the time evolution of hundreds of individually tracked satellite galaxies and relate their physical properties to the differing halo environmental conditions. Our main science driver is to understand how satellite galaxies lose their mass and react to tidal stripping. Unlike previous work our results are performed in a fully self-consistent cosmological context. The preliminary results demonstrate that while environment may vary considerably with respects to formation time and richness of substructure, the satellites evolve similarly. ",
        "introduction": "Four standard  \\LCDM\\ cosmological simulations with side length 64\\hMpc\\ were run at low mass resolution ($128^3$ particles), with 8 halos selected sampling a range of environments. In this contribution we present results for our oldest (8.3 Gyrs) and youngest (3.4 Gyrs) dark matter halos, halo \\# 1 and \\# 8 respectively.  These halos were re-simulated at 64 times higher mass resolution giving an effective mass per particle of $10^8$ \\hMsun, and a force resolution of approximately 1 \\hkpc.  The \\nbody\\ simulations were carried out using the publicly available adaptive mesh refinement code \\mlapm\\ (Knebe, Green~\\& Binney 2001). The code solves Poisson's equation on a hierarchy of nested grids: the whole computational volume is covered by one cubic domain grid whereas refined regions are of arbitrary shape and adjusted to the actual density field at each major step in order to follow the real distribution of particles at all times.  An example of \\mlapm\\ in action is shown in the left panel of Figure 1. These two images show the density field of our youngest halo superimposed by the refinement grids. The top image shows the 5th refinement level with grid resolution $4096^3$, the bottom the 6th level of refinement.  As the grids follow the over-densities in the simulations they by definition encompass the satellites.  Hence, we can build a hierarchy of nested isolated \\mlapm\\ grids into a \"grid tree\", and then generate a list of potential satellite centers  by treating the densest grid at the end of each branch as a \"potential\" halo center. We further check the validity of the center by assigning physical properties to the adjacent radial density distribution. The details can be found elsewhere (Gill, Knebe~\\& Gibson 2003).  \\begin{figure} \\vspace{88mm} \\special{psfile=Fig01a.eps vscale=27 hscale=27 voffset=-10 hoffset=10} \\special{psfile=Fig01b.eps vscale=27 hscale=27 voffset=-10 hoffset=200} \\caption{\\footnotesize {(Left Panel) Consecutive refinement levels of \\mlapm 's grid structure for halo \\# 8 superimposed upon the density projection of the particle distribution. (Right Panel) Orbits of all objects in the vicinity of the host halo; orbital lines graduate from dark at \\zform\\ time to light $z=0$.}} \\end{figure} ",
        "conclusions": ""
    },
    "0311/astro-ph0311564_arXiv.txt": {
        "abstract": "{We have derived element abundances in 310 emission-line galaxies from the Early Data Release of the Sloan Digital Sky Survey (SDSS) for which the [O {\\sc iii}] $\\lambda$4363 emission line was detected, allowing abundance determination by direct methods. We found no extremely metal-deficient galaxy ($Z$ $\\la$ $Z_\\odot$/12)\\thanks{12+log(O/H)$_\\odot$=8.69 $\\pm$ 0.05 (Lodders \\cite{L03}).}, probably as a consequence of selection effects in the SDSS sample. The oxygen abundance 12 + log O/H of the SDSS galaxies sample lies in the range from $\\sim$ 7.6 ($Z_\\odot$/12) to $\\sim$ 8.4 ($Z_\\odot$/2). This sample is merged with a sample of $\\sim$100 blue compact dwarf galaxies with high quality spectra containing some very low-metallicity objects to study the abundance patterns of low-metallicity emission-line galaxies. We find that the $\\alpha$ element-to-oxygen abundance ratios do not show any significant trends with the oxygen abundance, in agreement with previous studies. The Fe/O abundance ratio is smaller than the solar value, which we interpret as an indication that type Ia supernovae have not yet appeared in these galaxies, implying an age of less than 1 -- 2 Gyr. However, a slight decrease of the Fe/O abundance ratio with increasing metallicity suggests some depletion of iron onto dust in the galaxies with higher metallicities. The N/O abundance ratio  ranges from log N/O= --1.6 to --0.8. The fact that no galaxy with log N/O $\\la$ --1.6 was discovered implies that local low-metallicity emission-line galaxies are of a different nature than  high-redshift damped Ly$\\alpha$ systems with log N/O of $\\sim$ --2.3 and that their ages are probably larger than 100 -- 300 Myr. Our data indicate  the existence of a gradual nitrogen enrichment on a time-scale of a few Myr. ",
        "introduction": "The study of element abundances in  low-metallicity emission-line galaxies is important for our understanding of the chemical evolution of galaxies and for constraining models of stellar nucleosynthesis and the shape of the initial mass function. The optical spectra of H {\\sc ii} regions in these galaxies show strong narrow emission lines superposed on a stellar continuum that is rising toward the blue, allowing abundance determinations of such heavy elements as nitrogen, oxygen, neon, sulfur, argon, and iron.  The oxygen abundance  in the most metal-poor emission-line galaxies is around 1/30 times the solar value, which makes these objects among the least chemically evolved galaxies in the universe. It was proposed (Sargent \\& Searle \\cite{SS70}) that the most metal-poor galaxies were actually experiencing their first burst of star formation. However, subsequent photometric and spectroscopic studies have shown that the majority of metal-poor galaxies also possess several Gyr old stellar populations. In the case of the two extremely metal-deficient galaxies, I Zw 18 and SBS 0335--052, there is however no clear evidence of such stellar population and these galaxies may be genuinely young (e.g., Hunt, Thuan \\& Izotov \\cite{H03}; Guseva et al. \\cite{G03c}).  Studies of local metal-poor emission-line galaxies are also crucial in the context of galaxy formation. The proximity of these metal-poor galaxies allows studies of their structure, metal content, and stellar populations in a nearly pristine environment with a sensitivity, precision, and spatial resolution that faint, small angular size, distant high-redshift galaxies do not allow. In particular, the comparison of the abundance patterns in the local low-metallicity dwarf galaxies and distant damped Ly$\\alpha$ systems  (DLA) can shed light on the properties and the evolution of the primeval high-redshift galaxies.  The abundances of CNO  play an  especially important role in the understanding the chemical evolution of the galaxies. It is well established (e.g., Maeder \\cite{M92}) that oxygen and likely carbon are produced mainly by massive stars. The origin of nitrogen is more controversial.  In local metal-poor star-forming galaxies the nitrogen-to-oxygen abundance ratio  is found to be constant (e.g., Garnett \\cite{G90}) with a very low dispersion around the mean value of log N/O = --1.6 for galaxies with oxygen abundance 12 + log O/H $\\la$ 7.6 (Thuan, Izotov \\& Lipovetsky \\cite{TIL95}; Izotov \\& Thuan \\cite{IT99}). This led Thuan et al. (\\cite{TIL95}) and Izotov \\& Thuan (\\cite{IT99}) to infer that in those galaxies,  nitrogen has been produced by massive stars as a primary element and hence the lowest-metallicity galaxies are young systems. However,  Henry, Edmunds \\& K\\\"oppen (\\cite{H00}) proposed another explanation involving  a star formation occuring at a  very low rate, with massive stars producing oxygen and intermediate-mass stars producing primary nitrogen with a typical time-lag of 250 Myr.   In high-redshift low-metallicity damped Ly$\\alpha$ systems, on the other hand, the spread of the N/O abundance ratio appears to be very large (e.g., Lu, Sargent \\& Barlow \\cite{L98}; Centuri\\'on et al. \\cite{C03}). The N/O abundance ratio in some DLAs is found to be one order of magnitude lower than that in blue compact dwarf galaxies  (BCD) of same metallicity. If real, such differences suggest that DLAs and BCDs are systems in different evolutionary stages. This also raises doubts on the young age of the lowest-metallicity BCDs and suggests that the production of nitrogen by intermediate-mass stars is actually important in the low-metallicity BCDs and DLAs. This is in line with predictions of new models of rotating stars (e.g., Meynet \\& Maeder \\cite{MM02}), although these models also predict the production of a large amount of primary nitrogen by low-metallicity massive stars.  Emission-line galaxies with very low metallicity are rare. Only two dozens of BCDs with an oxygen abundance 12 + log O/H $\\la$ 7.6 are known up to now. The aim of this paper is to produce a large sample of low-metallicity emission-line galaxies with reliably derived element abundances and to compare them with those in the high-redshift DLAs. This should allow a better understanding of the nitrogen production and of the evolutionary status of both types of objects. We have used the early data release (EDR) of the Sloan Digital Sky Survey (SDSS) (Stoughton et al. \\cite{S02}) which includes the spectra of $\\sim$ 50000 galaxies, and searched for low-metallicity emission-line galaxies. The procedure and the resulting sample are described in section 2.  The derivation of the elemental abundances is presented and discussed in section 3. The main conclusions of this study are presented in section 4.  \\begin{figure}[t] \\hspace*{-0.0cm}\\psfig{figure=0101.f1.ps,angle=0,height=8cm,clip=} \\caption{Typical spectra of the (a) low-excitation and (b) high-excitation H {\\sc ii} regions from the SDSS EDR. The spectra are transformed to a linear wavelength scale and zero redshift.} \\label{Fig1} \\end{figure} ",
        "conclusions": "We have discussed in this study the abundances of the heavy elements in the 310 emission-line galaxies from the sample of the Sloan Digital Sky Survey (SDSS) Early Data Release (EDR) with a detected [O {\\sc iii}] $\\lambda$4363 emission line, allowing a direct element abundance determination based on the electron temperature. We have shown that the abundances are biased for galaxies with spectra of low excitation, and we attribute this bias mainly to a false detection of [O {\\sc iii}] $\\lambda$4363 or to a wrong placement of the continuum due to a low signal-to-noise ratio and adjacent stellar absorption lines. However, we cannot exclude that an additional heating mechanism different from the stellar radiation may be important for low-excitation galaxies. We have corrected for this bias by requiring that the Ar/O ratio be equal to the expected value. After correcting for this bias and merging the SDSS sample with the sample of the galaxies with large equivalent widths of emission lines (HeBCD sample), we obtain a sample of $\\sim$ 400 emission-line galaxies in total. Our main conclusions are as follows:  1. Despite an examination of the entire SDSS EDR sample of galaxies, we found no galaxy with extremely low metallicity (12 + log O/H $<$ 7.6, i.e. $Z$ $<$ $Z_\\odot$/12) while the HeBCD sample which consists of $\\sim$100 galaxies contains about 15 such galaxies. We suspect that the spectroscopic data base of the SDSS actually discriminates against the most metal-poor galaxies.  2. Emission-line galaxies from the SDSS sample in general follow the distributions of the N/O, Ne/O, S/O, Ar/O, Fe/O abundance ratios vs oxygen abundance found earlier for the galaxies in the HeBCD sample (Thuan et al. \\cite{TIL95}; Izotov \\& Thuan \\cite{IT99}).  3. The $\\alpha$ element-to-oxygen abundance ratios do not show any significant trend with oxygen abundance.  4. The Fe/O abundance ratio shows a significant underabundance of iron relative to oxygen as compared to solar, suggesting  either depletion of iron onto dust grains or a dominant production of iron in the massive supernovae of the type II. We favour the latter explanation, which implies an age of less than $\\sim$ 1 -- 2 Gyr.  5. No galaxy with log N/O $\\la$ --1.6 was found suggesting that local low-metallicity emission-line galaxies are different from those as  high-redshift DLAs with the considerably lower log N/O of $\\sim$ --2.3. These DLAs are considered to be truly young galaxies, with nitrogen produced only by massive stars. If this interpretation is correct, then our sample of $\\sim$ 400 dwarf emission-line galaxies contains no extremely young galaxy (i.e. with an age $\\la$ 100 -- 300 Myr).  6. Our data indicate a gradual enrichment of the galaxies in nitrogen by massive stars from the most recent starburst. The effect is best seen among galaxies from our sample that have intermediate metallicities."
    },
    "0311/astro-ph0311087_arXiv.txt": {
        "abstract": "Investigation of spectra from the Sloan Digital Sky Survey reveals the presence of a region of ionized gas of $>2\\arcdeg$ diameter centered approximately at $\\alpha = 10^{h} 37^{m}$ $\\delta = -00\\arcdeg 18'$ (J2000) (Galactic coordinates $l=248, b=+48$). [\\ion{O}{3}]~4959,5007 emission is particularly strong and emission from H$\\alpha$ and [\\ion{N}{2}]~6548,6583 is also detectable over a substantial area on the sky.  The combination of emission line ratios, the close to zero heliocentric radial velocity and the morphology of the structure are consistent with an identification as a very nearby planetary nebula.  The proximity of the hot, DO white dwarf PG~1034+001 further strengthens this interpretation.  The object is: i) the largest planetary nebula on the sky, ii) certainly closer than any planetary nebula other than Sh 2--216, iii) the first to be unambiguously associated with a DO white dwarf. A parallax distance for PG~1034+001 would establish whether the structure is in fact the closest, and one of the physically largest, planetary nebula known. ",
        "introduction": "\\label{Intro} The availability (Abazajian \\etal 2003) of a significant fraction of the spectra and images from the Sloan Digital Sky Survey (SDSS) (York \\etal 2000) provides a unique resource for the investigation of a wide variety of astrophysical phenomena.  The quality and scale of the database are such that a number of serendipitous discoveries can be expected in the coming years. In this {\\em Letter} we report the discovery, using SDSS spectra of unrelated objects, of the largest known planetary nebula (PN) on the sky. ",
        "conclusions": "Given that all observations are consistent with a PN nature for the newly discovered nebulosity, we designate it Hewett~1.  The proximity of the DO white dwarf PG~1034+001 (Wesemael, Green, \\& Liebert 1985) makes it probable that PG~1034+001 is responsible for the ionizing radiation.  Werner, Dreizler, \\& Wolff (1995) derived a spectroscopic distance for PG~1034+001 of $155^{+58}_{-43}\\,$pc and confirmation of the association would make Hewett~1 one of the closest PNs known (Napiwotzki 2001), the $2\\arcdeg$ diameter corresponding to a physical size of $3.5-7.0\\,$pc for the likely distance range of $100-200\\,$pc.  Again, Sh 2--216 provides the current benchmark, with a trigonometric distance of $130^{+9}_{-8}\\,$pc (Harris \\etal 1997) and a physical size of $\\simeq 3.5\\,$pc.  A parallax determination for PG~1034+001 would establish whether Hewett~1 is even closer.  Notwithstanding targeted searches (Werner \\etal 1997), very few PNs are known to be associated with non--DA (hydrogen poor) white dwarfs.  Hewett~1 may be the first PN to be discovered associated with a DO white dwarf, although PG~0108+101 (Reynolds \\etal 1987) and PG~0109+111 (Werner \\etal 1997) have candidate nebulae, neither detection is regarded as secure.  Since DO white dwarfs are thought to evolve from PG~1159 stars (see Dreizler \\& Werner 1996, and references therein), it will be of interest to compare the properties of Hewett~1 both to the PNs associated with PG~1159 stars (Napiwotzki 1999) and theoretical models (G{\\' o}rny \\& Tylenda 2000).  The derivation of a reliable age for the PN associated with PG~1034+001 will be of particular interest for constraining the time scales associated with the late--stages of evolution of post-asymptotic giant branch stars and the origin of PG~1159 stars and helium--rich WDs."
    },
    "0311/astro-ph0311034_arXiv.txt": {
        "abstract": "{We have carried out an in-depth study of three bright-rimmed clouds \\object{SFO 11}, SFO 11NE and \\object{SFO 11E} associated with the HII region IC 1848, using observations carried out at  the James Clerk Maxwell Telescope (JCMT) and the  Nordic Optical Telescope (NOT), plus archival data from IRAS, 2MASS and the NVSS. We show that the overall morphology of the clouds is reasonably consistent with that of radiative-driven implosion (RDI) models developed to predict the evolution of cometary globules. There is evidence for a photoevaporated flow from the surface of each cloud and, based upon the morphology and pressure balance of the clouds, it is possible that  D-critical ionisation fronts are propagating into the molecular gas. The primary O star responsible for ionising the surfaces of the clouds is the 06V star  \\object{HD17505}. Each cloud is associated with either recent or ongoing star formation: we have detected 8 sub-mm cores which possess the hallmarks of protostellar cores and identify  YSO candidates from 2MASS data.  We infer the past and future evolution of the clouds and demonstrate via a simple pressure-based argument that the UV illumination may have induced the collapse of the dense molecular cores found at the head of \\object{SFO 11} and \\object{SFO 11E}. ",
        "introduction": "\\object{IC 1848} is a large HII region, forming part of the   \\object{W5} HII region-molecular cloud complex in the Perseus Arm. The HII region component of  W5 (also known as S 199) is made up of two roughly circular thermal shells W5 East and W5 West, which are  separated by a dust lane. IC 1848 is located to the south of W5 West  (Braunsfurth 1983). W5 West is excited by the open cluster OCl 364, which comprises four O stars, whereas only one O star is  visible within W5 East (Normandeau, Taylor \\& Dewdney.~\\cite{normandeau}).  The whole complex lies at a distance of 1.9 kpc (Ishida \\cite{ishida}) and is a well known and well-studied star-forming region (e.g.~Val\\'ee, Hughes \\& Viner \\cite{vhv79}; Braunsfurth \\cite{braun}; Normandeau, Taylor \\& Dewdney \\cite{normandeau}; Heyer \\& Terebey \\cite{heyer}; Carpenter, Heyer \\& Snell \\cite{chs00}).  Numerous small bright-rimmed  clouds (sometimes also known as bright-rimmed globules) are found at the rims of W5 East and West,   which may be  star-forming regions triggered via the expansion of the HII regions (Sugitani, Fukui \\& Ogura \\cite{sfo}). The expansion of the HII regions drives shocks into the surrounding  molecular gas and these photoionisation-induced shocks are thought to trigger the collapse of sub-critical molecular cores within the clouds in a process known as radiative-driven implosion or RDI (Bertoldi \\cite{bertoldi}; Bertoldi \\& McKee \\cite{bk90}; Lefloch \\& Lazareff \\cite{ll94}, \\cite{ll95}).   Radiative-driven implosion of molecular cores at the periphery of HII regions may thus be responsible for a subsequent generation of star formation, amounting to a possible cumulative total of several hundred new stars per HII region (Ogura, Sugitani \\& Pickles \\cite{osp02}) and perhaps 15\\% or more of the low-to-intermediate mass stellar mass function (Sugitani, Fukui \\& Ogura \\cite{sfo}). Confirming bright-rimmed clouds as star-forming can provide important insights about the clustered mode of star formation and the overall star-formation efficiencies of molecular clouds.  Sugitani, Fukui \\& Ogura(\\cite{sfo}), hereafter referred to as SFO91, searched the Sharpless HII region catalogue (Sharpless \\cite{sharpless}) for bright-rimmed clouds associated with IRAS point sources, in order to identify potential star-forming clouds via their far-infrared emission.  Later, Sugitani \\& Ogura (\\cite{so94}) -- SO94 -- extended their search to  include bright-rimmed clouds from the ESO(R) Southern Hemisphere Atlas. At least 89 bright-rimmed clouds have been found to be associated with IRAS point sources. For brevity (and consistency with SIMBAD) we will refer to the combined  SFO91 and SO94 catalogues as the SFO catalogue.  Whilst a few individual clouds from the SFO catalogue have been studied in detail (e.g.~Lefloch, Lazareff \\& Castets~\\cite{llc97}; Megeath \\& Wilson \\cite{mw97}; Codella et al.~\\cite{cbnst01}) and shown to harbour protostellar cores, the vast majority of the SFO clouds  have not been associated with star-forming regions and the state of these clouds remains unknown.  We have carried out a star formation census of SFO bright-rimmed clouds (Thompson et al.~\\cite{thompsona}, \\cite{thompsonb}) to investigate the star-forming nature of the SFO bright-rimmed clouds and determine whether any star formation in the SFO sample was likely to have been triggered by the RDI process.  In this paper we report the results for three bright-rimmed clouds from our census.  The bright-rimmed cloud \\object{SFO 11} is found at the southern edge of IC 1848. It is positionally associated with the IRAS point source 02476+5950 and in optical images is double-lobed and cometary in appearance (see Fig.~\\ref{fig:dssnvss}). Two other bright-rimmed clouds that are not in the SFO catalogue are found within 6\\arcmin\\ of \\object{SFO 11}. Following the terminology of Ogura, Sugitani \\& Pickles (\\cite{osp02}) they are described in this paper as \\object{SFO 11NE} and SFO 11E. \\object{SFO 11NE} is a cometary cloud found 4\\arcmin\\ NE of \\object{SFO 11}, with a ``shoulder'' to the east side of the cloud. \\object{SFO 11E} is found 6\\arcmin\\  E of SFO11 and is associated with a bright ridge of nebular emission at the Southern ionisation boundary of IC 1848. A red optical image of the three clouds obtained from the Digitised Sky Survey is shown in Fig.~\\ref{fig:dssnvss}.  Neither \\object{SFO 11NE} or \\object{SFO 11E} are associated with any IRAS point sources.  This may be either due to  confusion caused by the limited IRAS resolution or  simply because SFO 11NE and E do not contain embedded protostars or stars. We included \\object{SFO 11} and the two neighbouring clouds in our study of the SFO catalogue  as they represent an ideal opportunity to study three possibly star-forming clouds close to each other on the sky and lying at a similar distance from the UV illumination source. In this paper  we present JCMT molecular line, SCUBA sub-mm  continuum, Nordic Optical Telescope narrowband H$\\alpha$ and archival observations (IRAS HIRES, 2MASS and VLA NVSS) to investigate the star-forming activity and general environment of these three clouds. These data are then modelled to try to understand their future evolution as potential star-forming regions.   In Sect.~\\ref{sect:obs} we describe the observational procedure. The data are analysed in Sect.~\\ref{sect:anal} where evidence for star formation, protostellar cores, embedded IR sources or molecular outflows is discussed. In Sect.~\\ref{sect:discuss} we explore implications for present and future star formation within the clouds. Finally in Sect.~\\ref{sect:conc} we present a summary of our conclusions.  \\begin{figure} \\begin{center} \\includegraphics*[scale=0.7,trim=255 70 270 80]{h4521f1_lowres.eps} \\label{fig:dssnvss} \\caption{A red Digitised Sky Survey image of the three clouds \\object{SFO 11}, \\object{SFO 11NE} and SFO 11E. The image is centred on \\object{SFO 11} and the other two clouds are the cometary bright-rimmed clouds found to the NE and E respectively. High pixel values have been excluded to emphasise the bright optical rims of the clouds. The arrow marks the direction toward the suspected exciting star of the clouds, HD 17505. The NVSS 20 cm radio emission is shown by the grey contours which start at 1.5 mJy and are spaced by 1.5 mJy.  The bright ridge to the south probably marks the southern ionisation boundary of IC 1848 and the broken appearance in the radio map may be partly a consequence of the limited sensitivity of NVSS to large angular size scales.} \\end{center} \\end{figure} ",
        "conclusions": "\\label{sect:conc}  We have carried out an in-depth study of three bright-rimmed clouds (SFO 11, \\object{SFO 11NE} and \\object{SFO 11E}) associated with the HII region IC 1848, in order to search the clouds for evidence of star formation and attempt to place the star formation in context with the likely past and future evolution of the clouds. We observed the three clouds using SCUBA on the JCMT to map their continuum emission at 450 and 850 $\\mu$m; the JCMT to map the J=2--1 lines of $^{12}$CO and $^{13}$CO; and the Nordic Optical Telescope to image the H$\\alpha$ emission from the cloud rims using a narrowband filter. We have also obtained archival data of the 20 cm radio emission from the NVSS; far and mid-infrared IRAS HIRES maps; and 2.2 $\\mu$m 2MASS images of the clouds. We draw the following conclusions from our observations:  \\begin{enumerate}  \\item The overall morphology of the clouds as seen in the H$\\alpha$ images,  SCUBA and CO maps is reasonably consistent with the predictions of the radiative-driven implosion (RDI) model of Lefloch \\& Lazareff (\\cite{ll94}). The cloud morphology resembles  the early collapse phase of the Lefloch \\& Lazareff (\\cite{ll94}) model, some 1--2 $\\times$ 10$^{5}$ years after the UV illumination began.  However, the cores found at the head of the clouds \\object{SFO 11NE} and \\object{SFO 11E} are not elongated in the direction of the UV illumination and this indicates that these cores may not have formed by the RDI process and could thus pre-date the UV illumination of the clouds. The H$\\alpha$ images show the typical radial striations expected for a photoevaporated flow of gas from the cloud surface.  \\item There is reasonable evidence for ongoing star formation within all three clouds at the class I or class II evolutionary stage. The SCUBA maps reveal 8 dust cores within the optical boundaries of the clouds which possess the same physical properties as protostellar cores. The dust cores found toward the head of the bright-rimmed clouds are typically larger and more massive than those found further away from the bright rims. 2MASS near-infrared images reveal several  candidate class I protostars and T Tauri stars associated with the clouds, although the majority of these objects are also consistent with reddened main sequence stars. More accurate infrared photometry is required to resolve ambiguities in classification. Molecular outflows may be associated with the two cores \\object{SFO 11NE SMM1} and  \\object{SFO 11E SMM1}, although higher spatial resultion is required to resolve the individual outflow lobes. A possible photoionised jet is seen in the H$\\alpha$ image of \\object{SFO 11NE} and optical spectroscopy to confirm the nature of this feature is a priority.  \\item The cometary morphology of the clouds suggests that their primary ionising star is \\object{HD17505}. The orientation of \\object{SFO 11E} is not aligned with the direction toward  \\object{HD17505} and this may indicate that \\object{SFO 11E} is also affected by another nearby OB star. The H$\\alpha$ emission from the bright rim of \\object{SFO 11E} is brighter on the  \\object{HD17505}-facing side of the cloud, suggesting that HD 17505 is the dominant exciting star. The 20 cm free-free emission from the clouds is consistent with a true distance from HD 17505 of approximately 22 pc. The star-cloud distance as seen in projection is approximately 11 pc.  \\item  All three clouds are in approximate pressure equilibrium between their ionised exterior and molecular interior, which is consistent with the RDI model prediction that D-critical ionisation fronts and photoionisation-induced shocks are currently propagating into the clouds. From the current data we cannot rule out the possibility that the clouds are over-pressured with respect to their ionised boundary layers  and in this scenario the photoionisation-induced shocks are stalled at the ionised boundary layer (Bertoldi \\cite{bertoldi}; Lefloch \\& Lazareff \\cite{ll94}). Nevertheless the close correspondence between the morphologies of the clouds and the RDI models suggests that the clouds are currently being shocked and on balance we conclude that it is likely that photoionisation-induced shocks are propagating into the clouds. We estimate induced shock velocities of $\\sim$ 1.4 km\\,s$^{-1}$, which have a crossing time of $\\sim 2 \\times 10^{5}$  years for the diameters of the cores found at the head of the clouds. The crossing times are similar to predictions of the cloud lifetimes determined from their morphology and indicates that there has been sufficient time since the clouds were illuminated for the shocks to propagate deep within the clouds.  \\item It is difficult to ascertain whether or not the current star formation observed within the clouds was induced by the photoionisation process. The duration of the cloud UV illumination, their shock crossing times and the estimated age of the protostars/YSOs that have formed within the clouds are all similar (typically a few times 10$^{5}$ years). We have shown via a simple pressure-based argument that the UV illumination may have rendered  \\object{SFO 11 SMM1} and \\object{SFO 11E SMM1} unstable against collapse but was not likely to have significantly affected  \\object{SFO 11NE SMM1}. A more detailed theoretical treatment of the cloud collapse, extending the existing RDI models in this direction, is required to make more concrete predictions in this area.  \\item Based upon the current mass loss from the SMM1 cores, predicted by the Lefloch \\& Lazareff (\\cite{ll94}) RDI model, we estimate their lifetimes as between 4 and 6 Myr. Potential  ionisation front instabilities seen in the H$\\alpha$ images of \\object{SFO 11} and SFO 11E may substantially reduce this lifetime. The ongoing star formation within the clouds should proceed relatively undisturbed by the dispersal of the clouds. In the absence of instabilities developing  within the ionisation front, which significantly accelerate the mass loss rate (Lefloch \\& Lazareff \\cite{ll94}), the dense core at the head of SFO 11NE may even outlive its ionising O star.  \\end{enumerate}"
    },
    "0311/astro-ph0311202_arXiv.txt": {
        "abstract": "Nearby clusters such as Virgo and Coma possess galaxy distributions which tend to be aligned with the principal axis of the cluster itself. This has also been confirmed by a recent statistical analysis of some 300 Abell clusters where the effect has been linked to the dynamical state of the cluster. Moreover, the orbits of satellite galaxies in galactic systems like our own Milky Way also demonstrate a high degree of anisotropy -- the so-called Holmberg effect, the origin of which has been the subject of debate for more than 30 years.  This study presents the analysis of cosmological simulations focusing on the orbits of satellite galaxies within dark matter halos.  The apocentres of the orbits of these satellites are preferentially found within a cone of opening angle $\\sim$40$^\\circ$ around the major axis of the host halo, in accordance with the observed anisotropy found in galaxy clusters.  We do, however, note that a link to the dynamical age of the cluster is not well established as both our oldest dark matter halos do show a clear anisotropy signal.  Further analysis connects this distribution to the infall pattern of satellites along the filaments: the orbits are determined rather by the environment of the host halo than some \"dynamical selection\" during their life within the host's virial radius. ",
        "introduction": "\\label{observations} Observations of the distribution of bright elliptical galaxies within the Virgo galaxy cluster show a remarkable collinear arrangement (West~\\& Blakeslee 2000).  Moreover, this axis also appears to be part of a filamentary bridge connecting Virgo and the rich cluster Abell 1367. This phenomenon has already been recognized (Arp 1968; Bingegeli, Tammann~\\& Sandage 1987) but only with accurate measurements of distance did it become possible to discriminate between this being a genuine three-dimensional structure or merely a chance alignment of galaxies. West~\\& Blakeslee (2000) based their distances upon the surface brightness fluctuations method (cf. Blakeslee, Ajhar~\\& Tonry 1999 for a recent review), concluding that not only were the brightest cluster members distributed anisotropically over the cluster, but the distribution of dwarf ellipticals was elongated in the same direction (Bingelli 1999). Plionis~\\ea (2003) investigated this substructure-cluster correlation statistically using 303 Abell clusters. They affirm that there is indeed such a signal and also that this signal is related to the dynamical state (and the environment) of the cluster.  Moreover, even on smaller scales -- in galactic systems -- there are clear observational indications that the distribution of the orbits of satellite galaxies is biased towards the galactic pole (Holmberg 1969) giving an anisotropic distribution.  Zaritsky~\\ea (1997) found that satellites of (isolated) disk galaxies are scattered asymmetrically about the parent galaxy and aligned preferentially with the disk minor axis. At first this signal was only prominent for distances out to $\\sim$50~kpc from the host (Holmberg 1969), but an analysis based on a much larger sample (cf.  Zaritsky~\\ea 1993) of satellites extends it to 200~kpc (Zaritsky~\\ea 1997).  This result is confirmed by a study of the satellites orbiting M31 (Hartwick 2000; Grebel, Kolatt~\\& Brandner 1999).  In addition, for the one galaxy where individual satellite orbits are known, the Milky Way, there is also evidence that that the orbits are preferentially polar (Zaritsky~\\& Gonzalez 1999). This is derived from information based upon the alignment of satellites on the sky (Kunkel~\\& Demers 1976; Lynden-Bell 1982), the orientation of the Magellanic Stream (Mathewson, Clearly~\\& Murray 1974), the three-dimensional distribution of satellites (Majewski 1994; Hartwick 1996), and their actual velocities (Scholz~\\& Irwin 1994).  All these observations clearly indicate that on scales spanning from galaxy clusters down to galaxies there is a signal indicating a correlation between the alignment of substructure and the shape of the gravitational potential this substructure moves in. However, the source of this alignment of satellites has been debated for a number of years, with two potential solutions suggested to explain this puzzling arrangement. If the distribution of infalling satellites is initially spherical, then dynamical selection may preferentially destroy or suppress the star-formation in those on more equatorial orbits\\footnote{With \\textit{equatorial} we mean \"perpendicular to the major axis\"}. For instance, Pe\\~narubbia, Kroupa \\& Boily (2002) have argued that in flattened dark matter halos, the timescale for dynamical friction is substantially longer for satellites that are on polar orbits.  Alternatively, the present non-isotropic distribution of satellite systems may reflect the fact that it has always been non-isotropic, with the satellites being accreted along preferred directions.  These hypotheses can be tested in numerical simulations of structure formation.  However, in previous attempts to decipher this signal from cosmological \\nbody\\ simulations there still remains a certain amount of uncertainty (Zaritsky~\\ea 1997). The first fully self-consistent simulations targeting the subject were performed by Tormen (1997). They though lacked resolution in time, space and mass to perform a detailed analysis of the satellite population for differing environments, and hence more refined simulations are required. Tormen (1997) could not follow the satellite distribution within the host's virial radius but rather tracked all progenitors prior to accretion. This allowed him to analyze the infall pattern rather than the orbital evolution of the satellites. The aim of this study is to investigate this subject with the latest state-of-the-art high-resolution \\nbody-simulations. We focus upon a detailed analysis of the temporal and spatial properties of satellite galaxies residing within host galaxy clusters that formed fully self-consistently within a cosmological framework.   \\begin{deluxetable}{ccccccccc} \\tablecaption{Summary of the eight host dark matter halos. \\label{DMhost}} \\tablewidth{0pt} \\tablehead{ \\colhead{Halo}                            & \\colhead{\\Rvir [\\hkpc]}                   & \\colhead{$v_{\\rm circ}^{\\rm max}$[\\kms]}  & \\colhead{\\Mvir [$10^{14}$\\hMsun]}         & \\colhead{\\zform}                          & \\colhead{age [Gyr]}                       & \\colhead{$N_{\\rm sat}(<\\!r_{\\rm vir})$}   & \\colhead{$T$} } \\startdata \\# 1 &  1349 & 1125 & 2.87 & 1.16 & 8.30 & 158 & 0.67 \\\\ \\# 2 &  1069 &  894 & 1.42 & 0.96 & 7.55 &  63 & 0.87 \\\\ \\# 3 &  1081 &  875 & 1.48 & 0.87 & 7.16 &  87 & 0.83 \\\\ \\# 4 &   980 &  805 & 1.10 & 0.85 & 7.07 &  57 & 0.77 \\\\ \\# 5 &  1356 & 1119 & 2.91 & 0.65 & 6.01 & 175 & 0.65 \\\\ \\# 6 &  1055 &  833 & 1.37 & 0.65 & 6.01 &  85 & 0.92 \\\\ \\# 7 &  1014 &  800 & 1.21 & 0.43 & 4.52 &  59 & 0.89 \\\\ \\# 8 &  1384 & 1041 & 3.08 & 0.30 & 3.42 & 251 & 0.90 \\\\ \\enddata  \\end{deluxetable} ",
        "conclusions": "\\label{conclusions} There are two possible scenarios that explain the substructure-cluster alignment found in observations of galaxy clusters (Plionis~\\ea 2003; West~\\& Blakeslee 2000) and the analysis of \\nbody-simulations presented in this study.  Either some dynamical process can be held responsible, meaning that initially the orbits were randomly distributed and only some satellites survive to give rise to the observed correlation in \\Fig{NdotE}, or alternatively, their orbital parameters are imprinted upon them at the time they enter the host halo as already pointed out by Tormen (1997).  ``Dynamical destruction'', whereby satellites on more equatorial orbits suffer more tidal disruption, can be rejected because a comparable (although not as pronounced) alignment signal is found when restricting the analysis to those satellites that are in fact disrupted: We placed mass-less tracer particles at the last centre of the satellites before they were classified \"disrupted\".  These ``disrupted systems'' are presented as the thin lines in \\Fig{NdotE}.  We therefore conclude that it must be the initial distribution of the satellite systems that is responsible for the present day alignment. For example, galaxies are \"funneling along the filaments\" (Kitzbichler~\\& Saurer 2003) which has also been confirmed by X-ray observations of the spatial distribution of substructure in galaxy clusters (West, Jones~\\& Forman 1995) and biases the types of possible orbits.  It would now be reassuring to confirm a link between the host halo shape and the surrounding environment, i.e. are the filaments that feed the halo with material preferentially angled with respect to the orientation of the host halo?  When projecting the positions of the satellites onto a sphere at the time they enter the virial radius of the host it is clear that they are not randomly distributed; they cluster in directions linked to the filamentary structure surrounding the host halo, as found before observationally (e.g.  West~\\& Blakeslee 2003; Plionis~\\& Basilikos 2002, and references therein) and in \\nbody-simulations (e.g.  Faltenbacher 2002; Hatton~\\& Ninin 2001; Onuora~\\& Thomas 2000; Colberg~\\ea 1999; Splinter~\\ea 1997; Tormen 1997).  Tormen (1997) already pointed out that there is a strong alignment between the distribution of infalling satellite galaxies and the shape of the dark matter host. However, Tormen's simulations did not have the spatial and mass resolution to investigate the orbits of satellites \\textit{within} the host halo. We are, for the first time, tracing extremely well-resolved dynamics of the satellites within the dark matter hosts confirming that the alignment for both, disrupted and surviving, satellites is maintained for several orbital periods. \\Fig{SatelliteOrbits} now presents a visual impression of the trajectories of all surviving satellites identified at formation time of host halo~\\#3 until redshift $z=0$, qualitatively supporting the scenario that these orbits are linked to the filamentary structure which will leave its imprint in observed anisotropy.  Plionis~\\ea (2003) also established a link between the alignment and the dynamical age of the clusters. Our simulations though show that this finding can not be generalized: both our oldest halos, in which the satellites had as many as 4--5 orbits, do show the correlation signal.  We do not observe a trend for a \"randomization\" of the orbits in older halos. The satellites actually preserve the alignment with the host they had when they first fell into the cluster.  In \\Sec{observations} we showed that the (observational) substructure-alignment signal spans from galaxy clusters down to galactic scales.  Our results are mostly applicable to the former, and any extrapolation to smaller scales has to be handled with care. In the case for galaxies, for instance, the interaction between the galactic disk and the incoming satellite could enhance destruction in the galaxy plane, i.e. satellite on prograde orbits decay faster than the ones on retrograde (or polar) orbits due to orbital resonances between the disk and the satellites (Pe\\~narubbia, Kroupa \\& Boily 2002).  Based on theoretical and numerical studies (cf. Lacey~\\& Cole 1993 and Moore~\\ea 1999) it should be possible to re-scale our simulations to a Milky Way sized object by requiring that the maximum of the circular velocity curves of our halos equals 220~\\kms.  The scaling factor lies in the range of about 4--5 (cf.~$v^{\\rm max}_{\\rm circ}$ values in \\Table{DMhost}) and hence our ``rescaled dark matter halos'' would correspond to the ones in the observational data with virial radii in the range 260--295\\hkpc. This scaling factor in length entails a scaling in mass of 64--125 (simply the length scale to the power of three). This brings the mass of our galaxy clusters down to $\\sim 10^{12}$\\hMsun\\ which agrees with the dark matter mass inferred for our Milky Way (Freeman 1996). More problematic, however, are the ages of our systems: our halos are only $\\leq$8.3~Gyrs old opposed to 12~Gyrs for the Milky Way. Satellites in the Milky Way had the chance to complete nearly twice as many orbits leaving more space for an explanation based upon the dynamical destruction scenario.  When re-scaling our data, putting aside the age issue, and trying to explain the Holmberg effect, another uncertainty comes into play: the orientation of the stellar disk. In a triaxial potential there are in general stable closed orbits about both the major and the minor axis (e.g., Binney \\& Tremaine 1987). So, in principal the disk plane could either be perpendicular to the major or perpendicular to the minor axis. Under the assumption it lies perpendicular to the major axis of the dark matter halo the satellites in our simulations will be on polar orbits. There are indications that this configuration results in the most stable disk configuration within a triaxial halo (Hayashi~\\ea 2003). Moreover, even though there are clear indications that the angular momentum of the dark matter is well aligned with the minor axis of the halo (e.g. Warren~\\ea 1992), van den Bosch~\\ea (2002) showed that the angular momentum of the baryonic component (i.e. gas) not necessarily follows that of the dark matter distribution. They found an average misalignment between $\\vec{L}_{\\rm gas}$ and $\\vec{L}_{\\rm DM}$ of the order of 40$^\\circ$ in their numerical simulations. Hence the orientation of the galactic disk with respects to the dark matter halo is not well determined. Turning to observations does not resolve this question, with studies of polar rings indicating strongly oblate halos (i.e.  Iodice~\\ea 2003), while others suggest dark matter halos are more spherical, or even oblate (Olling \\& Merrifield 2000; Ibata~\\ea 2001). The situation is not at all clear, but the assumption that disks are perpendicular to the major axis (which is in agreement with results presented by Hayashi~\\ea 2003) would provide an explanation for the Holmberg effect, even though this is a very speculative interpretation."
    },
    "0311/astro-ph0311458_arXiv.txt": {
        "abstract": "\\label{sec:abstract}  A new grid of theoretical color indices for the \\Stromgren $uvby$ photometric system has been derived from MARCS model atmospheres and SSG synthetic spectra for cool dwarf and giant stars having $-3.0\\leq\\FeH\\leq+0.5$ and $3000\\leq\\Teff\\leq8000\\,$K.  At warmer temperatures (i.e., $8000<\\Teff\\leq40000\\,$K), this grid has been supplemented with the synthetic $uvby$ colors from recent Kurucz atmospheric models without overshooting (Castelli, Gratton, \\& Kurucz, 1997, A\\&A, 318, 841).  Our transformations appear to reproduce the observed colors of extremely metal-poor turnoff and giant stars:  the various $uvby$ color-magnitude diagrams (CMDs) for the $\\FeH\\sim-2.2$ globular cluster M$\\,$92 can be matched exceedingly well down to $M_V\\approx6$ by the same isochrone that provides a very good fit to published $BV$ data (see Paper I), on the assumption of the same distance and reddening.  Due to a number of assumptions made in the synthetic color calculations, however, our color--$T_{\\rm eff}$ relations for cool stars fail to provide a suitable match to the $uvby$ photometry of both cluster and field stars having $\\FeH>-2.0$.  To overcome this problem, the theoretical indices at intermediate and high metallicities have been corrected using a set of color calibrations based on field stars having well-determined distances from {\\it Hipparcos}, accurate $\\Teff$ estimates from the infrared flux method, and spectroscopic $\\FeH$ values.  In contrast with Paper I, star clusters played only a minor role in this analysis in that they provided a supplementary constraint on the color corrections for cool dwarf stars with $\\Teff\\leq5500\\,$K.  They were mainly used to test the color--$T_{\\rm eff}$ relations and, encouragingly, isochrones that employ the transformations derived in this study are able to reproduce the observed CMDs (involving $u-v$, $v-b$, and $b-y$ colors) for a number of open and globular clusters (including M$\\,$67, the Hyades, and 47$\\,$Tuc) rather well.  Moreover, our interpretations of such data are very similar, if not identical, with those given in Paper I from a consideration of $BV(RI)_C$ observations for the same clusters --- which provides a compelling argument in support of the color--$T_{\\rm eff}$ relations that are reported in both studies.  In the present investigation, we have also analyzed the observed \\Stromgren photometry for the classic Population~II subdwarfs, compared our ``final\" $(b-y)$--$\\Teff$ relationship with those derived empirically in a number of recent studies, and examined in some detail the dependence of the $m_1$ index on $\\FeH$.\\footnote{Tabular versions of our \\Stromgren color-temperature relations can be retrieved via {\\it ftp} to the address ``uvphys.phys.uvic.ca\" using the login I.D. ``star\"  and ``vicmodel\" as the password. A simple FORTRAN code (``uvby.for\") is provided that interpolates within the low- and high-temperature color tables (``uvbylo.data\" and ``uvbyhi.data\", respectively) to yield $b-y$, $m_1$, $c_1$, and $BC_V$ indices for input values of $\\Teff$, $\\logg$, and $\\FeH$.} ",
        "introduction": "\\label{sec:intro}  Among the wide variety of photometric systems available today, the $uvby$ system of \\citet{Stromgren1963} still remains one of the most valuable for the study of stellar populations and Galactic structure.  Its usefulness stems from the fact that the four intermediate-width filters are designed to isolate and measure certain key features in a stellar spectrum which are highly sensitive to the underlying physical characteristics of the star itself.  For example, \\Stromgren $b-y$ provides an accurate indicator of effective temperature similar to broadband $B-V$ or $V-I$ while the two other \\Stromgren indices, $m_1$ and $c_1$, yield photometric estimates of the stellar metal abundance and surface gravity (or luminosity).  In this respect, the ability of the $uvby$ system to provide precise estimates of these fundamental parameters makes it much better suited than standard Johnson-Cousins $UBVRI$ photometry for the study of individual stars.  For years, \\Stromgren photometry has provided a wealth of information on the chemical and dynamical evolution of the Milky Way through its application to the field-star populations in both the Galactic halo and the solar neighborhood \\citep[see, for example, the works of][]{CleggBell1973, SchusterNissen1988, SchusterNissen1989a, SchusterNissen1989b, Olsen1984, Haywood2001}.  Moreover, it has also proven to be extremely valuable in its application to star clusters.  While the narrowness of the $uvby$ filters limited earlier investigations \\citep{CrawfordBarnes1969, CrawfordBarnes1970a, CrawfordPerry1966, CrawfordPerry1976} primarily to those stars in nearby open clusters (e.g., the Hyades, NGC$\\,$752, Praesepe, and the Pleiades) that were bright and isolated enough to be readily observed using traditional photomultipliers, the advent of modern CCD detectors meant that the \\Stromgren system could be extended to much fainter stars such as those found near the turnoff and main sequence regions in more distant clusters \\citep[see the pioneering studies of][]{Anthony-Twarog1987a, Anthony-Twarog1987b, Anthony-TwarogTwarog1987}.  This recent explosion in the amount of high-quality \\Stromgren data for a large number of open and globular clusters offers profound potential for refining our understanding of these systems.  For instance, it has already been shown that both the $m_1$ and $c_1$ indices can reveal the existence of carbon and nitrogen abundance variations in globular cluster RGB stars \\citep{Hilker2000, Grundahl2002a}.  Moreover, the $c_1$ index has also been used to derive distance-independent cluster ages using techniques akin to those developed by \\citet{SchusterNissen1989b} for field stars \\citep{Grundahl2000a} while the $m_1$ index has provided photometric $\\FeH$ estimates for individual turnoff and giant stars \\citep{Nissen1987, HughesWallerstein2000, HilkerRichtler2000}.  Despite the availability of new high-quality $uvby$ photometry for a number of open and globular clusters in the Galaxy, our ability to fully exploit these data using stellar evolutionary models still remains somewhat difficult due to the lack of accurate and reliable color--$\\Teff$ relations and bolometric corrections for the \\Stromgren system that are needed to transform these theoretical models to the observed cluster color-magnitude diagrams (CMDs).  While a number of empirical calibrations and analytical formulae have been derived over the years that serve to relate the \\Stromgren $b-y$, $m_1$, and $c_1$ indices to the fundamental stellar parameters of $\\Teff$, $\\FeH$, and $\\logg$, respectively, they are often very specific to certain types of stars that occupy limited regions of the H-R diagram and generally rely on only a small number of stars in the solar neighborhood that have well-determined properties from spectroscopic analysis.  As a result, these relations cannot be trusted if they are applied to situations beyond the range in which they were originally intended.  Alternatively, one may employ grids of theoretical colors computed from model stellar atmospheres and synthetic spectra to interpret the photometric data \\citep{Bell1988}.  These synthetic colors are very useful for not only confirming the empirical calibrations between various \\Stromgren indices and certain stellar properties, but also for characterizing the nature of stars in those regimes not covered by these relations.  Furthermore, they are ideal for transforming evolutionary models to the observational color-magnitude and color-color planes owing to their broad coverage of stellar parameter space.  Their accuracy, however, is ultimately limited by how well the models are able to reproduce the observed spectrum of an actual star.  Some examples of these synthetic color computations for the $uvby$ system include the grid of MARCS colors for cool dwarfs and subgiants presented by \\citet{VandenBergBell1985} and those computed by Kurucz from his atmospheric models \\citep{Kurucz1993}.  While the $uvby$ colors of the latter cover a much larger range in stellar parameter space than the \\citeauthor{VandenBergBell1985} results, they are known to have problems reproducing the observed colors of cooler dwarf and giant stars \\citep{Grundahl1998}.  Therefore, in order to accurately interpret and analyze cluster $uvby$ photometry using current stellar models, we must first check if these theoretical colors are in good agreement with the observed photometry for a collection of stars with well-determined physical properties.  Recently, \\citet[hereafter Paper I]{VandenBergClem2003} found that, by applying small corrections to synthetic color transformations for the $BV(RI)_C$ system towards cooler effective temperatures, it is possible to achieve good consistency with the observational data for cool dwarf and giant stars in both the metal-poor and metal-rich regimes.  Their so-called ``semi-empirical\" approach resulted in a set of $BV(RI)_C$ color--$\\Teff$ relations and bolometric corrections that were able to accurately and consistently interpret the observed $B-V$, $V-I$, and $V-R$ CMDs for both a sample of clusters (such as the Hyades, M$\\,$67, M$\\,$92, 47$\\,$Tuc, and NGC$\\,$6791) and {\\it Hipparcos} field stars.  In addition, their predicted solar metallicity $(B-V)$--$\\Teff$ relationship and computed $(B-V)_{\\odot}$ value agree exceedingly well with those derived from the empirical analysis of \\citet{SekiguchiFukugita2000}.  In theory, these same methods could also be used to overcome the problems with the theoretical colors computed for other photometric systems, provided that a suitable amount of data, both for clusters and field stars, are available in order to quantify what corrections are necessary to successfully place the synthetic indices onto the observational systems.  The goal of the present investigation is to develop a set of accurate and reliable semi-empirical color transformations for the \\Stromgren $uvby$ system.  In contrast with Paper I, which adopted previously published synthetic color grids as the initial framework for cool stars, we choose instead to compute an entirely new grid of $uvby$ colors using MARCS model atmospheres and SSG synthetic spectra.  These new $uvby$ colors effectively supersede those reported by \\citeauthor{VandenBergBell1985} since they are computed from more recent versions of the MARCS/SSG programs and cover a broader range in parameter space.  This new grid of $uvby$ colors is applicable to both dwarfs and giants having $3000\\leq\\Teff\\leq8000\\,$K with metal abundances extending from $\\FeH=-3.0$ to +0.5.  At $\\Teff>8000\\,$K our new grid is supplemented with the most recent Kurucz $uvby$ colors computed by \\citet[hereafter CGK97]{Castelli1997} from atmospheric models that neglect overshooting.  In the analysis that follows we will explain how these purely theoretical colors can be brought into better agreement with the observed $uvby$ data for a number of star clusters by correcting them against a sample of {\\it Hipparcos} field stars having accurate $\\Teff$ estimates.  Ultimately, the validity of our semi-empirical approach will be demonstrated in much the same way as in Paper I for the $BV(RI)_C$ transformations by showing that they yield consistent fits of model isochrones to the photometric data for a number of different clusters, regardless of which \\Stromgren color is considered.  More importantly, we will also show that the resultant interpretations of these $uvby$ CMDs are virtually identical to those obtained in Paper I, which analyzes $BV(RI)_C$ photometry of the same clusters, when reasonable estimates for their distances, reddenings, and metallicities are assumed. ",
        "conclusions": "\\label{sec:conclusions}  In this investigation, we have produced an extensive set of color-temperature relations for the \\Stromgren $uvby$ system that covers a broad range in $\\Teff$, $\\logg$, and $\\FeH$.  To be specific, at each $\\FeH$ value in the range from $-3.0$ to $+0.5$, in steps of 0.5~dex, $b-y$, $m_1$, and $c_1$ colors are provided at $-0.5\\leq\\logg\\leq5.0$ and $3000\\leq\\Teff\\leq6000\\,$K (in a ``low-temperature\" table), as well as at $2.0\\leq\\logg\\leq5.0$ and $6000<\\Teff\\leq40,000\\,$K (in a ``high-temperature\"  table).  The $v-b$ and $u-v$ colors are readily obtained from $m_1=(v-b)-(b-y)$ and $c_1=(u-v)-(v-b)$.  By convention, \\Stromgren $y$ and Johnson $V$ magnitudes are similarly normalized, with the result that the bolometric corrections to $y$ are nearly identical with those to $V$ and typically differ by $\\lesssim\\,$0.01$-$0.02~mag over most of parameter space.  As it is generally the case that observed $y$ magnitudes are calibrated to be on the $V$ magnitude scale, the $BC_V$ values reported in Paper I have been included in our color transformation tables in order for these tables and associated interpolation software to be a self-contained package.  The main thrust of this study has been to determine what corrections to purely synthetic color--$\\Teff$ relations are needed to ensure that the latter do the best possible job of satisfying empirical constraints.  The starting point for our analysis consisted of a set of synthetic colors that were computed from MARCS model atmospheres and SSG synthetic spectra at $\\Teff\\leq8000\\,$K, to which we appended the predictions for hotter stars as derived by \\citet{Castelli1997} from Kurucz non-overshooting model atmospheres.  There were no obvious problems with this composite grid (in conjunction with modern isochrones) to reproduce \\Stromgren observations for extremely metal-poor turnoff and giant stars (those having $\\FeH\\lesssim-2.0$), as in the globular cluster M$\\,$92.  Indeed, a completely consistent interpretation of the cluster data was found regardless of which color plane was examined, and that interpretation was virtually identical to that obtained from a consideration of $BV$ data in Paper I. Discrepancies between theory and observations became evident for cool dwarfs with $\\FeH\\approx-2.0$ via our comparisons to a sample of extremely metal-poor field stars. Moreover, as the metallicity increased, the predicted $uvby$ colors tended to be systematically bluer than observed ones.  We assume that at least part of these discrepancies, particularly in the \\Stromgren $u$ and $v$ filters, can be attributed to incomplete atomic and molecular line lists, not incorporating the Fe~{\\sc I} continuous opacity source \\citep{Bell2001} in the SSG computations, and neglecting CN processing in the atmospheres of our giant star models.  It is important to note, however, that the neglect of such CN variations in giants (or the improved treatment of Fe~{\\sc I} continuous opacity) is not a concern for the {\\it calibrated} color--$\\Teff$ relations that are presented in this paper because our transformations have been constrained to reproduce the colors of {\\it observed} stars.\\footnote{It is precisely for this reason that we have not put in the time and effort to do the best possible job of the synthetic color transformations.  There is no doubt that the latter could be significantly improved for cool, metal-rich stars (in particular) if we implemented a better treatment of the bound-free opacity due to neutral iron, allowed for CN processing in giants, etc., but the main results of this investigation do not depend on the accuracy of the synthetic colors.  Even if we had started with different model-atmosphere based transformations, we would still have obtained the same empirically constrained color--$\\Teff$ relations in the end.  While it is desirable to determine the extent to which our results can be reproduced solely from theory, this is outside the scope of the present project.  One concern that should be kept in mind, however, is the possibility that the observational equipment employed by some workers differs appreciably from that used to define the original $uvby$ system. This may account for the the wide variation in the colors sometimes found for the same star by different observers.  In such cases, the calculated colors would not be expected to match those observations very well.}  We first corrected the synthetic colors for $\\FeH=0.0$ using a fairly large sample of solar metallicity field stars having accurate $\\Teff$ determinations from the infrared flux method.  These adjustments tended to be rather large for cool stars, but they were nonetheless necessary to place the colors on the same system as the photometric data.  Important confirmation of the accuracy of the corrected transformations was provided by fitting isochrones to the CMD of the open cluster M$\\,$67 where models for the observed metallicity, and whose $\\Teff$ scale was precisely normalized to the Sun, yielded a superb match to the observations.  [As cluster data played a only limited role in the derivation of the color adjustments, and as the same interpretation of the data was obtained here as in Paper I, which analyzed $BV(RI)_C$ observations, we conclude that our transformations are accurate and fully consistent with those reported in Paper I for the Johnson-Cousins photometric systems.]  We found that the necessary adjustments to the synthetic colors at $\\Teff>5500\\,$K or $\\logg<3.5$ varied nearly linearly with $\\FeH$, whereas those outside this regime were constrained based on a sample of cool field and cluster dwarf stars.  Indeed, upon applying these corrections at $-1.5\\leq\\FeH\\leq-0.5$ and $\\FeH=+0.5$, it was possible to achieve very good fits of isochrones to the CMDs of M$\\,$3 and 47$\\,$Tuc, which have $\\FeH\\approx-1.55$ and $-0.75$, respectively, as well as to that of the Hyades ($\\FeH=0.12$).  Once again, there was excellent consistency with the results reported in Paper I concerning the $BV(RI)_C$ system.  Only at higher metal abundances did inconsistencies arise as demonstrated by the fact that our fit of the same isochrones used in Paper I to the CMD of the $\\FeH\\approx\\,+0.4$ open cluster NGC$\\,$6791 required a significant color offset to match the photometry if the same reddening used in Paper I was adopted.  As we have been unable to understand this discrepancy, which may require further observations to resolve it, our transformations at $\\FeH\\gtrsim0.15$ (probably those in Paper I as well) should be used with caution. More observational constraints are needed to put the color--$T_{\\rm eff}$ relations for super-metal-rich stars on a firm foundation.  We have found good consistency between our $(b-y)$--$\\Teff$ relation for dwarf stars having near solar abundances and those derived by \\citet{SaxnerHammarback1985}, Gratton et al.~\\citeyearpar{Gratton1996}, and Alonso et al.~\\citeyearpar{Alonso1996b}. Comparisons between the predicted and observed colors for the classic Population II subdwarfs have also offered encouraging support for our results.  Moreover, the dependence of $m_1$ on $b-y$ predicted from our color transformations agrees well with the relations for solar-metallicity dwarfs derived by \\citet{SchusterNissen1989a} at $b-y\\lesssim0.45$ and by \\citet{Haywood2002} at redder colors, and with the relations for giants as inferred from a sample RGB stars from a number of globular and open clusters that encompass a wide range in $\\FeH$.  It is our intention to update the transformations reported in this paper as new constraints become available.  Even at this stage, however, we believe that our results represent a considerable improvement over currently available color--$\\Teff$ relations for the \\Stromgren photometric system and that their use will lead to important refinements in our our understanding of stars and stellar populations.  The next paper in this series (Clem, VandenBerg, \\& Stetson, in preparation) will present color transformations for the Sloan $u^\\prime g^\\prime r^\\prime i^\\prime z^\\prime$ system."
    },
    "0311/astro-ph0311172_arXiv.txt": {
        "abstract": "{ The gamma-ray astronomical observatory INTEGRAL, succesfully launched on 17$^{\\rm th}$ October 2002, carries two large gamma-ray telescopes. One of them is the coded-mask imaging gamma-ray telescope onboard the INTEGRAL satellite (IBIS) which provides high-resolution ($\\approx$~12$'$) sky images of 29$^{\\circ}$~$\\times$~29$^{\\circ}$ in the energy range from 15 keV to 10 MeV with typical on-axis sensitivity of $\\approx$ 1~mCrab at 100 keV (3$\\sigma$, 10$^6$~s exposure). We report here the general description of the IBIS coded-mask imaging system and of the standard IBIS science data analysis procedures. These procedures reconstruct, clean and combine IBIS sky images providing at the same time detection, identification and preliminary analysis of point-like sources present in the field. Spectral extraction has also been implemented and is based on simultaneous fitting of source and background shadowgram models to detector images. The procedures are illustrated using some of the IBIS data collected during the inflight calibrations and present performance is discussed. The analysis programs described here have been integrated as instrument specific software in the Integral Science Data Center (ISDC) analysis software packages currently used for the Quick Look, Standard and Off-line Scientific Analysis. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0311/astro-ph0311491_arXiv.txt": {
        "abstract": "We propose an efficient and model independent method for reconstructing the primordial power spectrum from Cosmic Microwave Background (CMB) and large scale structure observations. The algorithm is based on a Monte Carlo principle and therefore very simple to incorporate into existing codes such as Markov Chain Monte Carlo.  The algorithm has been used on present cosmological data to test for features in the primordial power spectrum. No significant evidence for features is found, although there is a slight preference for an overall bending of the spectrum, as well as a decrease in power at very large scales.  We have also tested the algorithm on mock high precision CMB data, calculated from models with non-scale invariant primordial spectra. The algorithm efficiently extracts the underlying spectrum, as well as the other cosmological parameters in each case.  Finally we have used the algorithm on a model where an artificial glitch in the CMB spectrum has been imposed, like the ones seen in the WMAP data. In this case it is found that, although the underlying cosmological parameters can be extracted, the recovered power spectrum can show significant spurious features, such as bending, even if the true spectrum is scale invariant. ",
        "introduction": "The present cosmological concordance model has been shown to describe all present data surprisingly well with only about six parameters \\cite{Tegmark:2003ud,map2}. From a physics point of view it is based on reasonable assumptions, such as general relativity and forms of matter with constant equations of state, albeit the physical nature of both the dark matter and the dark energy remains unknown. The initial perturbation spectrum appears to be almost purely adiabatic, Gaussian and scale invariant. Furthermore a general framework for producing such fluctuations in the early universe, namely inflation, is known and well established. As long as the universe remains deeply within the slow-roll regime during inflation, and the inflaton remains weakly coupled, Gaussian quantum fluctuations are produced which eventually become the adiabatic metric fluctuations observed for instance in the cosmic microwave background. However, there are many reasons to believe that there are deviations from this behaviour during the inflationary epoch, and countless physical mechanisms have been investigated. In the simplest models the deviations from scale invariance are caused by non-zero derivatives of the inflaton potential. Depending on the exact inflation model this can produce a tilt in the spectrum, and in general one can expand the produced spectrum of metric fluctuations in a Taylor series around the scale-invariant Harrison-Zeldovich spectrum  \\begin{equation} \\log P (k) = \\log P (k_0) + (n-1) \\log \\frac{k}{k_0} +\\left. \\frac{1}{2} \\frac{d\\, n}{d \\, \\log k} \\right|_{k_0} \\, \\log ^2 \\frac{k}{k_0} + \\cdots \\label{power} \\end{equation}  where the different terms in the expansion are related to the derivatives of the potential.  Within this framework it has to be assumed that the universe does not deviate significantly from slow-roll in the sense that the power series above converges, and higher order terms are very small.  However, the present cosmological data shows some indication that the second order term is not small compared with the first order term \\cite{map3,Leach:2003us,Lue:2003su,Barger:2003ym,Kinney:2003uw,Kinney:2000nc,Hannestad:2000tj,Hannestad:2001nu}, perhaps indicating that the above series expansion does not converge, and that slow-roll may not be a valid approximation.  If one does away with the assumption of slow-roll during the production of fluctuations then much more radical deviations from scale invariance can be achieved. If for instance there are phase transitions or resonant production of particles during inflation, then step-like features in the power spectrum can occur \\cite{Starobinsky:ts,Chung:1999ve,Elgaroy:2003hp,Barriga:2000nk,Adams:1997de}. It is in fact possible that such a feature is seen in the CMB data at the very largest scales, where the level of fluctuations is significantly smaller than expected.  On a different note it is also possible that effects from Planck scale physics can show up during inflationary slow-roll which occurs at much lower energies and perturb the spectrum of fluctuations \\cite{Martin:2000xs,Brandenberger:2000wr,Niemeyer:2000eh,Easther:2001fi,Easther:2001fz,Danielsson:2002kx,Easther:2002xe,Alberghi:2003am,Martin:2003kp,Bergstrom:2002yd,Elgaroy:2003gq,Martin:2003sg}, one possibility being a logarithmic modulation of the power spectrum. This can occur even if the slow-roll conditions are fulfilled.   It is of course possible to search for specific types of features predicted in a given type of model, such as steps, Gaussian peaks, or oscillations (see e.g.\\ \\cite{Elgaroy:2001wu}). However, with such a plethora of different possibilities it would be highly desirable with a model-independent and robust method for searching for deviations from scale invariance in the power spectrum. Such a method should be able to naturally encompass the power-law spectra predicted by standard slow-roll, as well as the more radical features such as steps and bumps.  In the present paper we describe exactly such a method which is able to extract features in the underlying spectrum, using either CMB or large scale structure data. Furthermore it is very easy to implement in existing software for cosmological parameter estimation.  In section 2 we describe general features of likelihood analysis, as well as the present cosmological data. In sections 3 and 4 we describe the spectrum reconstruction algorithm and use it on present cosmological data to search for features. In section 5 we go on to test the algorithm against mock data from future CMB experiments such as the Planck Surveyor. In section 6 we test how systematic errors in the $C_l$ data can affect parameter estimation and spectrum recovery. To that end we insert a Gaussian feature directly into the mock CMB data and do likelihood analysis on this modified data. Finally section 7 contains a discussion and conclusion. ",
        "conclusions": "We have developed an efficient and model-independent method for reconstructing the primordial power spectrum from cosmological observations of both CMB and large scale structure. This method has the major advantage that it can locate even narrow features in the spectrum very accurately. Furthermore the algorithm is very easy to implement in current likelihood calculators which are Monte Carlo based, although it can also be used in grid based algorithms.  We then tested the algorithm on current cosmological data and found, in accordance with numerous other studies, no significant evidence for any features in the primordial spectrum of fluctuations, $P(k)$. There is a slight preference for an overall negative curvature of the spectrum, corresponding to a negative running of the spectral index. However, the statistical significance is very low. Furthermore there is an indication of a break in the power spectrum at the very largest observable scales, corresponding to the dip in power at low $l$ seen by WMAP. Again, however, the statistical significance of this feature is very low.  In order to test the efficiency and robustness of the algorithm we then went on to use it on mock data from a future high-precision experiment such as the Planck Surveyor. Two cases were investigated, one with a step-like feature imposed, and one with a Gaussian suppression of the power spectrum. In both cases the algorithm was able to recover the true, underlying power spectrum very efficiently and with no biasing, something which would for instance be extremely difficult for algorithms with fixed bins.  Finally we tested the algorithm on a mock CMB spectrum, where a noise feature of Gaussian shape was imposed {\\it after} the data was generated. The reason is that such glitches are seen in the present WMAP data, and it is of interest to know whether the true power spectrum, as well as other cosmological parameters can be safely recovered even if there are such spurious features present.  Running the algorithm we found that the cosmological parameters of the underlying model could be recovered safely, but that the extracted primordial spectrum no longer resembled to scale-invariant spectrum used to generate the data. Rather it showed a distinct bending. This means that care should be taken when extracting information on the primordial power spectrum, including such parameters as the running of the spectral index, from power spectra which show spurious, localized features in $l$-space."
    },
    "0311/astro-ph0311344_arXiv.txt": {
        "abstract": "{We describe XMM-Newton Guaranteed Time observations of a sample of eight high redshift ({\\it 0.45$<$z$<$0.62}) clusters. The goal of these observations was to measure the luminosity and the temperature of the clusters to a precision of $\\sim$10\\%, leading to constraints on the possible evolution of the luminosity--temperature ($L_{\\rm x}-T_{\\rm x}$) relation, and ultimately on the values of the matter density, $\\OmM$, and, to a lesser extent, the cosmological constant $\\OmL$. The clusters were drawn from the SHARC and 160 Square Degree (160SD) ROSAT surveys and span a bolometric (0.0--20~keV) luminosity range of 2.0 to 14.4 $\\times 10^{44}$ \\ergs (H$_{o}$=50,$\\OmM$=1, $\\OmL$=0).  Here we describe our data analysis techniques and present, for the first time with XMM-Newton, a $L_{\\rm x}-T_{\\rm x}$ relation.  For each of the eight clusters in the sample, we have measured total ($r<r_{virial}$) bolometric luminosities, performed $\\beta$-model fits to the radial surface profiles and made spectral fits to a single temperature isothermal model. We describe data analysis techniques that pay particular attention to background mitigation. We have also estimated temperatures and luminosities for two known clusters (Abell 2246 and RXJ1325.0-3814), and one new high redshift cluster candidate (XMMU J084701.8+345117), that were detected off-axis. Characterizing the $L_{\\rm x}-T_{\\rm x}$ relation as $L_{\\rm x} = L_{6} (\\frac{T}{6 keV})^{\\alpha}$, we find $L_{6}=15.9 ^{+7.6}_{-5.2} \\times 10^{44}$\\ergs~ and $\\alpha$=2.7$\\pm$0.4 for an $\\OmL=0.0, \\OmM=1.0$, $H_0=50$ ~km~s$^{-1}$~Mpc$^{-1}$ cosmology at a typical redshift $z \\sim 0.55$. Comparing with the low redshift study by \\cite{markevitch}, we find $\\alpha$ to be in agreement, and assuming $L_{\\rm x}-T_{\\rm x}$ to evolve as $(1+z)^{A }$, we find $A$=0.68$\\pm$0.26 for the same cosmology and $A=1.52^{+0.26}_{-0.27}$ for an $\\OmL=0.7, \\OmM=0.3$ cosmology. Our $A $ values are very similar to those found previously by \\cite{Vikh_hiz} using a compilation of Chandra observations of $0.39<z<1.26$ clusters.  We conclude that there is now evidence from both XMM-Newton and Chandra for an evolutionary trend in the $L_{\\rm x}-T_{\\rm x}$ relation. This evolution is significantly below the level expected from the predictions of the self-similar model for an $\\OmL=0.0, \\OmM=1.0$, cosmology, but  consistent with self-similar model in an $\\OmL=0.7, \\OmM=0.3$ cosmology. Our observations lend support to the robustness and completeness of the SHARC and 160SD surveys. ",
        "introduction": "\\subsection{Motivation} As the most massive gravitationally bound objects in the universe, galaxy clusters are particularly sensitive to the evolution of the density perturbations responsible for their formation.  Cluster abundance as a function of mass and redshift is dictated by the mass function (\\cite{ps}; \\cite{jenkins}), which gives the comoving space density of collapsed objects as a function of mass and redshift.  In standard models, the mass function falls off as a Gaussian at the high mass end, reflecting the Gaussian nature of the density perturbations.  Galaxy cluster abundance is therefore exquistely sensitive to the amplitude of the perturbations and its evolution. Since this evolution is controlled by the underlying cosmological background, observations of cluster abundance offer an effective way to constrain certain cosmological parameters, such as the density parameter (\\cite{ob1};\\cite{bb}) or the dark energy parameter and equation--of--state (\\cite{Majumdar};\\cite{Wang}).  The present day cluster abundance is degenerate in the matter density and amplitude of the power spectrum; evolution of the abundance breaks this degeneracy.  Constraints obtained in this manner are complementary to others that essentially rely on determinations of cosmological distances; for example, those found by observations of supernovae type Ia or by measurements of cosmic microwave background anisotropies.  Cluster mass is, however, difficult to measure directly, and in practice one seeks a direct observable that is closely related to virial mass. Lensing surveys would seem the most suited to the task, as the effects of lensing are of course directly related to mass (albeit projected along the line--of--sight). However, cluster mass estimates from weak and strong lensing remain controversial as they suffer from several systematic uncertainties (e.g. projection effects).  Among X--ray observables, intracluster gas temperature is expected to be tightly correlated with virial mass, an expectation borne out by simple hydrostatic considerations as well as by numerical simulations~(\\cite{evrard}; \\cite{bn}). X-ray luminosity, on the other hand, is a much less robust mass indicator, despite being significantly easier to measure, because it depends on the density profile of the intracluster gas, the physics of which is currently difficult to model.  \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics{0687fig1.ps}} \\caption{The predicted X--ray temperature function at redshifts $z=0.05$ (black), 0.33 (light grey) and 0.5 (dark grey) for a flat high density model (left; $\\OmL=0.12$) and a concordance model (right; $\\OmL=0.7$), both fitted to the local ($z \\sim 0.05$) temperature function.} \\label{fig:dndT} \\end{center} \\end{figure}  With a calibrated $T_{\\rm x}-M$ (temperature--mass) relation, the mass function can be translated into an observable cluster temperature distribution function $dn/dT$.  The exact $T-M$ relation to use is of course a key ingredient, one that may be addressed, for example, using numerical simulations, or directly from detailed observations that determine both cluster mass and temperature (e.g.~\\cite{jukka}). Figure~\\ref{fig:dndT} compares temperature function predictions for a high matter density and a flat model, both normalized to the present--day, observed $dn/dT$ (e.g. \\cite{ha}; \\cite{edge}; \\cite{b2000}; \\cite{ikebe01}). Evolution toward higher redshift is strikingly different in two cosmological models {\\it differing by their matter content alone} and illustrating the above argument.  Despite this obvious promise, it has been difficult to use $dn/dT(z)$ to explore the cosmological model because direct measurements of $T_{\\rm x}$ at $z>0$ require long integration with satellite observations.  Prior to the launch of Chandra and XMM-Newton, the number of measured cluster temperatures at high redshifts was very small.  This is attested to by the fact that the most distant determination of $dn/dT$ (\\cite{h00}) to date was based on a sample of only fourteen $z>0.3$ clusters  (mean $\\sim$0.38) from the Einstein Medium Sensitivity Survey (EMSS; \\cite{emss}).  An alternative, but related approach is to apply a luminosity--temperature ($L_{\\rm x}-T_{\\rm x}$) relation to a flux--limited sample, thereby obtaining either the temperature function, or a redshift distribution at given temperature (\\cite{ob2};\\cite{sbo};\\cite{reichart};\\cite{borgani99};\\cite{borgani01}). The advantage of this approach is that it does not require detailed X-ray spectroscopy of all the clusters in a flux-limited sample to establish the $L_{\\rm x}-T_{\\rm x}$ relation and its evolution.  There are now several flux limited catalogs of medium to high redshift clusters, in addition to the EMSS, to which one can apply this technique. These include those based on data from the ROSAT All Sky Survey (e.g. \\cite{MACS}; \\cite{Zhang}), NEP (e.g. \\cite{nepg}; \\cite{neph}); and from the ROSAT pointing archive (e.g. \\cite{brightsharc2}; \\cite{southern}; \\cite{Vikh}; \\cite{Mullis03}; \\cite{WARPS}; \\cite{RDCS}).  \\subsection{The XMM--Newton $\\Omega$--Project} The goal of the XMM--Newton $\\Omega$--project (\\cite{Bartlett}) is to increase the number of high quality X-ray cluster temperature measurements at $z>0.3$, and thus enhance the scientific yield from the various ROSAT surveys. We describe below XMM-Newton observations of eight ROSAT clusters at $z>0.4$ (median of $z=0.54$) performed as part of the Guaranteed Time programme. Related observations conducted during the open time programmes of seven $0.3<z<0.4$ clusters are discussed elsewhere (\\cite{opentime1}; \\cite{opentime2}). Cosmological interpretation is presented in \\cite{XMM2}. In Sections 2 \\& 3 we describe the observations and data analysis techniques. In Section 4 we discuss each of the eight clusters in turn. In Section 5 we present our $L_{\\rm x}-T_{\\rm x}$ relation and compare it to previous work.  This paper serves also to introduce the Project, and represents an opportunity to provide detailed descriptions of analysis techniques used in the Project and which can be be used generally for XMM-Newton observations of cluster targets. Except where explicitly stated elsewhere we use a $\\OmM=1$, $\\OmL=0$, $q_0=0.5$ model with $H_0=50$~km~s$^{-1}$~Mpc$^{-1}$ that has been most frequently used in the past as the parameter set for X-ray cluster studies. When necessary for examining cosmological implications we correct our results to a concordance model. ",
        "conclusions": "\\subsection{Luminosity-Temperature Relations} \\label{sec:LT}  Table~\\ref{tab:L} summarizes our analysis of the eight cluster observations described above. The mean elemental abundance ($Z=0.28\\pm 0.08$), core radius ($r_c=195\\pm50 h_{50}^{-1}$ kpc) and $\\beta$ ( $\\beta=0.70\\pm0.05$), are consistent with typical values measured at low redshift (e.g. \\cite{fukazawa};\\cite{jonesforman};\\cite{Mohr}). The clusters were derived from well understood surveys (SHARC and 160SD) and thus should be representative of the cluster population as a whole at these redshifts. We have used our observations to determine the $L_{\\rm x}-T_{\\rm x}$ relation at $0.45<z<0.62$ and to investigate evolutionary effects. We plot luminosity versus temperature, with one sigma errors, for the 8 cluster targets in Figure~\\ref{fig:lxtx}, using the values T$^{cf}$, L$^{cf}$ of Table 5. The errors are clearly dominated by those of the temperature measurements.  Characterizing the $L_{\\rm x}-T_{\\rm x}$ relation as \\begin{center} \\begin{equation} L_{\\rm x} = L_{6} \\left(\\frac{T}{6 keV} \\right)^{\\alpha} % \\end{equation} \\end{center} and assuming an EdS cosmology, we find $\\alpha$=2.7 $\\pm$0.4, and $L_{6}=15.9 ^{+7.9}_{-5.5} \\times 10^{44}$\\ergs~ using the bisector variant of the BCES fitting package (\\cite{BCES}). This relation is shown as the solid line on Figure~\\ref{fig:lxtx}. The slope of the relation is similar to most previous measurements, e.g. for an EdS cosmology, $\\alpha=2.64\\pm0.27$ (\\cite{markevitch}); $\\alpha=2.33\\pm0.43$ (\\cite{allen98}); $\\alpha=2.88\\pm0.15$ (\\cite{MonGus}); $\\alpha=2.47\\pm0.14$ (\\cite{Ikebe}); $\\alpha=2.82\\pm0.32$ (\\cite{novicki}).    \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics{0687fig17.ps}} \\caption{\\label{fig:lxtx} $L_{\\rm x}-T_{\\rm x}$ relation for an EdS model ($q_0=0.5$) and $H_0$=50 ~km~s$^{-1}$~Mpc$^{-1}$.  The solid line is the best fit for our sample after cooling flow excision, the dotted line is the relation from Markevitch (1998) for low red-shift clusters.} \\end{center} \\end{figure}    To characterize the possible evolution with $z$, we assume \\begin{center} \\begin{equation} L_{\\rm x} = L_{6}\\left(\\frac{T}{6 keV} \\right)^{\\alpha} (1+z)^{A } \\end{equation} \\end{center} Adopting the low redshift $L_{\\rm x}-T_{\\rm x}$ relation from Markevitch (1998) ($\\alpha=2.64\\pm0.27$, $L_{6}=12.44\\pm 1.08 \\times 10^{44}$\\ergs) we performed a $\\chi^{2}$ minimization on our data points corrected by a $(1+z)^{A }$ term.  The Markevitch (1998) relation was derived from ROSAT and ASCA observations of 30 clusters at $0.04<z<0.1$ ($\\bar z=0.05$) and is plotted as a dotted line in Figure~\\ref{fig:lxtx}.  On average, the XMM-Newton data points lie away from the Markevitch (1998) relation, suggesting an evolutionary effect whereby clusters of the same temperature were more luminous in the past.  The best fit value for $A $ is 0.68$\\pm$0.26 ($\\chi_{red}$=1.2). We have also determined the BCES best fit for our sample when all the points are corrected with this evolution to a single redshift (low redshift $z=0.05$ of Markevitch) bin, and find $\\alpha=2.71\\pm$0.22, and $L_{6}=12.5^{+4.9}_{-3.5} \\times 10^{44}$\\ergs. We note that we have chosen to compare with Markevitch (1998), rather than with the recent study of 82 clusters with ASCA temperatures by Ikebe et al. (2002), because the  Markevitch (1998) method to derive cluster temperatures is closer to our own (Section~\\ref{sec:MarkVikh}), thus we try to minimise any systematic effects from different data treatment. In Figure~\\ref{fig:lxtx2} we re-plot our data and the Markevitch (1998) relation after converting them to a common cosmology of $H_0=70$ ~km~s$^{-1}$~Mpc$^{-1}, \\OmL=0.7, \\OmM=0.3$. In this cosmology we find $A $=1.52$^{+0.26}_{-0.27}$.  For comparison, we include the data points obtained by Vikhlinin et al. (2002) from a compilation of Chandra observations of high-$z$ clusters ($0.39<z<1.26$) on Figure~\\ref{fig:lxtx2} (after adjusting to $H_0=70$ ~km~s$^{-1}$~Mpc$^{-1}$).  From this Figure, it is clear that both the XMM-Newton and Chandra points lie away from the Markevitch (1998) relation and that there is no obvious systematic offset between the XMM and Chandra data (see Section~\\ref{sec:MarkVikh}). Vikhlinin et al. (2002) also used Markevitch (1998) as a low redshift benchmark to investigate $L_{\\rm x}-T_{\\rm x}$ evolution. From their data, Vikhlinin et al. (2002) found $A=0.6\\pm0.3$ for $\\OmL$=0, $\\OmM$=1 and $A=1.5\\pm0.3$ for a $\\OmL=0.7, \\OmM=0.3$ cosmology.       The value of $A $ has important implications for our understanding of structure formation and cluster evolution. In a self-similar model (Kaiser, 1981) cluster X-ray properties are driven by gravitational processes, such as shock heating, and a value of A$\\sim$1.5 would be expected in most cosmological models.  A self-similar EdS model predicts exactly 1.5, and the value is only slightly different in a low density concordance model, the correction due to the redshift dependence of the virial density (\\cite{bn})  being small. In a model where the cluster X-ray properties are influenced by non-gravitational processes, such as energy injection by AGN's or supernovae, lower values of $A $ are predicted (e.g. \\cite{tozzi}). The evolution deduced from our observations assuming an EdS model is significantly below the predicted value.  Adopting a low density concordance model, on the other hand, leads to a value consistent with predictions. We conclude that there is now evidence from both XMM-Newton and Chandra for an evolutionary trend in the $L_{\\rm x}-T_{\\rm x}$ relation.   Previous studies % of $L_{\\rm x}-T_{\\rm x}$ evolution based on either ASCA (e.g. \\cite{Mush97}; \\cite{allen98}; \\cite{novicki}; Sadat et al., 1998), ROSAT PSPC (e.g. \\cite{Fairley}) or Chandra (e.g. \\cite{borgani01}; \\cite{holden}) temperature measurements found results consistent with no evolution (\\cite{sbo}). When these studies quote values for $A$, these values are generally smaller than those measured by us or Vikhlinin et al. (2002), but are usually still at least one sigma away from $A =0$. Sadat et al. (1998) found a positive evolution $A =0.5\\pm 0.3$ for $\\OmM=1,\\OmL=0$ and Novicki, Sornig, \\& Henry (2002) $A =1.1\\pm1.1$ for $\\OmM=1,\\OmL=0$, and $A=2.1\\pm1.0$ for $\\OmL$=0.7, $\\OmM$=0.3 respectively.  It is clear that we are  beginning to probe evidence that  $L_{\\rm x}-T_{\\rm x}$ evolution, although many more clusters need to be studied, and systematic biases (see below) examined in detail, before solid conclusions regarding structure formation models can be drawn.  \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics{0687fig18.ps}} \\caption{$L_{\\rm x}-T_{\\rm x}$ relation for $H_{o}$=70 $\\OmL=0.7, \\OmM$=0.3. The dotted horizontal lines are the CHANDRA sample reported by \\cite{Vikh_hiz}, and the dashed diagonal line the Markevitch (1998)  \\LxTx relation corrected to this cosmology. The  solid crosses indicate the results from the clusters in our study, corrected for this cosmological model, and the diagonal solid line is the best fit for our data.}  \\label{fig:lxtx2} \\end{center} \\end{figure}  \\subsection{Data Treatment and Systematics} \\begin{figure*} \\begin{center} \\begin{tabular}{cc} \\includegraphics[scale=0.5]{0687fig19a.ps}&\\includegraphics[scale=0.5]{0687fig19b.ps}\\\\ \\includegraphics[scale=0.5]{0687fig19d.ps}&\\includegraphics[scale=0.5]{0687fig19c.ps}\\\\ \\end{tabular} \\caption{\\label{fig:sb_i}  Surface brightness distribution of clusters, compared with best $\\beta$--model fit (dotted line). Reading clockwise from top left: RXJ1120.1, RXJ1334.3, RXJ0337.7 and RXJ0505.3.}  \\end{center} \\end{figure*}  We have made every effort to ensure that the $L_{\\rm x}-T_{\\rm x}$ relation presented above is robust. We have adopted a uniform approach to the spatial and spectral analysis of the eight clusters in the sample. We have excised regions with clear signs of point source contamination before extracting spectra and surface brightness profiles. We have not used $\\beta$-models to calculate total cluster count rates, but have rather summed up the counts within a viral radius. We have tried to ensure that our approach to background subtraction does not bias the measured cluster parameters (see below).  We have used measured temperatures to make conversions between count rates and fluxes and, where possible, we have used data from all three EPIC cameras to derive $L_{\\rm x}$ and $T_{\\rm x}$ values. For RXJ0847.2 and RXJ1325.5, data were not available from all three cameras. However, we are confident that the derived quantities for these two clusters are robust, as we have shown, using the RXJ1120.1 observation, that PN-only or MOS-only spectra are in excellent agreement.  The background subtraction has been complicated by the extended nature of the targets and the high proton flare contamination in some of the observations. We used  background template files during the spatial analysis. These files covered a limited energy range (0.3--4.5~keV), to minimize any calibration uncertainties. They were filtered using the same rate filtering criteria as their respective cluster observation and were normalized to it using the particle background rates in the CCD regions outside the FOV. For various reasons (see Section~\\ref{sec:back}), we did not use the background templates for the spectral analysis. Instead, we adopted an in-field background subtraction technique. With this method comes the concern that the proton background can be over-vignetted. We have checked for this by applying a different method, the double subtraction technique, to the RXJ1120.1 and RXJ1334.3 observations. We found that the results from both techniques agree within the statistical errors. For the two observations with the worst proton flare contamination, RXJ1354.2 and RXJ1701.3, we applied a 2\\% scaling to the in-field background spectrum to compensate for potential over-vignetting of the proton background.  In Figure~\\ref{fig:protvig} we demonstrate that this over-vignetting should not present a problem for the other clusters in our sample.  Despite all these quality controls, we still cannot rule out the possibility of there being some systematic bias in our results.  We discuss possible sources of such bias and the potential impact on our claimed $L_{\\rm x}-T_{\\rm x}$ evolution below.  \\subsubsection{Impact of Cooling Core Clusters} \\label{sec:cflows} \\begin{figure*} \\begin{center} \\begin{tabular}{cc} \\includegraphics[scale=0.5]{0687fig20a.ps}&\\includegraphics[scale=0.5]{0687fig20b.ps}\\\\ \\includegraphics[scale=0.5]{0687fig20d.ps}&\\includegraphics[scale=0.5]{0687fig20c.ps}\\\\ \\end{tabular} \\caption{\\label{fig:sb_ii}  Surface brightness distribution of clusters, compared with best $\\beta$--model fit (dotted line). Reading clockwise from top left: , RXJ0847.2, RXJ1325.5, RXJ1701.1 and RXJ1354.2.} \\end{center} \\end{figure*}  Cooling  cores have the joint effect of boosting the cluster luminosity and of reducing the measured global cluster temperature. These effects have been shown to introduce scatter in the $L_{\\rm x}-T_{\\rm x}$ relation and to modify its slope (\\cite{markevitch}; \\cite{allen98}). Various approaches have been taken to alleviate the effects of cooling cores on measured $L_{\\rm x}-T_{\\rm x}$ relations, e.g. by excluding cooling core clusters from the analysis ,\\cite{MonGus}, by fitting two temperature models (\\cite{Ikebe}), by masking the central regions from obvious cooling core clusters (\\cite{Vikh_hiz}), or by masking the central regions from all clusters (\\cite{markevitch}). Based on the images (Figures~\\ref{fig:image_i} \\& \\ref{fig:image_ii}) and radial profiles (Figures~\\ref{fig:sb_i} \\& \\ref{fig:sb_ii} ) of the eight clusters in our sample, we do not expect cooling cores to have much, if any, impact on the measured $L_{\\rm x}$ and $T_{\\rm x}$ values; only RXJ1701.3 and RXJ1325.5 show evidence of a central surface brightness excess. We have investigated this by comparing the results before and after excising a circular region \\footnote {The radius of the excised region was 50 $h_{50}^{-1}$ kpc except in the case of RXJ1701.3 and RXJ1325 where $\\sim$120kpc was used, see sections 3.6 and 3.7} from the cluster cores.  The results of the re-analysis can be found in Table~\\ref{tab:L}, (rows 4 and 5).  Temperatures changed very little -- always within the statistical error -- after we excised the central $r=50 h_{50}^{-1}$ kpc region. In two cases, ( RXJ1120.1 and  RXJ1701.3), the luminosity changes by more than $1\\sigma$, but still by only 5\\% and 9\\% respectively. Without the excision the best fit \\LxTx relation is $L_{6}=17.6 ^{+5.3}_{4.1}\\times 10^{44}$\\ergs~ and $\\alpha$=2.66$\\pm$0.25 for an $\\OmL=0.0, \\OmM=1.0$, $H_0=50$ ~km~s$^{-1}$~Mpc$^{-1}$ cosmology at a typical redshift $z \\sim 0.55$.  In summary, we find the $L_{\\rm x}-T_{\\rm x}$ relations before and after the core excision are consistent within the errors. However, we caution that some residual cooling flow contamination may still remain after this excision; $50 h_{50}^{-1} kpc \\simeq 14''$, corresponding to an encircled energy fraction of a little under 50\\% of the on-axis mirror PSFs.  \\subsubsection{Incompleteness and Flux Errors in the ROSAT Catalogs} \\label{sec:sharc-flux}  Our $L_{\\rm x}-T_{\\rm x}$ evolution result implies that high redshift clusters in our sample are more luminous than clusters of the same temperature at lower redshift.  We interpret this as evidence for a general evolutionary trend in the \\LxTx relation, but it might also reflect an underlying selection bias in the cluster catalogs from which the sample was drawn, e.g. if the SHARC and 160SD surveys were biased toward merging systems. It is possible  that the luminosity and temperature of clusters can change dramatically during the process of a major merging event (\\cite{ricker}).  Based on the XMM-Newton images of the 8 clusters studied here (Figures~\\ref{fig:image_i} \\& \\ref{fig:image_ii}), such a bias does not appear to exist; only RXJ1354 has evidence for a possible double component, and excluding this from the best fit $L_{\\rm x}-T_{\\rm x}$ relation changes the luminosity normalisation only $\\sim$2\\%. As long as the angular resolution of observations allows sufficient discrimination against obvious mergers, then presumably an ensemble sample of local and distant clusters are similarly affected by merger boosting of luminosity.  Moreover, we find no evidence for any incompleteness in the SHARC and 160SD surveys, at least in the small area covered by these observations. No new clusters were found in the fields surrounding the seven SHARC clusters in our sample.  In the field of the 160SD cluster, RXJ0847.2, we have discovered a new X-ray cluster (XMMU J084701.8+345117). However this cluster was too faint to have been included in the original 160SD survey.  We estimate the ($r<r_v$) flux of this object to be $1.3\\times 10^{-14}$ \\ergcms (based on a fitted temperature of $T_x=1.8$ keV and an assumed redshift of $z=0.56$), compared to the 160SD flux limit of $3.7\\times 10^{-14}$ \\ergcms, so this cluster is too faint to have been included in the original 160SD survey. In summary we do not find evidence for any intrinsic bias in the SHARC or 160SD survey selection functions that could mimic $L_{\\rm x}-T_{\\rm x}$ evolution.  Another possible way to mimic the claimed evolution would be to systematically overestimate the $L_{\\rm x}$ values. We have investigated this by comparing our flux measurements to those published in several works (\\cite{brightsharc2}; \\cite{southern}; \\cite{Vikh}). We report here that a systematic offset does appear to exist, but in the opposite direction to that needed to mimic evolution.  The XMM-Newton fluxes quoted in Table~\\ref{tab:L} are all lower than previously determined values, typically by 20\\% but as much as by 40\\% in the case of RXJ0505.3.  We note that this discrepancy persists regardless of the ROSAT catalog from which the cluster was selected. The average discrepancy between the Southern SHARC (4 clusters), Bright SHARC (3 clusters) and 160SD (3 clusters) fluxes and the XMM-Newton fluxes is 22\\%, 21\\% and 16\\% respectively.  Improvements in point source rejection, count rate to flux conversions and the signal to noise have meant that the typical statistical error on the XMM-Newton determined fluxes in Table~\\ref{tab:L} is $\\simeq 3\\%$, compared to $\\simeq 10\\%$ for the ROSAT determinations. Therefore, even accounting for a possible offset in the absolute cross calibration of the two observatories, it is clear that a systematic flux offset does exist.  There are likely to be several reasons for this offset. Of these, the use of an incorrect spatial model, to convert between an aperture flux and a total flux, is probably the most significant. As shown in \\cite{brightsharc2}, the use of a generic ($\\beta=0.67$, $r_c=250 h_{50}^{-1}$ kpc), rather than best fit, $\\beta$-model resulted in a typical enhancement of 10\\% in the total cluster flux. In our XMM-Newton analysis, we do not use the $\\beta$-model fits to determine the total flux. Instead we directly sum up the counts within a virial radius. For some clusters this difference in approach can explain the entire ROSAT to XMM-Newton discrepancy. For example, in the case of RXJ1334.3, our estimate of $r_v$ is almost identical to the $r_{80}$ radius within which 80\\% of the total flux from a generic $\\beta$-model would fall. \\cite{brightsharc2} divided the ROSAT counts within $r_{80}$ by 0.8 to estimate the total count rate. For RXJ1334.3 this division would artificially boost the flux by 25\\% which more than accounts for the 22\\% the discrepancy between the XMM-Newton and Bright SHARC values.  For the four Southern SHARC clusters in our sample, we have recalculated the total ROSAT PSPC fluxes using the $\\beta$ and $r_c$ values given in Table~\\ref{tab:L} (in \\cite{southern} generic values were used; $\\beta=0.67$, $r_c=250 h_{50}^{-1}$ kpc).  By doing so, we improve the consistency between the ROSAT and XMM-Newton fluxes to within 5\\% for RXJ1354.2 and RXJ1325.5 and within 15\\% for RXJ0337.7. For RXJ0505.3 there remains a large, 35\\%, discrepancy which requires further investigation.  For the three Bright SHARC clusters in our sample, we have gone back to the original \\cite{brightsharc2} data and measured count rates inside $r_v$. Doing so reduces the ROSAT determined flux by $\\simeq 30\\%$ in the case of both RXJ1120.1 and RXJ1334.3, i.e. so that they are $\\simeq10\\%$ lower than the values quoted in Table~\\ref{tab:L}. In the case of RXJ1701.3, the measured ROSAT flux actually increased slightly when measured inside $r_v$ rather than $r_{80}$. On closer examination of the ROSAT data, it was apparent that the nearby QSO and cluster (see Figure~\\ref{fig:image_ii}) were contaminating both the source and background apertures in the PSPC image. When the source and background were accumulated only from the (source free) region to the West of the cluster, the ROSAT flux within $r_v$ dropped to within $\\simeq10\\%$ of the XMM-Newton determined value. These examples demonstrate the importance of using XMM-Newton to re-calculate fluxes and luminosities for high redshift clusters detected at low signal to noise by ROSAT.  We stress that, despite these flux uncertainties, the SHARC and 160SD catalogs are still fair representations of the underlying cluster population and can still be used to probe the $L_{\\rm x}-T_{\\rm x}$ relation.  It is also important to note that the these flux uncertainties do not apply to the ROSAT observations used by \\cite{markevitch} to determine $L_{\\rm x}$ values for the 30 low redshift clusters in his $L_{\\rm x}-T_{\\rm x}$ analysis. Those observations have exquisite signal to noise, so that the $L_{\\rm x}$ values derived from them will be limited only by the absolute calibration of the instrument.  The fact that we find a \\LxTx relation consistent with the low redshift value argues against a bias in our selection towards the more luminous tail of the cluster population; for either \\begin{itemize} \\item all the clusters must be biased in the same way, or \\item they are biased in just such a way as to offset any real evolution in the value of the slope. \\end{itemize}  Nevertheless the reader might be cautioned that should there remain a common bias for all cluster samples, despite these cross-checks, it is possible the clusters selected represent the brightest portion of the intrinsic high-z sample, and thus remain on the high luminosity end of the true population, hence mimicking the luminosity evolution we see. Further validation of the ROSAT data with CHANDRA and XMM--Newton should help to close this issue.   \\subsubsection{Comparison with Markevitch (1998) and Vikhlinin et al. (2002)} \\label{sec:MarkVikh}  We come to very much the same conclusions regarding \\LxTx evolution as did \\cite{Vikh_hiz}. Given that these conclusions provide evidence for significant evolution, contrary to previous analyses, it is important to investigate whether this concordance is genuine or coincidental.  Both studies rely on \\cite{markevitch} for the low redshift comparison, so let us first consider how our XMM-Newton clusters might appear either too luminous, or too cool, compared to the \\cite{markevitch} sample:  \\begin{enumerate}      \\item  {  Markevitch 1998 measured cluster fluxes by summing the count rate inside circular apertures. However, rather than using a cluster specific virial radius, he used a fixed metric aperture of $r=1 h_{50}^{-1}$ Mpc and ignored any flux that might lie outside this radius. However, this would be a very small effect; Vikhlinin et al. (2002) estimate that the no more than 4\\% of the total flux would be missed by adopting a fixed $r=1 h_{50}^{-1}$ Mpc aperture.  }  \\item{  When Markevitch 1998 excised the central $r=50 h_{50}^{-1}$ kpc region from all the clusters in his low redshift sample before measuring the flux, he  applied a multiplication factor of 1.06 to account for flux falling inside the excised region. We carried out an excision and flux lost correction  on our eight clusters with a cluster-specific surface brightness model. (Section~\\ref{sec:cflows}, and Table~\\ref{tab:L}), we found that for six of the eight clusters, the measured $L_{\\rm x}$ values dropped slightly. The difference in our techniques could systematically affect cluster brightness, compared with  the low redshift counterparts, but only at the few percent level.  }   \\item {   Uncertainties in the cross calibration between XMM-Newton and the instruments used for the \\cite{markevitch} analysis (ROSAT and ASCA) may mimic evolution in the $L_{\\rm x}-T_{\\rm x}$ relation. At this time it is not possible to rule out that possibility; to date, a full comparison of XMM-Newton and ASCA determined cluster temperatures has not been carried out.  }    \\item {  Ikebe et al., 2002 have used a different approach to \\cite{markevitch} to analyse ASCA observations of low redshift clusters. They measure temperatures that are on average lower than those measured by \\cite{markevitch}, with the trend becoming more pronounced as $T_{\\rm x}$ increases. However, this should not impact our conclusions regarding $L_{\\rm x}-T_{\\rm x}$ evolution, given that we used a similar technique to \\cite{markevitch} to measure $T_{\\rm x}$.  }     \\item { Some estimate of the effect of external systematic error {\\em could} be taken by considering the value of $A $ in fitting to a different low-z \\LxTx sample. We have done so using the \\cite{MonGus} relation and find $A=0.95 \\pm 0.2$ (again in the Einstein de Sitter cosmology). Therefore, despite a possible contamination in that sample from cooling flows we still see evidence for evolution } \\end{enumerate}      We note that the first two concerns listed above do not apply to the Vikhlinin et al., (2002) study, as they deliberately adopted the same methodology as \\cite{markevitch} and the last three concerns are not likely to be important either; Vikhlinin et al., (2002) have carried out a cross comparison of Chandra and ASCA $T_{\\rm x}$ values for $\\simeq 20$ low redshift clusters and find them to be in agreement at better than the 5\\% level, with no systematic offset. Despite this, we should still consider the possibility that both studies (ours and \\cite{Vikh_hiz}) might be pointing to the wrong value of $A $. At present we cannot quantify how factors such as uncertainties in the cross-calibration of Chandra and XMM-Newton, especially at low energies, might impact our conclusions. We note in particular that Vikhlinin et al., (2002) applied an empirical factor of 0.93 to improve the cross-calibration of the ACIS front- and back-illuminated CCDs below 1.8keV. Other differences between the observatories may also be important. For example, it will be easier to account for point source and cooling core contamination using Chandra data, because of the improved spatial resolution (see for example the case of RXJ1701.3). Alternatively, it should be easier to correctly account for the various particle (cosmic ray and proton) backgrounds using XMM-Newton data, because its CCD's cover a larger area both inside and outside the FOV. In summary, although we cannot rule out the possibility that a combination of factors have lead us to measure an artificially large value for the $A $ parameter, we believe that our XMM measurements provide clear evidence for evolution in the $L_{\\rm x}-T_{\\rm x}$ relation.     \\small \\begin{table*} \\begin{center} \\begin{tabular}{l l l l l l l l l}\\hline\\hline &RXJ0337.7&RXJ0505.3&RXJ0847.2&RXJ1120.1&RXJ1325.5&RXJ1334.3&RXJ1354.2&RXJ1701.3\\\\ \\hline centroid RA & 03:37:45&05:05:19& 08:47:10&11:20:07&13:25:37&13:34:20&13:54:17&17:01:24\\\\ centroid Dec& -25:22:32&-28:48:50&+34:48:54&+43:18:04&-38:25:44&+50:31:02&-02:21:44&  +64:14:12 \\\\ \\hline &&\\multicolumn{5}{c}{Cooling Flow Excised}&&\\\\ \\hline $T_{x}^{cf}$ (keV)& 2.52 & 2.56 & 3.91 & 5.35 & 3.77 & 4.98 & 3.86 & 4.8\\\\ & $^{+0.36}_{-0.32}$ & $\\pm$0.3 & $^{+0.5}_{-0.35}$ & $^{+0.42}_{-.32}$ & $^{+0.4}_{-0.36}$ &$^{+0.26}_{-0.32}$ & $^{+0.62}_{-0.55}$ & $^{+1.9}_{-1.3}$\\\\ \\hline $L_{bol}^{cf}$~10$^{44}$ & 1.91 & 1.97 & 3.89 & 13.7 & 4.2 & 9.47 & 5.3 & 10.2\\\\ erg s$^{-1}$ & $\\pm$0.17 & $\\pm$0.16 & $\\pm$0.7 & $\\pm$0.23 & $\\pm$0.8 & $\\pm$0.24 & $\\pm$0.4 & $\\pm$1.1 \\\\\\hline &&\\multicolumn{5}{c}{No Cooling Flow Excision}&&\\\\ \\hline $T_{x}$ (keV)& 2.6 & 2.5 & 3.62 & 5.45 & 4.15 & 5.20 & 3.66 & 4.5\\\\ & $\\pm$ 0.35& $\\pm$0.3 & $^{+0.58}_{-0.51}$ & $\\pm$0.3 &$^{+0.4}_{-0.3}$ & $^{+0.26}_{-0.28}$ & $^{+0.6}_{-0.5}$ & $^{+1.5}_{-1.0}$\\\\ \\hline $L_{bol}$~10$^{44}$ & 1.97 & 2.03 & 4.07 & 14.4 & 4.5 & 9.59 & 5.18 & 11.1\\\\ erg s$^{-1}$ & $\\pm$0.1 & $\\pm$0.11 & $\\pm$0.2 & $\\pm$0.2 & $\\pm$0.7 & $\\pm$0.27 & $\\pm$0.3 & $\\pm$1.0 \\\\\\hline Flux (0.5--2)& 4.37 & 5.64 & 7.04 & 24.5 & 8.0 &14.1 &9.8 &24 \\\\ 10$^{-14}$cgs & $\\pm$0.2 & $\\pm$0.2 & $\\pm$0.3 & $\\pm$0.54 & $\\pm 0.03 $ & $\\pm$0.2 & $\\pm$0.5 & $^{+1.5}_{-1.2}$\\\\ \\hline \\nh (10$^{20}$ )&0.99&1.5&3.2&2.1&4.8&1.05&3.4&2.6 \\\\ (atom cm$^{-2}$)&&&&&&&& \\\\\\hline $\\beta$ & 0.76 & 0.66 & 0.81 & 0.77 & 0.64 & 0.66 & 0.68 & 0.64\\\\ & $^{+0.08}_{-0.04}$ &$^{+0.05}_{-0.04}$ &$\\pm$0.045&$\\pm$0.03&$^{+0.09}_{-0.07}$&$\\pm$0.02&$\\pm$0.055&$\\pm 0.05$\\\\\\hline r$_{c}$ (kpc)&145&164&307&209&115&154&248&204\\\\ h$_{50}^{-1}$&$\\pm$18&$\\pm$ 17&$\\pm$30&$^{+9}_{-8}$&$\\pm 20$&$\\pm$10&$^{+36}_{-26}$&{\\em fixed}\\\\\\hline Abundance &0.38&0.17&0.30&0.47&0.31&0.15&0.25&0.24\\\\ (Z)&$\\pm$0.09 &$\\pm$0.08&$\\pm$0.28&$\\pm$0.09&$^{+0.19}_{-0.17}$ & $\\pm$0.08 & $\\pm$0.14 & $\\pm 0.2$\\\\\\hline $z$ (optical) &0.577 &0.51&0.56&0.60&0.445&0.62&0.551&0.45\\\\ \\hline Normaln.& 1.68 & 2.12 &2.65 &9.42 & 2.7 &5.61 &4.01 & 5.58 \\\\ mekal 10$^{-4}$& $\\pm$0.10 & $\\pm$0.08&$^{+0.12}_{-0.11}$ &$^{+0.23}_{-0.17}$ &$^{+0.08}_{-0.07}$ &$\\pm$0.16 &$^{+0.14}_{-0.23}$ &$^{+0.50}_{-0.40}$\\\\ \\hline Radius ($''$)& 120 & 120 & 120 & 145 &90 & 120 & 120 & 120\\\\ Spect. extract.&&&&&&&&\\\\ \\hline Fract$_{r}$ counts  & 0.90 & 0.87 & 0.79 & 0.965 & 0.52 & 0.84 & 0.83 & 0.81\\\\ within $r_v$&&&&&&&&\\\\ \\hline  \\end{tabular} \\caption{\\label{tab:L} Summary of cluster parameters for EdS model H$_{o}$=50, q$_{o}$=0.5. Spectral fitting errors, L,$\\beta $ and r$_{c}$ are 1$\\sigma$ on one parameter. Fluxes are the measured, absorbed fluxes in ROSAT band. Fract$_{r}$ is the fractional correction made from the spectral extraction region, to the total counts within the $r_{v}$ radius after accounting for point source excision, inter-chip gaps loss etc.. The 3rd \\& 4th rows summarise the data for the case when the core region has been excised from the cluster core. (Except for RXJ1701.3, where following the CHANDRA data, we excise a cooling flow enhancement to 120kpc). } \\end{center} \\end{table*} \\normalsize"
    },
    "0311/astro-ph0311129_arXiv.txt": {
        "abstract": "A first analysis of a deep blind HI survey covering the southern Zone of Avoidance plus an extension towards the north ($196\\deg \\le \\ell \\le 52\\deg$) obtained with the Multibeam receiver at the 64\\,m Parkes telescope reveals slightly over a thousand galaxies within the latitude completeness limit of $|b| \\le 5\\deg$. The characteristics and the uncovered large-scale structures of this survey are described, in particular the prominence of the Norma Supercluster, the possible cluster around PKS\\,1343$-$601 (both in the Great Attractor region), as well as the Local Void and the clustering in the Puppis region.  In this blind HI survey, HIZOA\\,J0836$-$43, one of the most massive spiral galaxies known to date was discovered ($M_{\\rm HI} = 7.3 \\cdot 10^{10}{\\rm M_\\odot}$, $M_{\\rm T} = 1.1 \\cdot 10^{12}{\\rm M}_\\odot$; H$_0 = 75$\\,km/s/Mpc). Although of similar mass as Malin 1-like objects, this galaxy does not share their typical low-surface brightness properties but seems an exceptionally massive but normal, high-surface brightness, star-forming galaxy. ",
        "introduction": "Based on the intensive multi-wavelength surveys of the Zone of Avoidance (ZOA) in the last decade -- aimed at addressing cosmological questions about the dynamics of the Local Group, respectively the possible existence of nearby hidden massive galaxies, dipole determinations based on luminous galaxies, the continuity and size of nearby superclusters, the mapping of cosmic flow fields -- enormous progress has been achieved in uncovering the galaxy distribution in this obscured region of the sky (see Kraan-Korteweg \\& Lahav 2000 for a review).  In particular, the near infrared, whole-sky homogeneous surveys such as 2MASS (see Huchra et al., these proceedings) have resulted in a magnificent reduction of the ZOA. However, even this survey does become increasingly incomplete at low Galactic latitudes in the larger Galactic Bulge (GB) area ($\\ell \\approx 0\\deg \\pm 90\\deg$), including the Great Attractor (GA) region.  Even if obscured galaxies can be identified, redshifts are difficult if not impossible to obtain at the higher extinction levels.  Because of the transparency of the Galaxy to the 21\\,cm radiation of neutral hydrogen, systematic HI-surveys are particularly powerful in mapping large-scale structure (LSS) in this part of the sky. The redshifted 21\\,cm emission of HI-rich galaxies are readily detectable at lowest latitudes and highest extinction levels and the signal will furthermore provide immediate redshift and rotational velocity information. ",
        "conclusions": ""
    },
    "0311/astro-ph0311403_arXiv.txt": {
        "abstract": "The acceleration of the expansion of the universe has deep implications for structure formation, the composition of the universe, and its fate.  Roughly 70\\% of the energy density is in a dark energy, whose nature remains unknown.  Mapping the expansion history through supernovae, mapping the geometry of the universe and formation of structure through redshift surveys, and mapping the distance to recombination through the cosmic microwave background provide complementary, precise probes of the equation of state of the dark energy.  Together these next generation maps of the cosmos can reveal not only the value today, but the redshift variation, of the equation of state, providing a critical clue to the underlying physics. ",
        "introduction": "Observations of the distance-redshift relation of Type Ia supernovae have given firm evidence of an accelerated expansion of the universe (Knop et al.\\ 2003; Tonry et al.\\ 2003). As calibratable ``standardized candles'', these supernovae are excellently suited to map the expansion history $a(t)$ due to the direct relation between the measurements and the cosmological dynamics.  The redshift of the supernova, $z=a^{-1}-1$, measures the scale factor $a$, the size of the universe when the supernova exploded relative to its current size.  The calibrated peak magnitude gives the distance, which translates to the lookback time to the explosion.  The data clearly indicate an acceleration to the expansion rather than the slowing down under gravitational attraction that was previously expected.  This gravitational repulsion is generally interpreted in terms of an additional component to the energy density of the universe, and given the name dark energy.  Since the effective gravitating mass in general relativity depends on both the energy density $\\rho$ and pressure $p$ in the combination $\\rho+3p$, such a repulsion and hence acceleration could be induced by a component with strongly negative pressure.  Characterized in terms of the equation of state ratio $w=p/\\rho$, the condition for a single component to accelerate the expansion is $w<-1/3$.  In order to achieve the acceleration deduced from the distance-redshift measurements, in a flat universe with (decelerating) matter density as well, the energy density in dark energy must amount to $\\sim$70\\% of the total. Thus the majority of the universe is composed of dark energy, determining the cosmic dynamics and the fate of the universe.  Moreover, the equation of state must be substantially negative, $w\\approx -1$.  The physics underlying the dark energy sets the equation of state, so to understand this new gravitational or high energy physics requires precise and accurate measurement of this quantity.  Is the dark energy Einstein's cosmological constant ($w=-1$ exactly), some high energy physics scalar field (often called quintessence), or a sign of modifications to gravity or the presence of extra dimensions?  Except for the cosmological constant, almost all theories predict dynamical dark energy, with an equation of state evolving with the cosmic expansion.  In fact, this dynamics, in the form of the time variation $w(z)$, contains the main clue to the new physics.  Thus the goal is to bring together astrophysics and particle physics to map the dark energy equation of state. ",
        "conclusions": "\\label{sec.concl}  Dark energy poses a fundamental mystery as to what composes the majority of the universe, dominates its dynamics through a gravitational repulsion, and determines the fate of the universe.  Constraints on an averaged equation of state quantity from impressive efforts over the last five years show that it behaves roughly like a cosmological constant.  The precision of 15\\% is already good enough that continuing measurements along these lines are unlikely to be able to detect a deviation from cosmological constant properties at 3$\\sigma$.  But powerful next generation experiments are being designed that are sensitive not to a crude approximation of a constant EOS but that can map out the dynamical physics of a time varying equation of state $w(z)$.  These will also generate extraordinary astronomical data resources.  Mapping the expansion history to $z=1.7$ through the Type Ia supernova method offers great promise.  Even more valuable gains in accuracy and precision come from working together with crosschecking and complementary methods such as CMB measurements, weak gravitational lensing, and baryon oscillations measured in the matter power spectrum by large, deep galaxy surveys.  Mapping the dark energy equation of state will give us guideposts to high energy physics, the early universe, extra dimensions, or the theory of gravity, as well as reveal to us the true nature of the universe we live in and a picture of our fate."
    },
    "0311/astro-ph0311259_arXiv.txt": {
        "abstract": "We present observations and analysis of rotation curves and dark matter halo density profiles in the central regions of four nearby dwarf galaxies.  This observing program has been designed to overcome some of the limitations of other rotation curve studies that rely mostly on longslit spectra. We find that these objects exhibit the full range of central density profiles between $\\rho\\propto r^0$ (constant density) and $\\rho\\propto r^{-1}$ (NFW halo). This result suggests that there is a distribution of central density slopes rather than a unique halo density profile. ",
        "introduction": "The last few years have seen a flurry of activity in the field of precision measurements of central density profiles in dark matter halos, as demonstrated elsewhere in these proceedings. Most of this activity has concentrated on addressing two questions: is there a unique central density profile slope?, and, more to the point, do the measurements agree with the predictions of the simulations? In fact, the observations appear to point to a substantial disagreement between the central density profiles measured in low mass, low surface brightness dwarf galaxies (e.g., de Blok et al. 2001) and the predictions of most Cold Dark Matter simulations (e.g., Navarro, Frenk, \\& White 1996, hereafter NFW; Moore et al. 1999; Jing \\& Suto 2000). The significance of this discrepancy, however, is a matter of debate. Some authors ascribe it to intrinsic systematic problems in the observations, which conspire to poorly constrain the central density slopes (e.g., Swaters et al. 2003a), while others acknowledge these problems but argue that the data provide strong enough constraints to rule out universal slopes as steep as those predicted by the simulations (e.g., de Blok, Bosma, \\& McGaugh 2003).  Can observations be used to test the predictions of cosmological simulations? In the current era of ``precision cosmology'' the answer should be most emphatically {\\em yes!} Close attention needs to be paid to the potential systematic effects, however, in order to minimize their importance. In this paper we present a series of studies of high resolution velocity fields of dwarf galaxies which we have designed to remove as much as possible the impact of systematic uncertainties. Our overall conclusion is that (as also suggested by de Blok et al. 2003) these objects appear to exhibit a range of central density profiles rather than a unique value. The discrepancy between observations and simulations is thus real, and perhaps related to the absence of baryons and their related astrophysics in the simulations (although other explanations are also possible; e.g., Ricotti 2003). In this regard, these observations should be taken as a reality check on future simulations that incorporate all the physics relevant on the small spatial scales. ",
        "conclusions": "The four galaxies in this series of studies for which we were able to retrieve reliable central density profiles appear to span the full range of behaviors; from constant density cores to NFW halos. In particular, the data for NGC~2976 clearly do not allow an NFW halo. Because we have tried to eliminate most of the systematics present in this type of measurement, we believe that this range of central density slopes is real and reflects an underlying distribution, although the caveats associated with a small sample certainly apply. To study this distribution we plan to expand the sample of observed galaxies to 10--15 objects in the near future. This larger sample will also allow us to look for correlations between central density slopes and other parameters, such as galaxy mass or the magnitude of noncircular motions.  Perhaps the most important conclusion, however, is that measurements of the central density profiles of dwarf galaxies can be accurate. Multiwavelength imaging spectroscopy is key to minimize the vulnerability of these observations to potential systematic problems, such as erroneous positioning of the spectrograph slit. At the same time, 2D high--resolution data also provides a wealth of kinematic information that would be otherwise unavailable.  Finally, the presence of measurable radial motions in two out of the four galaxies studied presents a bit of a mystery; are we observing the remnants of the processes that erased the original central cusps? There does not appear to be a clear link to asymmetries or bars in our small sample, and maintaining such motions over long time periods appears very difficult in these small objects. Furthermore, it is unclear what the effect of these motions is on the inversion the rotation curve to obtain the density profile (we explicitly assumed that they provide no ``support'', which is equivalent to ignoring them). % Expanding the sample of studied objects will allow us to determine how general this phenomenon is.  \\begin{figure} \\plotone{bolatto_a_fig4.eps} \\caption{Rotation curve of NGC~5963. Panel (a) shows the minimum disk solution, together with an NFW fit to the measurements with a concentration parameter of 20. This rotation curve is better fit by the NFW functional form than by a power law. Panel (b) shows the maximum disk solution, with $ M_*/L_I=0.7 M_\\odot/L_{\\odot I}$. The resulting dark matter halo rotation curve is well--fit by the NFW solution with R$_{\\mbox{s}}=3.5$ kpc, V$_{200}=95$ km s$^{-1}$, and R$_{200}=50$ kpc. \\label{n5963}} \\end{figure}"
    },
    "0311/astro-ph0311290_arXiv.txt": {
        "abstract": "We present a method for detecting cosmic rays in single images. The algorithm is based on simple analysis of the histogram of the image data and does not use any modeling of the picture of the object. It does not require a good signal to noise ratio in the image data. Identification of multiple-pixel cosmic-ray hits is realized by running the procedure for detection and replacement iteratively. The tests performed by us, show that the method is very effective, when applied to the images with the spectroscopic data. It is also very fast in comparison with other single image algorithms found in astronomical data processing packages. Practical implementation and examples of application are presented. ",
        "introduction": "\\label{sec1}  Cosmic ray hits cause defects in all astronomical images obtained with CCD detectors. A relatively efficient approach to removing traces of cosmic rays from such images is to use multiple frames of the same object and then combine them with some algorithm for rejection of the outlying data. Methods of this type have been presented among others by \\citet{Shaw1992} and \\citet{Windhorst1994}. They may be found in most astronomical data processing packages.  In many cases we are not able to obtain multiple images of the same object or the required time resolution prevented us from using multiple image methods. A straightforward method that could be used in such cases would be sigma clipping. In practice it is almost impossible to find a good detection threshold for such an algorithm. In fact, it may leave some obvious cosmic rays untouched, while giving numerous spurious detections and rejecting valid data at the same time. Other types of single image algorithms may also be found in most astronomical data processing packages. Most of these methods rely on the sharpness of cosmic rays, relative to the atmospheric smoothing of real images. These methods are based on some interactive learning techniques \\citep{Murtagh1992,Salzberg1995} or involve a special fitting of a model representing a real image of a star with a superimposed cosmic ray \\citep{Rhoads2000}. These methods were designed to correct direct images of the stellar objects and we find them not suitable for spectral images.  The exposure times of the spectra are usually longer than direct images and range from 600~seconds to 1800~seconds. Considerable amounts of cosmic ray events accumulate in the images during such exposures. We have constructed a simple and straightforward algorithm for detecting cosmic rays in single images. Our method does not need any model of the shape of the image features themselves. We analyze the histograms of pixel counts in small sub-frames in order to detect pixels deviating by some factor from the bulk of the pixels under consideration.  We would like to note here, that low value bad pixels tend to hide weak cosmic rays by artificially making the local variance larger than would be caused only by usual sources of noise in the image: readout noise, dark current and the Poisson noise associated with background and object signal. The cosmic rays have random distributions of their positions in the images. On the other hand, bad pixels may be found as deviating pixels at the same position in every frame taken with a given instrument. Therefore we suggest to make bad pixel correction as a separate step of the reduction process prior to cosmic rays correction.  We describe details of our algorithm in paragraph 2. In paragraph 3. we describe an implementation and the results of our tests. In paragraph 4. we present examples of application to real astronomical data. Paragraph 5 contains the summary. ",
        "conclusions": "\\label{sec5}  We have presented a cosmic-ray rejection algorithm based on a simple analysis of the histogram of the image data. The most important advantage of our method is that it does not require modeling of the image data and may {\\it apriori} be applied to any type of well sampled image data. We have checked that for the spectroscopic images, it is very effective in detecting cosmic rays, while avoiding numerous spurious detections. Our method does not require advanced and extensive computations, so it is relatively fast.  The weak point of this approach, is that bright objects may shield cosmic rays in their neighborhood from detection. This is caused by two factors. First: the sensitivity to cosmic ray events is reduced at the locations of bright objects because of the Poisson noise associated with the image photons. Second: we look for bright spots and the legitimate object may be brighter than a nearby comic ray trace.  Our suggestion is to use the algorithm presented above for spectroscopic data whenever multiple image methods cannot be employed for cosmic rays removal.  Future improvement to the presented method could be an introduction of better interpolation for the replacement of pixels affected by cosmic rays."
    },
    "0311/astro-ph0311223_arXiv.txt": {
        "abstract": "An analytical description of the time structure of the pulses induced by muons in air showers at ground level is deduced assuming the production distance distribution for the muons can be obtained elsewhere.  The results of this description are compared against those obtained from simulated showers using AIRES. Major contributions to muon time delays are identified and a relation between the time structure and the depth distribution is unveiled. ",
        "introduction": "One of the highest priorities in Astroparticle Physics has become the understanding of the origin, composition and energy spectrum of Ultra High Energy Cosmic Rays.  Cosmic rays have been observed for a long time with energies known to exceed $10^{20}~$eV \\cite{NW} but the remarkably low flux at these energies makes their detection difficult. The air showers developing in the atmosphere from the interaction of these particles can be sampled at ground level with arrays of particle detectors.  Alternatively the fluorescence light emitted by the nitrogen in the atmosphere can be collected by optical systems \\cite{FE}.  Both types of experiments have confirmed the existence of cosmic rays of such high energies although there are discrepancies between the cosmic ray fluxes inferred by the different techniques \\cite{AGASA,HiRes}.  The study of extensive air showers of the highest energies is difficult because the atmosphere is part of the detector and the interpretation of measurements is very indirect. The air shower is a complex chain of interactions resulting in a very large number of particles from which only sparse data samples are taken. From this sample, the properties of the primary particle must be deduced.  We must rely on simplified models, or on complex numerical simulations, to infer these properties.  These simulations have intrinsic and unavoidable uncertainties arising from interaction models that require untested extrapolations of low energy accelerator data.  One of the complications arises because it is often very difficult to discern physical effects due to shower fluctuations or to primary composition from effects due to uncertainties in the interaction models. It is hoped that through redundancy in the observations, both the high energy models and the primary particle properties can be constrained.  Currently, the most important step in this direction is the development of the Auger Observatory. The southern part of the Observatory is now under construction in Argentina and will combine the fluorescence technique with a particle array of water \\v Cerenkov tanks covering a surface area of 3000~km$^2$ \\cite{Auger}. The signal from the tanks is recorded with Flash Analog to Digital Converters and relative timing is synchronized using GPS technology to 10-20~ns. The time structure of the shower front is of crucial importance because the shower arrival direction in an array of particle detectors is determined from the relative arrival times of the signals. This is relatively easy since the shower front can be approximated by a plane but clearly there are curvature corrections to the front and the time structure of the signal and its fluctuations will have an impact on the precision of the reconstruction procedure. Understanding the time structure of the signal is most important for controlling the angular resolution. Last but not least, preliminary observations at the Auger Observatory are already revealing that the time structure of the tank signals carries much information \\cite{ICRC} and it is hoped that this new information can be used to answer fundamental questions such as primary composition.  As most high energy cosmic ray detectors are, or have been, arrays of particle detectors, the time structure of the shower front has been studied before \\cite{Lapikens,watson,honda,battistoni,Kascade,time1}. However most work is either experimental or phenomenological and there is not a great deal of knowledge about the time structure of signals in air showers from the theoretical side or of its relation to the shower parameters. In this work we study the time structure of the muon signal in air showers. Our work is based on ideas that have been introduced for the description of muon densities in inclined showers \\cite{ModelPaper} and is a natural extension of these ideas. Similar ideas have been developed in \\cite{time1,time2,time3}.  Muons are naturally quite energetic when they reach the ground: otherwise they would have decayed in flight.  As they have long radiation lengths they do not suffer many interactions in their travel path from the point where they are created to ground level. As a result if muons are present in the signal they define the earlier part of the shower front, i.e. its onset which is usually chosen to define the arrival time to be used in the reconstruction of the direction. Also muons have a much flatter lateral distribution than photons or electrons and they always dominate the signal at ground level for sufficiently large distances from the shower axis. Moreover, for inclined showers, they dominate the signal everywhere at ground level because the electromagnetic part of the shower, generated by the decay of neutral pions, is absorbed high in the atmosphere.  Therefore, the description of the time structure of muons in air showers is of important practical interest for the reconstruction of air showers from measurements made in extensive air shower arrays.  Understanding the muon part of the shower is not the complete story because there are important contributions to the shower signal due to electrons and photons. We will leave the study of the time structure of electrons and photons for future work.  In this article we neglect the effect of the earth's magnetic field, although it is known that magnetic effects are important for the description of muon densities at high zenith angles. We have checked that the time structure of muons is not affected by the magnetic field up to zenith angles of 80 degrees and more, depending on the strength of the magnetic field. Although our results are general, and we expect that they will be valid at all zenith angles, we will concentrate on the large zenith angles range, where the muon signal is more important.  The article is organized as follows. In Section II we present the main assumption on which this work is based and we develop an analytical description of muon energy distributions in air showers as a function of transverse distance. These will be used in the next section when the time structure of the shower front is addressed. In Section III we describe the two main sources of time delays associated with muons and we present a model for their description, discussing how it compares with the results obtained directly using simulations and the relevance of the magnetic field effects. In Section IV we apply the method to give general results about time distributions in air showers and in Section V we summarize and conclude. ",
        "conclusions": "We have obtained a procedure that can be used to describe the time structure of the muon signal in extensive air showers. This procedure is based on two simple mechanisms, time delays due to the extra path length which are discussed as geometrical delays and those due to the sub-luminal velocity of the muons. The global delay can be expressed as a simple convolution of these two delays. The description combines the physical mechanisms with phenomenological descriptions of the energy, transverse momentum and position of the muons that are produced in an air shower with several approximations. The results obtained agree sufficiently well with those obtained with direct simulations for most distances of interest and for all zenith angles to be of a great importance for practical applications. This agreement allows us to state with some confidence that these two mechanisms dominate the time structure of the muons in air showers."
    },
    "0311/astro-ph0311015_arXiv.txt": {
        "abstract": "In loop quantum cosmology, the universe avoids a big bang singularity and undergoes an early and short super-inflation phase. During super-inflation, non-perturbative quantum corrections to the dynamics drive an inflaton field up its potential hill, thus setting the initial conditions for standard inflation. We show that this effect can raise the inflaton high enough to achieve sufficient e-foldings in the standard inflation era. We analyze the cosmological perturbations generated when slow-roll is violated after super-inflation, and show that loop quantum effects can in principle leave an indirect signature on the largest scales in the CMB, with some loss of power and running of the spectral index. ",
        "introduction": "The inflationary paradigm has been very successful, solving various problems in the big bang model of cosmology, and providing a framework for understanding structure formation in the Universe. Recent observational data on the cosmic microwave background (CMB) anisotropies, together with other data, are well accounted for by a simple inflationary model with suitable cold dark matter and dark energy content~\\cite{Spergel}.  However, this simple picture belies a number of fundamental puzzles. In particular, there is the problem of how to explain the initial conditions for successful inflation; for example, with the simple potential, \\beq \\label{po} V(\\phi)={1\\over2} m_{\\phi}^2\\phi^2\\,, \\eeq one requires an initial inflaton value $\\phi_i\\gtrsim 3\\mpl$. The ``eternal inflation\" scenario provides an answer to this, relying on the infinite extent of space, but its arguments have been challenged~\\cite{hawking}. In any event, whatever the merits of eternal inflation, it is useful to seek alternative explanations and to explore what quantum gravity may tell us about this. Quantum gravity effects above the Planck energy $\\mpl$ should be able to set the initial conditions for inflation--or to provide an alternative to inflation with adequate predictive power for structure formation. This is one of the challenges facing candidate theories of quantum gravity.  Results from WMAP also point to a possible loss of power at the largest scales and running of the spectral index. These effects, if confirmed, are difficult to explain within standard slow-roll single-field inflation, and may have roots in Planck-scale physics.  Loop quantum gravity (or quantum geometry) is a candidate quantum gravity theory that has recently been applied to early universe cosmology (see~\\cite{Bojo_review} for a recent review). It improves on a conventional Wheeler-DeWitt quantization, by discretizing spacetime and thus avoiding the breakdown of the quantum evolution even when the classical volume becomes zero~\\cite{Bojo1}.  The quantization procedure involves a fiducial cell corresponding to a fiducial metric to define a symplectic structure. Introduction of a fiducial cell, which is not required at a classical level, is a necessity for Hamiltonian quantization. This procedure has been developed for spatially flat and closed Friedmann geometries (and related Bianchi geometries). It is important to note that the scale factor in the quantum regime has a different meaning than in classical general relativity~\\cite{thanks}. For a flat geometry (which is the case we consider), in the quantum regime the scale factor is not subject to arbitrary rescaling, as in general relativity. Instead, the scale factor $\\tilde a$, obtained from loop quantization by redefining canonical variables, is related to the conventional scale factor $a$ via \\begin{equation} \\tilde {a}^3 = a^3 V_0\\,, \\end{equation} where $V_0$ is the volume of the fiducial cell. In this definition the quantization is independent of the rescaling of the fiducial metric and thus invariant under rescaling of the conventional scale factor. However, one can also quantize without redefining the canonical variables. In this case it has been shown that states in Hilbert space corresponding to two different fiducial metrics are unitarily related and the physics does not change with rescaling of the fiducial metric~\\cite{abl}. This is consistent with the background-independence of loop quantization.  What emerges in the quantization is a fundamental length scale, \\begin{equation}\\label{ls} \\tilde{a}_*=\\ell_*\\equiv\\sqrt{\\gamma j \\over 3}\\,\\lpl\\,, \\end{equation} where $\\gamma \\, (\\approx 0.13)$ is the Barbero-Immirzi parameter and $j $ is a half-integer (where $j>1$ is necessary in order to study evolution via an effective matter Hamiltonian~\\cite{Bojo2}). This length scale corresponds to a particular effective quantum volume (or the state of the universe) below which the dynamics of the universe is significantly modified by loop quantum effects. The crucial quantity that determines the dynamics and quantifies the state of the universe relative to the critical state, is \\begin{equation} q\\equiv {\\tilde {a}^2 \\over \\tilde {a}_*^2} = {a^2 \\over a_*^2}\\,. \\end{equation}  The Planck length $\\lpl$ is not put in by hand, but arises from the quantization procedure. The fundamental scale $\\ell_*$ is the same for a flat or closed classical geometry and has nothing to do with the topology. Thus for a flat non-compact spacetime, there is a fundamental length scale defined by quantization, unlike general relativity, where no length scale (other than the Hubble length) is defined by the geometry. If a compact topology is imposed on a flat geometry, then that introduces another (classical) length scale, independent of $\\ell_*$. Here we consider a flat non-compact universe. In the semi-classical regime, where spacetime is well approximated as a continuum but the dynamics is subject to non-perturbative loop quantum corrections, the conventional scale factor has the usual rescaling freedom. For convenience, we use the rescaling freedom to fix the classical scale factor at the critical epoch of transition from quantum to classical evolution: \\begin{equation} a_*=\\ell_*\\,. \\end{equation} It should be noted that any other choice could be used, or the rescaling freedom could be kept. The point is that the relevant physical quantity $q$ remains invariant. ",
        "conclusions": "In summary, we have shown that loop quantum effects can drive the inflaton from equilibrium up its potential hill, and can thus in principle set appropriate initial conditions for successful standard inflation. Loop effects may also leave an indirect imprint on the CMB on the largest scales, if the inflaton is not driven too far up its potential hill. This is possible in particular for very small initial fluctuations in the inflaton and its velocity, compatible with the uncertainty principle, Eq.~(\\ref{up}). We stress that the observational signature we found is an indirect effect of loop quantum gravity, rather than a direct consequence of perturbations generated in the regime $a<a_{*}$. In the loop quantum super-inflationary era, the standard quantization using the Bunch-Davies vacuum is no longer valid. Further analysis is required to understand the quantization of metric perturbations in loop quantum gravity.  By analyzing the dynamics and perturbations in the non-slow roll regime that follows the super-inflation, we showed that the loop quantum effects can indirectly lead to a running of the spectral index, with some loss of power on the largest scales. It should be noted that a similar loss of power on the largest scales can also be obtained via string theoretic scenarios (see Ref.~\\cite{other} and references therein). It would be interesting to distinguish between the different cosmological effects predicted by loop quantum gravity and string theories, using future high-precision observations.   \\vspace*{0.2cm} \\noindent{\\bf Acknowledgements}\\, We thank A. Ashtekar, B. A. Bassett, M. Bojowald, N. Dadhich, D. Jatkar and M. Sami for very helpful discussions. S.T. is supported by JSPS (No.~04942), P.S. by CSIR, R.M. by PPARC. P.S. thanks the ICG Portsmouth for support for a visit during which this work was initiated."
    },
    "0311/astro-ph0311365_arXiv.txt": {
        "abstract": "We explore the idea that the power-law tail in the mass function of protostellar condensations and stars arises from the accretion of ambient cloud material on to a condensation, coupled with a nonuniform (exponential) distribution of accretion lifetimes.  This model allows for the generation of power-law distributions in all star-forming regions, even if condensations start with a lognormal mass distribution, as may be expected from the central limit theorem, and supported by some recent numerical simulations of turbulent molecular clouds. For a condensation mass $m$ with growth rate $dm/dt \\propto m$, an analytic three-parameter probability density function is derived; it resembles a lognormal at low mass and has a pure power-law high-mass tail. An approximate power-law tail is also expected for other growth laws, and we calculate the distribution for the plausible case $dm/dt \\propto m^{2/3}$. Furthermore, any single time snapshot of the masses of condensations that are still accreting (and are of varying ages) also yields a distribution with a power-law tail similar to that of the IMF. ",
        "introduction": "\\subsection{IMF background}  Despite decades of detailed modeling of gravitational collapse and fragmentation, it is still not clear how certain physical conditions determine the mass of an individual star; a larger question is the origin of the {\\it distribution} of stellar masses at birth, the stellar initial mass function (IMF). Additionally, the observed mass spectrum of protostellar condensations in molecular clouds appears to follow a power-law decrease, similar to that of stars, in the mass range above $\\sim 0.5 \\Msun$ (Motte, Andr\\'e, \\& Neri 1998; Testi \\& Sargent 1998; Johnstone et al. 2000). Hence, the fundamental shape of the IMF may be determined at the molecular fragment scale, even though the transformation of a condensation to star(s) is not a simple process.  Starting with the expression for the Jeans (1928) mass, and working up to more sophisticated scenarios, there is a general tendency for multiplicative dependence on physical parameters (see, e.g., Adams \\& Fatuzzo 1996), leading to the claim that the central limit theorem should ensure that the IMF has a lognormal form. The formation of star-forming condensations in a molecular cloud is surely affected by multiple effects of a turbulent flow field, and statistical variations of ambient magnetic field, density, pressure, and temperature. Numerical models of turbulent molecular clouds typically yield a lognormal distribution of dense clump masses (e.g., Padoan, Nordlund, \\& Jones 1997; Ostriker, Stone, \\& Gammie 2001; Klessen 2001). Although much additional physics remains to be included in these models, it may be reasonable to believe that the stochastic feature of a lognormal distribution will continue to hold. Hence, one can ask a simple set of questions about star formation: Is the IMF lognormal? If not, and an approximate power-law tail is inferred, then how is it generated?  While Miller \\& Scalo (1979) did obtain a lognormal fit to the IMF, subsequent detailed studies find a region of power-law decrease in the intermediate and high mass ($M$) range that is similar to the original result of Salpeter (1955), i.e., $dN/d \\ln M \\propto M^{-1.35}$, where $dN$ is the differential number of stars in the logarithmic mass interval $d \\ln M$ (see e.g., reviews by Scalo 1998; Meyer et al. 2000; Elmegreen 2001). Scalo (1998) emphasizes a varying power-law index in the intermediate and high mass range, with significant region to region variations reflecting intrinsic variability and/or observational uncertainties.  \\subsection{Lognormal or power-law distribution?}  The generation processes of lognormal and power-law distributions are often intimately connected. This phenomenon has been studied extensively in social science, biology, and computer science, e.g., in the study of power-law distributions of income (the Pareto Law), city sizes (Zipf's Law), and Internet file sizes. Useful models and discussion can be found in Reed (2002, 2003), Reed \\& Hughes (2002), Mitzenmacher (2003), and references within. It turns out that many seemingly small effects in the generative process of a lognormal distribution can lead to a power-law tail instead. Such an explanation seems particularly relevant for the IMF, since the power-law tail is typically measured to span even less than two decades of mass, hardly qualifying as evidence for a truly scale-free process.   In this Letter, we propose a minimum-hypothesis explanation: even if the central limit theorem is relevant in determining the starting masses of condensations, the temporal effect of their subsequent interaction with the environment skews the mass distribution into a power-law form. This paper is based on the statistical model summarized by Reed \\& Hughes (2002), and builds upon some elements of the model of Myers (2000). The basic idea is that the growth of masses by accretion coupled with a {\\it distribution of times} for the growth yields a power-law tail in the mass function. It is a statistical approach, not dependent on the detailed values of physical variables in different environments. While the mass of any individual star is dependent on the particular initial values of various physical parameters, we are concerned here only with the {\\it shape} of the mass distribution function; specifically the intermediate and high mass tail. ",
        "conclusions": "We believe that the basic features of this model will be preserved in more complicated and realistic scenarios; that is, even a lognormal distribution of condensation masses at birth coupled with subsequent accretion growth (proportional to some power of mass) in which the accretion time obeys its own distribution (approximately exponential), will result in a power-law tail.  Zinnecker (1982) has presented another model of accretion growth in which a distribution of initial masses $m_0$ undergo Bondi accretion ($\\dot{m} \\propto m^2$); the exponent greater than 1 means that small differences in the $m_0$ values are magnified over time. After a fixed accretion time for all objects, when $m \\gg m_0$, $f(m)$ tends to a power-law form with index $-2$. For comparison, the mechanism of exponential distribution of lifetimes studied in this paper yields a power-law tail even if the initial masses are all the same. The general combination of accretion growth (due to the Bondi process or geometric accretion) and a distribution of accretion lifetimes (where both processes are controlled by the dynamical time of the parent system) may play a role in determining the IMF in detailed global simulations of star formation (e.g., Bate, Bonnell, \\& Bromm 2003) in which stars are formed throughout the duration of the simulation. Other ideas related to the {\\it spatial} structure of turbulence may also contribute to power-law fragment masses, e.g., thresholded regions in density fields with lognormal probability density functions may have power-law clump mass distributions (Pudritz 2002; Elmegreen 2002). However, we believe that the {\\it temporal} effect studied in this paper provides a compelling part of the explanation for the power-law tail in the IMF."
    },
    "0311/astro-ph0311153_arXiv.txt": {
        "abstract": "Thermal models for the quasar Big Blue Bump generally lead to bound-free continuum features, which may be in absorption or emission.  Searches for Lyman edges attributable to quasar atmospheres (in particular accretion disk atmospheres) have been ambiguous at best, but various relativistic and non-LTE effects may make them hard to detect.  The Balmer edge features tentatively predicted by such models might be easier to detect since they'd form farther out in the potential wells.  These can be sought in certain cases using spectropolarimetry to remove the effects of atomic emission (i.e. the Small Blue Bump) from the spectrum. We do find apparent Balmer edges in absorption in several quasars using this method! Although the features we see are believable and apparently common, more data and more modeling are needed to verify that they are best interpreted as atmospheric Balmer edges. ",
        "introduction": "Early models of quasar accretion disks\\footnote{It must be emphasized that almost all published accretion disk models are quasi-static and based on the Shakura-Sunyaev model.  This is known to be qualitatively and quantitatively inconsistent with the observations of variability, spectral energy distributions and other properties.  A sampling of references for this statement is the following:  Alloin et al 1985, Antonucci 1988, 1999; Krolik et al 1991;  Gaskell et al 2003; D vanden Berg 2003, these proceedings; O Blaes 2003, these proceedings. However, the models may indicate in a general way the behavior of optically thick thermal radiation in the case of energy deposition at large optical depth, especially if the emission follows the potential drop as in the Shakura-Sunyaev paradigm.}predicted strong Lyman edges in absorption, as expected for stars of moderately high temperature, due to source function gradients in the atmospheres (e.g. Kolykhalov and Sunyaev 1984).  Any such edges aren't easily seen (Antonucci, Kinney and Ford 1989 and later papers)!\\footnote{Also note that partial Lyman Limit Systems such as those of S5 0014+81 can give false positives for atmospheric edges (e.g. Fig 1a of Antonucci et al 1989 and the accompanying discussion). Thus no edges should be claimed to be atmospheric unless it has been shown that they lack accompanying narrow absorption lines. Candidates for broadened edges have generally failed this test, when the test has been made (Koratkar, Kinney and Bohlin (1992)).}  Changes in spectral slope do appear to be common in the vicinity of the Lyman limit (Zheng et al. 1995 and later papers).  Such spectral breaks might be consistent with smeared out and/or Comptonized Lyman edge features (e.g. Lee, Kriss, \\& Davidsen 1992, Hubeny et al. 2000).  However, detailed fits to the spectral break observed in 3C~273 with the latest disk atmosphere models still produce a local emission bump longward of the limit that is not in agreement with the data (Blaes et al. 2001), and other problem arise when fitting the reported spectral breaks as atmospheric edges (Kishimoto et al, in prep).  Accretion disk models place the $\\sim$4000A continuum around seven times farther out in the gravitational potential well for peak emissivity, compared with the Lyman edge region. This would greatly reduce relativistic smearing. Also since Balmer continuum absorption is not a resonance process, there is less chance of being fooled by intervening absorbing matter.  In fact the distinction between atmospheric and foreground absorption becomes in part a semantic one in that case.  The Balmer edge behavior has never been discussed as a diagnostic of the quasar continuum emission mechanism because very strong Balmer continuum EMISSION and FeII features, attributed to the Broad Line Region, contaminate this wavelength region.  Very fortunately there is a way around the problem of contamination by atomic emission in some objects.  These have slightly polarized BBBs, and UNPOLARIZED emission lines.  The likely explanation is that there is some Thomson scattering interior to the BLR, which polarizes the BBB only.\\footnote{We discuss the possibility that the observed polarization is due to a synchroton component (Schmidt and Smith 2002) in Kishimoto et al 2003.   Here we just note that the existence of the edge feature in polarized flux requires that at least some of the polarization result form a thermal process.}  Thus a plot of polarized flux can cleanly shave off the atomic emission.  Published data of moderate SNR confirm that in these cases the Small Blue Bump is indeed unpolarized (e.g. Antonucci 1988;  Schmidt and Smith 2000).  With recent higher-SNR data from larger telescopes, we can determine the behavior of the BBB alone at the Balmer edge postion. ",
        "conclusions": "We often find that the polarized flux DECREASES at wavelengths below ~4000A, strongly suggesting that the BBB has such a feature, and is thus thermal emission from an optically thick source. It may be that the total flux of the BBB has no feature at this wavelength, but that instead the BBB percent polarization has the feature. This type of behavior occurs in some models (Laor, Netzer and Piran 1990;  a rigorous calculation which gives rather different behavior in some part of parameter space was later published as Blaes and Agol 1996).  Still the implication is optically thick thermal emission.  Although the observed Balmer edge behavior is qualitatively consistent with thermal model predictions, more data and more modeling are needed to ensure that this is the correct interpretation.  On a slightly different topic, a high point of the SDSS conference for us was a lunchtime discussion of the polarization of 2MASS quasars with Gary Schmidt.  His group has observed these objects extensively (e.g. Smith et al 2003).  The objects were selected for red J-K colors. Considering the nonzero redshifts, this means in practice that they were selected for LARGE RATIOS OF HOT DUST EMISSION TO BBB EMISSION. The polarization data are very remarkable!  We'd have expected the 2MASS objects to be mostly reddened quasars, with polarization due to dust scattering or absorption.  Dusty far-IR selected quasars behave this way in general (e.g. Hines et al 2001).  Dust scattering typically produces very blue polarized flux spectra in the objects studied in that paper and in other similar objects, at least in the optical part of the spectrum.  (The shorter wavelengths are often cut off by dust absorption.)  None of the 2MASS objects behave that way. However many 2MASS quasars are highly polarized:  the percent polarization is often $\\sim10\\%$, rather than $\\sim1\\%$ as for our objects!  In general the broad lines ARE polarized in objects selected this way, and so that this high polarization arises on BLR scales of larger.  These properties make the interpretion of the spectropolarimetry very different from that of the low-polarization quasars described above.  Modest foreground reddening suppressing the BBB wouldn't necessarily produce very red J-K colors since that wavelength interval is small and at relatively long wavelength;  rather as noted above it could indicate a large ratio of dust emission to BBB emission.  Actual optical/NIR SEDs will help to distinguish the two possibilities.  Over the lunch able we considered the possibility that the BBB emission is suppressed by foreground Thomson scattering.  The (often very high) polarizations may be due to Thomson scattering as well.  Large electron-scattering optical depths are very reasonable for quasars and luminous Seyfert galaxies.  A powerful argument is that the scatterers in objects like NGC1068 provide an optical depth of 1$\\%$ due to electron scattering (Miller, Goodrich, and Mathews 1991; Antonucci et al 1994; Ogle et al 2003) \\footnote{This paper actually finds an angle-averaged optical depth on $\\sim$100pc scales of order a percent, but within the ionization/scattering cone, a small covering factor of sight lines with much higher Thomson depth.} and this polarization arises  on $\\sim$100pc scales (Antonucci et al 1994; Capetti et al 1995). The outflowing scattering region is likely to be a wind which originates near the sublimation (or ''ablation\") radius of $<$1pc (Krolik and Begelman 1986).  Thus given the density law for winds, THE OPTICAL DEPTH OF THE WIND IS OF ORDER UNITY on pc scales. (Ferland et al, this meeting, also argue for such electron-scattering optical depths near the torus inner edge.) Thus NGC1068 would look like a luminous Seyfert 1 or weak quasar if seen from the polar regions, and would present a very substantial electron-scattering optical depth to the BBB-emitting region. Similarly many quasars show large X-ray absorption columns with little or no optical reddening and absorption. (e.g. Salvati and Maiolino 2000; Wilkes et al 2002).  These observations are interpretable in the context of the scattering wind (Krolik and Kriss 2001).  A delightful aspect of the 2MASS spectropolarimetric observations is that the DIRECT, unpolarized line profile can often be seen simultaneously with the Thomson-scattered line profile!  See for example 2MASS 222202.2 +195231 in Fig 2 of Smith et al 2003. Thus we can get an accurate temperature and outflow velocity, as well as a good idea of the electon- scattering optical depth.   For the object mentioned, the electron temperature comes out to one million degrees (FWZI = 18,000 km/s) and the outflow velocity is around 4000 km/sec (G. D. Schmidt, 2003 p.c.). However, demonstration of large electron-scattering optical depth in front of many quasar BBBs will entail many observational tests.  To summarize, the polarization property is underexploited in this field, and it can provide many unique insights.  We thank Gary Schmidt and Pat Ogle for sharing their thoughts.  The work was funded in part by NSF AST-0098719."
    },
    "0311/astro-ph0311147_arXiv.txt": {
        "abstract": "This Letter presents the first results of an observational campaign to study the Galactic center with \\integ, the {\\it Internaltional Gamma-Ray Astrophysics Laboratory}. Mosaicked images were constructed using data obtained with ISGRI, the soft gamma-ray instrument of the coded aperture IBIS imager, in the energy ranges 20--40 and 40--100\\un{keV}. These give a yet unseen view of the high-energy sources of this region in hard X-rays and gamma-rays with an angular resolution of 12\\am (FWHM). We report on the discovery of a source, IGR J1745.6--2901, coincident with the Galactic nucleus \\sgra to within 0\\hspace{1pt}\\am\\hspace{-5pt}.9. Located at R.A.(J2000.0) = $\\mathrm{17^{h}45^{m}38^{s}\\hspace{-2pt}.5}$, decl.(J2000.0) = $-29^{\\circ}01'15''$, the source is visible up to about 100\\un{keV} with a 20--100\\un{keV} luminosity at 8\\un{kpc} of $(2.89 \\pm 0.41) \\times 10^{35}$\\un{ergs\\; s^{-1}}. Although the new \\integ source cannot unequivocally be associated to the Galactic nucleus, this is the first report of significant hard X-ray emission from within the inner 10\\am of the Galaxy and a contribution from the Galactic supermassive black hole itself cannot be excluded. ",
        "introduction": "From the prediction of the existence of a massive, compact source at the center of the Milky Way by Lynden-Bell \\& Rees (1971) % to the discovery of such a source in the radio domain 3 years later by Balick \\& Brown (1974), % and to the first detection of soft X-rays unmistakably attributable to it by {\\it Chandra} in 1999 \\citep{c:baganoff03a}, the supermassive black hole candidate \\sgra and the Galactic center (GC) region as a whole have been put under intense scrutiny for many years from radio wavelengths to gamma-rays. This has led to several discoveries, and advances in our understanding of the processes and interactions at the heart of the Milky Way. For example, it is now known that the soft X-ray emission ($<$10\\un{keV}) in the central 10\\am is heavily dominated by diffuse emission due primarily to hot gas \\citep{c:koyama96,c:sidoli-mereghetti99}, and that only about 10\\% of the total emission in this energy range can be accounted for by X-ray point sources brighter than $10^{31}\\un{ergs\\; s^{-1}}$ \\citep{c:muno03a}. Also, the contribution of point sources in this domain is about the same along the Galactic ridge as it is in the GC \\citep{c:ebisawa01}.   The Galactic nuclear region consists of six primary components that give rise to an array of complex phenomena through their mutual interactions. These constituents are: \\sgra, a supermassive black hole with a mass of around $3.5 \\times 10^6 \\un{M_{\\odot}}$ \\citep{c:schodel02,c:ghez03a}; the surrounding cluster of evolved and young stars; ionized gas streamers, some of which form a three armed spiral centered on \\sgra known as Sgr\\,A\\,West; a molecular dusty ring surrounding Sgr\\,A~West; diffuse hot gas; and a powerful supernova-like remnant known as Sgr\\,A~East \\citep{c:melia01}. Furthermore, both the IR and the X-ray point source populations decreases in spatial density approximately as $1/R^{2}$, where $R$ is the distance from the GC \\citep{c:serabyn96,c:muno03a}.   The first imaging observations of the GC in hard X-rays were performed by the X-ray telescope on {Spacelab\\,2} in the range 2.5--20\\un{keV} with an angular resolution of 3$'$. A source located within 1\\am\\hspace{-6pt}.1 of \\sgra was detected \\citep{c:skinner87}. Such a source was also detected by ART-P (8--20\\un{keV}) on the {\\it Granat} satellite in 1990-1991 \\citep{c:pavlinsky94}. At energies above 20\\un{keV}, only the coded mask instrument SIGMA/{\\it Granat}, sensitive to energies above 35\\un{keV}, provided imaging capability with a $\\sim$20\\am angular resolution, allowing an exploration of the dense GC region. In spite of the deep $9 \\times 10^{6}\\un{s}$ SIGMA survey of the central parts of the galaxy performed between 1990 and 1997, only upper limits were set for the hard X- and gamma-ray emission from the neighborhood of \\sgra at energies above 35\\un{keV} \\citep{c:goldwurm94,c:goldoni99}. The derived low bolometric luminosity of the Galactic nucleus (GN), in contrast with the powerful output from active galactic nuclei or black hole binaries, has motivated the developement of several models for radiatively inefficient accretion onto or ejection from \\sgra. These models have been widely applied to other accreting systems \\citep{c:melia01}.   A recent breakthrough discovery by {\\it Chandra} and {\\it XMM-Newton} is that \\sgra is the source of powerful X-ray flares \\citep{c:baganoff01,c:goldwurm03a,c:porquet03} during which the soft X-ray luminosity can increase by factors of 50--180 over a period of up to 3 hr. Some of these X-ray flares feature a significant hardening of the spectrum up to photon indices of $\\sim$1. Furthermore, a very recent discovery with the Very Large Telescope NAOS/CONICA imager \\citep{c:genzel03} and the Keck telescope \\citep{c:ghez03b}, is that \\sgra is also the source of frequent IR flares. This activity could indicate the presence of an important population of non-thermal electrons in the vicinity of the black hole. These results have raised great interest in the possibility of observing hard X-rays from the GN, a measure of which may particularly shed light on the relative role of accretion and ejection in the \\sgra system.  This Letter presents preliminary results of an observation campaign to study the GC at high energies with \\integ, the {\\it Internaltional Gamma-Ray Astrophysics Laboratory}, performed during the \\integ Galactic Center Deep Exposure (GCDE) program in the spring of 2003. The focus is set on the inner 10\\am--15\\am of the Galaxy and, more particularly, on the detection of a significant excess in this region. The results are based on data obtained with the \\integ Soft Gamma-Ray Imager (ISGRI), the low-energy camera of the main Imager on Board the \\integ Satellite (IBIS), proven to be accurate between 20 and 200\\un{keV}. The morphology of the central 2$^{\\circ}$ in the 20--100\\un{keV} band and the derived positions of closeby high-energy sources are briefly discussed to elucidate the experimental context from which the results on the central source are drawn.  \\begin{figure}[htb] \\epsscale{1.1} \\plotone{f1.eps} \\caption{\\footnotesize Milky Way's center seen in a $2^{\\circ}\\hspace{-2pt}\\times 2^{\\circ}$ field by the IBIS/ISGRI instrument in the energy ranges 20--40\\un{keV} (top) and 40--100\\un{keV} (bottom). Each image pixel size is equivalent to about 5\\un{arcmin}. Ten contour levels mark iso-significance linearly from about 4$\\sigma$ to 15$\\sigma$. \\label{f:sgra}} \\end{figure} ",
        "conclusions": "\\begin{deluxetable}{llcc} \\tablecolumns{4} \\tablewidth{0pc} \\tablecaption{Measured and extrapolated fluxes} \\tablehead{ \\colhead{Instrument} & \\colhead{Energy}   & \\multicolumn{2}{c}{Energy Flux} \\\\ \\colhead{} & \\colhead{(keV)}\t& \\colhead{($10^{-11}\\un{ergs\\; cm^{-2}\\; s^{-1}}$)} & \\un{(mCrab)} } \\startdata \\integ\t\t& 20--40$^{\\mathrm{a}}$\t& \t$1.92 \\pm 0.36$\t\t& $3.2 \\pm 0.6$ \\\\ (IGR 1745.6--2901)& 40--100$^{\\mathrm{a}}$&\t$1.86 \\pm 0.40$\t\t& $3.4 \\pm 0.7$ \\\\ \\hline & 2--10$^{\\mathrm{a}}$\t&\t\t10.0\t\t& 5.78\t\t\\\\ {\\it XMM-Newton}& 20--40$^{\\mathrm{b}}$\t&\t\t0.61\t   \t& 1.03\t\t\\\\ & 40--100$^{\\mathrm{b}}$&\t\t0.031\t   \t& 0.04\t\t\\\\ \\hline & 2--10$^{\\mathrm{a}}$\t&\t\t11.0\t   \t& 6.36\t\t\\\\ {\\it BeppoSAX}\t& 20--40$^{\\mathrm{c}}$\t&\t\t0.89\t   \t& 1.49\t\t\\\\ & 40--100$^{\\mathrm{c}}$&\t\t0.061\t\t& 0.09 \\enddata \\tablenotetext{a}{Measured} \\tablenotetext{b}{Extrapolated using kT = 7.2 keV} \\tablenotetext{c}{Extrapolated using kT = 8.0 keV} \\label{t:fluxes} \\end{deluxetable}  In attempting to identify the excess at $\\mathrm{17^{h}45^{m}38^{s}\\hspace{-2pt}.5}$, $-29^{\\circ}01'15''$, we first consider which sources of high-energy emission can be excluded based solely on ISGRI's error radius of 4\\am. There are four such sources: the radio arc located about 15\\am from the GN; the non thermal X-ray filaments associated with it and the dense molecular cloud {G0.13--0.13} \\citep{c:yusef02}; GRS 1743--290, observed by SIGMA in 1991 \\citep{c:goldwurm94}; and GRS 1741.9-2853 mentioned above and clearly unable to account for the bulk of the excess.  Second, we argue that this detection cannot be attributed to both diffuse emission and point sources as thus far observed at low X-ray energies within 8\\am--10\\am of \\sgra. This is so because it is incompatible with the extrapolation at higher energies of the flux integrated over this region as measured by {\\it XMM-Newton} (A. Decourchelle 2003, private communication) and {\\it BeppoSAX} \\citep{c:sidoli99}. The analysis of 2001 {\\it XMM-Newton} data in the 2--10\\un{keV} range consisted of integrating the total X-ray flux over a radius of 10\\am around \\sgra. The resulting spectrum is best fitted by a two-temperature plasma model with $\\mathrm{kT_{1}}$=0.6\\un{keV} for the soft component and $\\mathrm{kT_{2}}$=7.2\\un{keV} for the harder one. 1997 {\\it BeppoSAX} observations of the Sgr\\,A complex over a radius of 8\\am around \\sgra in the same energy band agree with {\\it XMM-Newton}'s results. In this case the data are also best described by a two-temperature model with $\\mathrm{kT_{1}}$=0.6\\un{keV} and $\\mathrm{kT_{2}}$=8\\un{keV}. The extrapolation of the {\\it XMM-Newton} and {\\it BeppoSAX} results to higher energies was done using kT=7.2 and 8\\un{keV}, respectively. Both instruments' direct measurements of the diffuse X-ray flux are consistent with those of {\\it Chandra} \\citep{c:baganoff03a} and ART-P \\citep{c:sunyaev93}. Table~\\ref{t:fluxes} reports the {\\it XMM-Newton} and {\\it BeppoSAX} measurements as well as their respective extrapolations. Note that the variability timescale of the \\integ source equally excludes a major diffuse contribution and also strongly suggests that accelerated charged particles in the expanding shell of Sgr\\,A East \\citep{c:melia98} cannot significantly contribute to the detected signal.   Third, since the detected excess cannot be fully accounted for by diffuse emission and is variable, it is reasonable to consider it a point source, IGR~J1745.6--2901, of 20--100\\un{keV} luminosity $\\sim$$3 \\times 10^{35}\\un{ergs\\; s^{-1}}$. One known point source found within 4\\am of IGR~J1745.6--2901 that could contribute to the signal is the eclipsing burster AX J1745.6--2901 \\citep{c:maeda96,c:sakano02}. This source lies about 1\\am.3 from \\sgra and was detected in a high state only once, in 1994 \\citep{c:sakano02}. The extrapolation of the flux measured in that state is marginally compatible with that of IGR J1745.6--2901, and this only in the low-energy range. The last confirmed detection of AX J1745.6--2901 was in 1997 by ASCA, when it was found in a much lower flux state. It may have recently been detected by {\\it Chandra}. A timing analysis of the ASCA data on AX J1745.6--2901 revealed a modulation with a 40\\% decrease in flux and period of $8.356 \\pm 0.008 \\un{h}$ \\citep{c:maeda96}. 1997 ASCA data folded with that period show a slight dip, but since phase information could not be preserved over the 2.5 yr that separated the two observations, no unique solution for the period was found \\citep{c:sakano02}. Folding the \\integ data for IGR~J1745.6--2901 with a period of 8.356 hr does not reveal any modulation.   {\\it Chandra} observations of the GC performed on 2003 June 19 indicate that two sources, CXOGC J174540.0--290005 and CXOGC J174540.0--290014, found within 0\\am\\hspace{-6pt}.4 of \\sgra, and another, CXOGC J174535.5--290124, lying 1\\am\\hspace{-6pt}.3 from \\sgra and possibly the same as AX~J1745.6--2901, had a combined intensity in the 2--8\\un{keV} band that was more than 30 times that of \\sgra in its quiescent state on that day (F.~K. Baganoff 2003, private communication). The GN was not particularly bright on June 19, but it must have been somewhat more active in May when an intense IR flare was observed to come from a region possibly very near the supermassive black hole's event horizon \\citep{c:genzel03}.  It is noteworthy that the fairly hard X-ray source that ART-P detected in 1990 at the position of \\sgra, had a 8--20\\un{keV} flux that varied between about 5 and 11\\un{mcrab} \\citep{c:pavlinsky94}. This is consistent with our detection in the 20--40\\un{keV} energy range. We also recall the 2~$\\sigma$ upper limits set by SIGMA on the 40--80\\un{keV} and 80--150\\un{keV} luminosity of the GN: $3.4\\times 10^{35}$ and $2.8\\times 10^{35}\\un{ergs\\; s^{-1}}$ respectively \\citep{c:goldoni99}, both of which are compatible with our detection.  Finally, \\integ has detected a very hard X-ray source at a position coincident to within 1\\am with \\sgra. Because of IBIS's 12\\am angular resolution, the emission from IGR~J1745.6--2901 cannot be attributed to one specific object and could be made up of several contributing sources found within a couple of arcminutes of \\sgra. Nonetheless, this remains the first detection of emission at energies greater than 20\\un{keV} from the very close vicinity of the GN, and a contribution from \\sgra itself cannot be excluded.  More INTEGRAL data are needed to better constrain the position, spectral shape, variability properties and the possibly multiple nature of IGR~J1745.6--2901. We expect that the most constraining results will be provided by simultaneous observations in hard and soft X-rays with INTEGRAL and {\\it XMM-Newton} or {\\it Chandra}."
    },
    "0311/astro-ph0311001_arXiv.txt": {
        "abstract": "We examine the possibility that hyper-accretion onto newly born, black holes occurs in highly intermittent, non-asymmetric fashion favorable to gravitational wave emission in a neutrino cooled disk. This picture of near-hole accretion is motivated by magneto-rotationally induced, ultra-relativistic disk dynamics in the region of the flow bounded from below by the marginally bound geodesic radius $r_{\\rm mb}$. For high spin values, a largely coherent magnetic field in this region has the dynamical implication of compact mass segregation at the displacement nodes of the non-axisymmetric, MRI modes. When neutrino stress competes favorably for the disk dynamical structure, the matter clumps may be rather dense and sufficiently long-lived to excite the Quasi-Normal Ringing (a.k.a. QNR) modes of the Kerr geometry upon their in-fall. We find that such accretion flow may drive bar-like, quadrupole $(l,m=2,2)$ modes in nearly resonant fashion for spin parameters $a \\geq .9$. The ensuing build up in strain amplitude of the undamped oscillations warrants a brisk rate of energy deposition into gravitational waves. A detectability assessment for the LIGO interferometers through the match filtering technique is given by integrating the energy flux over a one second epoch of resonant hyper-accretion at $1 \\msun \\sec^{-1}$. Thus, a $15 \\msun$ Kerr black hole spinning at $a \\simeq .98$ ($f_{\\rm QNR} \\simeq 1677$ Hz), and located at 27 Mpc (e.g., GRB980425), will deliver a characteristic strain amplitude, $h_{\\rm char} \\simeq 2.2_{-21}$, large enough to be detectable by LIGO II. If resonant hyper-accretion were sustainable for a longer period (or at higher rates) possibly associated with a second broad hump in a GRB light-curve, these objects could be detected by LIGO I at very low redshifts. ",
        "introduction": "\\label{sec:Intro}  It seems only fitting for ultra-relativistic, black hole accretion to be the leading contender to energize and explode the massive stars associated with gamma-ray bursts and hyper-nov\\ae. One must recognize, however, that such a scenario resides on a very exotic front: stellar-mass black holes accreting at twelve orders of magnitude above the Eddington limit! On the other hand, since GRB engines are effectively hidden from view, gravitational wave emission models may constitute the most effective probes to such enigmatic events. Indeed, there is a great deal of impetus to address such extreme accretion scenario and its associated, non-standard energy deposition channels. In this paper we are chiefly concerned with the possibility that the dynamical structure of hyper-accreting flows is highly intermittent on large scales, with the in-fall of large mass over-densities leading to prolific gravitational wave emission coincident with the gamma-ray stage.  We envisioned the accretion disk setting following the hyper-accreting black hole models of Popham, Woosley \\& Fryer (1999, hereafter PWF). For a fiducial scenario: ${\\rm M}_{h} = 3 \\msun, ~\\alpha_{_{\\rm SS}} = 0.1, ~{\\rm and}~ \\dot{\\rm M} = 0.1 \\msun \\sec^{-1}$, these authors found that the onset of photo-disintegration and of neutrino cooling at radii $r \\ltaprx \\,70 ~[GM/c^2]$ yield a mildly advective, semi-thin, $\\cH_\\Theta / r \\ltaprx .4$, accretion disk structure with a slightly sub-Keplerian rotation profile. An improved account of neutrino transport (Di Matteo \\etal 2002) has found that the innermost disk portion should become optically thick to neutrinos for accretion rates $\\dot{\\rm M} \\gtaprx 0.1 \\msun \\sec^{-1}$. This being the case, cooling by radial advection of energy will compete with that from local neutrino emission and advection will overwhelm neutrino cooling for $\\dot{\\rm M} \\gtaprx 1 \\msun \\sec^{-1}$ (Di Matteo \\etal 2002). Yet, such a high accretion rate is rather unlikely (MacFadyen \\etal 2001), so we restrict our analysis to $\\dot{\\rm M} \\leq \\,1 \\msun \\sec^{-1}$  The structure of this paper is a follows. In \\S 2, we assess the relevance of physical processes occurring in the relativistic region of the flow: $r_{\\rm mb} \\leq r \\leq r_{\\rm ms}$. Magneto-rotationally-induced, relativistic disk dynamics is addressed in \\S 3, where we also discuss the role of compressibility and estimate the size of mass over-densities from a neutrino cooled disk. In \\S4 we build an idealized, analytical model for gravitational wave emission and, in \\S 5, estimate the detectability of the GW signal for the LIGO interferometers with the match filtering technique. ",
        "conclusions": "We have laid out a physical--if simple minded--model for resonant driving of the quasi-normal ringing (QNR) wave modes of the Kerr geometry. A nascent black hole hyper-accreting at rates $\\dot{\\rm M} \\approx 1 \\msun \\sec^{-1}$ from a neutrino cooled disk is reckoned to oscillate near resonance of its ($l,m= 2,2$) quadrupole QNR frequency due to the in-fall of compact mass over-densities from the cusp in effective potential on a dynamical time scale. This model is based on large-scale, magneto-rotationally--induced fluid dynamics in the ultra-relativistic region of the flow bounded from below by the marginally bound orbit radius: $r_{\\rm mb}$ (assumed to coincide with the aforementioned cusp). The exact location of the cusp will change the dynamical time scale for in-fall but not by much when the spin parameter $a \\geq .9$. Furhermore, since the MRI time scale is suitable fast, a small delay to reach the free-fall stage is all that is needed for MRI dynamics to act. Heat conduction from neutrino diffusion out of mass over-densities on the time scale of the instability warrants that compressive perturbations behave non-adiabatically. This leads to sufficiently long-lived clumps in a highly intermittent medium. When the spin parameter of the hole exceeds 90\\%, large-scale magneto-rotational wave modes will segregate compact matter over-densities from large-scale magnetic field domains. The subsequent in-fall of these clumps from $r_{\\rm mb}$ will drive the QNR modes of the geometry in resonant fashion if the clumps arrive to $r_{\\rm mb}$ steadily, with a small spread around time intervals of $\\cO[\\Omega^{-1}_+]$, and for long enough to build the strain amplitude up to a saturation point.  It is difficult to address the probability distribution for clump arrival times from the linear theory alone but one anticipates that it will involve a co-moving time scale of $\\cO[\\Omega^{-1}_+]$ and the integrated, spin dependent WKB wavelength on the inspire trajectory toward $r_{\\rm mb}$. We have argued that for large mass clumps, $\\delta M \\simeq \\dot{M}\\Omega_+^{-1}$, the arrival times to $r_{\\rm mb}$ will not be entirely stochastic since ``the bucket needs time to fill-up\", nor entirely coherent since the seed density perturbations are stochastic but large-scale MRI dynamics will work to minimize the impact of this initial small-scale randomness. If the arrival times were entirely stochastic, the enhancement in energy deposition would be roughly proportional to the Q value of the mode instead of to Q$^2$ and the corresponding signal would be smaller by about an order of magnitude. A detailed probability distribution of arrival times requires an accurate account of the radial velocity profile and of the effective, spin-dependent WKB wavelengths on in-spiral trajectories which is beyond the scope of this study.  If the QNR modes are fed resonantly for a few seconds of hyper-accretion, the enhanced amplitude of the oscillations yields a very high rate of energy deposition into gravitational waves. Indeed, the integrated energy deposition is large enough to ``evaporate\" the equivalent of a factor of a few times the total rest-mass--energy of a single clump into gravitational waves, Yet, the characteristic frequency of emission resides on a less than optimal range with respect to the advanced LIGO noise curve: $f\\simeq 8000-1600$Hz for $M \\simeq 3-15 \\msun$. On the other hand, the high quality and coherence of the expected signal should make this model a prime candidate for narrow-band LIGO searches. For gamma-ray bursts at low redshifts, this may ultimately lead to a signal strong enough to be seen by LIGO II.  A 15$\\msun$ black hole accreting at $1 \\msun \\sec^{-1}$ may come about from the death throes of a very massive star. If such an explosion is associated with a long GRB, and the broad properties of the prompt gamma-ray light curve, to inner engine activity; hyper-accretion onto this heavy weight, fast spinning hole should be associated with a second broad hump in the light curve. We anticipate that gravitational waves with the predicted properties may be present at this stage of the GRB. The proposed mechanism is thus a direct probe of the elusive inner engine of gamma-ray bursts.   \\vfill"
    },
    "0311/astro-ph0311237_arXiv.txt": {
        "abstract": "We report on the analysis of high-resolution infrared spectra of the newly discovered brown dwarf $\\epsilon$ Indi B.  This is the closest known brown dwarf to the solar system, with a distance of 3.626 pc.  Spectra covering the ranges of $\\lambda$2.308--2.317$\\mu$m and $\\lambda$1.553--1.559$\\mu$m were observed at a resolution of $\\lambda$/$\\Delta$$\\lambda$=R=50,000. The physical parameters of effective temperature and surface gravity are derived for $\\epsilon$ Ind Ba by comparison with model spectra calculated from atmospheres computed using unified cloudy models.  The results are T$_{\\rm eff}$= 1500$\\pm$100K, log g= 5.2$\\pm$0.3 (in units of cm s$^{-2}$), placing it in the critical boundary between the late-L and early-T dwarfs.  The high spectral resolution also allows us to measure an accurate projected rotational velocity, with vsin($\\iota$)= 28$\\pm$3 km s$^{-1}$. Combined with a published luminosity for $\\epsilon$ Ind Ba (with log(L/L$_{\\odot}$)= -4.67), the derived parameters result in a ``spectroscopic'' mass estimate of $\\sim$32M$_{\\rm Jupiter}$, a radius of $\\sim$ 0.07R$_{\\odot}$, and a maximum rotational period of $\\sim$ 3.0 hours.  A compilation and comparison of effective temperatures derived from spectroscopy using model atmospheres versus those derived from luminosities and theoretical M$_{\\rm bol}$--radius relations reveals a systematic disagreement in the T$_{\\rm eff}$ scale. The source of this disagreement is unknown. ",
        "introduction": "The recently discovered nearby T-dwarf companion to $\\epsilon$ Indi (Scholz et al. 2003) will be important in improving our understanding of the behavior of the brown dwarfs that fall within the newly defined L and T spectral types (e.g., Kirkpatrick et al. 1999; Geballe et al. 2002). The T-dwarf reported by Scholz et al. (2003) was discovered by Volk et al. (2003) to be a close optical double consisting of an early T dwarf ($\\epsilon$ Indi B) and a late T dwarf ($\\epsilon$ Indi C) separated by 0.6 arcseconds. With an accurately known distance of 3.626$\\pm$0.009 pc, $\\epsilon$ Indi Ba and Bb are the nearest known brown dwarfs.  They share a common proper motion with the K5V star $\\epsilon$ Indi, lying at a projected distance of $\\sim$1460 AU from their presumed primary star.  The age of $\\epsilon$ Indi itself has been estimated by Lachaume et al. (1999) to be $\\sim$0.8 to 2.0 Gyr: this age estimate is based upon its rotational velocity and Ca II K-line emission.  $\\epsilon$ Indi Ba and Bb are thus brown dwarfs with very well-defined luminosities and approximate ages.  We present the first high-resolution infrared spectroscopic observations of $\\epsilon$ Indi B.  Synthetic spectra computed from unified cloudy models by Tsuji (2002) are compared to the observed high-resolution spectra; these comparisons are used to derive the stellar parameters of effective temperature (T$_{\\rm eff}$), surface gravity (defined as log g), and projected rotational velocity (vsin(i)). ",
        "conclusions": "We have used high-resolution infrared spectra of the nearest brown dwarf, $\\epsilon$ Ind B, to derive its physical parameters using comparisons to synthetic spectra calculated from model atmospheres. The spectroscopic T$_{\\rm eff}$=1500K, with log g=5.2, and an estimated mass of M=32M$_{\\rm Jupiter}$.  The projected rotational velocity is vsin($\\iota$)=28 km s$^{-1}$, indicating a maximum rotational period of $\\sim$3.0 hours.  The comparison between the T$_{\\rm eff}$ scales derived from spectroscopic plus model atmosphere analyses against those derived from M$_{\\rm bol}$--radius relations reveals a significant systematic difference that is unexplained."
    },
    "0311/hep-ph0311228_arXiv.txt": {
        "abstract": " ",
        "introduction": "In addition to their pivotal role in elucidating non-perturbative aspects of string theory, in recent years D-branes have become ubiquitous in particle phenomenology and cosmology. In the brane world scenario \\cite{HW,ADD,TeV,Lykken,Ovrut,RS}, the Standard Model gauge bosons and matter fields are localized on some space-filling branes whereas only gravity and closed string modes (such as the radions) can propagate in the bulk, thus offering the possibility of a lower fundamental energy scale. It is worth noting that in this framework the D-branes are part of the string theory vacuum. The Standard Model particles and the bulk modes are excitations above the vacuum described by {\\it perturbative} string states in the background of D-branes. Therefore, even though the branes themselves are non-perturbative objects, the physics of the brane world scenario can be studied using standard techniques of string perturbation theory.  In this paper, we explore yet a different phenomenological feature of D-branes. Unlike the background D-branes in the context of brane world, the D-branes we consider are not space-filling but rather behave as point particles in four dimensions. These particle-like states arise when all spatial dimensions of $Dp$-branes (irrespective of $p$) are wrapped around the compact space. We refer to them collectively as {\\it D-matter}. As we will discuss in more detail, these D-matter states appear quite naturally in four-dimensional string models. They are non-perturbative objects with mass $m \\sim M_s/g_s$ where $g_s$ is the string coupling and so at weak coupling they are heavier than the perturbative string states. However, in order to produce the observed gauge and gravitational couplings, $g_s$ is not arbitrarily small but is typically of order one. Hence, the D-matter states can be sufficiently light to be phenomenologically and cosmologically relevant. The D-matter states are dynamical degrees of freedom and so, in contrast to the background D-branes, we will treat them as excitations above the vacuum.   The stability of D-branes is due to the charges that they carry. Depending on their types (BPS or non-BPS), D-branes could carry integral or torsion (discrete) charges. In fact, it is because of the torsion charges that the complete spectrum of D-branes  should be classified by K-theory \\cite{Moore,Witten,Horava}. Among the D-matter states, the lightest D-particle (LDP) is stable because it is the lightest state carrying its specific charges. Therefore, just like the lightest supersymmetric partners (LSPs) in supersymmetric models, the LDPs are possible candidates for cold dark matter. As we will show, an important difference between the D-matter states and other non-perturbative objects (such as magnetic monopoles \\footnote{Magnetic monopoles are common in string models and their cosmological implications have been explored \\cite{Witten:2002wb}. As we will discuss, there are similarities and differences between D-matter and magnetic monopoles.}) is that they could have {\\bf {\\it perturbative} }couplings. % Hence, the LDPs are weakly interacting and so they could be candidates for wimps or wimpzillas depending on the string scale. We also comment on the case where the D-matter states are unstable (i.e., the spectrum of stable D-branes does not contain four-dimensional particle-like states). Although they are not long-lived to be cosmologically relevant, they could be produced at colliders and give rise to interesting signatures.   In addition to D-matter which are particle-like states (from a four-dimensional perspective), other types of stable defects (such as cosmic strings and domain walls) could also exist in a string vacuum. Given a specific string model, it is straightforward to deduce the complete spectrum of such defects. In this paper, we focus on the general properties of D-matter and point out their relevance to phenomenology. We also stress that the cosmological constraints on the various kinds of stable defects could provide important additional criteria on the viability of a string model (and hence in estimating the number of realistic string vacua \\cite{Douglas:2003um}). For example, we can rule out a string vacuum if there exist stable electrically charged D-matter states.   This paper is organized as follows. In Section \\ref{properties}, we describe the  properties of D-matter including its stability, mass, and interactions.   In Section \\ref{phenomenology}, we discuss the phenomenological implications and some possible signatures of D-matter. Finally, we end with some discussions and conclusion in Section \\ref{conclusion}. Some details of the four-point heterotic string amplitude which is dual to the annihilation cross-section of D-matter states are relegated to the appendix. ",
        "conclusions": "In this paper, we study the phenomenology of D-matter, i.e., particle-like states originated from D-branes whose spatial dimensions are all wrapped around the compact space. An interesting feature of D-matter is that although they are non-perturbative objects, they could have perturbative couplings with each other and with the Standard Model fields. Therefore, the D-matter states are weakly interacting and the lightest stable D-matter states (whose stability is due to integral or torsion charges they carry) are candidates for wimps or wimpzillas depending on the string scale. In the case of a low string scale, there are also potentially rich collider phenomenologies. Exploring the mass pattern of the lower lying D-matter states could reveal additional information about the structure of the compact dimensions.For illustrative purposes, we consider the simple case where all the Standard Model particles are localized on a single set of $Dp'$-branes. It would be interesting to explore the phenomenology of D-matter in more realistic brane world models such as models with background intersecting D-branes \\cite{CSU1,CSU2,CSU3,CIS,Berkooz,Berlin1,Sagnotti,Madrid1,Berlin2,Madrid2,Berlin3}. In these models, the gauge bosons and matter fields are localized differently in the extra dimensions (e.g., in different codimensions) and hence they will couple with different strengths to the D-matter.   Another interesting direction motivated by this work is to compute the spectra of stable D-branes (and their corresponding K-theories) in some realistic brane world models, e.g., certain four-dimensional $N=1$ supersymmetric orientifold models whose perturbative string spectra contain the Standard Model or Grand Unified Theories \\cite{CSU1,CSU2,CSU3,CIS}. The K-theories which classify the D-brane spectra have been computed for some concrete orientifold models, e.g., orientifolds with $N=2$ supersymmetries (in a four-dimensional sense) in  \\cite{Quiroz,Braun}. It would be interesting to extend these analyses to models with $N=1$ supersymmetry where chiral fermions can appear. Such analyses could also help to better understand the types of defects (from a four-dimensional perspective) that are formed toward the end of brane inflation \\cite{DvaliTye}, besides the cosmic strings that are generically produced  \\cite{Jones,Sarangi,ShiuTyeWasserman,Shiu,Pogosian}. Therefore, in addition to the conditions one often imposes (such as appropriate gauge groups, three generations of chiral matter, etc)  in estimating the statistics of realistic string theory vacua \\cite{Douglas:2003um}, the cosmological constraints on the existence of stable defects (such as cosmic strings and domain walls) may further restrict the space of viable string models.   \\appendix"
    },
    "0311/nucl-ex0311002_arXiv.txt": {
        "abstract": "A multi-detector array has been constructed for the simultaneous measurement of energy- and angular correlation of electron-positron pairs produced in internal pair conversion (IPC) of nuclear transitions up to 18 MeV. The response functions of the individual detectors have been measured with mono-energetic beams of electrons. Experimental results obtained with 1.6 MeV protons on targets containing $^{11}$B and $^{19}$F show clear IPC over a wide angular range. A comparison with GEANT simulations demonstrates that angular correlations of $e^+e^-$ pairs of transitions in the energy range between 6 and 18 MeV can be determined with sufficient resolution and efficiency to search for deviations from IPC due to the creation and subsequent decay into $e^+e^-$ of a hypothetical short-lived neutral boson. ",
        "introduction": "\\label{sec:intro} Spectroscopy of internal pair conversion (IPC), has a long tradition \\cite{wilson-siegbahn}. It has been an essential tool to identify nuclear transitions with energies in the few MeV region. This paper describes an experimental arrangement which has been used in a search for elusive neutral bosons. In the investigations with this equipment we obtained positive signals for deviations from IPC which may be of basic interest.  The IPC-branching ratio relative to the competing process of ${\\gamma}$-ray emission and the angular correlation of the pairs depend on the multipolarity of the transition and measurements of IPC have enabled an accurate determination of energy, spin and parity of many nuclear levels. Magnetic spectrometers have been used to identify coincident $e^+e^-$ pairs in order to measure the multipolarity and, with high precision, the energy of the corresponding transition. However, the excellent energy resolution of magnetic spectrometers can only be achieved at the expense of a small acceptance and hence long counting times were needed for such measurements.  For high-statistics measurements of the angular correlation of $e^+e^-$ pairs, primarily detector telescopes with large solid angles were used. Each telescope consists of two detectors, usually a combination of a thin ($\\Delta$E) and a thick (E) detector operated in coincidence to suppress singles ${\\gamma}$-ray events (e.g. from Compton processes). The angular correlation was then measured stepwise with two telescopes placed at different angles or, preferably, simultaneously by a multi-telescope assembly.  Following the early experience obtained in the fifties and sixties, several new large acceptance $e^+e^-$ pair spectrometers have been constructed in the nineties in particular at KVI Groningen \\cite{budaThesis,budaNIM,buda-nucl-phys,budaphysrev,budaphysrevlett} and Stony Brook \\cite{montoya,schadmand1,schadmand2}.  Some years ago, a search for neutral bosons \\cite{fokke0}, possibly emitted in nuclear transitions and decaying into $e^+e^-$ pairs, has been started offering a new challenge to old techniques. Within the constraints of energy and spin-parity conservation, such a neutral boson would compete with ${\\gamma}$-ray and IPC decay in certain nuclear transitions. The signature of such a boson would be its two-body decay into an electron-positron pair.  IPC occurs when a nucleus emits an $e^+e^-$ pair instead of a ${\\gamma}$-ray quantum. Together with the recoil of the emitting nucleus this forms a planar three-body decay with usually a small $e^+e^-$ opening angle, E0 transitions forming a noticeable exception to this rule. However, in the conceivable case when the nucleus emits a neutral boson, instead of IPC, this happens by two-body decay. When the boson subsequently decays into an $e^+e^-$ pair, again a two-body decay occurs, resulting in back-to-back emission of the $e^+e^-$ pairs in the centre of mass system. In the laboratory system the momentum of the decaying boson makes this angle smaller than 180$^{o}$ but leaves the ${\\Theta}$ distribution totally different from IPC: the ${\\Theta}$ distribution becomes truly instrumental in searches for such elusive neutral particles.    Various experimental studies have provided constraints on the possible existence of light, elusive neutral bosons \\cite{gunion,rev}. Short-lived neutral bosons have been searched for in beam dump experiments \\cite{riordan,davier,bross} for the mass region between the $e^+e^-$ threshold at 1.022 MeV/c$^2$ and the mass of the muon (105.7 MeV/c$^{2}$). It is interesting to note that these searches still allow a mass-lifetime window for masses between 5 and 100 MeV/c$^{2}$ at lifetimes shorter than 10$^{-13}$\\, s.  A decade ago our interest in the possible existence of a hitherto undiscovered neutral boson was also triggered by results from emulsion studies \\cite{elnadibadawy,deboervandantzig} of relativistic heavy ion reactions, in which a distinct cluster of $e^+e^-$ pairs was observed at short distance from the interaction vertex. Interpreted in terms of the emission and subsequent decay of a light $X$ boson, the relevant events show an average invariant mass $m_X$ of ${\\sim}$ 9 MeV/c$^{2}$ and lifetime ${\\tau}_X$ $\\sim 10^{-15}$ s.   The possibility to investigate the open mass-lifetime window provided the motivation \\cite{froehlich} to build an apparatus to perform high-statistics IPC measurements of E1 and M1 transitions and to search for anomalies in the angular correlations of $e^{+}e^{-}$ pairs with invariant mass between 5 and 15 MeV/c$^{2}$ and a lifetime shorter than {\\mbox{$10^{-10}$ s.}  Here, we report on the design and construction of the apparatus at the Institut f\\\"ur Kernphysik, Frankfurt (IKF) and describe some experiences with the equipment during a number of experiments \\cite{fokke0,fokke1,fokke2,fokke3} which demonstrate its performance. ",
        "conclusions": "\\label{conclusion}  A multi-detector array has been constructed to measure energy and angular correlations of IPC from high-energy nuclear transitions. The information of all possible combinations of the eight lepton detector telescopes allows the simultaneous detection of $e^{+}e^{-}$ correlation angles between 20 and 130 degrees. The energy resolution of the lepton telescopes is sufficient to identify and separate the nuclear transitions by means of the sum energy and energy asymmetry for lepton pairs from transitions in the energy range between 5 and 18 MeV. By comparing the results of various calibration experiments to extended Monte Carlo simulations with the program code GEANT, it has been shown that the spectrometer is well suited to perform searches for an elusive short-lived neutral boson with a mass between 6 and 15 MeV/c$^{2}$.  A series of dedicated experimental studies have already been performed with this setup. In one experiment, an excellent agreement with conventional IPC was found for the isovector E1 transitions at 12.2 and 17.2 MeV in $^{12}$C. These measurements result in an upper limit of $2.4{\\cdot}10^{-5}$ for the branching ratio of a hypothetical boson with a mass of 9 MeV/c$^{2}$. The latter value is significantly below the theoretical pair conversion coefficient \\cite{rose} of $3.9{\\cdot}10^{-3}$ for a 17.2 MeV E1 transition, indicating the high sensitivity of the experimental device."
    },
    "0311/astro-ph0311598_arXiv.txt": {
        "abstract": "We discuss the submillimeter properties of Extremely Red Objects (EROs) in the two   Canada-UK Deep Submillimeter Survey (CUDSS) fields. We measure the mean submillimeter flux of the ERO population (to {\\kp}$<$ 20.7) and find  0.40 $\\pm$ 0.07 mJy for EROs selected by $(I-K)>$ 4.0  and 0.56 $\\pm$ 0.09 mJy for EROs selected by  $(R-K)> $ 5.3 but,  these measurements are dominated by discrete, bright submillimeter sources. We estimate that EROs produce  7-11\\% of the far-infrared background at 850\\micron. This is substantially less than a previous measurement by \\citet{weh02} and we discuss possible reasons for this discrepancy.    We show that ERO counterparts to bright submillimeter sources lie within the starburst region of the  near-infrared color-color plot of  \\citet{poz00}.  Finally, we claim that pairs or small groups of EROs with separations of $\\lesssim$ 10 arcsec often mark regions of strong submillimeter flux. ",
        "introduction": "A complete accounting of all the systems which produce the extragalactic background light is key to our understanding of galaxy assembly and evolution. Observations of the early universe  have revealed a  wide variety of objects  whose relationship to each other,  importance to galaxy formation, and  contribution to the extragalactic background energy remains unclear. The fact that these systems have been selected over different wavelength ranges and using different selection criteria has made it is difficult to determine if we are studying different aspects of the same objects, or completely orthogonal populations.  Many of these objects are found over similar redshift ranges and are co-evolving, while others may be linked through time by evolutionary processes.  At near-infrared wavelengths a population of red objects begins to emerge at faint, $K\\gtrsim$ 18, magnitudes \\citep{els88,mcc92,hu94}. This class of systems,  labeled Extremely Red Objects (EROs), is photometrically defined by NIR color, typically $(I-K)>$ 4 or  $(R-K)> $  5 $-$ 5.3.  Two different galaxy populations are selected  by this color definition: an old stellar population at $z\\sim$1, or alternatively, a young starburst at high-redshift, deeply enshrouded in dust.  Low-mass stars and brown-dwarfs within the Galaxy also have similar NIR colors but have such low surface densities relative to faint galaxies that they do not significantly contaminate ERO samples.  Determining the fraction EROs belonging to each category has been difficult work, hindered by the faintness of these objects at optical and NIR wavelengths. Spectroscopic observations indicate that  evolved stellar populations and dusty starbursts are present in roughly equal numbers \\citep{dun96,gra96,coh99,cim02a} and this picture is generally supported by larger studies employing  less expensive, but more ambiguous,  techniques such as color \\citep{poz00,mar01,sma02a}, radio emission \\citep{sma02a}, or morphology \\citep{yan03,gil04}. There is substantial variation between the results however, stemming from different survey depths, varying color-selection criteria, and clustering effects \\citep{dad00,mccar01,smi02,sma02a,dad02,yan03}.   The detection of submillimeter emission from an ERO is a direct confirmation of a starburst nature (ignoring for the moment the problem of confusion within the beam). At current observational limits however submillimeter imaging can only detect the most luminous star forming EROs, those that overlap with the discrete submillimeter-bright population  at high-redshift  discovered with SCUBA (SMGs)  \\citep{sma97, bar98, hug98, eal00, sco02} which  are sites of  perhaps the most substantial star formation in the high-redshift universe.  Though still suffering from small number statistics $\\sim$20-30\\% of the systems detected deep submillimeter surveys  are EROs (when observed in the NIR)  \\citep{sma99,ivi02,cha03a,web03c} though this fraction obviously depends on the relative depths of the the 850{\\micron} and $K$-band imaging and the uncertainty in the in the identifications of the counterparts of non-radio detected SMGs.  What is not clear, however, is how many of the $K$-selected EROs are discrete, bright submillimeter sources and the submillimeter propeties of  EROs below the current detection limits of submillimeter surveys. These less luminous systems, though not individually submillimeter-bright could together produce a substantial fraction of the FIR background light. Two previous studies have attempted to address  these issues through statistical studies EROs \\citep{moh02,weh02} and the results of both groups indicate that as a population these systems  emit substantial submillimeter flux. In particular, \\citet{weh02} (hereafter W02) found $K^{\\prime} <$ 21.6 EROs produce $\\sim$50\\% of the FIR background light at 850$\\mu$m, with  this contribution peaking at $K^{\\prime}\\sim$20.    However, as with all ERO studies, these results may depend on  ERO selection criteria and field-to-field variance.  Thus, the relationship between EROs and bright SMGs, and the submillimeter properties of  EROs below the detection limit of SCUBA, remain  open questions, as does the ERO  contribution to the far-infrared background energy. Those EROs which are extremely luminous in the submillimeter are a minority of both the ERO population and the SMGs (at least to current $K$ depths) but must represent an important phase of galaxy evolution. Understanding their natures will provide a key link between these two populations. Isolating  the properties of EROs which are individually or statistically submillimeter bright which  differentiate them from the rest of the ERO population is an important step toward this understanding.  Such properties may include  the NIR colors and magnitudes \\citep{poz00,weh02},  morphologies, redshifts and the local environment.  Environment plays a key role in galaxy evolution and there is growing evidence of a correlation between submillimeter emission and high-density regions or galaxy merging and interactions \\citep{ivi00,sma03,ale03,web03c,ner03,cha03c} and this may be  be an important driver of star-formation in the ERO population. Additionally, the different redshift distributions of the EROs and SMGs (as far as we know them) \\citep{cha03b} hint at an evolutionary relationship between these populations perhaps with SMGs fading into less luminous star-forming EROs at later times.  Addressing these issues will not only illuminate the natures of these two populations but the process of galaxy evolution in general.  The paper is organized as follows. In \\S2 and \\S3  we present the data and define our ERO samples. In \\S 4 we measure the measure  submillimeter flux of EROs. We discuss the ERO contribution to the extragalactic background light in \\S 5. In \\S 6  we present the \\citet{poz00} color-color plot for these systems. In \\S 7 we  discuss the effect of ERO companions in the analysis. We finish in \\S 8 with some thoughts difference between submillimeter-bright and faint EROs.  Throughout we have used a flat, $\\Lambda$=0.7 cosmology and H$_{\\circ}=$ 72 km/s/Mpc. ",
        "conclusions": "We investigated the submillimeter properties of EROs in the CUDSS fields and found the following:  \\begin{enumerate} \\item{ We detect significant submillimeter flux from EROs in the CUDSS fields of  0.40 $\\pm$ 0.07 mJy or 0.56 $\\pm$ 0.09 mJy for $(I-K)>$ 4.0 and $(R-K)>$ 5.3 ERO selection criteria, respectively.  This implies EROs to a depth of $K<$ 20.7 produce 7-11\\% of the extragalactic background light at 850\\micron. However, the flux  is dominated by a minority  of systems which are very bright at submillimeter wavelengths and overlap with the discrete population of submillimeter-bright galaxies discovered with SCUBA.} \\item{EROs with  strong submillimeter flux  lie on the expected starburst side of the  NIR color-color diagram, as suggested by \\citet{poz00}. } \\item{Pairs or small groups of EROs mark, on average, regions of strong submillimeter flux, indicating that star formation rates are enhanced within these high-density areas.  This implies, some fraction of strong submillimeter emitters trace over-dense regions in the early universe.  We are unable to tell from these data if more than one object within these groups is responsible for the submillimeter emission}. \\item{We tentatively suggest that EROs which show submillimeter emission  are the higher-redshift and higher-luminosity tail of the dusty starburst EROs within these data,  and may have a history of significant and extended  merger activity.} \\end{enumerate}   {\\it ACKNOWLEDGMENTS}  We thank the anonymous referee for many helpful comments. Research by Tracy Webb is supported by the NOVA Postdoctoral Fellowship program.  Research by Steve Eales is funded through the Particle Physics and Astronomy Research Council and Leverhulme Trust.  We wish to thank Natascha F\\\"orster Schreiber, Doug Johnstone, and Kirsten Kraiberg Knudsen for very helpful discussions.  The JCMT is operated by the Joint Astronomy Center on behalf of the  Particle Physics and Astronomy Research Council in the UK, the National Research Council of Canada and the Netherlands Organization for Scientific Research.  The Canada-France-Hawaii Telescope is jointly operated by the National Research Council of Canada, the Centre National de la Recherche Scientifique of France and the University of Hawaii."
    },
    "0311/astro-ph0311551_arXiv.txt": {
        "abstract": "\\vspace{-6mm} \\centerline{\\bf Abstract}   We report on the results of eighteen observations of the core, or peak, of the \\fekalfa emission line at $\\sim 6.4$~keV in fifteen Seyfert~I galaxies using the \\chandra High Energy Grating (HEG). These data afford the highest precision measurements of the peak energy of the \\fekalfa line, and the highest spectral resolution measurements of the width of the core of the line to date. We were able to measure the peak energy in seventeen data sets, and, excluding a very deep observation of NGC~3783, we obtained a weighted mean of $6.404 \\pm 0.005$~keV. In all fifteen sources the two-parameter, 99\\% confidence errors on the line peak energy do not exclude fluorescent $K\\alpha$ line emission from Fe~{\\sc i}, although two sources (Mkn~509 and 3C~120) stand out as very likely being dominated by $K\\alpha$ emission from Fe~{\\sc xvii} or so. We were able to measure the line core width in fourteen data sets and obtained a weighted mean of $2380 \\pm 760 \\ \\rm km \\ s^{-1}$ FWHM (excluding the NGC~3783 deep exposure), a little larger than the instrument resolution ($\\sim 1860 \\ \\rm km \\ s^{-1}$ FWHM). However, there is evidence of underlying broad line emission in at least four sources. In fact, the width of the peak varies widely from source to source and it may in general have a contribution from the outer parts of an accretion disk {\\it and} more distant matter. For the disk contribution to also peak at 6.4~keV requires greater line emissivity at hundreds of gravitational radii than has been deduced from previous studies of the \\fekalfa line.  {\\bf Keywords:} accretion disks -- galaxies: active -- line: profile -- X-rays: galaxies  \\begin{center} {\\it Accepted for Publication in the Astrophysical Journal 21 November 2003} \\end{center} ",
        "introduction": "\\label{intro}  \\begin{table*}[h] \\caption{\\chandra Parameters of the Core Fe K Line Emission from $\\it Chandra$ (HEG) Data} \\begin{tabular*}{\\textwidth}{@{}l@{\\extracolsep{\\fill}}cccccc} \\hline Source, & $z^{a}$ & $E^{b}$ & $I^{c}$ & $  {\\rm EW}^{d}$ & $ {\\rm FWHM}^{e}$ & $ F^{f}$  \\\\ \\& Refs. $^{g}$ & & (keV) & & (eV) & ($\\rm km \\ s^{-1}$) & $ L^{f}$ \\\\ \\hline & & & & & & \\\\  NGC 7314 & 0.004760 & $6.412^{+0.010}_{-0.015}$ & $1.7^{+0.6}_{-0.8}$ & $48^{+17}_{-23}$ & $100$ f & 3.0 \\\\ \\vspace{-1mm} $[1]$ & & ($6.395 - 6.430$) & ($0.7 - 2.7$) & ($20 - 76$) & \\ldots & $0.15$ \\\\ NGC 3516(1) & $0.008836$ & $6.398^{+0.017}_{-0.007}$ & $4.0^{+1.3}_{-1.2}$ & $140^{+46}_{-42}$ & $1290^{+1620}_{-1290}$ & $2.3$ \\\\ \\vspace{-1mm} $[2]$ &    & ($6.389 - 6.408$) & ($2.5 - 5.8$) & ($88 - 203$) & ($0 - 3630$) & $0.40$   \\\\ NGC 3516(2) & $0.008836$ & $6.401^{+0.015}_{-0.017}$ & $3.8^{+1.5}_{-1.3}$ & $143^{+56}_{-49}$ & $3630^{+2350}_{-1540}$ & $2.1$  \\\\ $[2]$  &    & ($6.378 - 6.422$) & ($2.2 - 5.8$) & ($83 - 218$) & ($1560 - 7640$) & $0.35$ \\\\ Mkn 509 & $0.034397$ & $6.430^{+0.024}_{-0.023}$ & $3.1^{+1.9}_{-1.6}$ & $54^{+33}_{-28}$ & $2820^{+2680}_{-2800}$ & $5.3$ \\\\ \\vspace{-1mm} \\ldots   &   & ($6.394 - 6.461$) & ($0.9 - 5.7$) & ($16 - 99$) & ($0 - 7710$) & $14.0$ \\\\ NGC 5548(1) & $0.016760$ & $6.397^{+0.019}_{-0.023}$ & $3.1^{+1.5}_{-1.2}$ & $115^{+56}_{-45}$ & $3750^{+2590}_{-1890}$ & $2.4$  \\\\ \\vspace{-1mm} $[3]$  &    & ($6.365 - 6.422$) & ($1.5 - 5.1$) & ($56 - 189$) & ($1250 - 7610$) & $1.5$ \\\\ NGC 5548(2) & $0.016760$ & $6.400^{+0.010}_{-0.010}$ & $2.1^{+0.8}_{-0.7}$ & $67^{+26}_{-22}$ & $1780^{+1420}_{-1220}$ & $2.9$  \\\\ \\vspace{-1mm} $[4]$   &    & ($6.386 - 6.414$) & ($1.2 - 3.3$) & ($38 - 105$) & ($0 - 3970$) & $1.8$ \\\\ 3C 120 & $0.033010$ & $6.415^{+0.017}_{-0.017}$ & $3.0^{+1.8}_{-1.5}$ & $61^{+37}_{-31}$ & $2000^{+2950}_{-2000}$ & $4.4$ \\\\ \\vspace{-1mm} \\ldots    &   & ($6.392 - 6.450$) & ($1.1 - 5.5$) & ($22 - 112$) & ($0 - 6770$) & $10.8$ \\\\ NGC 4593 & $0.008301$ & $6.403^{+0.012}_{-0.038}$ & $3.4^{+3.2}_{-1.3}$ & $80^{+75}_{-31}$ & $2140^{+8370}_{-1230}$ & $4.1$ \\\\ \\vspace{-1mm} \\ldots  &    & ($6.349 - 6.434$) & ($1.8 - 8.1$) & ($42 - 191$) & ($310 - 15920$) & $0.63$ \\\\ NGC 3783(1) & $0.009730$ & $6.400^{+0.014}_{-0.016}$  & $5.0^{+2.4}_{-2.1}$ & $73^{+35}_{-31}$ & $2550^{+2420}_{-1560}$ & $6.0$ \\\\ \\vspace{-1mm} $[5]$   &   & ($6.376 - 6.420$) & ($2.4 - 8.3$) & ($35 - 121$) &  ($500 - 6290$) & $1.2$ \\\\ NGC 3783(2) & $0.009730$ & $6.397^{+0.003}_{-0.003}$ & $4.9^{+0.6}_{-0.5}$ & $70^{+9}_{-7}$ & $1700^{+410}_{-390}$ & $6.2$ \\\\ \\vspace{-1mm} $[6,7]$    &   & ($6.393 - 6.401$) & ($4.2 - 5.6$) & ($60 - 80$) & ($1180 - 2250$) & $1.3$ \\\\ MCG -6-30-15 & $0.007749$ & $6.408^{+0.030}_{-0.025}$ & $1.6^{+1.1}_{-0.9}$ & $49^{+34}_{-28}$ & $3250^{+5230}_{-3250}$ & $3.2$  \\\\ \\vspace{-1mm} $[8]$  &    & ($6.349 - 6.454$) & ($0.4 - 3.3$) & ($12 - 101$) & ($0 - 12670$) & $0.42$ \\\\ Mkn 279 & $0.030451$ & $6.415^{+0.047}_{-0.027}$  & $1.9^{+1.0}_{-0.9}$ & $132^{+69}_{-63}$ & $5010^{+6550}_{-2810}$ & $1.3$  \\\\ \\vspace{-1mm} \\ldots  &   & ($6.379-6.521$) & ($0.7-5.2$) & ($49-361$) &  ($1100-36980$) & $2.7$ \\\\ NGC 4051 & $0.002336$ & $6.419^{+0.038}_{-0.033}$ & $3.2^{+1.7}_{-1.4}$ & $191^{+101}_{-84}$ & $6330^{+7740}_{-3330}$ & $1.6$ \\\\ \\vspace{-1mm} $[9]$   &   & ($ 5.910 - 6.475$) & ($1.4 - 25.0$) & ($84 - 1492$) & ($1780 - 137470$) & $0.019$ \\\\ IC 4329A & $0.016094$ & $6.309^{+0.089}_{-0.099}$ & $11.3^{+7.2}_{-6.4}$ & $62^{+40}_{-35}$ & $15090^{+12430}_{-9950}$ & $16.4$  \\\\ \\vspace{-1mm} \\ldots  &    & ($6.171 -6.468$) & ($1.8 - 21.3$) & ($10 - 117$) & ($0 - 37150$) & $9.4$ \\\\ F 9 & $0.047016$ & $6.373^{+0.254}_{-0.092}$ & $4.9^{+8.7}_{-3.3}$ & $216^{+384}_{-145}$ & $17040^{+55960}_{-14270}$ & $2.1$  \\\\ \\vspace{-1mm} \\ldots  &  & ($6.203 - 6.753$) & ($1.2 - 17.0$) & ($53 - 749$) & ($2160 - 92940$) & $10.6$ \\\\ Mkn 766 & $0.012929$ & $6.423^{+0.018}_{-0.016}$ & $0.7^{+0.6}_{-0.5}$ & $34^{+29}_{-24}$ & $100$ f & $2.1$  \\\\ \\vspace{-1mm} \\ldots   &    & ($6.398 - 6.456$) & ($0.1 - 1.5$) & ($5 - 73$) & \\ldots & $0.79$ \\\\ NGC 3227 & $0.003859$ & $6.384^{+0.015}_{-0.016}$ & $1.1^{+1.2}_{-0.9}$ & $39^{+43}_{-32}$ & $100$ f & $2.3$  \\\\ \\vspace{-1mm} \\ldots   &    & \\ldots & ($0 - 2.8$) & ($0 - 99$) & \\ldots & $0.075$ \\\\ Akn 564 & $0.024684$ & $6.400$ f & $0.3^{+0.5}_{-0.3}$ & $16^{+27}_{-16}$ & $100$f & $2.5$  \\\\ \\vspace{-1mm} $[10]$   &    & \\ldots & ($0.0 - 1.1$) & ($0 - 59$) & \\ldots & $3.4$ \\\\ & & & & & & \\\\ \\hline \\end{tabular*} {\\small {\\it Chandra} HEG data, fitted with a power law plus Gaussian emission line model in the 2--7 keV band (for NGC 3516 photoelectric absorption was also included in the model). All parameters (except redshift) were free in the fits, except in the cases of NGC~7314, Mkn~766, NGC 3227 (line width fixed in these cases), and Akn 564 (line energy and width fixed). All parameters are quoted in the source rest frame. Statistical errors are for the 68\\% confidence level, whilst parentheses show the 90\\% confidence level ranges of the parameters. The number of interesting parameters assumed for calculating the statistical errors was equal to the number of free parameters in the Gaussian component of the model. For 3, 2, and 1 interesting parameter(s), the corresponding values of $\\Delta C$ for 68\\% confidence are  3.506, 2.279 and 0.989 respectively. For 3, 2, and 1 interesting parameter(s), the 90\\% confidence values of $\\Delta C$  are 6.251, 4.605 and 2.706 respectively. $^{a}$  Redshifts obtained from NASA Extragalactic Database (NED). $^{b}$  Gaussian line center energy. $^{c}$  Emission-line intensity in units of $\\rm 10^{-5} \\ photons \\ cm^{-2} \\ s^{-1}$. $^{d}$  Emission line equivalent width. $^{e}$  Full width half maximum, rounded to $10 \\ \\rm km \\ s^{-1}$. $^{f}$  $F$ is the estimated 2--10~keV observed flux in units of $10^{-11} \\ \\rm ergs\\ cm^{-2}\\ s^{-1}$. The power-law continuum was extrapolated to 10 keV. $L$ is the estimated 2--10~keV source-frame luminosity (using the 2--10 keV estimated flux), in units of $10^{43} \\ \\rm ergs\\ s^{-1}$, assuming $H_{0} = \\rm 70 \\ km \\ s^{-1} \\ Mpc^{-1}$ and $q_{0}$. $^{g}$ Key for previous publications on the same HEG data: [1] Yaqoob \\etal 2003a, [2] Turner \\etal 2002, [3] Yaqoob \\etal 2001, [4] Kaastra \\etal 2002, [5] Kaspi \\etal 2001, [6] Kaspi \\etal 2002, [7] Netzer \\etal 2003, [8] Lee \\etal 2002, [9] Collinge \\etal 2001, [10] Mastumoto, Leighly, \\& Marshall 2001, http://www.pha.jhu.edu/groups/astro/workshop2001/papers/matsumoto\\_c.ps}. \\end{table*}  \\begin{figure*}[h] \\vspace{10pt} \\centerline{\\psfig{file=tyf1a.eps,width=6.5in,height=6.5in,angle=180}} \\caption{\\footnotesize \\chandra HEG spectra in the Fe K band from each of the eighteen observations of fifteen type~I AGN (see \\tablefitsp). The data are binned at $0.01\\AA$ for the first sixteen spectra shown, and $0.02\\AA$ for the last two spectra (NGC~3227 \\& Akn~564). This can be compared to the HEG spectral resolution, which is $0.012\\AA$ FWHM. The data are combined from the $-1$ and $+1$ orders of the grating. The spectra have been corrected for instrumental effective area and cosmological redshift. Note that these are {\\it not} unfolded spectra and are therefore independent of the model that is fitted. The statistical errors shown correspond to the $1\\sigma$ Poisson errors, calculated using equations (7) and (14) in Geherls (1986) to approximate the upper and lower errors respectively. The solid line corresponds to the continuum model fitted over the 2--7~keV range (extrapolated to 7.5~keV), as described in the text (\\S\\ref{hegspec}). The vertical dotted lines represent (from left to right), the rest energies of the following transitions: Fe~{\\sc i}~$K\\alpha$, \\fexxv forbidden, two intercombination lines of \\fexxvp, \\fexxv resonance, and \\feklyap. The spectrum for NGC~7314 corresponds to that during a low continuum state, as defined and described in Yaqoob \\etal 2003a. Note that an apparent narrow feature at $\\sim 6.35$~keV in MCG~$-$6-30-15 and at $\\sim 6.55$~keV in NGC~3783(2) are detected in one arm of the grating only and are narrower than the spectral resolution, so are not real.} \\end{figure*}  \\setcounter{figure}{0} \\begin{figure*}[tbh] \\vspace{10pt} \\centerline{\\psfig{file=tyf1b.eps,width=6.5in,height=6.5in,angle=180}} \\caption{ -- {\\it continued}} \\end{figure*}  \\setcounter{figure}{0} \\begin{figure*}[tbh] \\vspace{10pt} \\centerline{\\psfig{file=tyf1c.eps,width=6.5in,height=6.5in,angle=180}} \\caption{ -- {\\it continued}} \\end{figure*}   At least part of the \\fekalfa fluorescent emission line in type~I active galactic nuclei (AGNs) is believed to originate in a relativistic accretion disk around a black hole (e.g. see reviews by Fabian \\etal 2000; Reynolds \\& Nowak 2003). The dominant peak energy of the \\fekalfa line at $\\sim 6.4$~keV appears to be ubiquitous and this core of the line carries a substantial fraction of the total line flux (e.g. Nandra \\etal 1997a; Sulentic \\etal 1998; Lubi\\'{n}ski \\& Zdziarski 2001; Weaver, Gelbord, \\& Yaqoob 2001; Yaqoob \\etal 2002; Perola \\etal 2002; Reeves 2002). Often, the broad part of the \\fekalfa line is absent, leaving only the narrow core. It has been traditional to associate such narrow \\fekalfa lines with an origin in distant matter, at least several thousand gravitational radii from the putative black hole (e.g. the optical broad-line region (BLR), the putative obscuring torus, or the optical narrow-line region (NLR)). However, Petrucci \\etal (2002) recently reported a {\\it variable}, narrow \\fekalfa line in Mkn~841, supporting an accretion-disk origin. Moreover, rapidly variable, narrow Fe K line emission has been observed in the Seyfert~I galaxy NGC~7314 (Yaqoob \\etal 2003a). Thus, even narrow \\fekalfa lines may have a significant contribution from the accretion disk (Lee \\etal 2002; Yaqoob \\etal 2003a). While such lines may be interpreted in terms of a truncated disk (e.g. Done, Madejski, \\& \\.{Z}ycki 2000), they could be due to low-inclination angle disks with a flat radial line emissivity (i.e. intensity per unit area falling off with radius more slowly than $r^{-2}$).  In this paper we address two very specific questions, using the \\chandra high energy grating transmission spectrometer ({\\it HETGS}; see Markert \\etal 1995), which affords the best spectral resolution currently available at 6.4~keV ($\\sim 39$~eV, or $\\sim 1860 \\ \\rm km \\ s^{-1}$ FWHM). Namely, for a sample of type~I AGNs, what are the energies of the peaks of the Fe K line emission and are these line cores resolved by {\\it Chandra}? This information can give important clues about the ionization state of Fe responsible for the line emission, and its origin. With \\chandra we can measure the peak energies with better precision than {\\it ASCA}, at least by a factor of four.  We note that, due to the small throughput of the \\hetg (especially above $\\sim 7$~keV), it is very difficult for the grating data to constrain the parameters of any underlying broad \\fekalfa line emission. However, in a different study we shall systematically compare the total \\fekalfa line emission observed with \\chandra and \\asca data and show that, aside from variability in some sources, there is good agreement between the two sets of data. We emphasize that even though the \\chandra data for most type~I AGN show narrow \\fekalfa line peaks, this by no means indicates that there is no broad \\fekalfa line emission. ",
        "conclusions": ""
    },
    "0311/astro-ph0311284_arXiv.txt": {
        "abstract": "Recent data have radically altered the X-ray perspective on cooling flow clusters.  X-ray spectra show that very little of the hot intracluster medium is cooler than about 1 keV, despite having short cooling times.  In an increasing number of cooling flow clusters, the lobes of a central radio source are found to have created cavities in the hot gas.  Generally, the cavities are not overpressured relative to the intracluster gas, but act as buoyant bubbles of radio emitting plasma that drive circulation as they rise, mixing and heating the intracluster gas.  All this points to the radio source, i.e. an active galactic nucleus, as the heat source that prevents gas from cooling to low temperatures.  However, heating due to bubbles alone seems to be insufficient, so the energetics of cooling flows remain obscure.  We briefly review the data and theory supporting this view and discuss the energetics of cooling flows. ",
        "introduction": "The radiative cooling time of the hot intergalactic gas close to the centres of about 70 percent of rich clusters of galaxies is significantly shorter than the Hubble time.  These systems are known as cooling flows (Fabian 1994).  Over the lifetime of a cooling flow cluster, radiative losses have a significant impact on the gas unless the radiated heat is replaced.  Cooling gas is compressed in order to maintain hydrostatic equilibrium, causing inflow near to the cluster centre and the deposition of large quantities of cool gas.  X-ray data from {\\it Chandra} and {\\it XMM-Newton} confirm central cooling times as short as $10^8$ -- $10^9$ y in many clusters (e.g. David et {al.} 2001), highlighting the significance of radiative losses.  However, 100 -- $1000\\rm\\ M_\\odot\\ y^{-1}$ of cold gas should be deposited by cooling flows (e.g.~White, Jones \\& Forman 1997) and very little evidence is found of this gas.  Many forms of cool gas (e.g. Crawford et {al.} 1999; Edge 2001) and recent star formation (e.g. Mittaz et {al.}  2001) have been found at the centres of cooling flow clusters, but the amounts fall well short of those expected.  Cooling flows also occur in groups and isolated elliptical galaxies (Mathews \\& Brighenti 2003).  X-ray spectra from the Reflection Grating Spectrograph (RGS) on {\\it XMM-Newton} show that there is very little gas cooler than about 1 keV in cluster cooling flows.  If the late stages of cooling are isobaric as expected, then the luminosity in a line is $L_{\\rm line} = \\dot M (5k/(2\\mu m_{\\rm H}) \\int_0^{T_{\\rm max}} \\Lambda_{\\rm line}(T) / \\Lambda(T) \\, dT$, where $\\dot M$ is the deposition rate of cooled gas, $\\Lambda$ is the cooling function and $\\Lambda_{\\rm line}$ the part of the cooling function due to the line.  This prediction is quite robust for low temperature lines, but RGS data show that some low energy lines are at least an order of magnitude weaker than expected (e.g. Peterson et {al.}  2003). ",
        "conclusions": "There is good evidence of AGN heating in cooling flows, but it remains unclear whether it is significant for cooling flows as a whole.  Since AGN heating is linked to cooling by feedback, this process can plausibly explain how gas can be maintained with short cooling times in cooling flows.  AGN heating probably needs to be augmented in order to account for global energetics of cooling flows.  Many details of the heating process remain obscure.  In particular, it is unclear how a cluster can be heated from its centre without producing a constant entropy core and mixing out observed abundance gradients.  To end on a speculative note, AGN outbursts also occur in isolated elliptical galaxies, where they can prevent the cooling of hot gas more readily than in clusters.  If so, AGN feedback inhibits cooling, hence star formation, in almost any system dominated by hot gas.  In that case, the effect of AGN feeback is imprinted on the galaxy luminosity function."
    },
    "0311/astro-ph0311417_arXiv.txt": {
        "abstract": "High resolution 7 mm observations of the W49N massive star forming region have detected recombination line emission from the individual ultracompact (UC) HII regions on 50 milliarcsecond (600 AU) scales. These line observations, combined with multifrequency, high-resolution continuum imaging of the region at 7 mm (VLA) and at 3 mm and 1 mm (BIMA), indicate that five to seven of the eighteen ultracompact sources in W49N are broad recombination line objects (BRLOs) as described by Jaffe \\& Martin-Pintado (1999). BRLOs have both broad radio recombination lines ($\\Delta$V$>$60 \\kms) and rising spectra (S$_{\\nu}\\sim\\nu^{\\alpha}$), with $\\alpha$ values greater than 0.4. The broad line widths of the H52$\\alpha$ line are probably related to motions in the ionized gas rather than pressure broadening. A number of models have been proposed to explain the long lifetime of UC HII regions, including the photoevaporated disk model proposed by Hollenbach et al. (1994). This model can also explain the broad lines, rising spectra and bipolar morphologies of some sources. We suggest$-$based on line and continuum observations as well as source morphology$-$that in a subset of the W49N ultracompact sources we may be observing ionized winds that arise from circumstellar disks. ",
        "introduction": "The W49A star-forming region is among the most luminous in the Galaxy (10$^7$ L$_\\odot$) and is located within one of its most massive giant molecular clouds (10$^6$ M$_\\odot$). While W49A is a distant region (located at about 11.4 kpc), the large number, high luminosity and wide morphological variety of its HII regions make it a valuable region to study high mass star formation, as long as observations have sufficient angular resolution. The W49 North (W49N) cloud core within W49A contains over a dozen ultracompact (UC) HII regions powered by OB-type stars arranged in a 2 pc diameter ring (Dreher et al. 1984; Dickel \\& Goss 1990). The discovery of so many compact sources within W49N was one of the origins of the UC HII region \"lifetime problem\". That is, the observed number of UC sources seems to exceed the numbers expected based on a free expansion of the ionized gas (Dreher et al. 1984; Wood \\& Churchwell 1989). A variety of models have been proposed to account for this discrepancy, and were recently summarized in Jaffe \\& Martin-Pintado (1999; hereafter JMP).  A recent millimeter survey of UC HII regions (JMP) indicates that about 30\\% of them are characterized by (1) rising spectral indices ($\\alpha>$0.4), and (2) broad radio recombination lines ($\\Delta$V$>$60~\\kms), much broader than would be expected if the width were thermal ($\\sim$25 \\kms).  JMP refer to these sources as Broad Line Recombination Objects (BRLOs). In their analysis, JMP discuss the apparent correlation between spectral index values and radio recombination line (RRL) widths, summarize the existing models that could explain the correlation, and conclude that the disk-wind model (Hollenbach et al. 1994) is the most promising. Sewilo et al. (2003) have recently carried out a VLA survey of seventeen hypercompact (HC) HII regions, and report the detection of eight sources with H92$\\alpha$ or H76$\\alpha$ line widths greater than 40 \\kms. They propose hierarchical clumping in the nebular gas as an explanation of the spectral indices observed. Here we describe new H52$\\alpha$ line observations that have uncovered a number of BLROs in W49N consistent with the 30\\%-50\\% of such sources found in these surveys. The scales probed by these observations ($\\sim$600 AU) may further clarify the nature of BRLOs, and provide new constraints for models of these young sources. ",
        "conclusions": "High resolution, high frequency line and continuum observations of the W49N region indicate that broad radio recombination line widths detected on small scales at 1.3 cm persist at 7 mm. In addition, we have detected line emission from four sources (A1, B1, B3 \\& G1S) unresolved in previous H52$\\alpha$ RRL observations of W49N (De Pree et al. 1995a). Six ($\\pm$1) of the 18 continuum sources in the W49N core have characteristics that would classify them as BLROs, roughly in agreement with the 30\\% value found in the JMP survey of UC HII regions. Broad lines in these regions detected at 7mm are likely to have a kinematic origin. The new RRL observations clearly provide a discriminating factor between BRLOs and normal UC HIIs in this massive star forming region."
    },
    "0311/astro-ph0311621_arXiv.txt": {
        "abstract": "{In this lecture I review the theory of magnetized accretion discs driving jets, with a focus on Young Stellar Objects (YSOs). I first introduce observational and theoretical arguments in favor of the ``disc wind'' paradigm. There, accretion and ejection are interdependent, requiring therefore to revisit the standard picture of accretion discs. A simple magnetostatic approach is shown to already provide some insights of the basic phenomena. The magnetohydrodynamic equations as well as all usual assumptions are then clearly listed. The relevant physical mechanisms of steady-state accretion and ejection from Keplerian discs are explained in a model independent way. The results of self-similar calculations are shown and critically discussed, for both cold and warm jet configurations. I finally provide observational predictions and the physical conditions required in YSOs discs. The necessity of introducing a magnetospheric interaction between the disc and the protostar is briefly discussed.} ",
        "introduction": "Collimated ejection of matter is widely observed in several astrophysical objects: inside our own galaxy from all young forming stars and some X-ray binaries, but also from the core of active galaxies. All these objects share the following properties: jets are almost cylindrical in shape; the presence of jets is correlated with an underlying accretion disc surrounding the central mass; the total jet power is a sizeable fraction of the accretion power.  \\subsection{Why should jets be magnetized~?}  There are basically two kinds of jet observations: spectra (blue shifted emission lines) and images (in the same lines for YSOs, or in radio continuum for compact objects). Most of these images show jets that are extremely well collimated, with an opening angle of only some degrees. On the other hand, the derived physical conditions show that jets are highly supersonic. Indeed, emission lines require a temperature of order $T\\sim 10^4$ K, hence a sound speed $C_s \\sim 10$ km/s while the typical jet velocity is $v_j\\sim 300$ km/s. The opening angle $\\theta$ of a ballistic hydrodynamic flow being simply $\\tan \\theta = C_s/v_j$, this provides $\\theta\\sim 5^o$ for YSOs, nicely compatible with observations. Thus, jets could well be ballistic, showing up an inertial confinement. But the fundamental question remains: {\\bf how does a physical system produce an unidirectional supersonic flow~?} Naively, this implies that confinement must be closely related to the acceleration process.  This has soon been recognized in the AGN community (where jets were being observed for quite a long time, see Bridle \\& Perley 1984), leading Blandford \\& Rees to propose in 1974 the ``twin-exaust'' model. In this model, the central source is emitting a spherically symmetric jet, which is confined and redirected into two bipolar jets by the external pressure gradient. Indeed, the rotation of the galaxy would probably produce a disc-like anisotropic distribution of matter around its center. Thus, in principle, the ejected plasma could be focused towards the axis of rotation of the galaxy (where there is less matter there) and thereby accelerated like in a De Laval nozzle. Such an idea was afterwards applied by Cant{\\'o} (1980) and K\\\"onigl (1982) for YSOs.  However, this model had severe theoretical drawbacks, related to Kelvin-Helmholtz instabilities quickly destroying the nozzle. Besides, it does not explain the very origin of the spherical flow.  But the strongest argument comes from recent observations, like those of HH30 (see eg. Ray et al. 1996). We know now that jets are already clearly collimated close to the central star, with no evidence of any relevant outer pressure. This implies that jets must be self-collimated. In my opinion, such an observation also rules out the proposition that jets are collimated by an outer poloidal magnetic field (Spruit et al. 1997).  The only model capable of accelerating plasma along with a self-confinement relies on the action of a large scale magnetic field carried along by the jet. In fact, Lovelace and Blandford proposed independently in 1976 that if a large scale magnetic field would thread an accretion disc, then it could extract energy and accelerate particules (electron positron pairs in their models). Then, Chan \\& Henriksen (1980) showed, using a simplified configuration, that such a field could indeed maintain a plasma flow collimated. But it was Blandford \\& Payne who, in 1982, produced the first full calculation of the interplay between a plasma flow (made of electrons and protons) and the magnetic field, showing both acceleration and self-collimation.  \\subsection{What are the jets driving sources~?}  To make a long story short, there are three different situations potentially capable of driving magnetized jets from young forming stars\\footnote{An alternative model is based on some circulation of matter during the early infall stages (see Lery et al. 1999 and references therein). However, by  construction, such a model is only valid for Class 0 sources and cannot be used to explain jets from T-Tauri stars. See also Contopoulos \\& Sauty (2001).}. \\begin{itemize} \\item {\\bf the protostar alone}: these purely stellar winds extract their energy from the protostar itself (eg. Mestel 1968, Hartmann \\& McGregor 1982, Sauty et al. 1999). \\item {\\bf the accretion disc alone}: ``disc winds'' are produced from a large radial extension in the disc, thanks to the presence of a large scale magnetic field (eg. Blandford \\& Payne 1982, Pudritz \\& Norman 1983). They are fed with both matter and energy provided by the accretion process alone. \\item {\\bf the interaction zone between the disc and the protostar}: these ``X-winds'' are produced in a tiny region around the magnetopause between the disc and the protostar (eg. Shu et al. 1994, Lovelace et al. 1999, Ferreira et al. 2000). \\end{itemize}  Purely stellar wind models are less favoured because observed jets carry far too much momentum. In order to reproduce a YSO jet, a protostar should be either much more luminous or rotating faster than observations show (DeCampli 1981, K\\\"onigl 1986). This leaves us with either disc-winds or X-winds\\footnote{Camenzind and collaborators proposed an ``enhanced'' version of stellar winds, related to an old idea of Uchida \\& Shibata (1984). In this picture, a magnetospheric interaction with the accretion disc is supposed to strongly modify both jet energetics and magnetic configuration, leading to enhanced ejection from the protostar (Camenzind 1990, Fendt et al. 1995, see Breitmoser \\& Camenzind 2000 and references therein).}. From the observational point of view, it is very difficult to discriminate between these two models. In a nice review, Livio (1997) gathered a number of arguments for the disc wind model. The main idea is to look for a model able to explain jets from quite a lot of different astrophysical contexts (YSOs, AGN, X-ray binaries). The only ``universal'' ingredient required is an accretion disc threaded by a large scale magnetic field. Such a paradigm naturally explains (qualitatively) all accretion-ejection correlations known and is consistent with every context: young stars (see Cabrit's contribution, this volume), microquasars (eg. Mirabel \\& Rodriguez 1999) and AGN (see eg. Serjeant et al. 1998, Cao \\& Jiang 1999, Jones et al. 2000).  \\subsection{Where does this magnetic field come from~?}  Let's face it: we don't know. There are two extreme possibilities. The first one considers that the field has been advected by the infalling material, leading to a flux concentration in the inner disc regions. The second one relies on a local dynamo action in the disc. Most probably, the answer lies between these two extreme cases.  If the interstellar magnetic field has been indeed advected, the crucial issue is the amount of field diffusion during the infall. Indeed, if we take the fiducial values $n \\sim 1\\mbox{ cm}^{-3}$ and $B\\sim 4\\ \\mu$G of dense clouds and use the law $B\\propto n^{1/2}$ (Crutcher 1999), we get a magnetic field at 1 UA ranging from 10 to $10^3$ G! One must then consider in a self-consistent way the dynamical influence of the magnetic field, along with matter energy equation and ionization state (see Lachieze's contribution). This is something extremely difficult and, as far as I know, no definite result has been yet obtained (see however Ciolek \\& Mouschovias 1995).  How exactly disc dynamo works is also quite unclear. Dynamo theory remains intricate, relying on the properties of the turbulence triggered inside the disc. In our current picture of accretion discs, these are highly turbulent because of some instability, probably of magnetic origin (see Terquem's contribution). Such a turbulence is believed to provide means of efficient transport inside the disc, namely anomalous viscosity, magnetic diffusivity and heat conductibility. Obviously, small scale (not larger than the disc vertical scale height), time dependent magnetic fields will then exist inside accretion discs. But we are interested in the mean flow (hence mean field) dynamics. So, in practice what does ejection require~? To produce two opposedly directed jets, there must be a large scale magnetic field in the disc, which is open and of one of the following topologies:  \\indent {\\bf Dipolar}: the field threads the disc, with only a vertical component at the disc midplane, matter being forced to cross the field lines while accreting (eg. Blandford \\& Payne 1982).  \\indent {\\bf Quadrupolar}: field lines are nearly parallel to matter, entering the disc in its plane and leaving it at its surfaces, with only a non zero azimuthal component at the disc midplane (eg. Lovelace et al. 1987).  Most jet models and numerical simulations assume a dipolar magnetic configuration, with no justification. In fact, it turns out from the analysis of disc physics that only the dipolar configuration is suitable for launching jets from Keplerian accretion discs (see appendix A in Ferreira 1997). This has been recently confirmed using $\\alpha$ dynamo-generated magnetic fields (Rekowski et al. 2000). This is due to a change of sign of the $\\alpha$ effect across the disc midplane, as observed in numerical simulations of MHD turbulence in the shearing box approximation (Brandenburg \\& Donner 1997). But a realistic situation requires to treat both turbulence and the backreaction of the magnetic field in a self-consistent way. Indeed, jet production is a means of flux leakage, hence of possible dynamo self-regulation (Yoshizawa \\& Yokoi 1993). Anyway, this severe issue of dynamo and turbulence lead theorists to simply assume the existence of a large scale magnetic field. Its' value and distribution are then either imposed or obtained as conditions for stationarity in Magnetized Accretion-Ejection Structures (MAES).  \\subsection{The MAES paradigm}  A large scale (mean) magnetic field of bipolar topology is assumed to thread an accretion disc, allowing ejected plasma to flow along open field lines. This field extracts both angular momentum and energy from the underlying disc and transfers them back to the ejected plasma. There have been numerous studies of magnetized jets (e.g. Blandford \\& Payne 1982, Heyvaerts \\& Norman 1989, Pelletier \\& Pudritz 1992, Contopoulos \\& Lovelace 1994, Rosso \\& Pelletier 1994, Lery et al. 1999 to cite only a few), but they all suppose that the underlying disc would support the jets. In all these works, the disc itself was treated as a boundary condition, people usually assuming a standard viscous accretion disc (Shakura \\& Sunyaev 1973). However, if jets are to carry away the disc angular momentum, they strongly influence the disc dynamics.  The investigation of MAES, where accretion and ejection are interdependent, requires a new theory of accretion discs. The relevant questions that must be addressed by any realistic model of stationary magnetized accretion-ejection structures are the following:  \\hspace{1cm} {\\bf (1)} What are the relevant physical mechanisms inside the disc~?  \\hspace{1cm} {\\bf (2)} What are the physical conditions allowing accretion and ejection~?  \\hspace{1cm} {\\bf (3)} Can we relate jet properties to those of the disc~?  To answer these questions, one must take into account the full 2D problem (not 3D, thanks to axisymmetry) and {\\bf not} treat the disc as infinitely thin as in a standard disc theory. As a consequence, no toy-model has been able yet to catch the main features of these accretion-ejection structures. There, the disc accretion rate exhibits a radial variation as matter is being ejected. The link between accretion and ejection in a MAES can therefore be measured by the quantity $ \\xi = \\frac{d \\ln \\dot M_a}{d \\ln r} $ called the ``ejection index''. This parameter (which can vary in the disc) measures a local ejection efficiency. In a standard accretion disc, $\\xi=0$ everywhere leading to a constant accretion rate. A complete theory of MAES must provide the allowed values of $\\xi$ as a function of the disc properties. ",
        "conclusions": ""
    }
}