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0710 | 0710.4549_arXiv.txt | We investigate the chemical abundances of NGC\,3603 in the Milky Way, of 30\,Doradus in the Large Magellanic Cloud, and of N\,66 in the Small Magellanic Cloud. Mid-infrared observations with the Infrared Spectrograph onboard the Spitzer Space Telescope allow us to probe the properties of distinct physical regions within each object: the central ionizing cluster, the surrounding ionized gas, photodissociation regions, and buried stellar clusters. We detect [S\3], [S\4], [Ar\3], [Ne\2], [Ne\3], [Fe\2], and [Fe\3] lines and derive the ionic abundances. Based on the ionic abundance ratio (Ne\3/H)/(S\3/H), we find that the gas observed in the MIR is characterized by a higher degree of ionization than the gas observed in the optical spectra. We compute the elemental abundances of Ne, S, Ar, and Fe. We find that the $\alpha$-elements Ne, S, and Ar scale with each other. Our determinations agree well with the abundances derived from the optical. The Ne/S ratio is higher than the solar value in the three giant H\2\ regions and points toward a moderate depletion of sulfur on dust grains. We find that the neon and sulfur abundances display a remarkably small dispersion (0.11\,dex in 15 positions in 30\,Doradus), suggesting a relatively homogeneous ISM, even though small-scale mixing cannot be ruled out. | \label{sec:intro} Giant H\2\ regions are ideal laboratories to understand the feedback of star-formation on the dynamics and energetics of the interstellar medium (ISM). Supernov{\ae} and stellar winds arising in such regions are reponsible for producing shocks, destroying dust grains and molecules, while compressing molecular clouds and triggering subsequent star-formation. They also allow the release of newly synthetized elements into the ISM, altering its metallicity. In order to study the star-formation properties as a function of the environment, we observed three giant H\2\ regions spanning a wide range of physical conditions (gas density, mass, age) and chemical properties (metallicity) with the Spitzer Space Telescope (Werner et al.\ 2004). Observations are part of the GTO program PID\#63. The regions are NGC\,3603 in the Milky Way, 30\,Doradus (hereafter 30\,Dor) in the Large Magellanic Cloud (LMC), and N\,66 in the Small Magellanic Cloud (SMC). The scope of this program is to address crucial issues such as the destruction of complex molecules by energetic photons arising from massive stars, the polycyclic aromatic hydrocarbon (PAH) abundance dependence on metallicity, or conditions that lead to the formation/disruption of massive stellar clusters. Photometry with Spitzer/IRAC (Fazio et al.\ 2004) has been performed and will be discussed in Brandl et al.\ (in preparation). The brightest mid-infrared (MIR) regions (knots, stellar clusters, shockfronts, ...) were followed spectroscopically with the Infrared Spectrograph (IRS; Houck et al.\ 2004). In Lebouteiller et al.\ (2007), we analyzed the spatial variations of the PAH and fine-structure line emission across individual photodissociation regions (PDRs) in NGC\,3603. The two other regions will be investigated the same way in follow-up papers (Bernard-Salas et al.\ in preparation; Whelan et al.\ in preparation). In this paper, we introduce the full IRS dataset (low- and high-resolution) of the giant H\2\ regions and we derive their chemical abundances. A subsequent paper will be focused on the study of molecules and dust properties (Lebouteiller et al.\ in preparation). Elemental abundances in H\2\ regions are historically derived from optical emission-lines. Large optical telescopes, together with sensitive detectors makes it possible to determine the chemical composition of very faint H\2\ regions. Because of dust extinction, optical spectra only observe ionized gas toward sighlines with low dust content. In this view, the MIR range allows analyzing denser lines of sight, with possibly different chemical properties because of small-scale mixing and/or differential depletion on dust grains. MIR emission-lines constitute the only way to measure abundances in more obscure regions, and these abundances ought to be compared to abundances from the optical range. Although the optical domain gives access to some of the most important elements to constrain nucleosynthethic and stellar yields (C, N, O, Ne, S, Ar, Fe), it does not include some essential ionization stages necessary for abundance determinations of certain elements, such as S\4\ or Ne\2. The MIR range enables the abundance determination of Ne, S, and Ar, with the most important ionization stages observed. Iron abundance can be also determined from MIR forbidden emission-lines, but with considerably larger uncertainty due to ionization corrections. Finally, it must be stressed that abundance determinations in the optical are more sensitive to the electronic temperature ($T_e$) determination as compared to the MIR range. The effect of $T_e$ on abundances determinations is a significant source of error in optical abundance results. Wu et al.\ (2007) recently studied a sample of blue compact dwarf galaxies (BCDs) with the IRS and found a global agreement between abundances derived from the optical and those derived from the MIR. This suggests that the dense lines of sight probed in the MIR have a similar chemical composition as unextincted lines of sight and/or dense regions with possibly peculiar abundances do not contribute significantly to the integrated MIR emission-line spectrum. MIR abundances of the BCDs were calculated using mostly H$\beta$ or H$\alpha$ lines from the optical as tracer of the hydrogen content, with significant uncertainties from aperture corrections, or different observed regions because of extinction. The present sample of giant H\2\ regions provides the unique opportunity to measure accurate abundances, with a signal-to-noise ratio sufficiently high to observe directly the H\1\ recombination line at 12.37\mic. We provide abundances of Ne, S, Ar, and Fe toward lines of sight with different physical properties (PDRs, ionized gas, embedded source, stellar cluster, ...) within each giant H\2\ region. We first present the sample of the three giant H\2\ regions in \S\ref{sec:pres}. The data reduction and analysis are discussed in \S\ref{sec:observations}. We infer the ion abundances in \S\ref{sec:ionicab}. Elemental abundances are determined in \S\ref{sec:eleab} and are discussed in \S\ref{sec:discussion}. | We analyzed the chemical abundances in the ISM of three giant H\2\ regions, NGC\,3603 in the Milky Way, 30\,Dor in the LMC, and N\,66 in the SMC using the MIR lines observed with the IRS onboard Spitzer. \begin{itemize} \item Our observations probe the ISM toward various physical regions, such as stellar clusters, ionized gas, photodissociation regions, and deeply embedded MIR bright sources. The spectra show the main ionization stages of neon, sulfur, and argon in the ionized gas. We also detect [Fe\2] and [Fe\3] lines. \item Ionic abundances of Ne\2, Ne\3, S\3, S\4, Ar\2, Ar\3, Fe\2, and Fe\3\ were derived. The internal variation of electron density across a region has no impact on the ionic abundance determination. On the other hand, we find that electron temperature uncertainties and/or intrinsic variations could be responsible for an error of 20\% at most on the abundance determinations. Based on the (Ne\3/H)/(S\3/H) ionic abundance ratio, we find that the optical spectra probe a gas with a degree of ionization equal to or higher than the gas probed in the MIR. \item Elemental abundances were determined from the ionic abundances. No ionization corrections were needed, except for iron. We find that neon, sulfur, and argon scale with each other, which is expected from stellar yields. Abundances do not show any dependence on the physical region (PDR, stellar cluster, embedded region, ...). \item The Ne/S ratio is larger than the solar value, and suggests that sulfur could be depleted onto dust grains. The sulfur abundance in the MIR agrees best with the lowest optical determinations, which is likely due to uncertainties in the abundance determinations. \item Iron abundance shows a larger uncertainty than Ne/H, S/H, and Ar/H. The comparison of iron and neon abundances hints at significant depletion of iron onto dust grains at large metallicities. The agreement with the optical determination of Fe/H indicates however that there is no differential depletion on dust grains between the gas probed in the MIR and in the optical. \item Fe/H is found to be spectacularly large in one position, corresponding to a supernova remnant. This strongly suggest that iron atoms have been released from dust grains due to schocks from the SN. \item The metallicity of NGC\,3603 agrees with the Galactic abundance gradient. The metallicities of 30\,Doradus and N\,66 agree well with those of the PNe in their respective host galaxies. These findings suggest that the giant H\2\ regions did not experience a significant metal enrichment for at least 1\,Gyr. If enrichement occured, the metallicity was altered by less than a factor of two. \item Neon and sulfur abundances show remarkably little dispersion in the three H\2\ regions (e.g., 0.11\,dex dispersion in 15 positions in 30\,Dor). Small-scale mixing is apparently effective, abundance fluctuations are smaller than $\sim$55\%. However, internal variations of the abundances are likely to be on the order of $\lesssim$5\%, and determining their existence would require a significant improvement of the data quality and of the method to be evidenced. \end{itemize} | 7 | 10 | 0710.4549 |
0710 | 0710.1193_arXiv.txt | \small Small perturbations in spherical and thin disk stellar clusters surrounding massive a black hole are studied. Due to the black hole, stars with sufficiently low angular momentum escape from the system through the loss cone. We show that stability properties of spherical clusters crucially depend on whether the distribution of stars is monotonic or non-monotonic in angular momentum. It turns out that only non-monotonic distributions can be unstable. At the same time the instability in disk clusters is possible for both types of distributions. | The study of the gravitational loss-cone instability, a far analog of the plasma cone instability, has begun with the work of V. Polyachenko (1991), in which a simplest analytical model of thin disk stellar cluster has been treated. The interest to the problem of stability of stellar clusters has been revived recently by detailed investigation by Tremaine (2005) and Polyachenko, Polyachenko, Shukhman, (2007; henceforth, Paper I) of low mass clusters around massive black holes. The both papers have considered stability of small amplitude perturbations of stellar clusters of disk-like and spherical geometry. Tremaine (2005) has shown using Goodman's (1988) criterion that thin disks with symmetric DFs over angular momentum and empty loss cone are generally unstable. By contrast, analyzing perturbations with spherical numbers $l=1$ and $l=2$, he deduced that spherical clusters with monotonically increasing DF of angular momentum should be generally stable. Later we demonstrated (see Paper I) that spherical systems with non-monotonic distributions may be unstable for sufficiently small-scale perturbations $l \ge 3 $, while the harmonics $l=1,2$ are always stable. For the sake of convenience, we have used two assumptions. The first one is that the Keplerian potential of the massive black hole dominates over a self-gravitating potential of the stellar cluster (which does not mean that one can neglect the latter). Then the characteristic time of system evolution is of the order of the orbit precessing time, which is slow, compared to typical dynamical (free fall) time. Since a star makes many revolutions in its almost unaltered orbit, we can regard it as to be ``smeared out'' along the orbit in accordance with passing time, and study evolution of systems made of these extended objects. The second assumption is a so called {\it spoke approximation}, in which a system consists of near-radial orbits only. This approximation was earlier suggested by one of the authors (Polyachenko 1989, 1991). The spoke approximation reduces the problem to a study of rather simple analytical characteristic equations controlling small perturbations of stellar clusters. There are two questions that naturally arise in this context. First: Does the instability remain when abandoning the assumption of strong radial elongation of orbits? Second: Does the instability occur in spheres with monotonically increasing distributions in angular momentum if one consider smaller-scale perturbations with $l \ge 3$? The aim of the paper is to provide answers to these questions. To achieve the task we use semi-analytical approach based on analysis of integral equations for slow modes elaborated recently in Polyachenko (2004, 2005) for thin disks, and in Paper I for spherical geometry. Following Paper I, we shall restrict ourselves to studying monoenergetic models with DFs in the form \begin{align}\label{eq:1.1} F(E,L)=A\,\delta(E-E_0)\,f(L). \end{align} The models specified by function $f(L)$ are suitable for studying the effects of angular momentum distribution on gravitational loss-cone instability. On the other hand, the Dirac $\delta$-function permits one to reduce the integral equations for slow modes to one-dimensional integral equations, and to advance substantially in analytical calculations. Several arguments can be brought in favour of our simplified approach. First of all, the Lynden-Bell derivative (see Paper I, eq. 4.7) of the DF with respect to angular momentum $L$, keeping $J = L + I_1$ constant (here $I_1$ is the radial action) in the limit where the slow mode approximation is applicable, can be replaced by a derivative, keeping energy $E$ constant: $$ \left(\frac{\p F}{\p L} \right)_{LB} = \Omega_\textrm{pr} \left(\frac{\p F}{\p E} \right)_L + \left(\frac{\p F}{\p L} \right)_E \approx \left(\frac{\p F}{\p L} \right)_{E}, $$ because $\Omega_\textrm{pr}$ is small. Thus, the derivative over energy is not included into the slow integral equation, and one can loosely say, that dependence on energy is only parametric. Another argument is that the results of independent study by Tremaine (2005), who used a non-monoenergetic DF, are in agreement with our conclusions. Section 2 is devoted to spheres, Section 3 -- to thin disks with symmetric DFs. The sections are organized alike. In the beginning we derive integral equations for initial distribution functions in the form (\ref{eq:1.1}). Then follow analytical and numerical investigations of these equations. We demonstrate that by contrast to the case of near-Keplerian sphere, the loss-cone instability in disks takes place even for the monotonic DF, $df/d|L|>0$, provided the precession is retrograde and the loss cone is empty: $f(0)=0$. Sec. 2 is complimented by stability analysis of models with circular orbits, which of course doesn't belong to the class of monoenergetic models of (\ref{eq:1.1}) type. In the last, Section 4, we discuss the results and some perspectives of further studies. | We have studied the stability of the spherically-symmetric and thin disk stellar clusters around a massive black hole. We conclude that stability properties of spherical clusters depend crucially on monotonity of initial distribution functions, while thin disk clusters are almost always unstable. If the initial distribution of the spherical cluster is monotonic, the cluster is most likely to be stable. This conclusion was first made in Tremaine (2005), where stability of $l=1$ mode was generally proved, and $l=2$ was tested numerically. We confirm this conclusion by considering a number of monotonic distributions for modes with arbitrary $l$. Besides, we have checked distributions obtained from monotonic ones by making them vanish quickly but smoothly at circular orbits. These models were also stable. However, a general proof of stability for any monotonic distributions was not yet found. Spherical clusters with the non-monotonic DFs should be generally affected by the gravitational loss-cone instability. The instability was first found in our Paper I using a simplification of systems with near-radial orbits. In the Sec. 2 we show that this instability is due to just non-monotony of distributions over angular momentum, the orbits may not necessary be near-radial. In our opinion, both monotonic and non-monotonic distributions are important for possible applications to real stellar clusters around black holes. The DFs monotonically increasing from the loss cone radius up to circular orbits are formed naturally due to two-body collisions of stars. It follows from numerical experiments (see, e.g., Cohn and Kulsrud, 1978), which predict establishment of such distributions after a characteristic time for collisional relaxation. These distributions may be approximated by the formula $F\propto \ln \bigl(L/L_{\rm min}\bigr)$. Such a slowly increasing function is, in fact, predetermined by the boundary conditions imposed in the cited numerical study and some other investigations. Indeed, the vanishing condition at $L=L_{\rm min}$, and the matching condition to isotropic (Maxwellian) distribution, $F=F(E)$, at the boundary $E=E_{\rm bound}=0$ of the phase space $(E,L)$ (boundary separates stars which is gravitationally coupled to the black hole from the others) is required. The last condition means the asymptotic (when $E \to E_{\rm bound}$) independence of the function $F(E,L)$ on the momentum $L$. So monotonic, or logarithmic, dependence of type of (\ref{eq:3.2}) is quite reasonable. The non-monotonic distributions are also real. If the cluster, is formed, for example, as a result of the collisionless collapse (several free fall times), then it remains collisionless for a long timescale of collisional relaxation (see, e.g., Merritt \& Wang, 2005). In principle, the system can have almost arbitrary DF both in the energy and in the angular momentum. During the collapse, a typical non-monotonic distribution of stars over the angular momentum, with empty loss cone and maximum at some value $L=L_{\ast}$, is formed. In Paper I we argued that stability properties of such a distribution is effectively analogous to one of typical plasma distributions of the ``beam-like'' type. But they can readily become unstable, as it is well-known in plasma physics (and also confirmed by direct stability study of corresponding stellar systems in Paper I). It is possible (as it is often so in plasma) that for the time of collisionless behavior, DF can undergo a dramatic change from its initial form. In particular, the collective flux of stars into the loss cone caused by the instability could, in principle, lead to the formation of a considerable part of the black hole. Checking of such possibilities is the most urgent task for future studies of unstable {\it non-monotonic} models. Since spherically-symmetric models with the {\it monotonic} DF are apparently stable, but analogous disk systems are unstable (see Tremaine 2005 and Sec. 3), a critical flatness of ellipsoid models at which the instability begins is expected. Study of such systems, as well as systems with more complex triaxial ellipsoids can be performed using numerical simulations. | 7 | 10 | 0710.1193 |
0710 | 0710.1700_arXiv.txt | The general world model for homogeneous and isotropic universe has been proposed. For this purpose, we introduce a global and fiducial system of reference (world reference frame) constructed on a $5$-dimensional space-time that is embedding the universe, and define the line element as the separation between two neighboring events that are distinct in space and time, as viewed in the world reference frame. The effect of cosmic expansion on the measurement of physical distance has been correctly included in the new metric, which differs from the Friedmann-Robertson-Walker metric where the spatial separation is measured for events on the hypersurface at a constant time while the temporal separation is measured for events at different time epochs. The Einstein's field equations with the new metric imply that closed, flat, and open universes are filled with positive, zero, and negative energy, respectively. The curvature of the universe is determined by the sign of mean energy density. We have demonstrated that the flat universe is empty and stationary, equivalent to the Minkowski space-time, and that the universe with positive energy density is always spatially closed and finite. In the closed universe, the proper time of a comoving observer does not elapse uniformly as judged in the world reference frame, in which both cosmic expansion and time-varying light speeds cannot exceed the limiting speed of the special relativity. We have also reconstructed cosmic evolution histories of the closed world models that are consistent with recent astronomical observations, and derived useful formulas such as energy-momentum relation of particles, redshift, total energy in the universe, cosmic distance and time scales, and so forth. It has also been shown that the inflation with positive acceleration at the earliest epoch is improbable. | \label{sec:intro} The main goal of modern cosmology is to build a cosmological model that is consistent with astronomical observations. To achieve this goal, tremendous efforts have been made both on theories and on observations since the general theory of relativity was developed. So far the most successful model of the universe is the Friedmann-Robertson-Walker (FRW) world model \cite{fri22,fri24,rob29,wal35}. The FRW world model predicts reasonably well the current observations of the cosmic microwave background (CMB) radiation and the large-scale structures in the universe. The precisely determined cosmological parameters of the FRW world model imply that our universe is consistent with the spatially flat world model dominated by dark energy and cold dark matter ($\Lambda\textrm{CDM}$) with adiabatic initial condition driven by inflation \cite{spergel07,tegmark06}. Although the flat FRW world model is currently the most reliable physical world model, one may have the following fundamental questions on the nature of the FRW world model. First, mathematically, if a space-time manifold is flat, then the Riemann curvature tensor should vanish, and vice versa. However, the Riemann curvature tensor of the flat FRW world model does not vanish unless the cosmic expansion speed and acceleration are zeros, which implies that the physical space-time of the flat FRW world is not geometrically flat but curved. Only its spatial section at a constant time is flat. Secondly, the cosmic evolution equations of the FRW world model can be derived from an application of the Newton's gravitation and the local energy conservation laws to the dynamical motion of an expanding sphere with finite mass density \cite{milne34,mccrea34}. Besides, the Newton's gravitation theory has been widely used to mimic the non-linear clustering of large-scale structures in the universe even on the horizon-sized $N$-body simulations \cite{colberg2000,park05}. On large scales, the close connection between the FRW world model and the Newton's gravitation law is usually attributed to the fact that the linear evolution of large-scale density perturbations satisfies the weak gravitational field condition. Recently, Hwang and Noh \cite{hwang06} show that the relativistic fluid equations perturbed to second order in a flat FRW background world coincide exactly with the Newtonian results, and prove that the Newtonian numerical simulation is valid in all cosmological scales up to the second order. However, one may have a different point of view that the Newton's gravitational action at a distance appears to be valid even on the super-horizon scales in the FRW world just because the world model does not reflect the full nature of the relativistic theory of gravitation. Thirdly, according to the FRW world model, the universe at sufficiently early epoch ($z \gtrsim 1000$) is usually regarded as flat since the curvature parameter contributes negligibly to the total density. The present non-flat universe should have had the density parameter approaching to $\Omega = 1$ with infinitely high precision just after the big-bang (flatness problem). On the other hand, if we imagine the surface of an expanding balloon with positive curvature, then the curvature of the surface is always positive and becomes even higher as the balloon is traced back to the earlier epoch when it was smaller. This prediction from the common sense contradicts the FRW world model. Observationally, the flat $\Lambda\textrm{CDM}$ universe is favored by the recent joint cosmological parameter estimation using the Wilkinson Microwave Anisotropy (WMAP) CMB \cite{hinshaw07,page07}, large-scale structures \cite{cole05,tegmark04}, type Ia supernovae (SNIa; \cite{riess07,wood07}), Hubble constant \cite{freedman01,macri06,sandage06}, baryonic oscillation data \cite{eisen05}, and so on. However, the WMAP CMB data alone is more compatible with the non-flat FRW world model ($\S7.3$ of \cite{spergel07} and Table III of \cite{tegmark06}). Besides, some parameter estimations using SNIa data or angular size-redshift data of distant radio sources alone suggest a possibility of the closed universe \cite{clocchi06,jackson06}. The combinations of the WMAP plus the SNIa data or the Hubble constant data also imply the possibility of the closed universe, giving curvature parameters $\Omega_k = -0.011\pm 0.012$ and $\Omega_k = -0.014\pm 0.017$, respectively \cite{spergel07}, although the estimated values are still consistent with the flat FRW world model. The questions above and the observational constraints on the cosmological model may bring about possibilities of non-flat or non-FRW world models. Interestingly, Einstein claimed that our universe is spatially bounded or closed \cite{ein22}. The primary reason for his preference to the closed universe is because Mach's idea \cite{mach93,misner73} that the inertia depends upon the mutual action of bodies is compatible only with the finite universe, not with a quasi-Euclidean, infinite universe. According to Einstein's argument, an infinite universe is possible only if the mean density of matter in the universe vanishes, which is unlikely due to the fact that there is a positive mean density of matter in the universe \footnote{However, in the appendix to the second edition of his book \cite{ein22}, Einstein summarized Friedmann's world models and discussed a universe with vanishing spatial curvature and non-vanishing mean matter density, which differs from his original argument.}. In this paper, we propose the general world model for homogeneous and isotropic universe which supports Einstein's perspective on the physical universe. The outline of this paper is as follows. In Sec. \ref{sec:metric}, we consider the effect of cosmic expansion on the physical space-time distance between neighboring events and describe how to define the line element for homogeneous and isotropic universes of various spatial curvature types. The metric and the cosmic evolution equations for flat, closed, and open universes are derived in Sec. \ref{sec:nspace}. It will be shown that our universe is spatially closed. In Sec. \ref{sec:some}, we reconstruct cosmic evolution histories of the closed world models, and derive interesting properties of the closed universe. In Sec. \ref{sec:inflation}, we discuss whether the inflation theory is compatible with the closed world model or not. Conclusion follows in Sec. \ref{sec:conc}. Throughout this paper, we adopt a sign convention $(+,-,-,-)$ for the metric tensor $g_{ik}$, and denote a 4-vector in space-time as $p^{i}$ ($i=0,1,2,3$) and a 3-vector in space as $p^{\alpha}$ ($\alpha=1,2,3$) or $\mathbf{p}$. The Einstein's field equations are \begin{equation} R_{ik} - \frac{1}{2} g_{ik}R = 8\pi G T_{ik}+\Lambda g_{ik}, \label{eq:einstein} \end{equation} where $R_{ik}=R^{a}_{~iak}$ is the Ricci tensor, $R=R^i_{~i}$ the Ricci scalar, $T_{ik}$ the energy-momentum tensor, $G$ the Newton's gravitational constant, and $\Lambda$ the cosmological constant. The Riemann curvature tensor is given by $R^{a}_{~ibk} =\partial_b \Gamma^{a}_{ki}-\partial_{k}\Gamma^{a}_{bi} +\Gamma^{a}_{bn} \Gamma^{n}_{ki}-\Gamma^{a}_{kn}\Gamma^{n}_{bi}$, with the Christoffel symbol $\Gamma^{a}_{ik}={1\over 2} g^{ab} (\partial_i g_{kb}+\partial_k g_{ib} -\partial_b g_{ik})$. The energy-momentum tensor for perfect fluid is \begin{equation} T_{ik} = (\varepsilon_\textrm{b} + P_\textrm{b})u_i u_k - P_\textrm{b} g_{ik}, \label{eq:emtensor} \end{equation} where $\varepsilon_\textrm{b}$ and $P_\textrm{b}$ are background energy density and pressure of ordinary matter and radiation, and $u_{i}$ is the 4-velocity of a fundamental observer. We assume that the cosmological constant acts like a fluid with effective energy density $\varepsilon_\Lambda = \Lambda/8\pi G$ and pressure $P_\Lambda = -\varepsilon_\Lambda$. The limiting speed in the special theory of relativity is set to unity ($c\equiv 1$). | \label{sec:conc} In this paper, the general world model for homogeneous and isotropic universe has been proposed. By introducing the world reference frame as a global and fiducial system of reference, we have defined the line element so that the effect of cosmic expansion on the physical space-time separation can be correctly included in the metric. With this framework, we have demonstrated theoretically that the flat universe is equivalent to the Minkowski space-time and that the universe with positive energy density is always spatially closed and finite. The open universe is unrealistic because it cannot accommodate positive energy density. Therefore, in the world of ordinary materials, only the spatially closed universe is possible to exist. The naturalness of the finite world with positive energy density comes from the Mach's principle that the motion of a mass particle depends on the mass distribution of the entire world. The principle is consistent only with the finite world because the dynamics of a reference frame cannot be defined in the infinite, empty world. The closed world model satisfies the Mach's principle and supports Einstein's perspective on the physical universe. We have reconstructed evolution histories of the closed world models that are consistent with the recent astronomical observations, based on the nearly flat FRW world models (Model I and II; Sec. \ref{sec:some}). The present curvature radius of the universe is $a_0 = 25.7$ Gpc ($a_0=40.2$ Gpc) for Model I (Model II). The expansion histories of both models imply that the closed universe dominated by dark energy expands eternally. However, the currently favored flat FRW world exists only as a limiting case of the closed universe with infinite curvature radius that is expanding with the maximum speed ($\dot{a}_0=1$, $\ddot{a}_0=0$). From the local nature of the FRW metric (Sec. \ref{sec:metric}) and of the proper time of a comoving observer (Sec. \ref{sec:rel_friedmann}), it is clear that the FRW world model describes the local universe as observed by the comoving observer. Since the Newton's gravitation law can be derived from the Einstein's field equations in the weak field and the small velocity limits, the gravitational action at a distance usually holds at a local region of space on scales far smaller than the Hubble horizon size (e.g., \cite{peebles80}). The proper Hubble radius $d_\textrm{H}$ (Fig. \ref{fig:horizon}, top) may provide a reasonable estimate of the characteristic distance scale where the Newton's gravity applies. The cosmic structures simulated by the Newton's gravity-based $N$-body method will significantly deviate from the real structures on scales comparable to $d_\textrm{H}$. The variation of the comoving Hubble radius $\chi_\textrm{H} = d_\textrm{H} /a$ also implies that in the past (future) the Newtonian dynamics was (will be) applicable on smaller region of space compared to the size of the universe (Fig. \ref{fig:horizon}, bottom). In this paper, the history of the universe has been tentatively reconstructed based on cosmological parameters of non-flat FRW world models. The more general cosmological perturbation theory and parameter estimation are essential for accurate reconstruction of the cosmic history. | 7 | 10 | 0710.1700 |
0710 | 0710.4914_arXiv.txt | We present 2D simulations of the cooling of neutron stars with strong magnetic fields ($B \geq 10^{13}$ G). We solve the diffusion equation in axial symmetry including the state of the art microphysics that controls the cooling such as slow/fast neutrino processes, superfluidity, as well as possible heating mechanisms. We study how the cooling curves depend on the the magnetic field strength and geometry. Special attention is given to discuss the influence of magnetic field decay. We show that Joule heating effects are very large and in some cases control the thermal evolution. We characterize the temperature anisotropy induced by the magnetic field for the early and late stages of the evolution of isolated neutron stars. | The observed thermal emission of neutron stars (NSs) can provide information about the matter in their interior. Comparing the theoretical cooling curves with observational data\cite{Yakovlev2004,Page2006} one can infer not only the physical conditions of the outer region (atmosphere) where the spectrum is formed but also of the poorly known interior (crust, core) where high densities are expected. There is increasing evidence that most of nearby NSs whose thermal emission is visible in the X-ray band have a non uniform temperature distribution\cite{Zavlin2007,Haberl2007}~. There is a mismatch between the extrapolation to low energy of the fits to X-ray spectra, and the observed Rayleigh Jeans tail in the optical band ({\it optical excess flux}), that cannot be addressed with a unique temperature (e.g. \rxdieciocho\cite{Pons2002}~, \rbdoce\cite{Schwope2007}~, and \rxcerosiete\cite{Perez2006}~). A non uniform temperature distribution may be produced not only in the low density regions\cite{Greenstein1983}~, but also in intermediate density regions, such as the solid crust. Recently, it has been proposed that crustal confined magnetic fields with strengths larger than $10^{13}$ G could be responsible for the surface thermal anisotropy \cite{Geppert2004,Azorin2006}~. In the crust, the magnetic field limits the movement of electrons (main responsible for the heat transport) in the direction perpendicular to the field and the thermal conductivity in this direction is highly suppressed, while remains almost unaffected along the field lines. Moreover, the observational fact that most thermally emitting isolated NSs have magnetic fields larger than $10^{13}$ G implies that a realistic cooling model must include magnetic field effects. In a recent work\cite{Aguilera2007}~, first 2D simulations of the cooling of magnetized NSs have been presented. In particular, it has been stated that magnetic field decay, as a heat source, could strongly affect the thermal evolution and the observations should be reinterpreted in the light of these new results. We present the main conclusions of this work next. | The main result of this work is that, in magnetized NSs with $B> 10^{13}$~G, the decay of the magnetic field affects strongly their cooling. In particular, there is a huge effect of Joule heating on the thermal evolution. In NSs born as magnetars, this effect plays a key role in maintaining them warm for a long time. Moreover, it can also be important in high magnetic field radio pulsars and in radio--quiet isolated NSs. As a conclusion, the thermal and magnetic field evolution of a NS is at least a two parameter space (Fig~\ref{fig_coupled}), and a first step towards a coupled magneto-thermal evolution has been given in this work. \begin{figure}[htb] \begin{center} \psfig{file=BT_v1_1.ps,width=0.5\textwidth,angle=-90} \end{center} \caption{Coupled magneto-thermal evolution of isolated neutron stars\cite{Aguilera2007a}~: $T_s$ for the hot component as a function of $B$ and $t$. Observations: squares for magnetars ($B>10^{14}$~G), triangles for intermediate-field isolated NSs ($10^{13}$~G$<B<10^{14}$~G) and circles for radio pulsars ($B<10^{12}$~G). Corresponding cooling curves in solid, dashed and dotted lines, respectively. } \label{fig_coupled} \end{figure} | 7 | 10 | 0710.4914 |
0710 | 0710.5123_arXiv.txt | We investigate the effect of spiral structure on the Galactic disk as viewed by pencil beams centered on the Sun, relevant to upcoming surveys such as ARGOS, SEGUE, and GAIA. We create synthetic Galactic maps which we call Pencil Beam Maps (PBMs) of the following observables: line-of-sight velocities, the corresponding velocity dispersion, and the stellar number density that are functions of distance from the observer. We show that such maps can be used to infer spiral structure parameters, such as pattern speed, solar phase angle, and number of arms. The mean line-of-sight velocity and velocity dispersion are affected by up to $\sim35$ km/s which is well within the detectable limit for forthcoming radial velocity surveys. One can measure the pattern speed by searching for imprints of resonances. In the case of a two-armed spiral structure it can be inferred from the radius of a high velocity dispersion ring situated at the 2:1 ILR. This information, however, must be combined with information related to the velocities and stellar number density in order to distinguish from a four-armed structure. If the pattern speed is such that the 2:1 ILR is hidden inside the Galactic bulge the 2:1 OLR will be present in the outer Galaxy and thus can equivalently be used to estimate the pattern speed. Once the pattern speed is known the solar angle can be estimated from the line-of-sight velocities and the number density PBMs. Forthcoming radial velocity surveys are likely to provide powerful constraints of the structure of the Milky Way disk. | It has been well established by now that the Milky way is not axisymmetric with both a central bar and spiral structure perturbing its disk. Due to our location in the Galactic plane both spiral and bar structure is impossible to observe directly. Galactic bar parameters such as orientation and pattern speed have been inferred indirectly from both asymmetries around the Galactic center (e.g., \citealt{blitz91,weinberg92}) and its effect on the local velocity distribution of old stars, i.e., the Hercules stream \citep{dehnen99,dehnen00,fux01,mq07b}. Spiral structure parameters, however, are much more uncertain. Current spiral density wave models \citep{fux01,lepine01,desimone04,qm05} strongly disagree on the strength of the spiral structure, the number of arms, and the pattern speed. These models differ in their predictions of the induced velocity streaming at different angular positions in the Galaxy. For example, a four-armed density wave with velocity perturbations of $\sim20$ km/s will exhibit rapidly varying radial and tangential velocity components with azimuth across distances of a few kpc, and we could expect to detect $\sim20-50$ km/s variations in the mean line-of-sight stellar velocity as a function of the distance from the Sun. However, the strength of the spiral arm perturbation remains controversial. Based on the OGLE number counts, \cite{paczynski94} estimated that the Sagittarius-Carina arm has a factor of two increase in density compared to the underlying disk. This model is inconsistent with COBE studies which find a much smaller contrast ($\sim15\%$) and show that the Perseus and Scu-Cru arms are more dominant \citep{drimmel01}. HI, CO, Cepheid, and far-infrared observations suggest that the Galactic disk contains a four-armed tightly wound structure. On the other hand, \cite{drimmel01} have shown that the near-infrared observations are consistent with a dominant two-armed structure. \citet{lepine01} suggest that locally the Milky Way can be modeled by the superposition of a two- and four-armed structure moving at the same pattern speed. By studying the nearby spiral arms, \cite{naoz07} find that the Sagittarius-Carina arm is a superposition of two features, moving at different pattern speeds. The effect of a two- and four-armed structure, moving at different angular velocities, on the velocity dispersion of a galactic disk has been explored numerically by \cite{mq06}. Estimates for the pattern speed of the Milky Way spiral structure, or equivalently, the Sun's position with respect to resonances associated with spiral structure, span a large range of values. Reviewing previous work, \citet{shaviv03} finds a clustering of estimates for the pattern speed of local spiral structure near $\Omega_s \sim 20 {\rm km s^{-1} kpc}^{-1}$, though other studies suggest $\Omega_s \sim 13 {\rm km s^{-1} kpc}^{-1}$. The model by \citet{lepine01} places the Sun near the corotation resonance $\Omega_s \sim 28 {\rm km s^{-1} kpc}^{-1}$), and was fit to Cepheid kinematics. The recent gas dynamical studies \citep{martos04, bissantz03} match the properties of the gas in nearby arms with a spiral pattern speed of $\sim 20 {\rm km s^{-1} kpc}^{-1}$. \citet{martos04} propose that a two-armed stellar structure consistent with the stellar distribution inferred from COBE could cause four-arms in the gas distribution near the Sun. The gas dynamical model proposed by \citet{bissantz03} with a similar spiral pattern speed matches HI and CO kinematics. The pattern speed of a spiral density wave can be tightly constrained from the location of its resonances. For example, \cite{qm05} associated stellar streams in the solar neighborhood with the 4:1 ILR resonance of a two-armed pattern and were then able to tightly constrain the pattern speed of the driving spiral density wave to within 5\%. Independent constraints on the pattern speed come from recent surveys of nearby open clusters (e.g. \citealt{dias05}) where the older clusters are found to have drifted further from their original density wave location. These authors concluded that the Sun is located near the CR. A solar circle near the CR is also favored by \cite{lepine01} and \cite{naoz07}. In this paper we investigate how spiral structure parameters can be inferred from velocity and density maps resulting from pencil-beam and large-scale surveys of the Galaxy. At present the influence of spiral arms on the observed kinematic properties of the Galactic disk is very poorly understood. With the advent of future Galactic all-sky (GAIA, SEGUE) and pencil-beam (ARGOS, BRAVA) radial velocity surveys, large amounts of kinematic data will be collected. The types of dynamical constraints made possible with these new data sets is not currently known. We address that issue here with synthetic models for the purpose of exploring how spiral structure might be constraint from these data. | Upcoming Galactic disk surveys will reveal the age, composition and phase space distribution of stars within the various Galactic components. These stellar excavations will provide essential clues for understanding the structure, formation and evolution of our Galaxy. To facilitate the interpretation of the huge amounts of data resulting from these surveys, Galactic disk models, such as the one presented here, are needed to interpret the observations. We have investigated how the Milky Was spiral structure parameters, such as pattern speed and solar phase angle, can be estimated in a deep all-sky survey. We performed a series of test-particle simulations of a warm galactic disk approximating the disk kinematics of the Milky Way. We considered both two- and four-armed spiral structure and suggested a way to distinguish between the two using velocity and number density maps. We found that the axisymmetric potential needs to be known to $\sim10\%$, line-of-sight velocities to $\sim20$ km/s, and distance uncertainties need to be less than $\sim30\%$. The mean line-of-sight velocity and the velocity dispersion are affected by up to $\sim35$ km/s which is well within the detectable limit for forthcoming radial velocity surveys. Pattern speed can be constrained by a hot ring at the 2:1 ILR in both two- and four-armed spiral structure. To distinguish between the two, however, we also need information related to the velocities and stellar number density. If the pattern speed is such that the 2:1 ILR is hidden inside the Galactic bulge the 2:1 OLR would be present in the outer Galaxy and thus can equivalently be used to estimate the pattern speed. Once the pattern speed is known the solar angle can be estimated from the number density variation with heliocentric distance; $\phi_0$ is also reflected in the $v_d$ PBMs. Future work needs to address the issue of how to obtain the axisymmetric background potential needed to subtract from the observational data as discussed at the end of Section \ref{sec:pbm}. Also, it is important to know what type of tracer stars are needed that would allow the estimation of photometric parallaxes with errors less than $\sim30\%$, and the distribution of those stars. While here we only considered steady state spiral stricture, other theories of spiral structure, such as transient and swing-amplified spirals, need to be investigated as well. It has also been suggested that the Galaxy harbors two sets of spiral structure moving at the same \citep{lepine01} or different \citep{naoz07,mq06} pattern speeds. We expect in all those cases it will be again resonant features to relate to the pattern speed and solar angle. In the case of non-steady state spirals, however, the structure in the PBMs will vary with integration time and interpretation will become more complicated. We refer these cases to a future study. It is also known that the Milky Way is a barred galaxy. The simulations performed here do not include the influence of the bar. This is not necessarily a shortcoming since most of the features in the PBMs we use to infer spiral structure parameters are caused by resonances and, unless a resonance overlap with the central bar exists in the same location, those would not be different when a bar is included in the simulations. Future work should also look at this problem. | 7 | 10 | 0710.5123 |
0710 | 0710.2190_arXiv.txt | The long-slit spectra obtained along the minor axis, offset major axis and diagonal axis are presented for 12 E and S0 galaxies of the Coma cluster drawn from a magnitude-limited sample studied before. The rotation curves, velocity dispersion profiles and the $H_3$ and $H_4$ coefficients of the Hermite decomposition of the line of sight velocity distribution are derived. The radial profiles of the \Hb , Mg, and Fe line strength indices are measured too. In addition, the surface photometry of the central regions of a subsample of 4 galaxies recently obtained with Hubble Space Telescope is presented. The data will be used to construct dynamical models of the galaxies and study their stellar populations. | This is the fourth of a series of papers aimed at investigating the stellar populations and the kinematics of early-type galaxies in the Coma Cluster. Spanning about 4 dex in the observed radial variation of the surface density of cluster members \citep[e.g.][]{kent1982}, the Coma Cluster is the ideal place to investigate these galaxy properties as a function of the environmental density in order to test the theories for galaxy formation and evolution. The sample of 35 E and S0 galaxies of the Coma Cluster is presented in the first paper of the series \citep[hereafter Paper I]{mehlert2000} along with the photometry and long-slit spectroscopy along their major axis. From these spectra the rotation curves, velocity dispersion profiles and the $H_3$ and $H_4$ coefficients of the Hermite decomposition of the line-of-sight velocity distribution (LOSVD) were measured out to 1--3 effective radii with high signal-to-noise ratio ($S/N$). Moreover, the radial profiles of the \Hb\, Mg, and Fe line strength indices were measured too. Subsequently, the spectroscopic database was complemented with the long-slit spectra obtained along the minor axis, an offset axis parallel to the major one and one diagonal axis for 10 objects \citep[hereafter Paper II]{wegner2002}. The central values and major-axis logarithmic gradients for the line strength indices were derived by \citet[][hereafter Paper III]{mehlert2003}. This allowed the estimation of the average ages, metallicities and \aFe\ ratios in the center and at the effective radius by using stellar population models with variable element abundance ratios from \citet{dthomas2003}. There is a dichotomy among the population of S0 galaxies. Some of them are dominated by old stellar populations and are indistinguishable from E galaxies. The remaining ones host very young stellar populations; hence they must have experienced relatively recent star formation episodes. Most massive galaxies had the shortest star formation timescales and were the first to form. The absence of age gradients implies that the stellar populations at different radii formed at a common epoch. The \aFe\ enhancement is not restricted to galaxy centers but it is a global phenomenon. Finally, negative metallicity gradients were measured to be significantly flatter than what is expected from gaseous monolithic collapse models. This suggests the importance of mergers in the galaxy formation history. Here the spectroscopic database of Paper I and II is completed with the long-slit spectra obtained along the minor axis, offset major axis and one diagonal axis for another 12 objects. As done in Paper II, these galaxies were selected from the sample of Paper I as the objects with the most extended and precise major-axis kinematics and therefore best-suited for dynamical modelling, balancing between the number of E and S0 galaxies. Moreover, the surface photometry of the central regions of a subsample of 4 galaxies recently obtained with Hubble Space Telescope (HST) is presented. The data shown here and in Paper I will allow the study of the stellar population gradients for a large number of early-type galaxies \citep{dthomas2008} in order to investigate possible systematic differences between the disk and bulge components of S0 galaxies. The photometric and kinematic data of the combined dataset allowed the construction of dynamical models of the objects to study the properties of the dark matter halos of flattened and rotating E and S0 galaxies \citep{jthomas2005,jthomas2007}. In fact, the implementation of Schwarzschild's orbit superposition technique for axisymmetric potentials by \citet{jthomas2004} was used to derive the stellar mass-to-light ratios and dark matter halo parameters for a subsample of 17 galaxies. About 10--50 percent of the mass inside the effective radius is dark with a central density which is at least one order of magnitude lower than the luminous mass density. The orbital system of the stars is reasonably close to isotropy, but the distribution function shows a lot of fine structure. This study was complementary to the one presented by \citet{gerhard2001} focusing on round and non-rotating ellipticals. The HST photometry is described in Sect. \ref{sec:photometry}. The spectroscopic galaxy sample, relative observations and data reduction are described in Sect. \ref{sec:spectroscopy}. The measured stellar kinematics, and line indices are given in Sect. \ref{sec:results}. Conclusions are drawn in Sect. \ref{sec:conclusions}. | \label{sec:conclusions} New radially resolved spectroscopy of 12 E and S0 galaxies of the Coma cluster was presented. The rotation curves, velocity dispersion profiles and the $H_3$ and $H_4$ coefficients of the Hermite decomposition of the line-of-sight velocity distribution were derived along the minor axis, offset major axis and one diagonal direction. Moreover, the line strength index profiles of Mg, Fe and \Hb\ line indices were measured too. In addition, the surface photometry of the central regions of a subsample of 4 galaxies recently obtained with HST/WFPC2 was presented. The data complement the existing set (Paper I, Paper II) and have a precision and radial extent sufficient to construct flattened and rotating dynamical models of the galaxies and study their radially resolved stellar populations. Other papers address these issues (Paper III, \citealt{jthomas2005,jthomas2007}). | 7 | 10 | 0710.2190 |
0710 | 0710.2473_arXiv.txt | In order to study the process of cooling in dark-matter (DM) halos and assess how well simple models can represent it, we run a set of radiative SPH hydrodynamical simulations of isolated halos, with gas sitting initially in hydrostatic equilibrium within Navarro-Frenk-White (NFW) potential wells. Simulations include radiative cooling and a scheme to convert high density cold gas particles into collisionless stars, neglecting any astrophysical source of energy feedback. After having assessed the numerical stability of the simulations, we compare the resulting evolution of the cooled mass with the predictions of the classical cooling model of White \& Frenk and of the cooling model proposed in the {\sc morgana} code of galaxy formation. We find that the classical model predicts fractions of cooled mass which, after about two central cooling times, are about one order of magnitude smaller than those found in simulations. Although this difference decreases with time, after 8 central cooling times, when simulations are stopped, the difference still amounts to a factor of 2--3. We ascribe this difference to the lack of validity of the assumption that a mass shell takes one cooling time, as computed on the initial conditions, to cool to very low temperature. Indeed, we find from simulations that cooling SPH particles take most time in traveling, at roughly constant temperature and increasing density, from their initial position to a central cooling region, where they quickly cool down to $\sim10^4$ K. We show that in this case the total cooling time is shorter than that computed on the initial conditions, as a consequence of the stronger radiative losses associated to the higher density experienced by these particles. As a consequence the mass cooling flow is stronger than that predicted by the classical model. The {\sc morgana} model, which computes the cooling rate as an integral over the contribution of cooling shells and does not make assumptions on the time needed by shells to reach very low temperature, better agrees with the cooled mass fraction found in the simulations, especially at early times, when the density profile of the cooling gas is shallow. With the addition of the simple assumption that the increase of the radius of the cooling region is counteracted by a shrinking at the sound speed, the {\sc morgana} model is also able to reproduce for all simulations the evolution of the cooled mass fraction to within 20--50 per cent, thereby providing a substantial improvement with respect to the classical model. Finally, we provide a very simple fitting function which accurately reproduces the cooling flow for the first $\sim10$ central cooling times. | Understanding galaxy formation is one of the most important challenges of modern cosmology. The rather stringent constraints on cosmological parameters now placed by a number of independent observations \citep[e.g.,][ for a recent review]{Springel06} allows us to precisely set the initial conditions from which the formation of cosmic structures has started. As a consequence, understanding the complex astrophysical processes, related to the evolution of the baryonic component, represents now the missing link toward a successful description of galaxy formation and evolution. So far, two alternative approaches have been pursued to make quantitative predictions on the observational properties of the galaxy population and their evolution in the cosmological context. The first one is based on the so-called semi-analytical models (SAMs, hereafter; e.g., \citealt{Kauffmann93,Somerville99,Cole00,Menci05,Monaco07}). In this approach, the background cosmological model predicts the hierarchical build-up of the Dark Matter (DM) halos, where gas flows in, cools and gives rise to the formation of galaxies, while the complex interplay between gas cooling, star formation, chemical enrichment and release of energy from supernova (SN) explosions and Active Galactic Nuclei (AGN) is modeled through a set of simplified or phenomenological models, which are specified by a number of free parameters. A posteriori, the values of the relevant parameters should then be constrained by comparing SAM predictions to observational data. The rather low computational cost of this approach makes it a quite flexible tool to explore the model parameter space. The second approach is based on resorting to full hydrodynamical simulations, which include the processes of gas cooling and suitable sub-resolution recipes for star formation and feedback. The obvious advantage of this method, with respect to SAMs, is that galaxy formation can be described by following in detail the gas-dynamical processes which determine the evolution of the cosmic baryons during the shaping of the large-scale cosmic structures. However, its limitation lies in the high computational cost, which makes it difficult to explore in detail the parameter space describing the process of galaxy formation and evolution. For this reason, following galaxy formation with full hydrodynamical simulations in a cosmological environment of several tens of Mpc is a challenging task for simulations of the present generation \citep[e.g., ][]{Nagamine04,Nagai05,Romeo05,Saro06}. This discussion shows that SAMs and hydrodynamical simulations provide complementary approaches to the cosmological study of galaxy formation. The ability of hydrodynamical simulations of accurately following gas dynamics calls for the need of a close comparison between these two approaches, in order to test the basic assumptions of the SAMs. The best regime to perform this comparison is when one excludes the effect of all those processes, like feedback in energy and metals, whose different modeling in SAMs and simulation codes would make the comparison scarcely telling. Since gas cooling is the most basic ingredient in any model of galaxy formation \citep[e.g., ][]{White91}, an interesting comparison would be performed when cooling is the only process turned on. In this spirit \cite{Benson01}, using a hydrodynamical simulation of a cosmological box and a stripped-down version of SAM, compared the statistical properties of ``galaxies'' found in the two cases. They discovered that SPH simulation and SAM give similar results for the thermodynamical evolution of gas and that there is a very good agreement in terms of final fractions of hot, cold and uncollapsed gas. Similar conclusions were reached by \cite{Helly03} and \cite{Cattaneo07}. They improved the comparison performed by \cite{Benson01} by giving to the down-stripped SAM the same halo merger histories extracted from the cosmological simulation. In this way they were able to compare cooling in DM halos not only statistically but on an object-by-object basis. The result was again that the two methods provide comparable ``galaxy'' populations. \cite{Yoshida02} performed a similar comparison for a simulation of a single galaxy cluster, obtaining similarly good results. While the general agreement between the two methods is encouraging, still all the above analyses generally concentrated on comparing the statistical properties of the galaxy populations. Furthermore, if one wants to test the reliability of the cooling model implemented in the SAMs, the cleanest approach would be that of turning off the complications associated to the hierarchical merging of halos, thereby allowing gas to cool in isolated halos. The purpose of this paper is to present a detailed comparison between the predictions of cooling models, as implemented in SAMs, and results of hydrodynamical simulations in which gas is allowed to cool inside isolated halos. Our controlled numerical experiments will be run for halos having the density profile (for DM particles) of \cite{NFW} (NFW hereafter), with a range of masses, concentration parameters and average densities (related to the halos' redshift). As a baseline model for gas cooling, we consider the classical one, as originally proposed by \cite{White91}. In this model, the cooling radius is defined as the radius at which the cooling time equals the time elapsed since radiative cooling is turned on. As a result, the growth rate of the cooled gas mass is simply related to the growth rate of the cooling radius. This model was claimed by \cite{White91} to be close the exact self-similar solutions of cooling flows presented by \cite{Bert89}. Simulation results will also be compared to another model of gas cooling, which has been recently proposed by \cite{Monaco07} in the context of the {\sc morgana} SAM, and is based on a ``dynamical'' definition of cooling radius. As a result of our analysis, we will show that the gas cooling rate in the simulations is initially faster than predicted by the classical cooling model. When the simulations are stopped, after about 8 central cooling times, this initial transient causes the classical cooling model to underestimate the cooled mass by an amount which can be as large as a factor of three, depending on the halo concentration, density and mass. A much better agreement with simulations is achieved by the alternative {\sc morgana} model. The plan of the paper is as follows. In section 2 we describe first the ``classical'' analytic model for cooling \citep{White91}, and the alternative {\sc morgana} cooling model \citep{Monaco07}. In section 3 we present numerical simulations, performed with the {\tt GADGET-2} code and in section 4 we discuss the results obtained by comparing simulations to analytical cooling models. We discuss our results in section 5 and draw our final conclusions in section 6. A more technical discussion on the differences between the classical and the {\sc morgana} cooling models is provided in Appendix A, while Appendix B gives a very simple fitting formula for the cooling flows. | We have presented a detailed analysis of cooling of hot gas in DM halos, comparing the predictions of semi-analytic models with the results of controlled numerical experiments of isolated NFW halos with hot gas in hydrostatic equilibrium. Simulations have been performed spanning a range of masses (from galaxy- to cluster-sized halos), concentrations and redshift (from 0 to 2). Smaller halos at higher redshift have not been simulated because the validity of the assumption of a hydrostatic atmosphere is doubtful when the cooling time is much shorter than the dynamical time. We have considered the ``classical'' cooling model of \cite{White91}, used in most SAMs, and the model recently proposed by \cite{Monaco07} within the {\sc morgana} code for the evolution of galaxies and AGNs. The main features of these models can be summarized as follows. The density and pressure profiles of the gas are computed by solving the equation of hydrostatic equilibrium in an NFW halo \citep{Suto98}. The classical cooling model assumes that each mass shell cools to low temperature exactly after one cooling time $t_{\rm cool}(r)$, computed on the initial conditions. The cooling radius $r_{\rm C}$ is then the inverse of the $t_{\rm cool}(r)$ function, and the cooled fraction is the fraction of gas mass within $r_{\rm C}$. The ``unclosed {\sc morgana}'' cooling model computes the cooling rate of each mass shell, then integrates over the contribution of all mass shells and follows the evolution of the cooling radius assuming that the transition from hot to cold phases is quick enough so that a sharp border in the density profile of hot gas is always present. This determines the evolution of the cooling radius $r_{\rm M}$. Moreover, to mimic the closure of the ``cooling hole'' due to the lack of pressure support at $r_{\rm M}$, the cooling radius (now called $r_{\rm M,ch}$) is induced to close at the sound speed. This defines the ``closed {\sc morgana}'' model. Our main results can be summarized as follows. \noindent {\bf (i)} The classical cooling model systematically underestimates the fractions of cooled mass. After about two central cooling times, they are predicted to be about one order of magnitude smaller than those found in simulations. Although this difference decreases with time, after 8 central cooling times, when simulations are stopped, the difference still amounts to a factor 2--3. This disagreement is ascribed to the lack of validity of the assumption that each mass shell takes one cooling time, computed on the initial conditions, to cool to low temperature. Seen from the point of view of a mass element, the time required by it to cool to low temperature is shorter than the initial cooling time when density increases and temperature is constant during cooling. This is what happens to gas particles in the simulations: they take most of time to travel from their initial position toward the cooling region, at roughly constant temperature and increasing density. The disagreement is stronger when the cooling gas comes from the shallow central region, in which case the cooling flow is markedly not self-similar. \noindent {\bf (ii)} The unclosed {\sc morgana} model gives a much better fit of the cooled mass fraction. This is mostly due to the relaxation of the assumption on the cooling time, mentioned in point (i). This model correctly predicts cooling flows which are stronger than the classical model, by a larger amount for flatter gas density profiles. In the Appendix we show that the solution is not self-similar if the slope of the density profile is shallower than $r^{-3/2}$. In this case cooling is not dominated by the the shells just beyond the cooling radius but the whole region for which the density profile is shallow contributes. \noindent {\bf (iii)} The closed {\sc morgana} model further improves the fit to the simulation results on the evolution of the cooled mass fraction, giving accurate results to within 20--50 per cent in all the considered cases, after about 8 central cooling times. This agreement is a good reward for the increase of physical motivation of this model, obtained at the modest price of letting the cooling radius close at the local sound speed. However, the closure of the cooling radius must be halted at later times for the model to give realistic results. In general, we consider the closed {\sc morgana} as a successful effective model of cooling, rather than as a rigorous physical model. \noindent {\bf (iv)} The cooling flow is well approximated by a constant flow, for which we give a fitting formula in Appendix B, and which is valid up to $\sim10$ central cooling times. In the context of models of galaxy formation, cooling of hot virialized gas is the starting point for all the astrophysical processes involved in the formation of stars (and supermassive black holes) and their feedback on the interstellar and intra-cluster media. We find that the classical model, used in most SAMs, leads to a significant underestimate of the cooled mass at early times. This result is apparently at variance with previous claims, discussed in the Introduction, of an agreement of models and simulations in predicting the cooled mass, (\citealt{Benson01,Helly03,Cattaneo07,Yoshida02}). Given the much higher complexity of the cosmological initial conditions used in such analyses, it is rather difficult to perform a direct comparison with the results of our simulations of isolated halos. Here we only want to stress the advantage of performing simple and controlled numerical tests in order to study in detail how the process of gas cooling takes place. It is well possible that the different behaviour of simulations and the classical cooling model is less apparent when the more complex cosmological evolution is considered. However, there is no doubt that the results of our analysis are quite relevant for the comparison between SAM predictions and observations. For instance, \citep{Fontanot07} have recently shown that the {\sc morgana} model of galaxy formation is able to reproduce the observed number counts of sources in the sub-mm band by using the standard Initial Mass Function (IMF) by \cite{Salpeter55}, without any need to resort to a top-heavier IMF \citep{Baugh06}. As argued by these authors, the bulk of starbursts are driven by massive cooling flows, so this difference is mostly due to the different cooling models. | 7 | 10 | 0710.2473 |
0710 | 0710.5253_arXiv.txt | Nine supergiant shells (SGSs) have been identified in the Large Magellanic Cloud (LMC) based on H$\alpha$ images, and twenty-three SGSs have been reported based on \ion{H}{1} 21-cm line observations, but these sets do not always identify the same structures. We have examined the physical structure of the optically identified SGSs using \ion{H}{1} channel maps and P-V diagrams to analyze the gas kinematics. There is good evidence for seven of the nine optically identified SGSs to be true shells. Of these seven H$\alpha$ SGSs, four are the ionized inner walls of \ion{H}{1} SGSs, while three are an ionized portion of a larger and more complex \ion{H}{1} structure. All of the H$\alpha$ SGSs are identified as such because they have OB associations along the periphery or in the center, with younger OB associations more often found along the periphery. After roughly 12 Myrs, if no new OB associations have been formed a SGS will cease to be identifiable at visible wavelengths. Thus, the presence and location of ionizing sources is the main distinction between shells seen only in \ion{H}{1} and those also seen in H$\alpha$. Based on our analysis, H$\alpha$ observations alone cannot unambiguously identify SGSs, especially in distant galaxies. | \label{sec:intro} Observations of the interstellar medium (ISM) of spiral and irregular galaxies have revealed complex filamentary structures indicative of a violent ISM \citep{MS79}. The largest of these structures are the supergiant shells (SGSs), with diameters approaching 1 kpc \citep{gm}. SGSs are thought to be formed by the fast stellar winds and supernova explosions of multiple OB associations; they require $10^{52}$--$10^{53}$ ergs for their creation, the equivalent of tens to hundreds of supernova explosions \citep{M80}. The diameter of these shells often exceeds the scale height of the galactic gas disk, allowing them to puncture the gas disk and vent their hot interior gas into the galactic halo. The expansion of SGSs may also cause further star formation. It has been suggested that the compression of the ISM on the rims of these shells and its subsequent gravitational collapse may be a significant cause of self-propagating star formation in galactic disks, thereby providing important insight into galactic evolution. To examine the impact of SGSs on their host galaxy, it is necessary to determine the physical structure of SGSs that have been identified based on their morphology alone, since some may be merely chance superpositions of unrelated filaments. The Large Magellanic Cloud (LMC) is an excellent site to explore the nature of SGSs. With its low inclination there is little line-of-sight confusion, and its proximity (50 kpc; Feast 1999) allows us to probe its ISM in a detail impossible for most other galaxies. SGSs in the LMC were first identified by \citet{gm}, based on visual inspection of H$\alpha$ images of the ionized ISM. This original paper tabulated four SGSs, a number that was later expanded to nine by \citet{M80}. Later work on LMC SGSs, however, has also been done based on observations of \ion{H}{1}, the neutral atomic ISM. \citet{Kim99} examined the LMC in \ion{H}{1} and tabulated 23 SGSs that have diameters greater than 360 pc. Curiously, there is not a one-to-one correspondence between the largest \ion{H}{1}-selected SGSs and the nine H$\alpha$-selected SGSs. To investigate the physical structure of the H$\alpha$-selected SGSs and their relationship with the \ion{H}{1}-selected SGSs, we have examined the \ion{H}{1} environment of the H$\alpha$-selected SGSs using the H$\alpha$ images from the Magellanic Cloud Emission Line Survey \citep[MCELS;][]{MCELS05} and the \ion{H}{1} synthesis maps made with combined observations from the Australia Telescope Compact Array and Parkes Telescope \citep{Kim03}. In this paper we describe these data sets in \S 2 and our methodology in \S 3, and critically examine the physical structure of the nine H$\alpha$-selected SGSs in \S 4. We discuss the nature of the SGSs and compare the \ion{H}{1}-selected and H$\alpha$-selected shells in \S 5. A summary is given in \S 6. | \label{sec:disc} Comparison of neutral and ionized hydrogen images has allowed us to gain an understanding of the structure and kinematics of the optically identified SGSs. We have determined that these shells can be separated into three categories: \begin{itemize} \item {\bf Simple Coherent Shells:} These are neutral shells whose inner walls are photoionized. While the stars in the interior of the SGSs are responsible for the expanding structure, most of the ionizing sources are distributed along the periphery. Of the optically identified LMC SGSs, LMC 1, 4, 5 and 6 fall within this category. \item {\bf Complex Shells:} These are structures in which the ionized shell delineates only part of a larger or more complex neutral shell. The complex shells include LMC 2, 3 and 8. In the cases of LMC 2 and LMC 3, the optical shell illuminates only the northern portion of an larger neutral structure, while the single optically identified shell LMC 8 appears to contain multiple expanding neutral structures, of which only the eastern part is associated with ionized gas. \item {\bf False Shells:} These are those optically identified structures whose neutral \ion{H}{1} counterpart does not show any of the following characteristics to indicate a shell structure: (1) expansion within the shell boundary, as expected from an expanding shell; (2) raised column densities along shell rims at the systemic velocity, showing a shell morphology; and (3) central cavity at the systemic velocity, or for an \ion{H}{1} hole. A physical shell, even if it has been stalled, can still be recognized by the latter two characteristics. We found no evidence in neutral hydrogen to confirm the shell structure of LMC 7 and LMC 9, and characterize them as false shells. LMC 7 lacked any corroborating \ion{H}{1} structure, while the optically identified shell LMC 9 was seen to be a chance superposition of ionized filaments and OB associations, with a neutral hydrogen structure in no way indicating a shell. \end{itemize} The nature of the optically identified SGSs can be understood as the result of the interplay between OB associations and the ISM. The fast stellar winds and supernova explosions from OB associations clear away the interstellar gas and form a large expanding shell structure. The initial OB associations will lose their ionizing power after $\sim$12 Myr, the lifetime of a 15 $M_{\odot}$ star, the least massive star whose UV flux still produces detectable \ion{H}{2} gas. The recombination timescale of ionized gas is $7.6\times10^4 N_{\rm e}^{-1}$ yr, where $N_{\rm e}$ is the electron density in cm$^{-3}$. The interstellar gas density in a SGS is most likely in the range of 0.1--1 cm$^{-3}$, and the recombination timescale is much less than 1 Myr. Thus, a SGS with no new OB associations will cease to be observable in H$\alpha$ soon after 12 Myrs, and only an \ion{H}{1} shell will be detected. A summary of the diameters ($D$), average expansion velocities ($V_{\rm exp}$), and maximum expansion velocities ($V_{\rm max}$) of the optically identified SGSs in the LMC is given in Table \ref{tab:SGS}. We estimate their ages from these data using the formula $t \simeq \eta (D/2) V_{\rm exp}^{-1}$, where $\eta$ is 0.4 for a momentum-conserving bubble \citep{S75}, and 0.6 for an energy-conserving bubble \citep{C75, W77}. We adopt an intermediate value $\eta=0.5$. Examining the ages derived in Table \ref{tab:SGS}, we find that all but one, SGS 8, are less than 12 Myr old. SGS 8, however, contains the OB associations LH18, LH24 and LH26 (see section \ref{sec:sgs8}), of which LH24 and LH26 are still young enough to contain \ion{H}{2} gas. Those that are younger than 12 Myr all have young OB associations along their peripheries. These young OB associations explain the presence of an optical counterpart to this \ion{H}{1} shell, regardless of their ages. Of the \ion{H}{1}-identified SGSs \citep{Kim99}, 15 have average diameters $\geq$ 600 pc, the threshold used for the optical SGSs: SGS 1, 3, 4, 5, 6, 7, 9, 10, 11, 12, 17, 18, 19, 20, and 23. Some of these \ion{H}{1}-identified SGSs are associated with the H$\alpha$-identified SGSs: SGS 3 with LMC 1, SGS 4 with LMC 8, SGS 7 with LMC 5, SGS 9 with LMC 9, SGS 11 with LMC 4, SGS 12 with LMC 3, and SGS 19 and 20 with LMC 2. As described in \S 4 and discussed earlier in this section, the H$\alpha$-identified SGSs may correspond to partial or complete walls of \ion{H}{1}-identified SGSs, depending on the distribution of OB associations that provide ionizing fluxes. Among the other \ion{H}{1} SGSs that are not associated with H$\alpha$ shells, SGS 1 and 6 do not contain any OB associations, and hence have no ionized counterparts; SGS 5 and 23 have no OB associations or \ion{H}{2} regions along the most well-defined shell rims; SGS 10 and 18 do not show obvious shell structures in either column density or channel maps, so they may not be physical shells; SGS 17's northern rim is coincident with a long H$\alpha$ arc that is likely to be ionized by massive stars in the 30 Doradus complex in the south \citep[the H$\alpha$ arc is sketched in dashes in Fig.\ 1 of][]{M80}. These comparisons adequately illustrate the importance of the presence of young OB associations to photoionize the gas in \ion{H}{1} SGSs to produce H$\alpha$-identified shell morphology. Supergiant shells have been identified in distant galaxies based on H$\alpha$ images \citep[e.g.,][]{HG97,M98}. However, from our analysis of the LMC H$\alpha$ SGSs, it is clear that the H$\alpha$ morphology alone is not sufficient to identify physically expanding shells and that the H$\alpha$ SGSs represent only the SGS population that possess young OB associations. High-resolution \ion{H}{1} position-velocity datacubes are needed to verify the shells' structure and to reveal the entire population of SGSs. | 7 | 10 | 0710.5253 |
0710 | 0710.3063_arXiv.txt | {The high energy neutrino telescope NT200+ is currently in operation in Lake Baikal. We review the status of the Baikal Neutrino Telescope, and describe recent progress on key components of the next generation kilometer-cube (km3) Lake Baikal detector, like investigation of new large area phototubes, integrated into the telescope. } \begin{document} | The Baikal Neutrino Telescope is operated in Lake Baikal, Siberia, at a depth of {1.1~km}. Deep Baikal water is characterized by an absorption length of $L_{abs}(480 $nm$) =20\div 24$ m, a (geometric) scattering length of $L_s =30\div 70$ m and a strongly anisotropic scattering function with a mean cosine of scattering angle $0.85\div 0.9$ \cite{APP1}, and by a level of bioluminescence and other natural backgrounds that are well below seawater sites. The first stage telescope, NT200, started full operation in spring 1998 and contained 192 Optical Modules (OMs). The favorable water properties, and a relatively simple and reliable design led to the physics success of this comparably small telescope. Low light scattering allows for a sensitive volume of a few Mtons at PeV shower energy scale, well beyond the geometric detector limits. For a review of the high sensitivity limits on UHE astrophysical neutrino's as well as best so far obtained limits on relativistic magnetic monopoles and other results, see \cite{ICRC07_B}. The upgrade to NT200+ was a logical consequence of the large external sensitive volume, now to be fenced by sparsely instrumented external strings of OMs. In this paper, we review the current status of the Baikal Neutrino Telescope as of 2007, and the activities towards the km3-scale detector. Results on a prototype device for acoustic neutrino detection, obtained with a stationary setup in 2006/2007, are reported elsewhere in these proceedings \cite{ICRC07_C}. | The Baikal Neutrino Telescope is taking data currently in it's NT200+ configuration - an upgrade of the original NT200 telescope for improved high energy shower sensitivity. For a km3-detector in Lake Baikal, R\&D-activities have been started. The NT200+ detector is, beyond its better physics sensitivity, used as an ideal testbed for critical new components. Modernization of the NT200+ DAQ allowed to install a prototype FADC PM readout. Six large area hemispherical PMs have been integrated into NT200+ (2 Photonis XP1807/12" and 4 Hamamatsu R8055/13"), to facilitate an optimal PM choice. A prototype new technology string will be installed in spring 2008; and a km3-detector Technical Design Report is planned for fall 2008. {\it This work was supported by the Russian Ministry of Education and Science, the German Ministry of Education and Research and the Russian Fund of Basic Research (grants 05-02-17476, 05-02-16593, 07-02-10013, 07-02-00791), by the Grant of the President of Russia NSh-4580.2006.2 and by NATO-Grant NIG-9811707 (2005).} \vspace*{-3mm} | 7 | 10 | 0710.3063 |
0710 | 0710.3313_arXiv.txt | A multi-dimension, time-dependent Monte Carlo code is used to compute sample $\gamma$-ray spectra to explore whether unambiguous constraints could be obtained from $\gamma$-ray observations of type~Ia supernovae. Both spherical and aspherical geometries are considered and it is shown that moderate departures from sphericity can produce viewing-angle effects that are at least as significant as those caused by the variation of key parameters in one-dimensional models. Thus $\gamma$-ray data could in principle carry some geometrical information, and caution should be applied when discussing the value of $\gamma$-ray data based only on one-dimensional explosion models. In light of the limited sensitivity of current $\gamma$-ray observatories, the computed theoretical spectra are studied to revisit the issue of whether useful constraints could be obtained for moderately nearby objects. The most useful $\gamma$-ray measurements are likely to be of the light curve and time-dependent hardness ratios, but sensitivity higher than currently available, particularly at relatively hard energies ($\sim 2$ -- $3$~MeV), is desirable. | \label{sect:intro} Although the paradigm that Type Ia supernovae (SNe~Ia) result from the explosions of carbon-oxygen white dwarfs is well established, many issues regarding the nature of the progenitors and the explosion mechanism remain unclear (see e.g. \citealt{hillebrandt00}). Achieving a clearer understanding of SNe~Ia is important because of their role in the chemical evolution of galaxies and as cosmological distance indicators. Considerable theoretical effort has gone into modelling of SNe~Ia explosions. Recently, fully three-dimensional (3D) modelling of the explosion \citep{reinecke02,gamezo03,roepke05,roepke06,jordan07} has become feasible, allowing a detailed treatment of the hydrodynamical instabilities and turbulence which play a pivotal role. The 3D structure predicted by these models has been shown to affect observables such as optical/ultra-violet/infrared ({\sc uvoir}) light curves and spectra (e.g. \citealt{kasen06c, sim07, sim07b}). Although they are primarily detected through their optical emission, SNe~Ia are also expected to be $\gamma$-ray sources owing to the large masses of radioactive isotopes synthesised in the explosion (e.g. \citealt{travaglio04}). In principle, measurements of $\gamma$-ray emission from SNe~Ia could provide important diagnostics since $\gamma$ rays trace almost directly the mass and velocity distribution of the products of nuclear burning. Therefore, there has been considerable theoretical work on SN~Ia $\gamma$-ray emission (e.g. \citealt{ambwani88, burrows90, mueller91, hoeflich92, kumagai97, hoeflich98, gomez98, milne04}). Unfortunately, owing to the low sensitivity of $\gamma$-ray observatories, to date only one SN~Ia (SN1991T) has been detected (see discussion by \citealt{milne04}). However, with current instrumentation such as the {\it SPI} \citep{vedrenne03} on-board {\it Integral} \citep{winkler03}, detection of SN~Ia within several Mpc would be possible (\citealt{gomez98}). Moreover, increased sensitivity in future missions should make detection feasible for more distant SNe~Ia. Most studies of SN~Ia $\gamma$-ray emission focus on making predictions from individual models representing specific explosion mechanisms. In general, these studies have been restricted to one-dimension (although see \citealt{hoeflich02}) and to Chandrasekhar-mass models (some sub-Chandrasekhar models have been considered, e.g. \citealt{hoeflich98}; \citealt{gomez98}). Such studies demonstrated that, for nearby SN~Ia, the prospects of obtaining useful data are fairly good. Here we adopt a different but complementary approach, motivated by the increasing variety of explosion conditions suggested (both in theoretical and semi-empirical studies): rather than determining the quality of $\gamma$-ray data that would be needed to distinguish specific models, we attempt to show what might be unambiguously determined solely from data and the physics of $\gamma$-ray radiation transport. We use a multi-dimensional, time-dependent code to compute $\gamma$-ray spectra for a set of parameterised geometries spanning a broad range in the relevant physical conditions. Using these reference spectra, we highlight the quantities that would be most useful diagnostics and therefore most worthy of consideration in the design of future instrumentation. In Section~\ref{sect:code}, we briefly describe the code used to compute the reference $\gamma$-ray spectra. The reference models we employ are motivated in Section~\ref{sect:models} and their spectra are discussed in Section~\ref{sect:spectra}. Guided by our reference spectra, in Section~\ref{sect:diagnostics} we examine observational diagnostics. Finally, in Section~\ref{sect:summ}, we summarise our results. | \label{sect:spectra} \begin{figure} \epsfig{file=fig3.eps, width=8.5cm} \caption{ Representative time series of $\gamma$-ray spectra computed with four spherically symmetric models (identified in the second panel) at times after explosion as indicated on each panel. The $^{56}$Ni and $^{56}$Co emission lines are identified in the top panel. The four flux-bands discussed in Section 5.3 are indicated in the bottom panel. Fluxes are for a source distance of 1~Mpc. \label{fig:spec} } \end{figure} Fig.~\ref{fig:spec} shows a time series of spectra computed from Model~SC and spectra from three other spherical models (SS, SFeR and SM) are over-plotted. (For clarity, spectra from the remaining four models are not plotted. They are, however, included in all the discussions of observable diagnostics in subsequent sections.) At times close to maximum light the $\gamma$-ray spectrum consists of strong emission lines, mainly due to $^{56}$Co, with significant continuum arising from Compton scattering of line photons. Since optical depths decrease with time in the expanding ejecta, the strength of the lines relative to the continuum increases with time. At early times, the ``mixed'' model (Model~SM) shows a harder continuum than the standard model and stronger Ni emission lines (particularly 0.158~MeV and 0.275~MeV). These lines are also fairly strong in Model~SNiS (not plotted) where they originate in the surface layer of $^{56}$Ni. As noted by \citet{gomez98}, these Ni lines can only form when the source Ni lies at small $\tau_C$; since $\sigma_C$ decreases with increasing photon energy, the soft-energy lines are degraded most easily, becoming swamped by photons down-scattered from harder energies. In Model~SS, the 0.158 and 0.275~MeV lines are almost completely buried by the strong continuum which persists until well after maximum light as a consequence of the high $\tau_C$'s in this model. Model~SFeR shows the effect of composition as intended. Above about 0.3~MeV, its spectra are indistinguishable from those of Model~SC. At soft energies, however, the Model~SFeR flux is lower by up to an order of magnitude. This is due entirely to the difference in the photoabsorption cross-section arising from the choice of composition. Although weaker, a similar effect appears in Model~SNiS because of the relatively large amounts of iron-group material in that model, particularly in the surface layer. Should energy-resolved data of sufficient sensitivity to measure both the line and continuum emission across the entire 0.1 -- 3.5~MeV spectrum be obtained, it would be used to evaluate SN models by direct comparison. However, given the rarity of very nearby SNe~Ia and the limited sensitivity of $\gamma$-ray observatories, in the next section we will consider what information could be extracted from low quality $\gamma$-ray data in the form of simple diagnostics. We will focus on line and continuum fluxes. Although the energy resolution of $\gamma$-ray instruments can be high enough to resolve spectral lines (e.g. with {\it SPI Integral}, see \citealt{roques03}), sensitivity limits make it practically impossible to measure detailed line shapes except for extremely nearby events \citep{gomez98}. \label{sect:summ} Using a Monte Carlo code we computed $\gamma$-ray spectra for a variety of models to explore whether unambiguous constraints could be obtained from $\gamma$-ray observations of SNe~Ia. Two aspherical toy geometries (a lop-sided distribution of Ni and an ellipsoidal ejecta) show that moderate departures from sphericity can produce viewing-angle effects at least as significant as those due to variations of key parameters in 1D models. Thus $\gamma$-ray data could carry some useful constraints on possible geometries, but caution must be applied when evaluating its potential usefulness in distinguishing specific explosion scenarios. Given the limited sensitivity of current $\gamma$-ray missions, we conclude, in agreement with previous studies (e.g. \citealt{gomez98}), that there is little prospect for obtaining useful constraints from line-ratio diagnostics except for fortuitously nearby objects. Instead, we suggest that the best prospects are offered by hardness ratios. In particular, owing to the simplicity of the physics underlying the $\gamma$-ray spectrum, a simple ratio of the total emission in a hard energy, line-dominated part of the spectrum to a soft energy, continuum-dominated region could discriminate between our more extreme models. We also emphasise the value of obtaining multiple observations over a wide time period given the diagnostic power of the light curve shape. In planning future $\gamma$-ray missions, greater sensitivity at harder energies ($\simgt 2$~MeV) should be given high priority since several of the most potentially useful diagnostics (e.g. our $R_{1}$- and $H_{4}$-ratios) require measuring the hardest energies in the spectrum. | 7 | 10 | 0710.3313 |
0710 | 0710.4019_arXiv.txt | Principal Component Analysis (PCA) is a well--known technique used to decorrelate a set of vectors. It has been applied to explore the star formation history of galaxies or to determine distances of mass--lossing stars. Here we apply PCA to the optical data of Planetary Nebulae (PNe) with the aim of extracting information about their morphological differences. Preliminary analysis of a sample of 55 PNe with known abundances and morphology shows that the second component (PC2), which results from a relation produced by the parameters log(N/O), initial and final mass of PNe, is depending on the morphology of PNe. It has been found that when log(N/O) $< -0.18$ the PNe's nitrogen is low independently on the oxygen abundance for either Bipolar ({\it B}), Elliptical ({\it E}) or Round ({\it R}) PNe. An interesting result is that both {\it E} and {\it R} PNe have log(N/O) $< 0$ while only {\it B} PNe show negative and positive values. Consequently, {\it B} PNe are expected to have higher nitrogen values than the {\it E} and {\it R} PNe. Following that and a second sample of 35 PNe, n$_{\rm e}$ is also found to be higher in {\it B} PNe. Also, in all PNe morphologies PC2 appears to have a minimum at 0.89 and PNe's initial mass at 2.6 M$_{\odot}$. 5--D diagrams between PCAs components and physical parameters are also presented. More results will follow while simple models will be applied in order to try to give a physical meaning to the components. | \label{sec:1} Planetary Nebulae (PNe) are powerful tools in the study of the evolutionary scenario of intermediate mass stars. They play an important role in the chemical enrichment history of the interstellar medium and many efforts have been devoted to determine the physical parameters of Galactic PNe (like T$_{\rm e}$, n$_{\rm e}$, T$^{\star}$, L$^{\star}$, distance, abundances; \cite{journal1} and references therein). For this study a number of methods have been used either by using the observational results directly or by developing simulation models like CLOUDY \cite{journal2}. Using the latter, the values of the physical parameters can be determined (i.e. \cite{journal3}) but not any possible correlation among them. So far, only statistical methods had been used in order to find correlations between the parameters. However, there is a well--known technique (PCA) which can be used in order to study the possible correlation between these parameters. PCA technique uses a sample of observed parameters and creates a new sample of independent components which are linear combination from the previous parameters. | 7 | 10 | 0710.4019 |
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0710 | 0710.3712.txt | We apply a wind model, driven by combined cosmic-ray and thermal-gas pressure, to the Milky Way, and show that the observed Galactic diffuse soft X-ray emission can be better explained by a wind than by previous static gas models. We find that cosmic-ray pressure is essential to driving the observed wind. Having thus defined a ``best-fit'' model for a Galactic wind, we explore variations in the base parameters and show how the wind's properties vary with changes in gas pressure, cosmic-ray pressure and density. We demonstrate the importance of cosmic rays in launching winds, and the effect cosmic rays have on wind dynamics. In addition, this model adds support to the hypothesis of Breitschwerdt and collaborators that such a wind may help explain the relatively small gradient observed in $\gamma$-ray emission as a function of galactocentric radius. | Large-scale galactic outflows are usually considered in the context of starburst galaxies or Active Galactic Nuclei \citep{VeCeBH2005}. These outflows are interesting not only intrinsically (what drives the outflow?) but for the interstellar and intergalactic media (how is the host galaxy affected, and what metals are ejected from the galaxy?). To examine these questions, we have built a thermal and cosmic-ray driven wind model. Our investigation into such models was first inspired by observational hints that the Milky Way may possess a kiloparsec-scale wind; this paper further explores that possibility. To motivate this study, we first introduce the observational evidence for a Galactic wind (\S\S\ref{diffuseXrays} and \ref{CRdensity}) and then introduce the cosmic-ray and thermally driven wind model (\S\ref{modelDef}). In \S\ref{compareToObs}, we calculate the X-ray emission from this wind, and then compare it to the observations. After finding the best-fit wind model, we then explore the parameter space around that model (\S\ref{ParameterSurvey}) to understand more about how the wind is modified by varying the input and fit parameters. Our conclusions are given in \S\ref{Conclusions}. But first, observational hints for an outflow from our own Galaxy. %To start applying this model, we look to a perhaps %unlikely first suspect: the Milky Way. \subsection{X-ray Observations}\label{diffuseXrays} The Milky Way does exhibit clues that it might drive a large-scale wind. The first of these is an enhancement in the diffuse soft X-ray emission, stretching over the longitude range $-20^\circ \la l \la 35^\circ$ with an emission scale height in the southern Galactic hemisphere of $b \sim -17^\circ$ (see Fig.~\ref{r45Intensity}). This emission was first noted by \citet{SnowdenEtAl95}, who modeled it with an isothermal plasma with a temperature $T = 4 \times 10^6$~K, a midplane electron density of $n_{\rm e, midplane} \sim 3.5 \times 10^{-3}$~cm$^{-3}$, and a midplane thermal pressure of $P_{\rm g, midplane}/k \sim 2.8 \times 10^4$~cm$^{-3}$~K. At approximately the same time, \citet{BreitschwerdtSchmutzler94} suggested that the average \textit{all-sky} X-ray emission (not only that emission in the region defined above) in all \textit{ROSAT} bands might be explained by delayed-recombination in a large-scale cosmic-ray and thermally driven wind \citep[see also][]{BreitschwerdtSchmutzler99}. Later, \citet{AlmyEtAl00} used intervening absorption to show that at least half of the central, enhanced X-ray emission lies more than 2~kpc from the sun \citep[see also][]{ParkEtAl97, ParkEtAl98}. Since that measurement was made in the Galactic plane, where the absorption is strongest, it was inferred that most of the emission observed at higher latitudes lies beyond that 2~kpc distance. \citet{AlmyEtAl00} also improved on previous modeling efforts: that work presents a model of the emission due to a static polytropic gas (with $\gamma = 5/3$), and very importantly, includes the effects of known background components, such as the stellar background, extragalactic background, and an additional isotropic background (to fit high-latitude emission). For comparison, their model had a central temperature of $T_0 = 8.2 \times 10^{6}$~K, a central electron density of $n_{\rm e} = 1.1 \times 10^{-2}$~cm$^{-3}$, and a central pressure of $P_{\rm g,0} = 1.8 \times 10^5$~cm$^{-3}$~K. We will compare our results with this static polytrope model to investigate whether a wind model for this emission is feasible. \begin{figure*}[] \begin{center} \includegraphics[width=14cm]{f1_color.ps} \caption{X-ray emission at $3/4$~keV (the ``R45 band'') as seen by \textit{ROSAT} \citep{SnowdenEtAl97}. These observations suggest a ``Galactic X-ray Bulge'', seen most clearly in the southern Galactic Hemisphere, and stretching over the Galactic longitude range, $l$, from $|l| \la 30^\circ$ and down to approximately $-15^\circ$ in Galactic latitude. This paper asks whether the X-ray bulge in the southern Galactic Hemisphere can be explained with a combined thermal and cosmic-ray driven wind. \label{r45Intensity}} \end{center} \end{figure*} \subsection{Cosmic Ray Source Density}\label{CRdensity} Another indicator of a Galactic wind comes from measurements of the density of cosmic rays as a function of Galactocentric radius, $R$. The source density of cosmic rays can be determined via $\gamma$-ray emission: the production of $\gamma$-ray photons with energies exceeding about 50~MeV is dominated by collisions of cosmic rays with gas in the interstellar medium \citep{BlBlHe84}. Since the galaxy is largely transparent to such high-energy photons, the $\gamma$-ray emissivity at those energies yields the cosmic ray source density. If cosmic rays are produced in supernovae remnants, then since the source density of supernovae remnants seems to increase with decreasing $R$, the cosmic-ray source density should increase as well. However, it has been known for some time \citep[e.g,][]{Bloemen89} that the inferred cosmic-ray source density is relatively flat, compared to the supernova density, as a function of $R$. There has, however, been some debate about whether supernovae remnants are an accurate tracer \citep[since those surveys are subject to various selection effects; see, e.g.,][and references therein]{StrongEtAl04}. Recent surveys of the pulsar population \citep{LorimerEtAl06} also show that the pulsar source density increases towards the center of the Galaxy, as shown in Figure~\ref{crSourceDensity}. This is true irrespective of the model of how $n_{\rm e}$ varies in the disk, although the magnitude of the pulsar population gradient with $R$ depends strongly on the $n_{\rm e}$ model. So, there remains a mismatch between the observed source density of cosmic ray ``producers'' and the cosmic rays themselves. It has already been pointed out that the observed slow rise in cosmic rays may be due to a wind emerging from the disk, advecting cosmic rays outwards \citep{BloemenEtAl93, BrDoVo02}. In the case of \citet{BloemenEtAl93}, a wind model was applied to the entire Galactic disk; as a result, only a very slow wind was found to be compatible with the inferred cosmic-ray source density. In contrast, \citet{BrDoVo02} applied their cosmic-ray and thermally driven wind model, where the wind velocity varied as a function of radius and height; they also took into account anisotropic diffusion. With this model, a small radial gradient in the cosmic ray source density could be explained. An alternate explanation for this slow change in the cosmic ray population with $R$ was proposed by \citet{StrongEtAl04}, who found that a radial variation in the $W_{\rm CO}$-to-$N(H_2)$ ratio by a factor of 5 to 10 could explain the $\gamma$-ray observations. In this paper we primarily address the question of the origin of the diffuse, soft X-ray background emission; we will, however, concentrate on a large-scale wind model, keeping in mind its possible application to the cosmic ray source density. \begin{figure}[h] \begin{center} \includegraphics[width=6cm,angle=-90]{f2.ps} \caption{Comparison of two different calculations of the pulsar population as a function of Galactocentric radius \citep{LorimerEtAl06} vs. the cosmic ray source density implied from the observed $\gamma$-ray emissivity \citep{StrongEtAl04}. The two different curves for the pulsar distribution result from assuming a smooth distribution of $n_{\rm e}$ in the Galaxy \citep{LyMaTa85}, or a clumped distribution, using \citet{CordesLazio02} and \citet{FaKa06}; for details, see \citet{LorimerEtAl06}. The fact that the cosmic-ray distribution does not seem to follow the pulsar population has been known for some time \citep{Bloemen89}, but there is no consensus on the reason. A cosmic-ray and thermal pressure-driven wind may help explain the cosmic-ray source population. \label{crSourceDensity}} \end{center} \end{figure} %Both the diffuse X-ray emission and the slow change in cosmic-ray %population towards the Galactic center hint at the possibility of a %large-scale wind driven by both thermal and cosmic-ray pressure. We %next outline a wind model including these components. | We have implemented a simplified cosmic ray- and thermally-driven wind and have used it to try to explain the soft, diffuse X-ray emission seen towards the Galactic Center. We find that such a wind can indeed match the observed averaged X-ray emission quite well, and in fact fits demonstrably better than the static polytrope model of \citet{AlmyEtAl00}. It is important to note that this wind is approximately equally powered by both cosmic rays and thermal pressure: cosmic rays are important in helping this relatively cool wind escape from the Galactic potential. It is also quite interesting that this wind does not require excessive thermal or cosmic-ray pressures (both pressures are not extreme compared to what has been estimated for the inner Milky Way), nor does this simple model require much more energy than the standard inferred supernova rate implies. Taking this result at face value, such a wind would be very important to the ``ecology'' of the Milky Way due to the high mass loss rate of 2~$M_{\odot}$~yr$^{-1}$. In addition, such a wind would also play an important role in removing angular momentum from matter in the Galactic disk and allowing matter to move radially inward \citep{Zirakashvili1996}. At the least, this shows that such wind models should be considered further for the Milky Way; they may be able to explain at least a substantial fraction of the observed soft X-ray emission. Further, as other researchers have already shown \citep{BrDoVo02}, such winds can also be used to explain the unexpectedly slow rise in $\gamma$-ray emission towards the center of the Galaxy; this work therefore gives independent support to the Galactic wind hypothesized in \citet{BrDoVo02}. We have not yet calculated the effect of the best-fit wind model on the cosmic-ray distribution, but a simple approximation will investigate if this wind is removing cosmic rays at too high a rate. We reason as follows. The best-fit wind model shown here has a cosmic-ray advective timescale of $\sim 4.3 \times 10^6$~years. Therefore, supernova in the disk must resupply the cosmic-ray pressure on that timescale. We know that the approximate total supernova energy in the disk (from observations, see \S\ref{totalEnergyFlux}) is $\sim 1.7 \times 10^{41}$~ergs~s$^{-1}$. So, if some fraction of this energy, $\epsilon_{\rm CR}$, is given to cosmic rays, and distributed over the volume occupied by hot gas where the wind is launched (over the Galactocentric radius range of $1.5$ to $4.5$~kpc, and a height range of $\pm$2~kpc, to be conservative), that energy density should be similar to the cosmic-ray pressure required to launch the wind. Calculating the resultant buildup of $P_{\rm c}$ over $4.3 \times 10^6$ years at the observed SN rate, we find $P_{\rm c} \sim \epsilon_{\rm CR} \cdot 6.9 \times 10^{-12}$~dyne~cm$^{-2}$ or $\sim \epsilon_{\rm CR} \cdot 5.0 \times 10^4$~K~cm$^{-3}$. This is actually of order the $P_{\rm c}$ that the best-fit model requires (see Table~\ref{bestFitParams}), although it would require $\epsilon_{CR} \sim 0.6$ to duplicate the best-fit $P_{\rm c}$, which is relatively high. Still, this simple, rather conservative calculation shows that the high $\dot{M}$ wind shown here does not remove cosmic rays much more quickly than they can be replenished by the normal SN rate in the Galaxy, although the removal rate of cosmic rays is certainly non-negligible, and would affect the density of cosmic rays towards the Galactic center. Of course, a more detailed calculation is required (with a more detailed wind model), but this again shows that the best-fit wind model is at least feasible, and would have a significant but not destructive effect on the Galaxy's cosmic-ray density. \subsection{Future Improvements} More detailed models are clearly needed; there are a few concerns about the current model that could be addressed with more realistic wind models. For instance, we have assumed uniform density at the base of the wind over the area of the disk from $R=1.5$~kpc to 4.5~kpc, which leads to a mass outflow rate of order 2~$M_{\odot}$~yr$^{-1}$. This seems quite high, but in the context of a more detailed model with variations in density within the wind, we might expect that the $n^2$ weighting of emissivity would favor overdensities, and allow an inhomogeneous wind to better reproduce the observations with a smaller mass outflow rate. In addition, it is possible that other effects limit the gas to velocities below those in this simple model; drag effects may slow down the wind \citep[e.g.,][]{EM07}, and could lead to some of the gas forming part of the Galactic fountain \citep{Bregman99}. On the other hand, turbulence may be an important additional source of energy for the wind, but we have not included such an input in this work. Also, the effects of distributed mass loading, which could be relatively easily incorporated into this model, have been ignored so far; such mass loading would be very important to the emission properties of the wind, and the potential observability of such winds in many galaxies. Finally, note that we have used \citet{AG89} abundances; if the wind starts out with super-solar abundances (particularly in oxygen), a smaller mass outflow rate would be required. The cosmic ray physics in this wind is still quite simple. As mentioned previously, we have assumed zero diffusivity of the cosmic rays throughout the wind \citep[c.f.,][]{BrMcVo1993}. We have also neglected thermal conductivity (our calculations show that conductivity may be important for $z \la 350$~pc). Including both of these effects would be essential to future progress for this wind model: as mentioned previously (see \S\ref{totalEnergyFlux}), including the effects of conduction may help lower the energy requirements in the wind, bringing the total power of the wind closer to that of the total inferred supernovae power in the Galactic disk below the wind. We also note that we are using a model of the gravitational potential that we know overestimates the potential in the Milky Way (see \S\ref{modelDef}). Adopting the more detailed potential model of \citet{DehnenBinney98} may allow a lower total-energy wind to duplicate the soft X-ray observations. Finally, we have also assumed the hydrodynamic model of \citet{McKenzieWebb84} and \citet{BrMcVo1991} for the interaction of cosmic rays with \alfven waves and the gas. To further examine this model, we will next consider the effects of higher cosmic-ray fluxes \citep{Zweibel2003} and apply it in other settings. \subsection{Future Tests} How can we further test this model? In analogy to early studies of the solar wind, this outflow may impact clouds in the vicinity of the galaxy, perhaps causing ``comet-tail'' extensions to high velocity clouds above the plane of the Milky Way. The formation of such ``tails'' would depend on the velocity of the wind. This has been studied in some detail before \citep{BenjaminCox02}, but should be reconsidered in the context of the predictions of these winds. In addition, it may be possible to study the kinematic impact of this wind on, for instance, the Magellanic Stream (A. Burkert, personal communication). Another way to test the model would be to compare the wind with absorption columns and emission spectra towards the center of the Galaxy. Concerning absorption measurements, recent \textit{Chandra} observations \citep{FutamotoEtAl04} towards the low-mass X-ray binary 4U 1820-303 show significant columns in \al{O VII}, \al{O VIII}, and \al{Ne IX}. As this X-ray binary is located within ten degrees of the Galactic center, at a distance of approximately 7.6~kpc, it seems an ideal target. However, similar absorption columns are found towards objects on sightlines that do not intercept the base of the wind; for instance, Mrk~421 shows similar absorption columns \citep{FutamotoEtAl04}, but is far out of the plane. This leads us to conclude that the observed absorption is somewhat local to the Sun's position, and, as such, this absorption does not constrain the wind model. However, recent Suzaku measurements of emission at various latitudes along different sightlines towards the Galactic center (Rocks, 2008, in preparation) may help constrain the properties of the wind. Finally, these winds, including their important cosmic-ray component, will also emit synchrotron radiation. We are now calculating the synchrotron emission expected from these wind models (Schiller et al., 2008, in preparation). This will allow exploration of the wind's synchrotron emission as compared to recent models of Galactic synchrotron which begin to map the three-dimensional cosmic-ray emissivity \citep{NordEtAl06}. Looking at a wider field of application, such wind models may also be quite important in application to starburst galaxies \citep[e.g.,][]{GallagherSmith05, SDR06} and dwarf galaxies. The general applicability of these kinds of models to starbursting galaxies has been shown by \citet{Breitschwerdt03} in fitting cosmic-ray and thermally driven wind models to NGC 3079. | 7 | 10 | 0710.3712 |
0710 | 0710.5535_arXiv.txt | We present new \emph{Chandra} observations of a high redshift ($z$$\sim$1) galaxy cluster discovered in the Red-Sequence Cluster Survey (RCS): \rcsB. X-ray luminosity measurements and mass estimates are consistent with $L_X$--$T_X$ and $M_\delta$--$T_X$ relationships obtained from low-redshift data. Assuming a single cluster, X-ray mass estimates are a factor of $\sim$10--100 below the red-sequence optical richness mass estimate. Optical spectroscopy reveals that this cluster comprises two components which are close enough to perhaps be physically associated. We present simple modeling of this two-component system which then yields an X-ray mass and optical richness consistent with expectations from statistical samples of lower redshift clusters. An unexpectedly high gas mass fraction is measured assuming a single cluster, which independently supports this interpretation. Additional observations will be necessary to confirm the excess gas mass fraction and to constrain the mass distribution. | There is considerable interest in measuring the evolution of the galaxy cluster mass function with redshift. In principle, such a measurement can test structure formation and constrain the cosmic expansion history. One of the observational challenges in such an undertaking is to identify all clusters in a particular volume; the other is to reliably infer cluster mass from observable cluster properties, over a range of cluster redshifts. Recent optical surveys, such as the Red-Sequence Cluster Survey \citep[RCS,][]{GY05}; and the Sloan Digital Sky Survey \citep[e.g.,][]{Kea07} have found large numbers of clusters at moderate to high redshift, with $0.2\lesssim z\lesssim 1$ and $0< z \lesssim 0.3$, respectively. At the high redshift extent, where spectroscopic information can be difficult to obtain on a survey scale, various optical properties, such as richness or optical luminosity, can be taken to be mass proxies, and have been shown to be strongly correlated with observables understood to be related to cluster mass, such as galaxy velocity dispersion and X-ray temperature \citep[e.g.,][]{YE03, Lea03, Pea04, Lea06}. In order to use optical richness (or any other optical property) as a mass proxy, understanding the evolution of the mass-richness relation is critical to extracting cosmological information, since any evolution in the relation must be separated from the evolution of the mass function. In particular, some care must be taken when attempting to extend the mass/observable relation to high redshift as many of these relations are computed using local data. Previous work has shown that optically-selected, high-redshift clusters at a fixed optical richness are typically lower in temperature and underluminous in X-rays when compared to local clusters \citep{Hea05,Gea04,Lea02,Dea01}. Since these clusters generally follow the local luminosity-temperature relationship, the low temperatures and luminosities are interpreted as an indication that the cluster galaxies are not in virial equilibrium with the X-ray emitting gas. However, it is important to account for the significant scatter in the relation between optical richness and X-ray luminosity, as an Eddington-like bias will scatter points of lower than expected X-ray luminosity into an optically-selected sample \citep[e.g.,][]{Bea94}. Observing these high redshift clusters is thus an avenue to study how clusters approach equilibrium. \rcsB\ is part of a larger sample of high redshift, high richness RCS clusters selected for X-ray follow up observations. X-ray properties of the sample as a whole are described in a forthcoming paper by \citet{Hea07}. Here we present \emph{Chandra} X-ray observations and optical spectroscopy of this apparently rich, X-ray underluminous object. In \S 2 we describe the X-ray data reduction and analysis. In \S 3 we compare mass estimates from the X-ray data to the known mass-temperature and mass-richness relations in literature. In \S 4 we describe the optical spectroscopy, in \S 5 we discuss our results and we summarize in \S 6. Throughout the paper we take a standard $\Lambda$CDM cosmology with $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_m = 0.3$ and $\Omega_\Lambda = 0.7$. At the cluster redshift, $z=0.9558$, 1\arcsec\ corresponds to a metric distance of 6.64 kpc. Errors are indicated with 90\% confidence intervals unless otherwise noted. | \rcsB\ has an inferred X-ray mass which is one to two orders of magnitude lower than the mass indicated by its red-sequence richness, assuming that the X-ray emission is from a single, spherical, isothermal gas distribution and that the richness is associated with a single cluster ($M_{200,X}=4.6\pm^{6.0}_{1.7}\times10^{13}$ M$_{\odot}$ versus $M_{200,B_{gcR}}=1.1\times10^{15}\pm0.46$ dex M$_{\odot}$). This is a significant discrepancy for two reasons. First, as mentioned before, richness is well correlated on average with other observable properties. Secondly, clusters with masses above $10^{15}$ M$_{\odot}$ are quite rare, especially at $z\sim1$. This means that if the richness mass estimate of \rcsB\ were truly indicative of its size, it would be among the largest overdensities in the observable universe, while also being very underluminous. Yet the X-ray properties, which suggest a much less massive object, are consistent with the expected $M_{\delta}$--$T_X$ and $L_X$--$T_X$ relations, meaning that the temperature of the plasma is consistent with its observed distribution. One possible interpretation of these results is that the red-sequence galaxies used to measure the optical richness and the X-ray emitting gas trace different volumes. Specifically, while the X-ray luminous plasma is confined within one or more deep gravitational potential wells, the galaxy population may occupy a more extended region which is not yet dynamically relaxed. The existence of such a structure in \rcsB\ is supported by its two-peaked radial velocity distribution. Merging and dynamically active clusters are expected and observed to be increasingly common at higher redshifts \citep[e.g.,][]{Jea05}. Thus any cluster sample will include an increasing fraction of dynamically young systems at higher redshift. Furthermore, N-body simulations indicate that a significant fraction of cluster samples selected using broadband color discrimination will be systems whose galaxy members are not predominantly associated with a single, large potential well, but rather are spread amongst a number of smaller, but still physically associated, dark matter halos distributed along the line-of-sight \citep{Cea07}. Simulations of the RCS technique on mock catalogs tuned to reproduce observables such as the observed galaxy color distribution and the two-point correlation function show that we expect false-positive cluster detections to occur at a frequency of $\sim$5\% in the RCS \citep{G02}. This agrees with initial results based on a small number of clusters \citep{Gilbank_ea07,Bea07}. Since the red-sequence galaxies used in the RCS richness measurements form very early ($z\gtrsim2$) in high density regions, overdensities of red-sequence galaxies can be associated with very large structures which may not be relaxed during the observed epoch. The X-ray luminous plasma, on the other hand, will be confined to one or more gravitational potential wells whose size is limited by the collapse timescale. In the hierarchical collapse paradigm of a $\Lambda$CDM universe, the large, virialized wells seen in local clusters do not develop until long after the galaxies have formed. In this scenario, for any evolving cluster, there is likely to be an epoch at which the apparent richness overestimates the virialized mass. Moreover, the misinterpretation of the X-ray emission from a complex, dynamically young cluster as a single, virialized structure can lead to an overestimate of the inferred gas mass fraction and an underestimate of the total mass, which can be seen as follows. Let us consider two models for the X-ray emission seen in \rcsB. Our one-component model (Model I) is the set of assumptions used in our X-ray analysis thus far: X-ray emission is from a single, spherically-symmetric distribution of gas in hydrostatic equilibrium with a gravitational potential which is well-described by a $\beta$-model. The two-component model (Model II) assumes that there are two nominally identical (i.e.~with the same values of $\beta$, $r_0$ and $T_X$), virialized gas distributions, separated along the line of sight, and that each of the two gas distributions is responsible for half of the measured X-ray flux. Each component is assumed to obey the local $M_{\delta}$--$T_X$ and $L_X$--$T_X$ relation. Because it assumes the flux is halved between the two components, the luminosity inferred per cluster in Model II is half that of Model I. The luminosity scales as the square of the gas mass, so since the spatial distribution is the same for each cluster, the inferred gas mass and thus also the electron density will be reduced by a factor of $\sqrt{2}$ in each of the clusters relative to Model I. Since our total mass estimate for each cluster depends only on the temperature of the X-ray emitting gas and its distribution in the plane of the sky, each of the two components in Model II will have the same total mass as the single component in Model I. This means that the inferred gas mass fraction will be a factor of $\sqrt{2}$ lower using Model II versus Model I. The total mass, total gas mass and gas mass fraction results using each of the two models can be found in Table \ref{masstable}. Similarly, $n$ identical components along the line of sight would reduce the inferred gas mass per component and the overall gas mass fraction by a factor of $\sqrt{n}$ in addition to reducing the luminosity per component by a factor of $n$. In principle, luminosity measurements and the $L_X$--$T_X$ relation constrain the number of components allowed, but our luminosity errors are too large to distinguish between Model I and Model II. \citet{ME01}, using a suite of hydrodynamic cluster simulations, found that during merger events the ICM spectral fit temperature will underestimate the mass-weighted ICM temperature by $\sim20\%$, because the cool, denser inflowing gas will dominate the emission over the gas already heated by the merger. This means that for a dynamically young system, such as \rcsB, where even the most virialized components are likely to have undergone recent mergers, X-ray mass estimates for those clusters may still be below the virialized masses by several tens of percent. In a recent detailed study of the cluster Cl 0024+17, \citet{Jea07} found a similar underestimation of the true mass distribution due to the assumption of a single virialized mass structure rather than two components extended along the line of sight and in a state of ongoing dynamic interaction. In summary, the assumption of spherical symmetry and virial equilibrium for a cluster system containing extended line-of-sight structure and dynamic evolution can lead to overestimation of the gas mass fraction and underestimation of the total mass. The baryon mass fraction in relaxed galaxy clusters is expected to be a universal quantity \citep[e.g.,][]{Wea93,Vea03}, suggesting that the comparison of cluster gas mass fractions inferred by assuming spherical symmetry to canonical cluster values might be used to infer the presence of line-of-sight structure. As noted by \citet{Wea93}, the cosmological baryon mass ratio, $\Omega_b/\Omega_m$, provides an upper limit to the total gas mass fraction. This ratio from the WMAP three-year data results is $\Omega_b/\Omega_m=0.175\pm0.012$ \citep{Sea07}. Galaxy cluster samples at high redshift are expected to have lower cluster gas mass fractions than local samples \citep[e.g.,]{Hea07, Sea05, Eea04}. Also, a positive correlation between temperature and gas mass fraction has been observed \citep[e.g.,][]{Vea06, Sea03}. This suggests that a high-redshift cluster with a low X-ray temperature, such as \rcsB, ought to have a correspondingly low gas mass fraction. Instead, we observe a very high gas mass fraction assuming that it is a single, spherical matter distribution, especially measured within a large radius. At $r_{500}$, we find $f_{gas,500}=0.17\pm^{0.07}_{0.08}$. We observe cluster emission to approximately this radius, so this gas mass fraction estimate is unlikely to have the extrapolation errors that the measurement at $r_{200}$ might. Assuming that Model II is more appropriate for the physical state of \rcsB\ lowers the gas mass fraction towards what is expected. Would this pair of clusters be physically associated? As an order-of-magnitude argument, we note that the richness mass ($\sim$1$\times10^{15}$ M$_{\odot}$) is roughly equal to that of a sphere of radius $\sim$13 Mpc with the cosmic density of $z=0.96$. Optical spectroscopy shows that the redshift separation of the two components in \rcsB\ is $\Delta z\approx 0.005$, or a physical separation of about 12 Mpc in the Hubble flow. This rough agreement supports the conclusion that \rcsB\ is an incompletely-virialized system which is still approaching equilibrium, and that the galaxy distribution traces unvirialized matter extended along the line-of-sight, as well as the virialized matter traced by the X-ray gas. The free-fall collapse time for the two components given by the richness mass is approximately equal to the lookback time, indicating that this cluster would be nearly virialized by about the present epoch. Additional observations are required to determine the physical state of \rcsB. Due to the low number of source photons from this cluster, there are large errors in the X-ray luminosity and temperature. Improvements on both would increase the precision of the X-ray mass estimates and the gas mass fractions. An independent mass estimate, such as might be obtained from weak lensing, could also help distinguish which mass estimate (richness or X-ray) is more appropriate, as well as giving additional insight into the distribution of matter. A weak lensing measurement, though difficult, would be particularly interesting for this system if it could trace the spatial distribution of cluster mass perpendicular to the line of sight, possibly revealing more substructure. Additional spectroscopy, currently underway, will help to further constrain line-of-sight substructure. | 7 | 10 | 0710.5535 |
0710 | 0710.0898_arXiv.txt | { We report the detection of transits by the $3.1 M_{\rm Jup}$ companion to the V=8.17 G0V star \object{HD 17156}. The transit was observed by three independant observers on Sept. 9/10, 2007 (two in central Italy and one in the Canary Islands), who obtained detections at confidence levels of 3.0~$\sigma$, 5.3~$\sigma$, and 7.9~$\sigma$, respectively. The observations were carried out under the auspices of the Transitsearch.org network, which organizes follow-up photometric transit searches of known planet-bearing stars during the time intervals when transits are expected to possibly occur. Analyses of the 7.9~$\sigma$ data set indicates a transit depth $d=0.0062 \pm 0.0004$, and a transit duration $t=186 \pm 5$ min. These values are consistent with the transit of a Jupiter-sized planet with an impact parameter $b=a \cos{i}/R_{\star} \sim 0.8$. This planet occupies a unique regime among known transiting extrasolar planets, both as a result of its large orbital eccentricity ($e=0.67$) and long orbital period ($P=21.2 {\rm d}$). The planet receives a 26-fold variation in insolation during the course of its orbit, which will make it a useful object for characterization of exoplanetary atmospheric dynamics. } | During the past several years, the discovery rate of transiting planets has begun to increase rapidly, and twenty transiting planets with secure characterizations are currently known \footnote{ Extrasolar Planets Encyclopedia {\tt http://exoplanet.eu}}. This aggregate consists mostly of short-period hot-Jupiter type planets, with prototypical examples being HD~209458b \citep{2000ApJ...529L..45C,2000ApJ...529L..41H} and HD~189733b \citep{2005A&A...444L..15B}. These planets tend to have $M\sim 1M_{\rm Jup}$, $2 {\rm d}<P<5{\rm d}$, and tidally circularized orbits. In the past year, two remarkable discoveries have significantly extended the parameter space occupied by known transiting planets. HD~147506b \citep{2007arXiv0705.0126B} with $M=8.04 \, M_{\rm Jup}$ is by far the most massive planet known to exhibit transits. It also has the longest orbital period (5.63 days) and a startlingly large orbital eccentricity, $e\sim 0.5$. At the other end of the mass scale, Gl~436b \citep{2004ApJ...617..580B,2007A&A...472L..13G} has $M=0.07 M_{\rm Jup}$, a 2.64 day orbital period, and an eccentricity $e=0.15\pm0.01$ \citep{2007arXiv0707.2778D}. These two planets straddle more than a hundred-fold difference in mass, and their significant non-zero eccentricities are also capable of imparting important information. At present, infrared observations of transiting extrasolar planets by Spitzer present an incomplete and somewhat contradictory overall picture. It is not understood how the wind vectors and temperature distributions on the observed planets behave as a function of pressure depth, and planetary longitude and latitude. Most importantly, the effective radiative time constant in the atmospheres of short-period planets remains unmeasured, and as a result, dynamical calculations of the expected planet-wide flow patterns \citep{2003ApJ...587L.117C,2005ApJ...629L..45C,2005ApJ...618..512B,2007ApJ...657L.113L,2007arXiv0704.3269D} have come to no consensus regarding how the surface flow should appear. This lack of agreement between the models stems in large part from the paucity of unambiguous measurements of the radiative time constant in the atmosphere. What is needed, is a transiting planet with both a long-period orbit and a large orbital eccentricity. If such a planet were known, then one could use Spitzer to obtain infrared time-series photometry of the planet during the periastron passage. The transit guarantees knowledge of both the geometric phase function and the planetary mass. This information would in turn allow a measured rate of increase in flux to inform us of the planet's atmospheric radiative time constant in the observed wavelength regime. The orbital periods of the known transiting planets are all significantly shorter than 6 days. This bias is due both to the intrinsically lower geometric probability of transit as one moves to longer periods, and also to the fact that ground-based wide-field transit surveys that rely on photometry folding become very significantly incomplete for planets with orbital periods longer than 5 days. If one wants to detect longer-period transiting planets from the ground, a more productive strategy is to monitor known RV-detected planet-bearing stars at the times when the radial velocity solution suggests that transits may occur. This strategy has the further advantage of producing transits around stars that tend to be both bright and well-suited for follow-up observations. Long-period transiting planets present an ideal observing opportunity for small telescope observers. \cite{2003PASP..115.1355S} have suggested that a global network of telescopes, all capable of $\sim1\%$ photometry can easily outperform a single large telescope in terms of efficiency of transit recovery. Since inception in 2002, the Transitsearch.org network has conducted follow-up searches on a number of intermediate-period planets (see e.g. \citealt{2006ApJ...653..700S}). The Doppler-based discovery of HD~17156b was recently published by the N2K consortium \citep{2007arXiv0704.1191F}. The planet has M$\sin i = 3.12{\rm M_{Jup}}$, with $P=21.22$ days and $e\sim0.67$. \cite{2007arXiv0704.1191F} report that the V=8.17 G0V host star has $M= 1.2 M_\odot$ and $R= 1.47 R_\odot$. The planet's semi-major axis $a=0.15 {\rm AU}$ thus indicates a periastron distance of $a_{min}=0.0495 {\rm AU}=7.2 R_{\star}$. A best fit to the radial velocities indicates longitude of periastron $\omega=121 \pm 11^{\circ}$. The orbital orientation is favorable, yielding an a-priori geometric transit probability of $P\sim13$\%. In their discovery paper, \cite{2007arXiv0704.1191F} reported 241 individual photometric measurements obtained over a 179 day interval, and with a mean dispersion $\sigma=0.0024$ mag. No significant rotation-induced periodicity was seen. Together, the observations sampled approximately 25\% of the $1-\sigma$ transit window, and no evidence for a transit was observed. After the \cite{2007arXiv0704.1191F} discovery paper was made public, the star was added to the Transitsearch.org candidates list\footnote {http://207.111.201.70/transitsearch/dynamiccontent/candidates.html} and observers throughout the Northern Hemisphere were repeatedly encouraged to obtain photometry of the star\footnote{see www.oklo.org}. The first available window of opportunity occurred on 9/10 Sept., 2007, with the transit midpoint predicted to occur at HJD $2454353.65 \pm 0.30$. | The detection of transits by a planet with a three-week orbital period demonstrates the utility of ad-hoc networks of small telescopes for obtaining photometric follow-up of planets whose orbital parameters have been determined via Doppler radial velocities. Indeed, the transits of HD~17156b offer a plethora of interesting opportunities for follow-up observations. With its high orbital eccentricity and small periastron distance, HD~17156b appears to bear a curious kinship to HD~80606b, HD~147506b, and HD~108147b. All three of these planets occupy a locus of the $a$--$e$ plane where they should actively be undergoing tidal dissipation, and therefore they should be generating significant quantities of excess interior heat. Our measurement indicates that tidal heating is not significantly inflating the planetary radius. The nominal $R=1.1 R_{\rm Jup}$ radius predicted by baseline models (e.g. those of \citealt{2003ApJ...592..555B}) is confirmed by our observations. Follow-up photometric measurements during future transits will allow a more accurate determination of the orbital inclination of HD~17156b. An improved value for $i$, in turn, will generate an accurate assessment of likelihood that the planet can be observed by Spitzer in secondary transit, and will enable a much-improved constraint on the still-uncertain radius of the parent star. In the event that secondary transits can be observed, a direct measurement of the excess tidally generated luminosity from the planet is a distinct possibility (see e.g. \citealt{2007arXiv0707.2778D}). As a consequence of its highly eccentric orbit, HD~17156b experiences a 26-fold variation in insolation during the 10.6 day interval between periastron and apoastron. This extreme radiative forcing may drive interesting, and potentially observable dynamical atmospheric flows on the planet \citep{2007ApJ...657L.113L}. The large tidal forces experienced during periastron have almost certainly forced the planet into pseudo-synchronous rotation (e.g. \citealt{1968ARA&A...6..287G,1981A&A....99..126H,2005MNRAS.364L..66P}). Rotationally induced modulation in the infrared light curve following periastron is potentially observable, and may be of great utility in selecting between the current divergent predictions for the actual value of the pseudo-synchronous spin frequency. HD~17156b is quite massive, as is often the case for planets orbiting one member of a binary pair \citep{2007A&A...462..345D}, and the eccentricity is large. These characteristics favor a formation scenario involving migration and/or dynamical evolution in the presence of a sufficiently close external perturber. Such a perturber could be either a companion star or an additional planet(s) in the system. For some of the close-in planets with $m \sin i > 1.5~M_{\rm Jup}$ orbiting single stars, there are already indications of a history of significant dynamical perturbations. For example, HD~118203b, HD~68988b and HIP~14810b all have anomalously high eccentricities that may be indicative of additional perturbing companions, perhaps with masses below (or periods longer than) the threshold of immediate radial velocity detection (this is certainly the case for HD~68988b and HIP~14810b, which both have long-period planetary companions \citealt{2006ApJ...646..505B}). Due to mutual perturbations, the eccentricities of the bodies in the precursor system may have grown to the point where crossing orbits were achieved. Repeated close encounters among the planets would have then generated a period of chaotic evolution that typically terminates with the ejection of one planet on a hyperbolic trajectory \citep{2002Icar..156..570M}. Alternately, a stellar companion could also effectively trigger dynamical evolution or instability in a precursor system, eventually leading to the current configuration (for some examples, see \citealt{2007A&A...467..347M,2007A&A...472..643M,2003ApJ...589..605W}). With reference to a stellar companion, a quick inspection of POSSI, POSSII, and 2MASS images do not reveal any clear association between faint field stars and HD~17156. The only potentially interesting source is 2MASS~02494068+7144583. It shows an appreciable proper motion, but in the opposite direction of the proper motion of HD~17156. The star lies 22.2\arcsec~from HD~17156; if they are at the same distance, the apparent separation is $\sim$1740 AU. A combination of continued radial velocity monitoring of HD~17156, in conjunction with accurate measurements of successive transit midpoints, gives hope for the detection and accurate characterization of additional bodies in the system via a novel set of constraints. | 7 | 10 | 0710.0898 |
0710 | 0710.5473_arXiv.txt | We identify a large sample of isolated bright galaxies and their fainter satellites in the 2dF Galaxy Redshift Survey (2dFGRS). We analyse the dynamics of ensembles of these galaxies selected according to luminosity and morphological type by stacking the positions of their satellites and estimating the velocity dispersion of the combined set. We test our methodology using realistic mock catalogues constructed from cosmological simulations. The method returns an unbiased estimate of the velocity dispersion provided that the isolation criterion is strict enough to avoid contamination and that the scatter in halo mass at fixed primary luminosity is small. Using a maximum likelihood estimator that accounts for interlopers, we determine the satellite velocity dispersion within a projected radius of 175~\hkpc. The dispersion increases with the luminosity of the primary and is larger for elliptical galaxies than for spiral galaxies of similar b$_{\rm J}$ luminosity. Calibrating the mass-velocity dispersion relation using our mock catalogues, we find a dynamical mass within 175~\hkpc\ of $\rm M_{175}/h^{-1}M_\odot \simeq 4.0^{+2.3}_{-1.5} \times 10^{12}\,(L_{\bjm}/L_*)$ for elliptical galaxies and $\rm M_{175}/h^{-1}M_\odot \simeq 6.3^{+6.3}_{-3.1} \times 10^{11}\,(L_{\bjm}/L_*)^{1.6}$ for spiral galaxies. Finally, we compare our results with recent studies and investigate their limitations using our mock catalogues. | The view that galaxies are surrounded by large dark matter halos dates back more than 30 years to the pioneering study of the rotation curve of M32 by Rubin \& Ford (1970). Extended galactic halos are, in fact, a generic feature of the cold dark matter model of galaxy formation (Blumenthal et al 1984, Frenk et al 1985), but this fundamental theoretical prediction has limited observational support. Zaritsky \etal~(1993) attempted to measure the mass and extent of dark matter halos by analysing the dynamics of satellite galaxies found around `isolated' galaxies. Since galaxies generally have only a few detectable satellites, they used a method that consists of stacking satellites in a sample of primaries of similar luminosity. In spite of the small size of their relatively inhomogeneous sample, Zaritsky \etal~(1993) were able to detect massive halos around isolated spiral galaxies extending to many optical radii. Having nearly doubled their satellite sample to 115 members, Zaritsky \etal (1997b) confirmed their earlier claims including a puzzling lack of correlation between the velocity dispersion of the (stacked) satellite system and the luminosity of the primary. More recently, McKay \etal~(2002) performed a similar analysis on data from the Sloan Digital Sky Survey (SDSS; York \etal~2000). They compared mass estimates derived from satellite dynamics to those derived from weak lensing analyses of the same data (McKay \etal 2001). With a much larger sample than that of Zaritsky et al (1997), they were able to detect a correlation between satellite velocity dispersion and primary luminosity. This trend was confirmed by Prada \etal~(2003) who also used SDSS data. Although they are both based on SDSS data, these two studies find results that, while consistent at first sight, are in, fact, somewhat contradictory. For example, although Prada et al (2003) found a strong dependence of satellite velocity dispersion on galactrocentric distance, their measured velocity dispersion within a radius of 125~\hkpc\ is similar to the values obtained by McKay \etal~(2002) at a radius of 275~\hkpc. Discrepant results were also found by Brainerd \& Specian (2003) who applied the same technique to the early, ``100k'' data release of the 2degree-Field Galaxy Redshift Survey (2dFGRS; Colless 2001) and derived satellite velocity dispersions which are in qualitative and quantitative disagreement with those of Zaritsky \etal (1997), McKay \etal~(2002) and Prada \etal~(2003). A more extensive analysis of the complete 2dFGRS (Colless \etal~2003) by Brainerd (2005) also led to disagreements with the results of earlier work. This somewhat confused picture of satellite dynamics is due in large part to different choices of primary galaxy samples and to differences in the modelling and analysis methods. This paper has multiple aims. Firstly, we carry out a new analysis of the dynamics of satellites around bright galaxies of different morphological types selected from the full 2dFGRS. The goal is to constrain the velocity dispersion and mass of their dark matter halos and we therefore select a sample of isolated galaxies chosen according to strict criteria. Secondly, we investigate the reliability and accuracy of commonly used dynamical analysis methods. For this, we make extensive use of realistic mock catalogues constructed from large cosmological N-body simulations and different semi-analytic galaxy formation models (Cole \etal~2000; \munichsambrak). A similar approach, but in a different context, was adopted by \vdbpapa. Finally, we attempt to understand the root cause of the differences found in previous work, again relying on the use of realistic mock catalogues. The paper is organised as follows. In Section~\ref{sec:data}, we briefly present some of the characteristics of the 2dFGRS data and simulations used for our analysis. In Section~\ref{sec:sample}, we describe the satellite sample selection scheme, together with its robustness to changes in the selection parameters. The analysis of the 'stacked' satellite velocity distribution is carried in Section~\ref{sec:anal} while, in Section~\ref{sec:vel_disp_est}, we present velocity dispersion estimates for our mock catalogues and for 2dFGRS primaries split according to luminosity and morphological type. Using a model for the relationship between dark halo mass and satellite velocity dispersion, we give, in Section~\ref{sec:mass_est}, an estimate of the mass of the halos around 2dFGRS galaxies. In Section~\ref{sec:discussion}, we compare our results with those of previous studies, and we conclude in Section~\ref{sec:conclusion}. | \label{sec:conclusion} We have developed, tested and applied a method to probe the properties of extended dark matter haloes around bright galaxies. We do this by carefully selecting isolated galaxies in the 2dFGRS and using their faint satellites as tracers of the gravitational potential. By stacking systems of similar primary luminosity to improve the signal-to-noise, we estimate the satellite velocity dispersion. Realistic mock galaxy catalogues, created from cosmological N-body simulations populated with a semi-analytic galaxy formation scheme, enable us to relate the measured velocity dispersion of the satellite system to the velocity dispersion and mass of the underlying dark halo. Fig.~\ref{fig:sigma_amag_morph} shows evidence for the existence of dark matter halos around typical galaxies. Our sample of satellites probes the potential well of the primaries out to several hundred kiloparsecs and demonstrates that the dark halos extend many times beyond the optical radius of the primary. This is in agreement with the current theoretical picture of galaxy formation in a cold dark matter universe (e.g. White \& Frenk 1991, Kauffmann \etal~1993, Cole \etal~2000). The satellite velocity dispersion increases with the luminosity of the primary and is much larger for elliptical galaxies than for spiral galaxies of similar b$_{\rm J}$ luminosity. The total extent of the dark halo is not constrained by our data. Although, the satellite distribution extends to $r_p \sim$~375~\hkpc, most of the signal comes from within $r_p \sim$~175~\hkpc. Within the errors, the velocity dispersion appears to be constant within $r_p \sim$~175~\hkpc\ and $r_p \sim$~375~\hkpc. In this range, the satellite velocity dispersion does not depend strongly on the luminosity of the primary. Our mock catalogues allow us to calibrate the velocity dispersion-mass relation for galaxies selected according to the isolation criterion of our 2dFGRS sample. Fig.~\ref{fig:amag_mass} then indicates that elliptical galaxies reside in haloes which are at least 4 times more massive than spiral galaxies of similar b$_{\rm J}$ brightness. Galaxy like the Milky-Way typical reside in dark matter halos of mass $\rm \sim 3.5^{+4.0}_{-2.1}\,10^{11}$~\hmsol\ within 175~\hkpc. A key assumption in our analysis is that isolated galaxies of similar luminosity reside in halos of similar mass. It is this that justifies the stacking procedure. Our semi-analytic models allow us to test the validity of this assumption. We find that in one of two semi-analytic models that we have investigated, the \munichsam\ model, there is very little scatter in the relation between central galaxy luminosity and halo mass. A different semi-analytic model, however, that by Cole \etal~(2000), predicts considerable scatter in this relation and this introduces large errors in the halo properties inferred from a stacking analysis. Mocks constructed from this model return an increasing satellite velocity dispersion as a function of primary luminosity which, however, deviates systematically from the velocity dispersion of the host dark halos. The reasons behind this difference in the galaxy formation models are not investigated in detail here but it serves to illustrate that significant theoretical uncertainties remain in the kind of analysis that we have presented here. | 7 | 10 | 0710.5473 |
0710 | 0710.5645_arXiv.txt | We describe the development and implementation of the SEGUE (Sloan Extension for Galactic Exploration and Understanding) Stellar Parameter Pipeline (SSPP). The SSPP derives, using multiple techniques, radial velocities and the fundamental stellar atmospheric parameters (effective temperature, surface gravity, and metallicity) for AFGK-type stars, based on medium-resolution spectroscopy and $ugriz$ photometry obtained during the course of the original Sloan Digital Sky Survey (SDSS-I) and its Galactic extension (SDSS-II/SEGUE). The SSPP also provides spectral classification for a much wider range of stars, including stars with temperatures outside of the window where atmospheric parameters can be estimated with the current approaches. This is Paper I in a series of papers on the SSPP; it provides an overview of the SSPP, and initial tests of its performance using multiple data sets. Random and systematic errors are critically examined for the current version of the SSPP, which has been used for the sixth public data release of the SDSS (DR-6). | The Sloan Extension for Galactic Understanding and Exploration (SEGUE) is one of three surveys that are being executed as part of the current extension of the Sloan Digital Sky Survey (SDSS-II), which consists of LEGACY, SUPERNOVA SURVEY, and SEGUE. The SEGUE program is designed to obtain $ugriz$ imaging of some 3500 square degrees of sky outside of the original SDSS-I footprint (Fukugita et al. 1996; Gunn et al. 1998, 2006; York et al. 2000; Stoughton et al. 2002; Abazajian et al. 2003, 2004, 2005; Pier et al. 2003; Adelman-McCarthy et al. 2006, 2007a). The regions of sky targeted are primarily at lower Galactic latitudes ($|b|~< ~35^{\circ}$), in order to better sample the disk/halo interface of the Milky Way. As of Data Release 6 (DR-6, Adelman-McCarthy et al. 2007b), about 85\% of the planned additional imaging has already been completed. SEGUE is also obtaining $R$ $\simeq$ 2000 spectroscopy, over the wavelength range $3800-9200$\, {\AA}, for some 250,000 stars in 200 selected areas over the sky available from Apache Point, New Mexico. The spectroscopic candidates are selected on the basis of $ugriz$ photometry to populate roughly 16 target categories, chosen to explore the nature of the Galactic stellar populations at distances from 0.5 kpc to over 100 kpc from the Sun. Spectroscopic observations have been obtained for roughly half of the planned targets thus far, a total of about 120,000 spectra. The SEGUE data clearly require automated analysis tools in order to efficiently extract the maximum amount of useful astrophysical information for the targeted stars, in particular their stellar atmospheric parameters, over wide ranges of effective temperature ($T_{\rm eff}$), surface gravity (log $g$), and metallicity ([Fe/H]). Numerous methods have been developed in the past in order to extract atmospheric-parameter estimates from medium-resolution stellar spectra in a fast, efficient, and automated fashion. These approaches include techniques for finding the minimum distance (parameterized in various ways) between observed spectra and grids of synthetic spectra (e.g., Allende Prieto et al. 2006), non-linear regression models (e.g., Re Fiorentin et al. 2007, and references therein), correlations between broadband colors and the strength of prominent metallic lines, such as the Ca~II K line (Beers et al. 1999), auto-correlation analysis of a stellar spectrum (Beers et al. 1999, and references therein), obtaining fits of spectral lines (or summed line indices) as a function of broadband colors (Wilhelm et al. 1999), or the behavior of the Ca~II triplet lines as a function of broadband color (Cenarro et al. 2001a,b). However, each of these approaches exhibits optimal behavior only over restricted temperature and metallicity ranges; outside of these regions they are often un-calibrated, suffer from saturation of the metallic lines used in their estimates at high metallicity or low temperatures, or lose efficacy due to the weakening of metallic species at low metallicity or high temperatures. The methods that make use of specific spectral features are susceptible to other problems, e.g., the presence of emission in the core of the Ca~II K line for chromospherically active stars, or poor telluric line subtraction in the region of the Ca~II triplet. Because SDSS stellar spectra cover most of the entire optical wavelength range, one can apply several approaches, using different wavelength regions, in order to glean optimal information on stellar parameters. The combination of multiple techniques results in estimates of stellar parameters that are more robust over a much wider range of $T_{\rm eff}$, log $g$, and [Fe/H] than those that might be produced by individual methods. In this first paper of a series, we describe the SEGUE Stellar Parameter Pipeline (hereafter, SSPP), which implements this ``multi-method'' philosophy. We also carry out a number of tests to assess the range of stellar atmospheric parameter space over which the estimates obtained by the SSPP remain valid. The second paper in the SSPP series (Lee et al. 2007b; hereafter Paper II) seeks to validate the radial velocities and stellar parameters determined by the SSPP by comparison with member stars of Galactic three globular clusters (M~15, M~13, and M~2) and two open clusters (NGC~2420 and M~67). A comparison with an analysis of high-resolution spectra for SDSS-I/SEGUE stars is presented in the third paper in this series (Allende Prieto et al. 2007; hereafter Paper III). Section 2 describes determinations of radial velocity used by the SSPP. The procedures used to obtain an estimate of the appropriate continuum, and the determination of line indices, is explained in \S 3. The methods that the SSPP employs for determining stellar parameters are presented in \S 4. Section 5 addresses the determinations of auxiliary estimates of effective temperature, based on both theoretical and empirical approaches; these methods are used for stars where adequate estimates of $T_{\rm eff}$ are not measured by our primary techniques. A decision tree that gathers the optimal set of parameter estimates based on the multiple methods is introduced in \S 6. Section 7 summarizes the conditions for raising various flags to warn potential users where uncertainties in parameter determinations remain. In \S 8, we discuss validation of the parameters determined by the SSPP based on SDSS-I/SEGUE stars for which we have obtained higher dispersion spectroscopy on various large telescopes, and also compare the parameters with those of likely member stars of Galactic open and globular clusters. Assignments of approximate MK spectral classifications are described in \S 9. In \S 10 we describe several methods (still under testing) for the determination of distance estimates used by the SSPP. Section 11 presents a summary and conclusions. In the following, the colors ($u-g$, $g-r$, $r-i$, $i-z$, and $B-V$) and magnitudes ($u,~g, ~r,~i,~z,~B,~\rm and~V$) are understood to be de-reddened and corrected for absorption (using the dust maps of Schlegel et al. 1998), unless stated specifically otherwise. | We have described the development and execution of the SEGUE Stellar Parameter Pipeline (SSPP), which makes use of multiple approaches in order to estimate the fundamental stellar atmospheric parameters (effective temperature, $T_{\rm eff}$, surface gravity, log $g$, and metallicity, parameterized by [Fe/H]) for stars with spectra and photometry obtained during the course of the original Sloan Digital Sky Survey (SDSS-I) and its current extension (SDSS-II/SEGUE). The use of multiple approaches allows for an empirical determination of the internal errors for each derived parameter, based on the range of the reported values from each method. From consideration of about 140,000 spectra of stars obtained during SDSS-I and SEGUE that have derived stellar parameters available in the range 4500~K $\le$ $T_{\rm eff}$ $\le$ 7500~K, typical internal errors obtained by the SSPP are $\sigma(T_{\rm eff})$ = 73~K (s.e.m), $\sigma(\log~g)$ = 0.19 (s.e.m), and $\sigma([\rm Fe/\rm H])$ = 0.10 (s.e.m). Paper III points out that the internal scatter estimates obtained from averaging the multiple estimates of the parameters produced by the SSPP underestimate the external errors, owing to the fact that several methods in the SSPP use similar same parameter indicators and atmospheric models. The results of a comparison with an average of two different high-resolution spectroscopic analyses of 124 SDSS-I/SEGUE stars suggests that the SSPP is able to determine $T_{\rm eff}$, log $g$, and [Fe/H] to precisions of 135~K, 0.26 dex, and 0.21 dex, respectively, after combining small systematic offsets quadratically for stars with 4500~K $\leq T_{\rm eff} \leq$ 7500~K. These errors differ slightly from the those obtained by Paper III ($\sigma(T_{\rm eff})$ = 130~K, $\sigma(\log g)$ = 0.21 dex, and $\sigma([\rm Fe/\rm H])$ = 0.11 dex), even though they share a common set of high-resolution calibration observations. This arises because Paper III derived the external uncertainties of the SSPP only taking into account the stars observed with the HET (on the grounds of internal consistency). The sample referred to as OTHERS in Paper III exhibits somewhat larger scatter in its parameters, when compared with those determined by the SSPP. Observation of several hundred additional stars from SDSS-I/SEGUE with HET is now underway. Thus, in the future, we will be able to use a homogeneous sample gathered by HET in our tests. Also, additional high-resolution data for stars outside of our adopted temperature range will enable tests for both cooler and warmer stars. Considering the internal scatter from the multiple approaches and the external uncertainty from the comparisons with the high-resolution analysis together, the typical uncertainty in the stellar parameters delivered by the SSPP are $\sigma(T_{\rm eff})$ = 154~K, $\sigma(\log~g)$ = 0.31 dex, and $\sigma([\rm Fe/\rm H])$ = 0.23 dex, over the temperature range 4500~K $\leq T_{\rm eff} \leq$ 7500~K. However, it should be kept in mind that the errors stated above apply for the very highest $S/N$ spectra obtained from SDSS ($S/N >$ 50/1), as only quite bright stars were targeted for high-resolution observations. In addition, outside of the quoted temperature range (4500~K $\le$ $T_{\rm eff}$ $\le$ 7500~K), we presently do not have sufficient high-resolution spectra to fully test the parameters obtained by the SSPP. The results of a comparison with likely member stars of a sample of Galactic open and globular clusters suggest that SSPP may slightly overestimate [Fe/H] (by $\sim 0.15$ dex) for stars with [Fe/H] $< -$2.0, and underestimate [Fe/H] (by $\sim 0.30$ dex) for stars with near-solar metallicities. Slight trends of [Fe/H] with $g-r$ are noticed for the higher metallicity clusters as well, although further data will be needed in order to verify this. Approximate spectral types are assigned for stars, based on two methods, with differing limitations. A preliminary set of distance determinations for each star is also obtained, although future work will be required in order to identify the optimal method. We conclude that the SSPP determines sufficiently accurate and precise radial velocities and atmospheric parameter estimates, at least for stars in the effective temperature range from 4500~K to 7500~K, to enable detailed explorations of the chemical compositions and kinematics of the thick-disk and halo populations of the Galaxy. | 7 | 10 | 0710.5645 |
0710 | 0710.2247_arXiv.txt | {The deeper and more extended survey of the central parts of the Galactic Plane by H.E.S.S. during 2005-2007 has revealed a number of new point-like, as well as, extended sources. Two point-like sources can be associated to two remarkable objects around ``Crab-like'' young and energetic pulsars in our Galaxy : G21.5-0.9 and Kes~75. The characteristics of each of the sources are presented and possible interpretations are briefly discussed.} \begin{document} | \begin{figure*}[!t]% \begin{center} \includegraphics*[width=0.46\textwidth,angle=0,clip]{g215exc008Fig1.eps} \includegraphics*[width=0.46\textwidth,angle=0,clip]{J1846m029exc008Fig1.eps} \end{center} \caption{ Smoothed excess maps ($\sigma=0.08^{\circ}$) of the 0.5$^{\circ}\times0.5^{\circ}$ field of view around the positions of HESS~J1833-105 (left) and HESS~J1846-029 (right). The white contours show the pre-trials significance levels for 4, 5, 6 $\sigma$, and 7, 8, 9 $\sigma$, respectively. The black triangle marks the position of the pulsars. The best-fit positions of the two sources are marked with an error cross (for HESS~J1846-029 the latter overlaps with the triangle).} \label{skymaps} \end{figure*} The standard candle of VHE astronomy, the Crab Nebula, has served for decades as a yardstick in almost all wavelengths, and yet it is a very peculiar object, harbouring the most energetic and one of the youngest pulsars of our Galaxy. Since the early days, where the similarities of the historical trio Crab/3C~58/G~21.5-0.9 were under debate~\cite{WilsonWeiler76}, radio and X-ray astronomy have provided a wealth of information by detecting and characterizing nebulae around rotation-powered pulsars. In the VHE domain, H.E.S.S. has revealed more than a dozen pulsar wind nebulae (PWN), either firmly established as such or compelling candidates~\cite{HESSpwnICRC07}, almost all of which are middle-aged (at least few kyrs up to $\sim$100 kyrs, except MSH~15-52) and exhibit an offset between the pulsar position and the nebula center. We report here on the VHE emission discovery of two remarkable objects, G~21.5-0.9 and Kes75, which also harbor very young and energetic pulsars and which on some aspects, especially their plerionic nebular emission due to an energetic pulsar, can be considered as Crab-like. \begin{figure*}[tb] \begin{center} \includegraphics[width=0.48\textwidth]{g21.5m09on011Pwl1Z.eps} \includegraphics[width=0.48\textwidth]{J1846on011Pwl1Zdec.eps} \end{center} \caption{Differential energy spectra above for HESS~J1833-105 (left) and HESS~J1846-029 (right). The shaded area shows the 1 $\sigma$ confidence region for the fit parameters.} \label{spectra} \end{figure*} G21.5-0.9 \cite{Altenhoff70}, recently revealed as a composite SNR consisting of a centrally peaked PWN and a 4{\arcmin} shell~\cite{Bocchino2005,MathesonSafi-Harb2005}, was previously classified as one of the about ten Crab-like SNR~\cite{Green2004}. Its flat spectrum PWN is polarised in radio~\cite{BeckerSzymkowiak1981} with a spectral break above 500 GHz~\cite{GallantTuffs1998}. The non-thermal X-ray PWN with radius $\sim$40{\arcsec} shows significant evidence of cooling ~\cite{Slane2000}, with the power-law photon index steepening from 1.43$\pm$0.02 near the pulsar to 2.13$\pm$0.06 at the edge of the PWN. There appears to be a synchrotron X-ray halo at a radius of 140{\arcsec} from the pulsar which could originate in the shell~\cite{Bocchino2005,MathesonSafi-Harb2005}, with a contribution of scattering off dust grains as proposed by Bocchino et al.~\cite{Bocchino2005}. The 61.8 ms pulsar PSR~J1833-1034, with a spin-down power of ${\dot E} = 3.3\times 10^{37}$erg/s and a characteristic age of 4.9 kyr was discovered only recently through its faint radio pulsed emission~\cite{Gupta2005,Camilo2006}. Given the derived distance of 4.7$\pm$0.4 kpc, the age of G~21.5-0.9 was revised downwards by a factor of $\sim$10 to force consistency with the freely expanding SNR shell~\cite{Camilo2006}. PSR~J1833-1034 in G~21.5-0.9 is the second most energetic pulsar known in the Galaxy. Kes 75 (SNR G29.70.3) is also a prototypical example of a composite remnant for which the distance of 19 kpc was estimated through neutral hydrogen absorption measurements~\cite{BeckerHelfand1984}. Its 3.5{\arcmin} radio shell surrounds a flat-spectrum highly polarized radio core, and harbors, at its center, the 325 ms X-ray Pulsar, PSR~J1846-258~\cite{Gotthelf2000}. The latter has the shortest known characteristic age $\tau_c= 723$ yr and a large inferred magnetar-like magnetic field of B$=4.9\times 10^{13}$G. The pulsar lies within a 25{\arcsec}$\times$20{\arcsec} X-ray nebula which exhibits an photon index of 1.92$\pm$0.04, but no evidence of cooling as a function of the distance to the pulsar. Like in G~21.5-0.9 there is an X-ray halo, in this case due mostly to dust scattering, but a non-thermal contribution from electrons accelerated in the shell remains possible~\cite{Helfand2003}. | It is remarkable that de Jager et al.~\cite{deJager1995} predicted that plerionic VHE $\gamma$-rays from G21.5-0.9 would be detectable at a level of $4\times10^{-13}$ cm$^{-2}$~s$^{-1}$ at 1~TeV with an electron spectral index of $\sim$2.8, which would give a photon index near 2.0 at VHE energies (after including KN effects given the contributions from dust and CMBR). Their prediction was based on an assumed equipartition field strength of 22 $\mu$G which is close to the value of $\sim$15$\mu$G implied from $\gamma$-ray observations reported here (assuming IC scattering on CMB photons only, and using the ratio of the X-ray to the $\gamma$-ray luminosities: $L_{\rm X}/L_{\gamma}\sim30$). The equipartition field strength was afterwards increased to 0.3 mG following the revision of the maximum spectral range of the radio PWN to 500 GHz \cite{GallantTuffs1998, Camilo2006}. However, the detection of VHE $\gamma$-rays by H.E.S.S. from PWN tends to confirm the suggestion of Chevalier \cite{Chevalier2004} that some PWN may be particle dominated, so that the true PWN field strength may be significantly lower than equipartition for some objects. In the case of Kes~75, $L_{\rm X}/L_{\gamma}\sim$10 yields also a lower than equipartition nebular magnetic field strength of $\sim$10 $\mu$G. It should be noted that Kes~75 shows the highest conversion efficiency in X-rays ($\sim15$\%) as compared to other ``Crab-like'' pulsars ($\sim3$\% and $\sim0.6$\% for the Crab and G~21.5-0.9, respectively) and a 100 times larger $\gamma$-ray efficiency ($\sim$2\%) than the Crab and G~21.5-0.9 which are similar in that respect ($\sim$0.02\%). However, the latter object's $L_{\rm X}/L_{\gamma}\sim30$ is 4 times smaller than that of the Crab Nebula $L_{\rm X}/L_{\gamma}\sim120$. These numbers together with the spin parameters and high surface magnetic field in the case of PSR~J1846-0258, show that these objects, although ``Crab-like'' in some aspects, do possess peculiar properities. Given the evidence for synchrotron emission in the SNR shell, an alternative interpretation of the VHE emissions of G~21.5-0.9 and Kes75 would be radiation from particles accelerated at the non-relativistic forward shock of the freely expanding SNR. However the required field strength in the shell to explain the H.E.S.S. detection in terms of IC scattering should be much lower than 10 $\mu$G, value which may be unreasonably low for typical expanding SNR shells. Deeper observations of both sources could help to constrain the size of the VHE emission region and to ascertain whether it is compatible with this scenario. | 7 | 10 | 0710.2247 |
0710 | 0710.0418_arXiv.txt | We are investigating mass fractions on the crust of a neutron star which would remain after one year of cooling. We use cooling curves corresponding with various densities, or depths, of the neutron star just after its formation. We assume the modified Urca process dominates the energy budget of the outer layers of the star in order to calculate the temperature of the neutron star as a function of time. Using a nuclear reaction network up to technetium, we calculate how the distribution of nuclei quenches at various depths of the neutron star crust. The initial results indicate that $^{28}$Si is the lightest isotope to be optically thick on the surface after one year of cooling. | The observed emission from a neutron star passes through a crustal layer of the neutron star. In order to fully interpret the observed emission from a neutron star, we need to understand what comprises this crustal layer of the neutron star \cite{Hern84b,heylmnras01}. Here, the mass fractions on the surface are calculated for what would exist on the crust after one year of cooling of the neutron star. These mass fractions are calculated using a 489 isotope reaction network which burns up to technetium written by F.X. Timmes \cite{timmesapjs99} and is available on Timmes's webpage \footnote{\url{http://www.cococubed.com/code_pages/burn.shtml}}. | We have examined the mass fractions that will result in an optically thick surface layer by looking at two cases: a density of $10^{12}\;$g/cm$^3$ and another of $10^{7}\;$g/cm$^3$. The mass fractions at these densities are calculated using a 489 isotope reaction network which burns up to technetium. The neutron star is cooled for a year, starting with a temperature of $10^{10}$K, assuming the modified Urca process dominates. Upon calculating the mass fractions the minimum mass fraction abundance required for enough of an isotope to form an optically thick layer such that it has a surface density of 1g/cm$^2$ is calculated. From the two cases we found that the deeper density resulted in a larger variety of lighter isotopes can form an optically thick layer. We found $^{28}$Si, $^{30}$Si, $^{31}$P, $^{32}$S, $^{33}$S, $^{34}$S, and $^{36}$Ar rise to the surface. Future work on this includes exploring more densities, comparing to analytic calculations and estimating the timescale for the light impurities to float to the surface. \begin{theacknowledgments} We would like to thank Ed Brown for helpful discussions during the conference. The Natural Sciences and Engineering Research Council of Canada, Canadian Foundation for Innovation and the British Columbia Knowledge Development Fund supported this work. This research has made use of NASA's Astrophysics Data System Bibliographic Services. \end{theacknowledgments} | 7 | 10 | 0710.0418 |
0710 | 0710.5521_arXiv.txt | I calculate the linear stability of a stratified low collisionality plasma in the presence of a weak magnetic field. Heat is assumed to flow only along magnetic field lines. In the absence of a heat flux in the background plasma, Balbus (2000) demonstrated that plasmas in which the temperature {\it increases} in the direction of gravity are buoyantly unstable to convective-like motions (the ``magnetothermal instability''). I show that in the presence of a background heat flux, an analogous instability is present when the temperature {\it decreases} in the direction of gravity. The instability is driven by the background heat flux and the fastest growing mode has a growth time of order the local dynamical time. Thus, independent of the sign of the temperature gradient, weakly magnetized low collisionality plasmas are unstable on a dynamical time to magnetically-mediated buoyancy instabilities. The instability described in this paper is predicted to be present in clusters of galaxies at radii $\sim 0.1-100$ kpc, where the observed temperature increases outwards. Possible consequences for the origin of cluster magnetic fields, ``cooling flows,'' and the thermodynamics of the intercluster medium are briefly discussed. | \label{sec:intro} Thermally stratified fluids are buoyantly unstable when the entropy increases in the direction of gravity, a result of considerable importance to the theory of stellar structure (Schwarzschild 1958). Remarkably, however, this well-known result changes in a low collisionality plasma in which i) the collisional mean free path of electrons is larger than the electron Larmor radius and ii) thermal conduction is the dominant mode of heat transport (Balbus 2000). In such a plasma, heat is transported primarily along magnetic field lines. For the simple problem of a horizontal magnetic field in a vertically stratified plasma, Balbus (2000) showed that the condition for the plasma to be buoyantly unstable becomes that the temperature (not entropy) increase in the direction of gravity. The resulting ``magnetothermal instability'' (MTI) has been studied with nonlinear simulations by Parrish \& Stone (2005; 2007). In a subsequent paper, Balbus (2001) generalized his initial result to rotating flows and magnetic fields of arbitrary orientation, but still under the assumption that there is no heat flux in the background plasma (i.e., that the field lines are initially isothermal). This latter assumption is unlikely to hold in many low collisionality astrophysical plasmas such as clusters of galaxies and hot accretion flows onto black holes. In this paper I extend Balbus's calculation and study the stability of weakly magnetized plasmas in the presence of a background heat flux. I show that the presence of a heat flux drives a buoyancy instability analogous to the MTI when the temperature {\it decreases} in the direction of gravity (a situation that is MTI stable according to Balbus's analysis). This instability is distinct from the heat flux driven overstabilities described in Socrates, Parrish, \& Stone (2007).\footnote{In my analysis below, I utilize the Boussinesq approximation to focus on nearly incompressible perturbations. In this limit, Socrates et al. predict that the slow mode is stable while the fast mode is unstable on a dynamical time. Because the Boussinesq approximation filters out fast waves, I do not expect any version of their overstabilities to be present in my analysis.} In the next two sections I summarize the equations and assumptions used in my analysis (\S \ref{sec:eqns}) and the results of the linear stability calculation (\S \ref{sec:results}). I then discuss possible applications of the heat flux-driven version of the MTI, in particular to the intercluster plasma in clusters of galaxies (\S \ref{sec:disc}). | \label{sec:disc} The above analysis shows that, regardless of the sign of the temperature gradient, a weakly magnetized low collisionality plasma in which heat flows primarily along magnetic field lines is buoyantly unstable. For $dT/dz < 0$, this instability is the magnetothermal instability (MTI) derived by Balbus (2000; 2001) and simulated by Parrish \& Stone (2005; 2007). Although a plasma with $dT/dz > 0$ is MTI stable according to Balbus (2001), I have shown that an analogous buoyancy instability in fact exists for $dT/dz > 0$ in the presence of a vertical magnetic field and a background heat flux. Physically, this new instability arises because perturbed fluid elements are heated/cooled by the background heat flux in such a way as to become buoyantly unstable (\S \ref{sec:new}). In many astrophysical plasmas, the sign of the temperature gradient is fixed to be $dT/dz < 0$ by basic principles. These systems may be MTI unstable, but they will be stable to the new buoyancy instability discussed in this paper. This is typically the case in cooling white dwarfs and neutron stars, where the flow of heat outwards requires $dT/dz < 0$. It also also the case in hot accretion flows onto compact objects, because the inflow of matter and the release of gravitational potential energy drives $dT/dz < 0$. The heat flux driven instability described in this paper may act in the transition region between cool dense gas and hot low density plasma in stellar coronae, accretion disks, and the multi-phase interstellar medium. However, these regions tend to be strongly magnetized, which will inhibit the instability (\S \ref{sec:strongB}). A more promising application is to the hot intercluster plasma in galaxy clusters. Plasma in hydrostatic equilibrium in a Navarro, Frank, \& White (1997) dark matter potential well has a temperature profile which is locally isothermal ($\rho \propto r^{-2}$) at a scale radius $R_s \simeq 100-400$ kpc. The temperature is predicted to decrease for radii both smaller and larger than $\sim R_s$. Such a radial variation in the temperature of the intercluster plasma is directly observed in many systems (e.g., Piffaretti et al. 2005). At radii larger than $\sim R_s$, the intercluster plasma is MTI unstable (e.g., Parrish \& Stone 2007), but it is MTI stable in the cores of clusters where the temperature decreases inwards. However, it is precisely these radii that are unstable to the buoyancy instability discussed in this paper. Thus, provided that the field is not too strong (see \S \ref{sec:strongB}), I conclude that the entire intercluster plasma in galaxy clusters is unstable to magnetically mediated buoyancy instabilities. The implications of this instability for the intercluster medium will be investigated in future papers using nonlinear simulations. Here I briefly comment on the possible consequences. Given the presence of exponentially growing instabilities that amplify magnetic fields {\it at all radii} in galaxy clusters, it is natural to suspect that these instabilities play a significant role in generating the observed (e.g., Federica \& Feretti 2004) $\mu G$ magnetic fields in clusters from smaller cosmological ``seed'' fields. In addition, just as the MTI is found to re-orient the magnetic field to be largely radial, allowing heat to flow down the temperature gradient (Parrish \& Stone 2007; Sharma \& Quataert, in preparation), I suspect that the heat flux driven buoyancy instability discussed here will generate a significant horizontal magnetic field if one was not present originally. This will act so as to decrease the net heat flux through the plasma (which is the origin of the instability in the first place). Given the close connection between the heat flux and magnetic field described in this paper, it is unclear whether current calculations of the effective conductivity and heat flux in cluster plasmas (e.g., Narayan \& Medvedev 2001; Chandran \& Maron 2004) are correct, since they do not capture this {dynamical} coupling. It may thus turn out that thermal conduction from large radii will prove to be less effective than previous authors have suspected (e.g., Bertschinger \& Meiksin 1986; Ruszkowski \& Begelman 2002; Zakamska \& Narayan 2003) at heating ``cooling flow'' cores in clusters. Regardless of the accuracy of this speculation, the results of this paper highlight the need for a proper treatment of the combined effects of thermal conduction and magnetically-mediated buoyancy instabilities on the plasma in galaxy clusters. In future work it will also be interesting to study the dynamics of the heat flux driven instability in the presence of cosmic rays, which may be energetically significant in cluster cores because of a central black hole, and which are known to modify the MTI (Chandran \& Dennis 2006). \ | 7 | 10 | 0710.5521 |
0710 | 0710.1302_arXiv.txt | {We use the techniques of effective field theory in an expanding universe to examine the effect of choosing an excited inflationary initial state built over the Bunch-Davies state on the CMB bi-spectrum. We find that even for Hadamard states, there are unexpected enhancements in the bi-spectrum for certain configurations in momentum space due to interactions of modes in the early stages of inflation. These enhancements can be parametrically larger than the standard ones and are potentially observable in future data. These initial state effects have a characteristic signature in $l$-space which distinguishes them from the usual contributions, with the enhancement being most pronounced for configurations corresponding to flattened triangles for which two momenta are collinear.} \preprint{PI-COSMO-64} \begin{document} | 7 | 10 | 0710.1302 |
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0710 | 0710.5071_arXiv.txt | The Laser Interferometer Space Antenna (LISA) mission will use advanced technologies to achieve its science goals: the direct detection of gravitational waves, the observation of signals from compact (small and dense) stars as they spiral into black holes, the study of the role of massive black holes in galaxy evolution, the search for gravitational wave emission from the early Universe. The gravitational red-shift, the advance of the perihelion of Mercury, deflection of light and the time delay of radar signals are the classical tests in the first order of General Relativity (GR). However, LISA can possibly test Einstein's theories in the second order and perhaps, it will show some particular feature of non-linearity of gravitational interaction. In the present work we are seeking a method to construct theoretical templates that limit in the first order the tensorial structure of some metric fields, thus the non-linear terms are given by exponential functions of gravitational strength. The Newtonian limit obtained here, in the first order, is equivalent to GR. | The extreme difficulties which arise if one tries to draw physically important conclusions from the basic assumptions of Einstein's theory are mainly due to the non-linearity of the field equations. Moreover, the fact that the spacetime topology is not given a priori, and the impossibility to integrate tensors over finite regions cause difficulties unknown in other branches of mathematical physics. Actually in this respect they are not so different from others fields, for example the electromagnetic field, the scalar field, etc., by themselves obey linear equations in a given spacetime, they form a non-linear system when their mutual interactions are taken into account. The distinctive feature of the gravitational field is that it is self-interacting (as the Yang-Mills field): it is non-linear even in the absence of other fields. This is because it defines the spacetime over which it propagates \citep{Hawking}. Linearized gravity is any approximation to General Relativity obtained from $g_{\mu\nu}=g_{\mu\nu}^{(0)}+h_{\mu\nu}$ (where $g_{\mu\nu}^{(0)}$ is any curved background spacetime) in Einstein's equation and retaining only the terms linear in $h_{\mu\nu}$ \citep{Wald}. The weakness of the gravitational field means in the context of general relativity that the spacetime is nearly flat. Small gravitational perturbations in Minkowski space can be treated in the simplest linearized version of General Relativity, \begin{equation} \label{linear} g_{\mu\nu}=\eta_{\mu\nu}+h_{\mu\nu}, \end{equation} as describing a theory of a symmetric tensor field $h_{\mu\nu}$ propagating on at background spacetime. This theory is Lorentz invariant in the sense of Special Relativity. If one wants to obtain a solution of the non-linear equations, it is necessary to employ an iterative method on approximate linear equations whose solutions are shown to converge in a certain neighbourhood of initial surface. It should be possible to avoid some of these difficulties of non-linearity by working in some spacetimes that shall be described in this paper. The proposal is that some metric fields can be separated into the parts carrying the dynamical information and those parts characterizing the coordinates system. In this proposal, terms of the coordinates system will have tensorial structure limited only in the first order. The tensors that will describe $g_{\mu\nu}$ have linear behavior. Naturally, there is a price that must be paid for the linear tensors. The dynamic terms that carry the gravitational strength have exponential structure. The principal idea came from the basic principle that one should interpret (\ref{linear}) as separation between pure mathematical and physical terms in metric field tensor. In spite of the fact that $\eta_{\mu\nu}$ plays a key role as empty flat and background spacetime of the Standard Model in the description of fundamental interactions, this background tensor metric $\eta_{\mu\nu}$ is an object wholly mathematical and entirely geometrical, while $h_{\mu\nu}$ contains the physical information. The strength of gravity is tied in the components of $h_{\mu\nu}$. The proposal of this paper is a working hypothesis to untie the strength of gravity from geometrical tensors. This proposal is valid for a family of metric field tensors $g_{\mu\nu}$, and some basic examples such Newtonian limit and gravitational plane waves of low amplitudes are described. This paper is outlined as follows: in Sec. II, we present the basic mathematic concepts of the (quasi-)idempotent tensors that compose the structure of metric fields approached in this work. In Sec. III, we propose how to link the strength of gravity with the tensors from Sec. II, then is defined a family of exponential metrics. In Sec. IV, we present some examples of these exponential metrics, such as: Yilmaz metric, circularly polarized wave and rotating bodies. In Sec. V, we present exponential metrics (`adjoint metric field') that are non-physical, but which help us to compute Christoffel symbols, and consequently the curvature tensors, Ricci tensors and determinant of metric field. In Sec. VI, we verify the Newtonian limit and also we obtain gravitational waves. In Sec. VII, we present a general conclusion. We assume spacetime $({\cal M},{\bf g})$ to be a ${C^{\infty}4-}$dimen\-sional, globally hyperbolic, pseudo-Riemannian manifold ${\cal M}$ with Lorentzian metric tensor ${\bf g}$ (whose components are $g_{\mu\nu}$) associated with the line element $$ds^2=g_{\mu\nu}(x)dx^{\mu}dx^{\nu},$$ assumed to have signature $(+---)$ \citep{Landau}. Lower case Greek indices refer to coordinates on ${\cal M}$ and take the values $0,1,2,3.$ The relation between the metric field $g_{\mu\nu}$ and the material contents of spacetime is expressed by Einstein's field equation, \begin{equation} \label{field equation} R_{\mu\nu}-\frac{1}{2}g_{\mu\nu}R=\frac{8\pi G}{c^4}T_{\mu\nu}, \end{equation} $T_{\mu\nu}$ being the stress-energy-momentum tensor, $R_{\mu\nu}$ the contract curvature tensor (Ricci tensor) and $R$ its trace. In an empty region of spacetime we have $R_{\mu\nu}=0$, such a region is called vacuum field. | This paper deals with a tensorial structure that assumes a (quasi-)idempotent feature to be able to improve at least the linear tensorial template of some tensor metric fields. It is clear that Einstein's field equations are non-linear, however, with these (quasi-)idempotent tensorial structure, without quadratic tensorial values, the non-linearity becomes more moderate although there is a price to pay. The part that carries the dynamical information, the strength of gravity is tied to the tensorial structure by exponential functions. In this approach the metric field can be characterized by a background spacetime conformally flat affected by a disturbance. We have approached some examples in this tensorial structure that results in exponential metric fields, we can point out as the main exponential metric obtained in this paper which has been extensively explored: the Yilmaz exponential metric \citep{Yilmaz1,Yilmaz2,Yilmaz3,Yilmaz4,Yilmaz5,Yilmaz6,Yilmaz7,Clapp,Robertson1,Robertson2,Ibison}. H. Yilmaz has argued that in his theory, the gravitational field can be quantized via Feynman's method \citep{Yilmaz9,Alley}. Further, it has been found that the quantized theory is finite. Incidentally in the exponential metric fields approached in this work just as in the Yilmaz theory there are no black holes in the sense of event horizons, but there can be stellar collapse \citep{Robertson1,Robertson2}. However, there are no point singularities. Interesting results obtained in this work from exponential metric fields are: circularly polarized wave; rotating bodies that in the first order is a deformation of Kerr metric and also we have a deformed static spherically symmetric spacetime. Many discussions around massive stellar objects have suggested, for example, that Kerr metric should be slightly deviated from Kerr. The possibility of discovering a non-Kerr object should be taken into account when constructing waveform templates for LISA's data analysis tools \citep{Glampedakis,White}. The technological development is ripe enough so much so in the years to come we might be able to test the second order relativist-gravity effects and may lead to answers to some important questions of gravity. In this work, we have obtained a simple and general expression for the volume element of a manifold in coordinates $(t,x,y,z)$ given in terms of strength of gravity and of traces of tensors $\bm \eta$ and $\bm \Upsilon$. It is possible that an analysis of any Lagrangian of field interacting with gravity will become easer. An interesting observation is the spacetime of circularly polarized plane wave, in this spacetime the volume element $\sqrt{-g}\,\,d^4x$ is the same of Minkowski spacetime, in this sense this gravitational radiation obtained from exponential metric field does not modify the volume element of background Minkowski spacetime where this plane wave travels onto. Moreover, it was purposed and verified the Newtonian limit as solution for Einstein's equation, since we can assume that the trace of stress-energy tensor is $T\approx \rho c^2$. Other important solution of Einstein's equation analysed in this paper was the plane gravitational wave for the empty space since we have considered the vanished stress-energy tensor to the first order in $\Phi$. Both solved Einstein's equations for Newtonian limit and plane gravitational wave propagating in the vacuum are cases that the strength of gravity is small, $\Phi\ll 1$. We have analysed the Newtonian limit in the case that $\partial_{\alpha}\bm\Upsilon=0$, and analysed the plane gravitational wave considering the strength of gravity as a constant term, thus we had two independent Ricci tensors, $R_{\mu\nu}^{(1)}$ which $\partial_{\alpha}\bm\Upsilon=0$ (for Newtonian limit) and $R_{\mu\nu}^{(2)}$ which $\partial_{\alpha}\Phi=0$ (for plane gravitational wave). In a forthcoming work, an analysis of Einstein's equations with both non-vanished $\partial_{\alpha}\bm\Upsilon$ and $\partial_{\alpha}\Phi$, will be considered. It is missing a discussion about quantities of physical interest in the solutions of Einstein's equations which describe the exterior and interior gravitational field. Yilmaz has argued the existence of the matter part in the right-hand side of the field equations correspondent to field energy in the exterior. This paper lacks a discussion about the interior and the exterior field energies denoted by a total stress-energy tensor. An analysis about the total stress-energy tensor will be the object of a forthcoming study, where the physical consequences of terms of deformity in Kerr and Schwarzschild solutions could be analysed. We know that the dark energy and the dark matter problems are challenges to modern astrophysics and cosmology; as a typical example, we could mention the galactic rotation curves of spiral galaxies, that probably, indicates the possible failure of both Newtonian gravity and General Relativity on galactic and intergalactic scales. To explain astrophysical and cosmological problems with arguments against dark energy and dark matter many works have been devoted to the possibility that the Einstein-Hilbert Lagrangian, linear in the Ricci scalar R, should be generalized. In this sense, the choice of a generic function $f(R)$ can be derived by matching the data and the requirement that no exotic ingredient have to be added \citep{Allemandi,Barrow,Capozziello1,Capozziello2,Capozziello3,Carroll1,Carroll2,Faraoni,Flanagan,Koivisto,Nojiri1,Nojiri2,Nojiri3,Sotiriou}. This class of theories when linearized exhibits others polarization modes for the gravitational waves, of which two correspond to the massless graviton and others such as massive scalar and ghost modes in $f(R)$ gravity \citep{Bellucci,Bogdanos}. In this way, analyses in any order to $f(R)$ gravity with `exponential metrics' proposed in the present work could give a positive contribution to the debate of astrophysical and cosmological questions. | 7 | 10 | 0710.5071 |
0710 | 0710.0979_arXiv.txt | We have obtained velocity-resolved spectra of the $\hoz ~ (\lambda = 2.1218 \micron)$ emission line at $2\arcsec$ angular resolution (or $\sim 0.08$~pc spatial resolution) in four regions within the central 10 pc of the Galaxy where the supernova-like remnant Sgr A East is colliding with molecular clouds. To investigate the kinematic, physical, and positional relationships between the important gaseous components in the center, we compared the H$_2$ data cube with previously published NH$_3$ data. The projected interaction-boundary of Sgr A East is determined to be an ellipse with its center offset $\sim1.5$~pc from Sgr A* and dimensions of 10.8~pc $\times$ 7.6~pc. This H$_2$ boundary is larger than the synchrotron emission shell but consistent with the dust ring which is believed to trace the shock front of Sgr A East. Since Sgr A East is driving shocks into its nearby molecular clouds, we can determine their positional relationships using the shock directions as indicators. As a result, we suggest a revised model for the three-dimensional structure of the central 10~pc. The actual contact between Sgr A East and all of the surrounding molecular material, including the circum-nuclear disk and the southern streamer, makes the hypothesis of infall into the nucleus and feeding of Sgr A* very likely. | In the central 10 pc of our Galaxy, the Sgr A region contains several characteristic objects; a candidate for super-massive black hole (Sgr A*) of about $4 \times 10^6~\msun$ (see \citealt*{ghe03,sch03} and references therein), a surrounding cluster of stars (the Central cluster), molecular and ionized gas clouds (the circum-nuclear disk (CND) and Sgr A West), supernova remnants (SNR G~359.92-0.09 and Sgr A East). They are surrounded by molecular structures including two giant molecular clouds (GMCs) M-0.02-0.07 and M-0.13-0.08 (also known as the `$50~\kms$ cloud' and the `$20~\kms$ cloud', respectively). In addition to the two GMCs, recent accurate radio observations have resolved several dense and filamentary molecular features around the Sgr A complex; the `molecular ridge', the `southern streamer', the `northern ridge', and the `western streamer' (see Figure~\ref{3D_model} of this paper and Figures~3, 9, and 10 of \citealt*{mcg01} and Figure~14 of \citealt*{her05}). These molecular features are believed to play important roles in feeding the central massive black hole \citep{ho91,coi99,coi00,mcg01}. The interaction between these various components is responsible for many of the phenomena occurring in this complicated and unique portion of the Galaxy. Developing a comprehensive picture of the primary interactions between the components at the Galactic center will also improve our understanding of the nature of galactic nuclei in general. As the complicated morphology of the central 10 pc is being unveiled thanks to the dramatic progress of radio technology, effort is being made to understand whether these features are really associated with the Galactic center or just seen along the line-of-sight in that direction, and to determine the relative positions of them along the line-of-sight, i.e., the three-dimensional (3-D) spatial structure of the Galactic center. Observations of 327~MHz absorption toward Sgr A West definitely place Sgr A East behind Sgr A West (\citealt*{yus87}; see also \citealt*{ped89}). \citet*{mez89} observed ring-shaped 1.3~mm dust emission surrounding Sgr A East across the $50~\kms$ cloud and the $20~\kms$ cloud, and suggested that Sgr A East has expanded into these molecular clouds. Based on these observational arguments, \citet{mez89} proposed a 3-D structure of the Sgr A complex and concluded that the event which created Sgr A East and the associated dust shell did not occur deep within the GMCs but close to their surfaces facing the sun. However, \citet*{geb89} found CO absorption toward a few Galactic center infrared (IR) sources and suggested that the $20~\kms$ cloud may be located in front of Sgr A West. They also found some evidence that the $50~\kms$ cloud lies partly in front of Sgr A West. Based on the NH$_3$ morphology and kinematics observed using the Very Large Array (VLA), \citet*{coi99,coi00} located the Galactic nucleus (defined to include Sgr A*, Sgr A West, and the CND throughout this paper) behind the southern streamer (and the $20~\kms$ cloud; see also \citealt*{gus80}) and the $50~\kms$ cloud (or the northern part of the molecular ridge) slightly behind Sgr A East. They also argued that the distance between Sgr A East and the $20~\kms$ cloud along the line-of-sight should be smaller than 8.4 pc, which is the size of the SNR G~359.92-0.09 in 20~cm radio continuum images \citep*{yus87,ped89}. \citet*{her05} updated and modified the 3-D model of \citet{coi00} based on their additional NH$_3$ line data and more recently published results (\citealt*{mae02} and references therein; \citealt{par04}), as follows. The nuclear region is placed just inside the leading edge of Sgr A East. Only some part of the $50~\kms$ cloud is located in front of the nucleus. The western streamer seen in NH$_3$ emission is highly inclined to the line-of-sight and is expanding outward with Sgr A East. The northern ridge is placed along the northern edge of Sgr A East and is expanding perpendicular to the line-of-sight. The southern streamer passes over the nucleus in projection but probably does not interact with it. Together, the 3-D models above agree on the following features: \begin{enumerate} \item The Galactic nucleus lies in front of Sgr A East but behind the southern streamer and part of the $20~\kms$ cloud along the line-of-sight. \item Sgr A East is expanding into the $50~\kms$ cloud, the northern ridge, and the western streamer. \item SNR G~359.92-0.09 is colliding with the southern part of the molecular ridge, the eastern edge of the $20~\kms$ cloud, and the southern edge of Sgr A East. \end{enumerate} On the other hand, contradictions among the models raise the following questions. \begin{enumerate} \item Is the nucleus in contact with or contained within Sgr A East? \item Is the southern streamer falling into the nucleus? \item Has Sgr A East expanded into the $50~\kms$ cloud significantly, or just started to contact it? \item Is Sgr A East colliding with the northern part of the molecular ridge? \item Is Sgr A East in contact with the $20~\kms$ cloud? \item Is the $20~\kms$ cloud located only in front of Sgr A East, or also extended further to the backside of it along the line-of-sight? \end{enumerate} It should be noted that the models above are all based on indirect evidence like morphology, kinematics of molecular clouds, or absorption of background radiation by these clouds, rather than on direct, physical interactions between the objects. To answer some of the above questions directly, we have observed molecular hydrogen (H$_2$) emission and constructed a 3-D picture of the Galactic center. H$_2$ emission is an excellent tracer of interactions between dense molecular clouds and other hot and powerful objects, like Sgr A East. In this paper we report observations of H$_2$ line emission from regions of interaction between Sgr A East and other gaseous components within the central 10 pc. Our observations were almost entirely of the H$_2$ 1-0 S(1) line. Unlike most previous work, we observed this line at sufficiently high spectral resolution to resolve the velocity profiles and at high enough angular resolution to obtain detailed information on the spatial structure of the emission. We also obtained measurements of the H$_2$ 2-1 S(1) line at 2.2477~$\mu$m at one location in order to investigate the excitation mechanism of the H$_2$. We describe the observations in Section~\ref{observation} and the reduction of the spectroscopic data in Section~\ref{reduction}. Based on the directions of the shocks derived from the direct comparison of radial velocities with those from the NH$_3$(3,3) data of \citet{mcg01}, we construct a 3-D model for the structure of the central 10~pc in Section~\ref{structure}. In a forthcoming paper we discuss the properties of the shocks and estimate the explosion energy and age of Sgr A East, from which we constrain its origin. | Based on the H$_2$ emission map, we determine the outer boundary of Sgr A East where it is driving shocks into the surrounding molecular clouds, to be approximately an ellipse with the center at ($+32\arcsec$, $+18\arcsec$) or $\sim1.5$~pc offset from Sgr A*, a major axis of length 10.8~pc, which is nearly parallel to the Galactic plane, and a minor axis of length 7.6~pc. This boundary is significantly larger than the synchrotron emission shell \citep*{eke83,yus87,ped89} but is closely consistent with the dust ring suggested by \citet*{mez89}. Since Sgr A East is in physical contact with all of its nearby molecular clouds (the $50~\kms$ cloud, the northern ridge, the molecular ridge, the southern streamer, the CND, and the western streamer), we are able to determine the positional relationships between Sgr A East and the molecular clouds along the line-of-sight using the shock directions as indicators. Based on the determined relationships and the strong evidence that Sgr A East is in contact with the nucleus, we suggest a revised model for the 3-D spatial structure of the central 10 pc of our Galaxy modifying the previous models of \citet*{mez89}, \citet*{coi00}, and \citet*{her05}. Our conclusions on the 3-D structure resolve most of the debates in previous studies as follows: \begin{enumerate} \item Is the nucleus in contact with or contained within Sgr A East? -- The Galactic nucleus is in physical contact with Sgr A East since the CND is pushed toward us by the expanding hot cavity of Sgr A East. \item Is the southern streamer falling into the nucleus? -- It is highly probable that the southern streamer is falling into the nuclear region and feeding the CND. Sgr A East is driving shocks into the northern-most part of this cloud where it meets the CND in projection. \item Has Sgr A East expanded into the $50~\kms$ cloud significantly, or just started to contact it? -- In the H$_2$ data of the northeastern field, we can see that most of this region is filled with shocked gas from the $50~\kms$ cloud. The area corresponds to at least one third of that of the entire cloud. Thus Sgr A East has significantly penetrated the cloud. If the cloud is wrapping around the Sgr A East shell with a casual morphological coincidence, however, the shocked layer might still be thin and on the surface. \item Is Sgr A East colliding with the northern part of the molecular ridge? -- Yes, we detected shocked H$_2$ emission from the northern-most part of this cloud. \end{enumerate} However, we cannot answer the questions related to the $20~\kms$ cloud. We know the branches (the southern streamer and the western streamer) from this GMC are interacting with Sgr A East but the main body of this cloud is located far to the south, beyond the scope of our observations. Therefore the position and extent along the line-of-sight of this GMC in Figure~\ref{3D_model} is uncertain. According to \citet*{coi99,coi00}, SNR G~359.92-0.09 is expanding into the molecular ridge, the $20~\kms$ cloud, and Sgr A East. Thus, if we observe the H$_2$ emission around SNR G~359.92-0.09 in a similar way to the 3-D observations for Sgr A East, the questions about the $20~\kms$ cloud could also be answered. | 7 | 10 | 0710.0979 |
0710 | 0710.5767.txt | In studies of star-forming regions, near-infrared excess (NIRX) sources---objects with intrinsic colors redder than normal stars---constitute both signal (young stars) and noise (e.g. background galaxies). We hunt down (identify) galaxies using near-infrared observations in the Perseus star-forming region by combining structural information, colors, and number density estimates. Galaxies at moderate redshifts (z = 0.1 - 0.5) have colors similar to young stellar objects (YSOs) at both near- and mid-infrared (e.g. Spitzer) wavelengths, which limits our ability to identify YSOs from colors alone. Structural information from high-quality near-infrared observations allows us to better separate YSOs from galaxies, rejecting 2/5 of the YSO candidates identified from Spitzer observations of our regions and potentially extending the YSO luminosity function below K of 15 magnitudes where galaxy contamination dominates. Once they are identified we use galaxies as valuable extra signal for making extinction maps of molecular clouds. Our new iterative procedure: the Galaxies Near Infrared Color Excess method Revisited (GNICER), uses the mean colors of galaxies as a function of magnitude to include them in extinction maps in an unbiased way. GNICER increases the number of background sources used to probe the structure of a cloud, decreasing the noise and increasing the resolution of extinction maps made far from the galactic plane. | The near-infrared (NIR) has been a valuable window into star-forming regions, providing us with a variety of tools to study the process by which dark clouds coalesce into stars: luminosity functions in the $K$ band (2.2$\mu$m) can provide an accurate estimate of the initial mass function \citep[e.g.][]{Muench:2002,Stolte:2005}; studies of embedded clusters reveal important details of the processes by which stars form either in groups or in isolation \citep[e.g.][]{Lada:1991,Roman-Zuniga:2007}; and extinction mapping in the near-infrared allows precise determination of the column density structure of the cloud \citep[e.g.][]{Alves:2001,Cambresy:2002}. Of particular use is the narrow range of intrinsic near-infrared colors of main-sequence stars (typically $H-K$ = 0 - 0.4 and $J-H$ = 0 - 1.0). This narrow range is due both to the near-infrared bands lying on the Wien tail of all stellar (hydrogen-fusing)-temperature blackbodies and to the relative paucity and uniformity of absorption feature\footnote{Principally H$^{-}$, but also CO features which are sensitive to surface gravity and produce a small split between dwarf and giant colors}. By assuming that all stars have the same intrinsic color we can do purely photometric extinction mapping. This is the heart of the Near Infrared Color Excess method (NICE) \citep{Lada:1994} and the Near Infrared Color Excess method Revisited (NICER) \citep{Lombardi:2001}. Additionally, we can use the narrow range of stellar near-infrared colors to identify young stars by their near-infrared excess. The thermal contribution from their disk or envelope changes their color from that of a plain photosphere and places them in a certain region of a near-infrared color-color ($J-H$ versus $H-K$) diagram. This is the CTTS locus defined by \citet{Meyer:1997}. Unfortunately, a number of other astronomical objects also have intrinsic red colors, similar to CTTS. We refer to all intrinsically red objects as near-infrared excess (NIRX) sources. Of particular concern for studies of star forming regions are objects which are not either young-stellar objects (YSOs) or T-Tauri type young stars (TTS). These objects---brown dwarfs, a variety of evolved stars, galaxies, and AGN, are often studied in the near-IR in their own right, but in studies of star forming regions where we have limited information about their nature (perhaps only $J$,$H$, and $K$ photometry) they are contaminants which must be understood and identified in order to avoid common biases in these studies. For example, including intrinsically red objects in NICER produces an overestimate of the extinction. In this paper we identify as galaxies a large number of NIRX sources at moderate redshifts (z = 0.1-0.5) in high-quality near-infrared images outside of the Perseus molecular cloud complex. After describing our observations and reduction in \S\ref{Obs} we present a detailed case that our NIRX sources are galaxies in \S\ref{ColorAndNumber}. In the heart of the paper we show how deep NIR images can help identify genuine YSOs in Spitzer observations (\S\ref{Clusters}) and how we can make use of these contaminating galaxies as additional probes for extinction mapping (\S\ref{GNICERSection}). | 7 | 10 | 0710.5767 |
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0710 | 0710.0712_arXiv.txt | We investigate molecular evolution in a star-forming core that is initially a hydrostatic starless core and collapses to form a low-mass protostar. The results of a one-dimensional radiation-hydrodynamics calculation are adopted as a physical model of the core. We first derive radii at which CO and large organic species sublimate. CO sublimation in the central region starts shortly before the formation of the first hydrostatic core. When the protostar is born, the CO sublimation radius extends to 100 AU, and the region inside $\lesssim 10$ AU is hotter than 100 K, at which some large organic species evaporate. We calculate the temporal variation of physical parameters in infalling shells, in which the molecular evolution is solved using an updated gas-grain chemical model to derive the spatial distribution of molecules in a protostellar core. The shells pass through the warm region of $10 -100$ K in several $\times$ $10^4$ yr, and fall into the central star $\sim 100$ yr after they enter the region where $T \gtrsim 100$ K. We find that large organic species are formed mainly via grain-surface reactions at temperatures of $20 -40$ K and then desorbed into the gas-phase at their sublimation temperatures. Carbon-chain species can be formed by a combination of gas-phase reactions and grain-surface reactions following the sublimation of CH$_4$. Our model also predicts that CO$_2$ is more abundant in isolated cores, while gas-phase large organic species are more abundant in cores embedded in ambient clouds. | In a last decade, great progress has been made in our understanding of the chemical evolution of low-mass star-forming cores \citep[][and references therein]{fra07, bt06}. Among the observational advances are the detection of chemical fractionation in several prestellar cores; emission lines of the rare isotopes of CO are weaker at the core center than at outer radii, while emission lines of nitrogen-bearing species (e.g., N$_2$H$^+$) show relatively good correlation with the centrally-peaked dust continuum \citep[e.g.][]{cas99, taf02}. Theoretical models show that the CO depletion is caused by adsorption onto grains \citep{bl97,aik01}, and that the CO depletion helps to temporarily maintain the N$_2$H$^+$ abundance at the core center. Theoretical models also predict that a fraction of the adsorbed CO will be hydrogenated to form H$_2$CO and CH$_3$OH on grain surfaces \citep{ar77,hhl92}, which is confirmed by laboratory experiments \citep{wkl02,fuchs07}. The existence of solid methanol in low-mass star formation regions has been confirmed observationally; \citet{ppp03} detected a high abundance of CH$_3$OH ice (15-25 \% relative to water ice) towards three low-mass protostars among $\sim 40$ observed protostars. Radio observations find gaseous CH$_3$OH to be abundant in the central regions of protostars \citep{sjv02}. Since the formation of CH$_3$OH is inefficient in the gas phase \citep{gpc06,gwh06}, it must be formed by the hydrogenation of CO on grain surfaces in the prestellar core stage, and then sublimated to the gas-phase as the core is heated by the protostar. Although CH$_3$OH ice is not detected towards the majority of low-mass protostars \citep{ppp03} and the background star Elias 16 \citep{chi96}, the upper limit for the CH$_3$OH ice abundance is a few $\%$ relative to water ice, which is not low enough to contradict the idea that gaseous CH$_3$OH around protostars is originally formed by grain-surface reactions and then sublimated. Two major questions, however, remain to be answered, the first being at what stage CO returns to the gas phase. The core remains nearly isothermal as long as cooling by radiation is more efficient than heating by contraction (compression). Eventually, though, the heating overwhelms the cooling, so that the core center becomes warmer. A newly-born protostar further heats the surrounding core. Laboratory experiments show although some fraction of CO can be entrapped in water ice \citep{cdf03}, a significant amount of CO sublimates at around 20 K \citep{sa88}. In order to predict if an observable amount of CO returns to the gas-phase during the prestellar core stage or after the birth of a protostar, the temperature distribution in a core should be calculated by detailed energy transfer. Such a prediction is important in order to observe very young protostars. Once it is sublimated, CO again becomes a useful observational probe. In addition, CO sublimation significantly affects the gas-phase chemistry; for example, it destroys N$_2$H$^+$. The second question concerns how large organic molecules are formed in protostellar cores. In recent years, diverse organic molecules, including methanol (CH$_3$OH), dimethyl ether (CH$_3$OCH$_3$), acetonitrile (CH$_3$CN), and formic acid (HCOOH), have been detected towards low-mass protostars \citep[][and references therein]{cec07}. They are still referred to as "hot-core species", because it was once thought that they are only characteristic of hot ($T\sim 200$ K) cores in high-mass star forming regions. A large number of modeling studies on hot-core chemistry show that sublimed formaldehyde (H$_2$CO) and CH$_3$OH are transformed to other organic species by gas-phase reactions within a typical timescale of $10^4-10^5$ yr \citep[e.g.][]{mh98}. In low-mass cores, however, the time scale for the cloud material to cross the hot ($T\sim 100$ K) region should be smaller than $10^4$ yr, considering the infall velocity and temperature distribution in the core \citep{bot04a}. Furthermore, theoretical calculations and laboratory experiments recently showed that gas-phase reactions are much less efficient in producing some hot-core species, such as methyl formate (HCOOCH$_{3}$) and dimethyl ether, than assumed in previous models \citep{hor04, gep06,gh06}. Several model calculations have been performed on the chemistry that occurs in low-mass protostellar cores. \citet{dot04} solved a detailed gas-phase reaction network assuming a core model for IRAS 16293-2422, and succeeded in reproducing many of the observed lines within 50 \%. The physical structure of the core; i.e., its density and temperature distribution, was fixed with time. They assumed gas-phase initial molecular abundances that pertain to the high-mass hot-core AFGL 2591. \citet{lbe04}, on the other hand, constructed a core model that evolves from a cold hydrostatic sphere to a protostellar core by combining a sequence of Bonnor-Ebert spheres with the inside-out collapse model by \citet{shu77}. They solved a chemical reaction network that includes gas-phase reactions along with adsorption and desorption of gas-phase/ice-mantle species. The resulting molecular distributions and line profiles are significantly different from those of the static core models \citep{leb05}. The chemical network of \citet{lbe04}, however, does not include large organic species. In the present paper, we re-investigate molecular evolution in star-forming cores with the partial goal of answering the two questions posed above. We adopt a core model by \citet{mi00}; it accurately calculates the radial distribution of temperature, which determines when and where the ice components are sublimated. The model also enables us to follow molecular evolution smoothly from a prestellar core to a protostellar core. The chemistry includes both gas-phase and grain-surface reactions according to \citet{gh06}; the surface reactions in particular are important for producing organic molecules in a warming environment. Here, we report a solution of the reaction network following the temporal variation of density and temperature in infalling shells to derive a spatial distribution of molecules, including complex organic ones, in a protostellar core. | \subsection{Physical structure of the core and sublimation radius} We have re-analyzed and adopted the results of \citet{mi00} to derive the sublimation temperatures of CO and large organic species, and to investigate molecular evolution in a star-forming core. The chosen conditions in the envelope can be different from those of other models because the temperature distribution in the envelope depends not only on the evolutionary state and mass of the central object (either the first or second core), but also on the mass distribution in the envelope, which should depend on the initial conditions. Recently, \citet{omu07} investigated the temperature distribution in the first-core envelope assuming mass distributions from the L-P model \citep{lar69} and Shu model \citep{shu77} similarity solutions. The former has a more massive envelope than the latter. When the first core has a mass of $0.05 M_{\odot}$, for example, the hydrogen number density at $r=10$ AU is several $10^{10}$ cm$^{-3}$ and $\sim 10^9$ cm$^{-3}$ with the L-P model and the Shu model, respectively. The L-P model gives a core luminosity of $10^{-1} L_{\odot}$, and sublimation radii of $r_{\rm 20K} \sim 100$ AU and $r_{\rm 100 K}\sim 10$ AU, while the Shu model yields a core luminosity of $\sim 10^{-3} L_{\odot}$ and an $r_{\rm 20 K}$ of $\sim 20$ AU. The latter model does not exceed 100 K at any radius. Our core model is warmer than the Shu model but colder than the L-P model. In the second core stage, on the other hand, the envelope structure can deviate from spherical symmetry and be accompanied by a circumstellar disk. Considering a typical angular velocity for molecular cloud cores \citep[$\sim 10^{-14}$ s$^{-1}$; e.g.,][]{ga85}, the centrifugal radius (i.e. the initial disk radius) is $\sim 100$ AU. Our results at $r\gtrsim$ several hundred AU are thus relatively robust, while the core structure would be significantly different from our model at smaller radii. For example, the large organic species could be sublimated by the accretion shock onto the forming disk rather than in the envelope, since the centrifugal radius $\sim 100$ AU coincides with the sublimation radius of large organic species in our model. \subsection{Simple molecules} \cite{lbe04} investigated the evolution of relatively simple molecules, such as HCN and N$_2$H$^+$, in a star-forming core by combining a sequence of Bonnor-Ebert spheres for the prestellar stage with the inside-out collapse model of \cite{shu77} for the core after the first-core formation. These simple molecules are often observed in star-forming cores and are of importance as observational probes, while we mainly discussed large organic species in \S 3. Figure \ref{cfLee} shows the radial distribution of simple molecules in our model at $t_{\rm final}$. Comparison with \cite{lbe04} is not easy because there are many differences in the physical core models and chemical reaction networks. First, the inside-out core model, which is adopted in \cite{lbe04}, yields lower temperatures than our core model, as discussed above. Secondly, we adopt different binding energies for molecules to the grain surface. While the difference is on the order of only a few hundred K for many species, the binding energy of NH$_3$ is significantly higher in our model (5534 K) than in \cite{lbe04} (1082 K); the latter value seems to originate from \cite{hh93}, in which the hydrogen bonding of NH$_3$ is not taken into account. Thirdly, grain-surface reactions, which are the main formation processes of saturated species such as NH$_3$ and H$_2$CO in our model, are not included in \cite{lbe04}. Hence, here we compare our results with \cite{lbe04} qualitatively rather than quantitatively. A main conclusion of \cite{lbe04} is that the molecular abundances vary significantly near and inside the CO sublimation radius. For instance, N$_2$H$^+$ is destroyed by CO, and thus declines steeply inwards across the CO sublimation radius, an effect which happens in our model as well. \cite{lbe04} also found that some molecules, such as H$_2$CO, have their peak abundance at the sublimation radius; the gas-phase abundance first increases inward via sublimation, but then decreases due to gas-phase reactions at smaller radii. Similar phenomena can be seen in our model, but the spatial variation of the gas-phase abundances is more moderate than in \cite{lbe04}, for which we can think of a few reasons. First, \cite{lbe04} calculated the molecular evolution in 512 shells, while we calculated the chemistry in only 13 shells. A larger number of shells are needed to resolve abundance fluctuations on a smaller radial scale. Secondly, higher infall speeds, in general, make the abundance distributions more uniform. Since the infall speed is higher at inner radii in the infalling envelope, it would be natural that the molecular abundances remain relatively uniform at $r\lesssim 100$ AU, which is not calculated in \cite{lbe04}. Thirdly, our chemical network includes a much larger number of species and reactions. \cite{lbe04} used a reduced reaction network with $\sim 80$ species and $\sim 800$ reactions to save computational time, while our model includes 655 species and 6309 reactions. In a small reaction network, a sudden increase of a species (e.g., as caused by sublimation) can easily change the abundances of other species though chemical reactions. In a large reaction network, on the other hand, a larger number of reactions contribute to the formation and destruction of each species, and thus the sudden abundance change of one species does not necessararily propagate to other species. Another noticeable difference in our results from those of \cite{lbe04} is that HCO$^+$ decreases more steeply inwards at $r\lesssim 1000$ AU in our model; H$_3$CO$^+$ and C$_3$H$_5^+$ are the dominant positive ions rather than HCO$^+$ at the inner radii. The oxygen-bearing species atomic oxygen (O), molecular oxygen (O$_2$), and gaseous H$_2$O are also of special interest because O is very reactive and because the two molecules have been intensively observed by {\it the Submillimeter Wave Astronomy Satellite (SWAS)} and {\it Odin} Satellite in recent years. Our model predicts that O reaches an abundance of $4\times 10^{-5}$ relative to hydrogen nuclei at $r=8000$ AU, while it steeply decreases inwards from $\sim 1000$ AU to 100 AU. Even at $r=8000$ AU, however, H$_2$O ice is the most abundant O-bearing species (Figure \ref{dist}). \cite{ber00} summarized the {\it SWAS} observations of cold molecular clouds by stating that the fractional abundance of O$_2$ lies under $10^{-6}$ and that of gaseous water lies in the range $10^{-9}$ to a few $\times$ $10^{-8}$. Molecular oxygen has recently been detected towards $\rho$ Oph by Odin with an abundance of $5\times 10^{-8}$ relative to hydrogen \citep{lar07}. These abundances are consistent with our predictions for the outermost radius $r=8000$ AU, where the density $n_{\rm H}\sim 10^4$ cm$^{-3}$ and temperature $T<20$ K are similar to those in molecular clouds. \subsection{Dependence on visual extinction at the core edge} So far we have assumed that the model core is embedded in ambient clouds of $A_{\rm v}= 3$ mag. In reality, some cores (e.g. Bok globules) are isolated, while others are embedded in clouds. Isolated cores are directly irradiated by interstellar UV radiation, which causes photo-reactions (photodissociation and ionization) in the gas-phase and ice mantles \citep[e.g.][]{lee96,rh01}. In order to evaluate the effect of ambient UV radiation, we recalculated molecular abundances in a core that is directly irradiated by interstellar UV radiation; i.e. $A_{\rm v}=0$ mag at the outer edge of the core ($r=4\times 10^4$ AU). The photodissociation rates of H$_2$ and CO were calculated by following \citet{lee96}, which gives the shielding factors as a function of $A_{\rm v}$ and column densities of CO and H$_2$ in the outer radii. The outermost shell for which we calculate molecular evolution is initially located at $r\sim 1.4\times 10^4$ AU and declines to $r=8000$ AU at $t_{\rm final}$. Column densities of CO and H$_2$ outside of this shell were estimated by assuming $n$(CO)/$n_{\rm H}=5\times 10^{-5}$ and $n$(H$_2$)/$n_{\rm H}=0.5$. Figure \ref{av0} shows the resultant distribution of molecular abundances in a protostellar core at $t_{\rm final}$. Compared with the embedded core model (Figure \ref{dist}), the fractional abundances of CH$_3$OH and H$_2$CO are lower by more than one order of magnitude, while the CO$_2$ abundance is higher. When the shells are still at $r\gtrsim$ several thousand AU and have relatively low $A_{\rm v}$ ($\lesssim 4$ mag), the photodissociation of H$_2$O ice efficiently produces OH, which reacts with CO to produce CO$_2$ ice. Methanol in the ice mantle is dissociated to H$_2$CO, which is further dissociated to CO. Species such as CH$_3$CHO, HCOOCH$_3$ and CH$_3$OCH$_3$ are also less abundant in the isolated model, because their formation path includes H$_2$CO in the ice mantle. Formic acid in the gas phase extends only up to $\sim 100$ AU, while it extends to several hundred AU in the embedded model. In the isolated model, it is formed mainly by OH + HCO on the grain surface. Our results may indicate that molecular abundances in hot cores and corinos depend on whether the core is isolated or embedded in clouds. The importance of photo-reactions on ice mantle abundances has also been investigated in a number of laboratory experiments. For example, \citet{wk02} and \citet{wm07} found that carbon dioxide is efficiently produced by UV irradiation on a binary ice mixture of H$_2$O and CO, a result that is consistent with our model. However, \cite{wm07} found that the UV irradiation also produces CH$_3$OH with a relative abundance of $n$(CH$_3$OH)/$n$(CO) $\sim 10$ \% in the ice mixture. Comparison of our model with their experiment is not straightforward because of differences in temperature and included reactions, but we may have underestimated the CH$_3$OH ice abundance in the irradiated core model. The discrepancy can arise from the H atom desorption rate in our model. We calculated that UV radiation penetrates into the ice mantle and dissociates H$_2$O to produce H atoms embedded in ice. Although we do not discriminate between H atoms on the ice surface and those embedded in the ice mantle, the embedded H atoms in reality would have a lower desorption rate and a better chance of reacting with neighboring CO, a reaction that initiates the formation of methanol in the mantles. Discrimination between surface and embedded hydrogen atoms should be included in future work. \subsection{Comparison with observation} Our model results can be compared with observational results of low-mass protostars. Comparison of the physical structure of the core has already been discussed in detail by \citet{mi00}. Here we concentrate on molecular abundances. First, we compare molecules in ice mantles. The observation of ices towards the low-mass protostar Elias 29 is summarized in \citet{es00}; the abundances of CO, CO$_2$, CH$_3$OH and CH$_4$ relative to water ice are 5.6 \%, 22 \%, $< 4$ \%, and $< 1.6$ \%, respectively. On the other hand, \citet{ppp03} detected a high abundance (15-25 \% relative to water ice) of CH$_3$OH ice towards 3 low-mass protostars among 40 observed protostars. Although the observation of ice features is difficult, it is probable that the composition of ice mantles depends on their environment and the history of the observed regions. From a theoretical point of view, the abundances of molecules on grains and their fractional abundances compared with H$_2$O ice depend on time and radius from the protostar. In addition, an embedded core and an isolated core lead to significantly different calculated abundances for solid CH$_3$OH, H$_2$CO and CO$_2$. Table~\ref{solid_theo} lists ratios of ice species with respect to water ice at $t_{\rm final}$ for local abundances at radii of 1000 AU and 8000 AU and for column densities towards the core center. Since most of the core material exists at $r\le 8000$ AU along the line of sight, and since the ices are present at $r>10$ AU, the column density is calculated by integrating the number density of ice species from 10 AU to 8000 AU. It is interesting to note that regardless of the distance from the protostar, the isolated core model gives smaller surface abundances of CH$_3$OH and H$_2$CO and a higher surface abundance of CO$_2$ than the embedded model. In both models, the surface abundance of CO at 1000 AU is much smaller than observed, but it increases towards larger radii up to 10 and 20\% in the embedded and isolated models respectively. Comparison with the observations is best done using our column density ratios, which are in reasonable agreement with Elias 29 for both models. Considering the variation among cores, our models show reasonable agreement with observation. The disagreement with CO doubtless results from our assumption concerning desorption rates. In the present work, we use a desorption energy for each species mainly referring to laboratory experiments on pure ice sublimation or theoretical estimates that sum up the van der Waals forces between adsorbed atoms and grain surface. But in reality, interstellar ice is a mixture, and hence volatile species can be entrapped by less volatile species; recent temperature-programmed desorption results show that in a mixture rich in water ice, much CO desorbs at considerably higher temperatures \citep{cdf03}. Table \ref{obs} lists the gas-phase abundances of large molecules towards the low-mass protostar IRAS 16293-2422 and in the central region ($r=30.6$ AU) of our core model. The physical and chemical structures of IRAS 16293-2422 have been intensively studied \citep{cec00a,cec00b, sch02,cau03, cha05}. The physical parameters derived by the model at $t_{\rm final}$ ($\sim 23 L_{\odot}$ and $r_{\rm 100K}\sim1.2\times 10^2$ AU) are very close to the ones of IRAS16293-2422; the observed luminosity of the source is 27 $L_{\odot}$ \citep{wal86} and the physical structure has been constrained by multi-line analysis of H$_2$O and H$_2$CO observations through a detailed radiative transfer code \citep{cec00b}. It should be noted that the emission lines of large organic molecules observed in this source are not spatially resolved, except for a few lines investigated by interferometric observations \citep{bot04b, kua04, rem06}. The density and temperature of the gas should vary both within the beam and along the line of sight. Hence the derived molecular abundances vary significantly depending on the assumptions concerning the physical structure of the core and the emitting region, which explains the difference in abundances determined by different investigators (see Table~ \ref{obs}). Although it is not obvious which core model, embedded or isolated, should be compared with IRAS16293-2422, we would prefer the embedded core model. While IRAS16293-2422 is in the Ophiuchus molecular cloud, which harbors several UV sources in the form of OB stars, the $^{13}$CO and C$^{18}$O observations indicate that the core is embedded in dense gas \citep{lor89,tac00}. The high molecular D/H ratios observed towards IRAS 16293-2422 and its neighboring starless core IRAS16293E \citep{cec07,vas04} indicate that these cores have been very cold and thus well-shielded from the ambient stellar radiation. Considering the uncertainties in observationally-estimated molecular abundances, HCOOCH$_3$, HCOOH, and CH$_3$CN in our embedded core model show reasonable agreement with the observations. The embedded model, however, underestimates CH$_3$OCH$_3$ and overestimates H$_2$CO and CH$_3$OH. The isolated core model, on the other hand, reproduces observed abundances of CH$_3$OH, HCOOH and CH$_3$CN, but strongly underestimates H$_2$CO, HCOOCH$_3$, and CH$_3$OCH$_3$. We can think of several possible solutions to improve the agreement with observation. First, it is noteworthy that the gaseous CH$_3$OH abundance estimated in IRAS 16293-2422 is much lower than the abundance of CH$_3$OH ice ($n$(CH$_3$OH ice)/$n_{\rm H} \sim 10^{-5}$, assuming n(H$_2$O ice)/ $n_{\rm H}\sim 10^{-4}$) detected by \citet{ppp03} towards some low-mass protostars. We may have missed or underestimated the reactions which transform CH$_3$OH to other large organic species. Secondly, the abundances of large organic species in the central region can vary with time. In the present work, we have concentrated on the distribution of molecules only at $t_{\rm final}$. Shells that reach the central region at different times should experience different temporal variations of physical conditions. Some shells may experience longer periods at $T\sim 20-40$ K, where large organic species start to be efficiently formed. This possibility will be pursued in a future work. Thirdly, core models with rotation could produce higher abundances of large organic species; because of the centrifugal force, core material migrates more slowly and stays longer in the temperature range preferable for the formation of large organic molecules (e.g. in a forming disk). Because of the beam size of the current radio observations, little is known concerning the spatial distribution of the large organic species; they are mostly confined within a few arc seconds from the core center \citep[e.g.][]{kua04}. But there are exceptions; \citet{rem06} found that HCOOH and HCOOCH$_3$ show extended emission of $\sim 5$ arcsec. Our embedded core model reproduces the extended emission of HCOOH, while HCOOCH$_3$ is confined to $r<100$ AU. \subsection{Carbon chains in a protostellar core} Recently, \cite{sshk07} detected strong emission lines of carbon chain species such as C$_4$H and C$_4$H$_2$ towards the low-mass protostar L1527, which is considered to be in a transient phase from class 0 to class I. In general, carbon-chain species are associated with the early stages of a cold cloud core when the dominant form of carbon changes from atomic carbon to CO. Hence, the existence of carbon-chain species towards L1527 is a surprise. Figure \ref{carbon_chain} (a-b) shows the temporal variation of CH$_4$ and carbon-chain abundances in the shell that reaches $r=2.5$ AU at $t_{\rm final}$. When the core starts to contract, methane (CH$_4$) and C$_3$H$_4$ are already abundant in ice mantles. Methane has been formed by the hydrogenation of carbon on grain surfaces, while C$_3$H$_4$ has been formed by a combination of gas-phase reactions (to form unsaturated carbon chains such as C$_{3}$H and C$_{3}$H$_{2}$) and grain-surface reactions (to hydrogenate them). When the CH$_4$ sublimates, some fraction reacts with C$^+$ to form C$_2$H$_3^+$, which is a precursor for ion-molecule reactions that lead to the production of larger unsaturated hydrocarbons. While gas-phase reactions tend to make the carbon chains longer, adsorbed species experience hydrogenation and dissociation by cosmic-ray induced UV radiation. Figure \ref{carbon_chain} (c-d) shows the distribution of carbon-chain species in our embedded core model. Methane and C$_3$H$_4$ are abundant in the ice mantle even at $r\gtrsim 10^3$ AU, while other hydrocarbons are abundant at a few 100 AU $\lesssim r \lesssim 10^3$ AU. Most of the carbon-chain species desorb at $r\lesssim$ a few 100 AU. In summary, our model indicates that carbon-chain species can be re-generated in the protostellar core by the combination of gas-phase reactions and grain-surface reactions. It should be noted, however, that large oxygen-containing organic species are not detected in L1527 \citep{sak07}, while they are abundant in the central region of our model. In future work, the temporal variation of the molecular distribution in model cores should be investigated to see if at certain evolutionary stages carbon-chain species are abundant while large organic species are not. For example, in their recent study of the chemistry of cold cores, \cite{gwh06} found that generation of gas-phase hydrocarbons from precursor surface methane occurs at very late times, after the abundance of surface methanol has essentially vanished. More detailed observations are also needed for a quantitative comparison with models. While our model predicts the column density of C$_4$H to be $2\times 10^{16}$ cm$^{-2}$ towards the central star, the observation gives $2\times 10^{14}$ cm$^{-2}$. The actual column density towards the central star can, however, be larger, because the emission is averaged within the beam (a few 10 arcsec) in the current observation. Methane ice, the precursor of the carbon-chain species in our model, is as abundant as 23 \% relative to water ice at the outer edge of our core model. A deep integration of ice features towards field stars is needed to constrain the CH$_4$ ice abundances at the outer edge of the core and/or quiescent clouds, while the column density ratio of CH$_4$ ice to H$_2$O ice (1 \%, see Table 2) in our model is consistent with the observation towards YSO's (\S 4.3), because of the relatively large sublimation radius of CH$_4$. | 7 | 10 | 0710.0712 |
0710 | 0710.1839_arXiv.txt | Black hole mass (\mbh) is a fundamental property of active galactic nuclei (AGNs). In the distant universe, \mbh\ is commonly estimated using the MgII, \hb, or \ha\ emission line widths and the optical/UV continuum or line luminosities, as proxies for the characteristic velocity and size of the broad-line region. Although they all have a common calibration in the local universe, a number of different recipes are currently used in the literature. It is important to verify the relative accuracy and consistency of the recipes, as systematic changes could mimic evolutionary trends when comparing various samples. At $z=0.36$, all three lines can be observed at optical wavelengths, providing a unique opportunity to compare different empirical recipes. We use spectra from the Keck Telescope and the Sloan Digital Sky Survey to compare \mbh\ estimators for a sample of nineteen AGNs at this redshift. We compare popular recipes available from the literature, finding that \mbh\ estimates can differ up to $0.38\pm0.05$ dex in the mean (or $0.13\pm0.05$ dex, if the same virial coefficient is adopted). Finally, we provide a set of 30 internally self consistent recipes for determining \mbh\ from a variety of observables. The intrinsic scatter between cross-calibrated recipes is in the range $0.1-0.3$ dex. This should be considered as a lower limit to the uncertainty of the \mbh\ estimators. | Understanding the growth of supermassive black holes along with their host galaxies is one of the fundamental questions in current astrophysics \citep[e.g.][]{DMS05,Cro++06b}. Black hole mass (\mbh) is a key parameter in revealing the nature of black hole-galaxy coevolution as well as the physics of active galactic nuclei (AGNs). However, direct mass measurements using the motions of gas and stars in the sphere of influence of a central black hole is limited to very nearby galaxies \citep[e.g.][]{K+G01,F+F05}. Beyond the very local universe, the so-called ``virial'' or ``empirically calibrated photo-ionization'' method based on the reverberation sample is popularly used for active galaxies \citep[e.g.][]{WPM99, Kas++00, Kas++05, Ben++06a}. This method utilizes broad line widths as velocity indicators and monochromatic continuum or line luminosities as indicators of broad-line region size, hence estimating virial \mbh. A combination of the MgII, \hb, or \ha\ broad emission line widths and the 3000\AA, 5100\AA, \hb, or \ha\ luminosities is typically used, depending on the redshift of the source and the observational setup. Several equations have been presented in the literature to estimate \mbh\ using various combinations of these indicators \citep[e.g.,][]{W+U02b,W+U02a,M+J02,TMB04,Kol++06,G+H05,V+P06,Woo++06,Sal++07,N+T07,Tre++07}. Although all three emission lines have a common calibration based on the reverberation sample in the local universe, it is important to verify that different recipes give consistent results; any systematic changes could mimic evolutionary trends given that different recipes are often used in various studies. At $z=0.36$, all three lines can be observed at optical wavelengths, providing a unique opportunity to cross-calibrate the different methods of \mbh\ estimation. Using data from the Keck Telescope and the Sloan Digital Sky Survey for a sample of nineteen AGNs at $z=0.36$, we compare the different methods of estimating \mbh, and derive a set of self-consistent equations for \mbh\ estimates using every combination of velocity scale (FWHM and line dispersion \sline\, of MgII, \hb, or \ha) and luminosity (3000\AA, 5100\AA - nuclear and total - \hb, or \ha). The paper is organized as follows. In \S~\ref{sec:data} we describe the sample selection, observations, and data reduction. In \S~\ref{sec:meas} we describe our line fitting process, based on expansion in Gauss-Hermite series, and the resulting luminosity and width measurements. In \S~\ref{sec:form} we review the various formulae adopted in the literature, and compare the various \mbh\ estimators. In \S~\ref{sec:res} we present our self-consistent recipes. Section~\ref{sec:sum} summarizes our results. Throughout this paper magnitudes are given in the AB scale. We assume a concordance cosmology with matter and dark energy density $\Omega_m=0.3$, $\Omega_{\Lambda}=0.7$, and Hubble constant H$_0$=70 kms$^{-1}$Mpc$^{-1}$. \begin{figure*} \epsscale{0.8} \plotone{f1.eps} \caption{Flux-calibrated spectra. The SDSS spectra are shown in black, and the Keck spectra are shown in blue and red. The MgII line can be seen on the far left of each wavelength range, while \hb\ is located in the center and \ha\ to the far right. } \label{spectra1} \end{figure*} \begin{figure*} \epsscale{0.8} \plotone{f2.eps} \caption{As in Figure~\ref{spectra1} for objects S11 to S29.} \label{spectra2} \end{figure*} | \label{sec:sum} In this paper we have used Keck and SDSS spectra of nineteen Seyferts at $z=0.36$, to perform a comprehensive study of ``virial'' black hole mass estimators for broad line AGNs. The main results can be summarized as follows: \begin{enumerate} \item We have fit Gauss-Hermite series to the data in order to measure the FWHM and \sline\, of MgII, \hb\,, and \ha\,, as well as \ha\, and \hb\, luminosities and continuum luminosities at 3000\AA\ and 5100\AA. Measurement errors are approximately 0.02 dex on the MgII and \hb\ line widths, 0.04-0.05 on the \ha\ line widths, 0.01 dex on the continuum luminosity, 0.01 dex on the \hb\ luminosity, and 0.06 dex on the \ha\ luminosity. \item We have compared twelve formulae taken from the literature, showing that \mbh\ estimates can differ systematically by as much as $0.38\pm0.05$ dex (or $0.13\pm0.05$ dex, if the same virial coefficient is adopted). Such differences should be taken into account when comparing data obtained with different methods. \item We have cross-calibrated a set of 30 empirical recipes based on all combinations of the velocity and luminosity indicators corresponding to the \ion{Mg}{2}, \hb, and \ha\ broad lines. Taking the masses measured by Treu et al. (2007) as our fiducial black hole masses, we find that: the absolute scale of the different indicators is calibrated to within $\sim$0.05 dex; the best agreement is found when using the line dispersion of \hb\, as a velocity estimator, with the residual 0.1 dex r.m.s. scatter resulting from the various continuum luminosity estimators; adopting the line dispersion of \ion{Mg}{2} raises the scatter to 0.2 dex; for the other estimators the intrinsic scatter is in the range 0.2-0.38 dex. This implies a lower limit of 0.1-0.2 dex on the validity of each estimator for each individual case. \end{enumerate} The newly calibrated recipes should be useful to reduce the sources of systematic uncertainties when comparing different studies. | 7 | 10 | 0710.1839 |
0710 | 0710.3461_arXiv.txt | We present the results of a deep, wide-field search for transiting `Hot Jupiter (HJ)' planets in the globular cluster $\omega$ Centauri. As a result of a 25-night observing run with the ANU 40-inch telescope at Siding Spring Observatory, a total of 109,726 stellar time series composed of 787 independent data points were produced with differential photometry in a 52$'$$\times$52$'$ (0.75 deg$^2$) field centered on the cluster core, but extending well beyond. Taking into account the size of transit signals as a function of stellar radius, 45,406 stars have suitable photometric accuracy ($\le$0.045 mag to V$=$19.5) to search for transits. Of this sample, 31,000 stars are expected to be main sequence cluster members. All stars, both cluster and foreground, were subjected to a rigorous search for transit signatures; none were found. Extensive Monte Carlo simulations based on our actual data set allows us to determine the sensitivity of our survey to planets with radii $\sim$1.5R$\rm{_{Jup}}$, and thus place statistical upper limits on their occurrence frequency $F$. Smaller planets are undetectable in our data. At 95$\%$ confidence, the frequency of Very Hot Jupiters (VHJs) with periods P satisfying 1d~$<$~P~$<$~3d can be no more than $F_{\rm VHJ} < 1/1040$ in $\omega$ Cen. For HJ and VHJ distributed uniformly over the orbital period range 1d~$<$~P~$<$5d, $F_{\rm VHJ+HJ} < 1/600$. Our limits on large, short-period planets are comparable to those recently reported for other Galactic fields, despite being derived with less telescope time. | The identification and subsequent study of extrasolar planets has become a subject of intense interest in recent years. To date, $\sim$250 giant planets have been discovered\footnote{http://exoplanet.eu/}, ranging in mass from Neptune to greater than Jupiter. These new worlds are altering our understanding of the formation and evolutionary processes of giant planets in the immediate solar neighborhood. Currently, the majority of them have been found using radial velocity (RV) techniques, which favours the detection of close-in, massive planets. Statistical analysis of current RV detections indicates that 1.2$\%$$\pm$0.3$\%$ of nearby F,G and K stars are orbited by `Hot Jupiter' (HJ) planets, those with orbital periods of only a few days ($\le$0.1 AU) and minimum masses approximately equal to that of Jupiter \citep{M2005}. Such discoveries have challenged traditional ideas of planetary evolution, implying that a rapid migration of the planet takes place soon after formation. Planet frequency from radial velocities appear to depend on the metallicity of the host star \citep{G1997,L2000,S2001,FV2005}. However, there is very little observational evidence for a lower planetary frequency at quite low metallicities, due to the bias of radial velocity detections in association with bright nearby high metallicity stars. Hence dedicated surveys in low metallicity environments, like globular clusters, can provide information to help understand this relationship in a more robust manner. Indeed, does low metallicity halt planet formation or just affect the planetary migration process? Of the currently known planets, $\sim$90$\%$ have only a minimum mass assigned to them, due to the unknown inclination of the planetary orbit, and unknown radii and densities. These quantities can be measured for transiting planets. Due to its short orbital period, each HJ has a non-negligible probability of transiting its host star that depends on the orbital separation and the ratio between the stellar and planetary radii. Typical transit depths and durations are $\sim$1.5$\%$ and $\sim$2 hours respectively. With precise photometry these transiting systems can be identified in the field, leading to direct measurements of the planetary radius (provided the stellar radius is known) from the depth of the transit dip. Studies of the planetary atmosphere may be attempted if transits occur, and when coupled with RV measurements, accurate mass and density determinations can be made. Currently, 28 transiting exoplanets are known \citep{C2000,He2000,K2003,K2005,B2004,B2005,A2004,P2004,S2005,MCC2006,Bakos2006,Bakos2007,Bakos2007c,OD2006,OD2007,CC2006,Sahu2006,Bu2007,G2007,Mand07,Barb2007,Kov2007,N2007}, only a handful of which were first discovered through radial velocity searches. Despite an increasing number, the detection rate of planets from transit searches is significantly lower than initially expected eg, \citet{H2003}. This lack of detections is due, in part, to observing strategy: a long observing window ($\ge$1 month equivalent) with a dedicated telescope coupled with a wide field and high temporal resolution are needed to sample enough stars frequently enough to allow the detection of a transit. Also the production of a large enough number ($\sim10^{5}\rightarrow10^{6}$) of high-quality ($\le$0.02 mag) lightcurves of dwarf stars coupled with a sensitive detection algorithm with low false alarm rates is non-trivial. Recently, \citet{Gouldetal2006} analyzed OGLE~III transit surveys in Galactic fields and concluded that the occurrence frequencies of the detected planets in these surveys is not statistically different from that found in RV surveys of nearby stars. However, they did conclude that the frequency of HJ planets with periods P satisfying 3$<$P$<$5 days ($F =$~1/320(1$^{+1.37}_{-0.59}$) at a 90$\%$ confidence) was statistically different from and larger than that of Very Hot Jupiter (VHJ) planets (1$<$P$<$3 days), which have $F = $1/710(1$^{+1.10}_{-0.54}$) at 90$\%$ confidence. Since the OGLE~III surveys detected no planets with radius larger than 1.3R$_{\rm{Jup}}$,they placed upper limits on the occurrence frequency of larger worlds. Transit searches, unlike RV, are not limited to the immediate solar neighborhood and can be used to measure relative planet frequencies in various regions of the Galaxy, providing information on the role environmental effects play on HJ planet formation. Early predictions for the success of transit searches in open clusters was presented by \citet{J1996}, indicating that with a large amount of telescope time planets could be detected via the transit technique in nearby clusters. More recently, \citet{PG2005} presented an analysis of the prospective harvest of cluster transit surveys by discussing the observational techniques and methods to maximise the chances of a detection. They concluded that due to their mass functions, the most populous, nearby and bright clusters have the greatest chance of yielding a planet. \citet{PG2006} then went on to discuss specifically the detection of short period `Hot Earth' and `Hot Neptune' planets. The detection yields for various nearby clusters with various instruments was estimated. \citet{Ai2007} agree that small-aperture wide field surveys targetting nearby clusters have the potential to discover transiting Hot Neptune planets. Transit searches have been undertaken in bright metal-rich open clusters including STEPPS \citep{Bu2005}, UStAPS \citep{S2003,Ho2005}, PISCES \citep{M2003,Mo2005,Moch2006}, $\it{Monitor}$ \citep{A2007}, EXPLORE-OC \citep{V2005} and in the Praesepe cluster (M44) with KELT \citep{Pepper07}. Searches have also been performed in the general Galactic field \citep{U2002,U2004,H2005,Ka2005,W2007a} and toward the Galactic Bulge \citep{Sahu2006}. If cluster candidates are confirmed as planetary in nature, difficult if fainter than V$\sim$17.0, they can provide information on the timescales of HJ formation and subsequent migration. Null results of high significance allow planet frequency upper limits to be estimated. Globular clusters provide an excellent opportunity to study the effects of environment on planetary frequency. Two bright, nearby southern clusters, 47 Tuc and $\omega$ Centauri, have stars in sufficient numbers ($\sim$10$^{5}$) and brightness (V$\leqslant$17) for meaningful statistics to be gained using ground-based telescopes of moderate aperture. 47 Tuc was previously sampled for planetary transits, resulting in a high significance ($>$3$\sigma$) null result in both the cluster core \citep{G2000} and in the outer halo \citep{W2005}. These two results strongly indicate that system metallicity - not crowding - is the dominant factor determining HJ frequencies in this cluster. This paper presents the results of a dedicated transit search in the second cluster, $\omega$ Centauri, in order to test further the dependence of planetary frequency on stellar metallicity and crowding. Omega Centauri has only 1/10$^{\rm{th}}$ the core density of 47 Tuc yet contains five times the total mass (5.1$\times$10$^6$M$_{\odot}$, \citet{Meylan1995}). Due to its low stellar density compared to other globular clusters and long stellar interaction timescale, a null result for $\omega$ Centauri can be used to test the relative importance of stellar metallicity over density in the formation of giant planets. Omega Centauri ($\omega$ Cen, NGC 5139) has been subjected to intense research over the years. The cluster is unique among globular clusters in that it displays a distinct spread of metallicity among its stars \citep{DW1967,NB1975,Lee1999,P2000,Sol2005}, due to an extended period of star formation and chemical enrichment. Using He abundances, \citet{N2004} has shown that the cluster has three distinct stellar populations, with metallicities of $-$1.7, $-$1.2 and $-$0.6 dex, corresponding to 0.80, 0.15 and 0.05 of the total population respectively. The cluster has a highly-bound, retrograde orbit \citep{D1999} and is by far the most massive of the globular clusters \citep{Meylan1995}. Indeed, these vagaries have led to the theory that the cluster had an external origin, being the left-over remains of a tidally disrupted dwarf galaxy \citep{BF2003,IM2004,BN2005}. With its relative proximity, $\omega$ Cen presents a statistically significant number of upper main sequence stars that can be searched for transiting HJ planets. Here we present the result from a vigorous search for the transit signatures of large planets on 45,406 lightcurves in a 0.75deg$^2$ field centered on $\omega$ Cen. The same set of observations yielded a total of 187 variable stars in the field, 81 of which are new discoveries, and are presented in a companion paper \citep{W2006}. Furthermore, we observed a control field in the Lupus Galactic Plane to test the data reduction and transit identification strategies. Analysis for this field is ongoing, but has led to the identification of several transit candidates, of which none similar were seen in the $\omega$ Cen dataset. One candidate in particular has excellent prospects for being a new Hot Jupiter planet \citep{W2007a}. Section 2 of this paper describes our observational strategy and data reduction details. Section 3 details how the photometry was obtained for both the crowded core regions and outer halo parts of the dataset. The cluster Color Magnitude Diagram dataset (along with astrometry) are also briefly discussed. Section 4 describes the stellar parameters of the cluster stars that were searched for transits, and the expected characteristics of the transits themselves. The total number of stars in the field (both in the cluster and the foreground galactic disk) is also calculated. Section 5 describes the transit detection algorithms used in our search and our removal of systematics in the photometry. Our Monte Carlo simulations to derive expected recovery of real transits and false alarms are described in Section 6, and their application to estimate our HJ sensitivity outlined in Section 7. The results of our transit search in $\omega$ Cen are presented in Section 8, with discussion, comparison to the literature and interpretation in Section 9. We conclude in Section 10. | We have presented the results of a wide-field, deep photometric search for transiting short-perod planets in the globular cluster $\omega$ Centauri, a region previously un-sampled for planetary transits. The cluster was observed with a 52$'$$\times$52$'$ field of view for 25 contiguous nights with the ANU 40-inch telescope at Siding Spring Observatory. From application of difference imaging analysis, a total of 109,726 time series were produced across the field, each being composed of 787 independent data points. A total database of 45,406 stars have photometric accuracy suitable for the search ($\le$0.045 mag scatter down to V$=$19.5), including 31,000 cluster stars extending 2.5 magnitudes down the main sequence. All of these were subjected to a rigorous (and vigorous) search for transit-like events; none were detected after variable stars and clear false-positives were removed. Simulations have shown that if large Hot Jupiters (HJs) formed in the cluster then dynamically speaking they would survive to be detectable in our data. Extensive Monte Carlo simulations via injection of transit signals into actual light curves were used to determine the sensitivity of the survey to R$\le$1.5R${\rm_{Jup}}$ planets over a range of orbital periods. Coupled with our null result, we are thus able to place strict, statistically significant upper limits on the occurrence frequency $F$ of large (R~$=$~1.5R), short-period planets in $\omega$ Centauri. We determine a limit of $F_{\rm VHJ} < 1/1040$ at 95$\%$ confidence for Very Hot Jupiter (VHJ) planets with periods distributed uniformly over 1d~$<$~P~$<$~3d. This upper limit for the cluster is less than that determined by \citet{Gouldetal2006} for smaller (1.3R${\rm_{Jup}} < $~R~$< 1.5$R${\rm_{Jup}}$) planets with the same period distribution in the Galactic fields surveyed by OGLE~III. The two results are consistent at the 90$\%$ confidence level, and more understandable if the low metallicity of $\omega$ Cen suppresses planet formation or planetary migration. Under the assumption that there is no difference in occurrence frequency for VHJ and HJ across the orbital period range 1d~$<$~P~$<$5d, we derive an upper limit of $F_{\rm VHJ+HJ} < 1/600$ in $\omega$ Cen. The corresponding result in the Galactic OGLE~III fields for comparably-sized planets is an upper limit of $F < 1/640$. Both results are quoted at 95$\%$ confidence. Our results are less dependent on model assumptions about the distance to the target population since the vast majority of stars in our fields are members of, and thus at the distance of, $\omega$~Cen. It is noteworthy that despite the fact that the OGLE~III campaigns monitored considerably more stars in better median seeing with more frames per field, our $\omega$ Cen study produces a comparable upper limit on the frequency of large, short-period planets. While part of the reason may lie in the longer exposure times, denser sampling, and the use of different cleaning and detection algorithms in our survey, a large part of the difference is due to the small fraction of the more distant and obscured Galactic bulge stars (which constituted about a third of the total OGLE~III sample) that can be meaningfully probed for transiting planets. This null result for VHJ and HJ planets in $\omega$ Cen, coupled with the null result of 47 Tucanae \citep{W2005} strengthens the evidence for the dominance of system metallicity over stellar interactions in determining short period planetary frequencies in globular clusters. At longer orbital periods stellar encounters may play a role in determining planetary frequencies. This is a result aligned with current work on the metallicity trend of planet-bearing host stars in the Solar Neighborhood and N-body simulations of planets in dense environments. Such a metallicity dependence is one of the main predictions of the core accretion model of planet formation. | 7 | 10 | 0710.3461 |
0710 | 0710.2128_arXiv.txt | We have resolved the classical nova V1663 Aql using long-baseline near-IR interferometry covering the period from $\sim$5--18 days after peak brightness. We directly measure the shape and size of the fireball, which we find to be asymmetric. In addition we measure an apparent expansion rate of $0.21 \pm 0.03$ ${\rm mas\,day^{-1}}$. Assuming a linear expansion model we infer a time of initial outburst approximately 4 days prior to peak brightness. When combined with published spectroscopic expansion velocities our angular expansion rate implies a distance of $8.9\pm3.6$ kpc. This distance measurement is independent of, but consistent with, determinations made using widely available photometric relations for novae. | Novae are violent stellar explosions exceeded in energy output only by $\gamma$-ray bursts and supernovae. They are erratic outbursts that occur in systems containing a white dwarf accreting mass from a late-type stellar companion (e.g Prialnik \& Kovetz\nocite{pk05} 2005). When the amount of accreted material on the surface of the white dwarf reaches some critical value a thermonuclear-runaway is ignited, giving rise to the observed nova outburst in which material enriched in heavy elements is ejected into the surrounding medium at high velocities. For certain elements, this ejected material may influence observed abundances in the ISM \citep{gehrz98,hernanz05}. Direct observations of the expansion of the nova shell provide an opportunity to accurately determine the distance to the nova. Such observations are usually only possible many months or years after the outburst, when the expanding shell can be resolved \citep{bode02}. However, several optical/IR interferometers are now capable of resolving novae and other explosive variables, allowing detailed studies of the initial "fireball", \citep{ches07,mon06} as well as later stages of development \citep{lane05,lane07}. Nova Aquilae 2005 (ASAS190512+0514.2, V1663 Aql) was discovered on 9 June 2005 by \citet{iauc8540}. At the time of discovery the magnitude was $m_V$ = 11.05; the source reached $m_V \sim 10.8$ the following day, and declined in brightness thereafter. A possible progenitor near the source coordinates (sep. $\sim$ 4.5 arcseconds) is seen on Palomar Optical Sky Survey plates (USNO-B1.0 0952-00410569, \citet{usno-b1}), with magnitudes $m_R \sim 18.1$ and $m_I \sim 16.45$. Soon after discovery \citet{iauc8544} obtained an optical spectrum with features indicating a heavily reddened, peculiar nova. H-$\alpha$ emission lines exhibited P Cygni line profiles and indicated an expansion velocity of $700 \pm 150 {\rm km\,s^{-1}}$ \nocite{iauc8544} (Dennefeld et al. 2005, confirmed via personal communication), somehwhat slow for a classical nova, but not outside the range of observed values. Recently, \citet{pog06} published spectra and analysis of published light-curves of this nova, deriving a distance in the range 7.3--11.3 kpc, and an expansion velocity of $\sim2000$~${\rm km\,s^{-1}}$. We have used the Palomar Testbed Interferometer (PTI) to resolve the $2.2 \mu$m emission from V1663 Aql and measure its apparent angular diameter as a function of time. We are able to follow the expansion starting $\sim 9$ days after the initial explosion; when combined with radial velocities derived from spectroscopy we are able to infer a distance to, and luminosity of, the object. We compare this result with values inferred by a maximum magnitude-rate of decline (MMRD) relation in the literature (see Poggiani 2006 for a summary). The Palomar Testbed Interferometer (PTI) was built by NASA/JPL as a testbed for developing ground and space-based interferometry and is located on Palomar Mountain near San Diego, CA \citep{col99}. It combines starlight from two out of three available 40-cm apertures and measures the resulting interference fringes. The high angular resolution provided by this long-baseline (85-110 m), near infrared ($2.2 \mu$m) interferometer is sufficient to resolve emission on the milli-arcsecond scale. | We have used long-baseline near-IR interferometry to resolve the classical nova V1663 Aql, starting $\sim 9$ days after outburst. We measure an apparent expansion rate of $0.21 \pm 0.03$ ${\rm mas day^{-1}}$, which can be combined with previously determined expansion velocities to produce a distance estimate to the nova of $8.9\pm3.6$ kpc; the precision is limited by the precision of the available spectroscopic radial velocities. Such a large distance is consistent with the large reddening ($E(B-V) \sim 1.1$) determined from photometry, as well as with distances found from MMRD relations. This represents only the third time a nova has been resolved using optical/IR interferometry, the previous cases being Nova V1974 Cyg 1992 \citep{q93} and more recently RS Oph \citep{mon06,lane07}. We anticipate that further instrumental improvements, in particular high spectral resolution interferometry such as that recently deployed on the Keck Interferometer \citep{eis07}, may break the inclination degeneracy and thus yield very precise expansion distances to these interesting objects. | 7 | 10 | 0710.2128 |
0710 | 0710.0407_arXiv.txt | The ``galactic shocks'' \citep{fujimoto68,roberts69} is investigated using a full three-dimensional hydrodynamic simulations, taking into account self-gravity of the ISM, radiative cooling, and star formation followed by energy feedback from supernovae. This is an essential progress from the previous numerical models, in which 2-D isothermal, non-self-gravitating gas is assumed. We find that the classic galactic shocks appears is unstable and transient, and it shifts to a globally quasi-steady, inhomogeneous pattern due to non-linear development of instabilities in the disk. The spiral patterns consists of many GMC-like dense condensations, but those local structures are not steady, and they evolves into irregular spurs in the inter-arm regions. Energy feedback from supernovae do not destroy the quasi-steady spiral arms, but it mainly contributes to vertical motion and structures of the ISM. The results and methods presented here are a starting point for more consistent treatment of the ISM in spiral galaxies, in which effects of magnetic fields, radiative transfer, chemistry, and dynamical evolution of a stellar disk are taken into account. | A standing spiral shock solution was first discovered in the 60s numerically in a rotating gas disk in the galactic potential with a tightly wound spiral perturbation (i.e. the ``pitch angle'' is very small). Since then in the most of theoretical studies on the galactic ``shock'' \citep[e.g.][]{fujimoto68,roberts69,wood75,johns86,lubow86}, the interstellar medium (ISM) is treated as a isothermal and homogeneous fluid with two-dimensional approximation, or it is modeled as a $N-$body system of small cloudlets \citep[e.g.][]{tomisaka86}. Global evolution of the spiral shock was studied in the 80s \citep{johns86} using time-dependent, two-dimensional (2D) hydrodynamic simulations show that spiral shocks are stable and long-standing for various pitch angles. However, the steady, smoothed galactic shock does not consistently explain the complicated distribution of the ISM around spiral arms and the inter-arm substructures akin to so-called `spurs' or `feathers'\footnote{In this paper, we refer the terms 'spurs' or 'feathers' as inter-arm gas structures, which often show quasi-periodic features associated with main spiral arms. See Paper I, \citet{shetty06}, and \citet{kim02} for numerical examples, and \citet{vigne06} for observations. On the other hand, `blanches' are sub-structures bifurcated from main spiral arms. They are smoother and longer azimuthal extent than spurs, which may be caused by resonances \citep{chark03}.} observed in real spiral galaxies \citep{elm80,scov01,calz05,vigne06}. Moreover, observed molecular clouds in galactic disks do not match the picture of hydrodynamic shocks in a uniform media. It was however shown by full 2D global simulations of a non-self-gravitating, isothermal gas disk that the spurs are in fact natural consequences of ``wiggle'' instability, which is caused by a purely hydrodynamic phenomenon, i.e. Kelvin-Helmholtz instability \citep{wadakoda04} (hereafter Paper I). This phenomenon was also found in \citet{shetty06}. They also pointed out that the features only grow in the inner-most several kpc regions if self-gravity and magnetic fields are ignored. More recently, three-dimensional (3D) response of the gas to the spiral potential was modeled using a local shearing box approximation in isothermal, MHD simulations taking into account self-gravity \citep{kim06}, and found that the wiggle instability is suppressed by radial flapping motion of the shock. These previous results suggest that hydrodynamic effects, self-gravity of the gas and magnetic fields play some important roles on the gas structures in a spiral potential. However, an important feature of the real ISM has been ignored; its inhomogeneous multi-phase structures. Apparently the ISM is not `isothermal fluid', nevertheless it is assumed in most HD and MHD simulations\footnote{\citet{dobbs07} recently studied gas dynamics in a spiral potential, taking into account multi-phase nature of the ISM. However, in their non-self-gravitating SPH simulations, an energy equation with realistic cooling and heating processes is not solved, alternatively warm ($10^4$ K) and cold ($100$ K) components are treated separately as isothermal gases without phase exchange.}. Effects of self-gravity and magnetic fields highly depend on gaseous temperatures and phases. In this sense, an energy equation with a realistic cooling and heating processes should be solved before taking into account those effects. Local box-approximation with a shearing periodic boundary \citep[e.g.][]{kim06} is not necessarily relevant for representing dynamics of the multi-phase ISM in galactic disks, because typical scales of the inhomogeneous structures of the ISM are not small enough compared to the disk size \citep[cf.][]{wada01, wada07, tasker06}. Moreover, since gravity is a long-range force, global, 3D simulations for the whole disk are essential if self-gravity of the gas is considered. This is also necessary for the multi-phase ISM, because the scale height and velocity dispersion are different for cold and hot gases. In 3D hydrodynamic simulations of self-gravitating gas disks with the radiative cooling, we should take into account energy feedback from stars, otherwise the cold gas disk becomes unrealistically thin and the gravitational instability is strongly affected. Here we show, for the first time, 3D evolution of the ISM in a galactic spiral potential, taking into account realistic radiative cooling and energy feedback from stars, especially from supernovae explosions. The simulations are performed for the whole gas disks without any assumptions for symmetry. In order to ensure a high spatial resolution (10 pc), which is essential to reproduce the multi-phase ISM, we here focus on a central part of a relatively small galaxy (radius is 2.56 kpc and the maximum circular velocity is 150 km s$^{-1}$). However, this is large enough to see the global effect of the spiral potential on the inhomogeneous ISM. Our simulations clarify apparent discrepancy between the steady solution of ``galactic shock'' and the complicated, non-steady structures of the ISM in real spiral galaxies. This is a major progress from the previous simulations, and it will be a starting point for more realistic numerical models of the ISM, taking into account `lived' stellar potential, magnetic fields, UV radiation from massive stars, and chemistry of molecules and atoms. | The high resolution hydrodynamic simulations, taking into account self-gravity of the gas and realistic cooling and heating processes for the ISM, reveal for the first time that the galactic spiral arms of the ISM are neither hydrodynamic shock waves nor an assembly of long-lived, bullet-like cloudlets. The global spiral arms are consist of complicated time-dependent substructures from which stars can be formed, but over a long time (at least 5-6 rotational periods), they stably exist under the influence of the spiral potential. The pseudo-spiral in the multi-phase ISM is robust for energy input from the supernovae, which mainly cause the vertical non-uniform structure of cold and hot gases. The pattern speed of the spiral potential and its strength are not key parameters to alter above features. The ISM in galactic disks has often been represented by isothermal, non-self-gravitating fluid or inelastic particles in many astrophysical simulations. Moreover, introducing a periodic boundary condition and reducing spatial dimensions were usual. However, those kinds of simplification do not necessarily represent the nature of spiral galaxies. Magnetic fields are not considered in the present simulations, but the present results suggest that even if the magnetic fields are weak, the spiral patterns can steadily exist on a global scale. Effect of magnetic fields and other important physical processes, such as UV radiation and chemistry, should be investigated based on 3D global models, as presented in this paper. The present treatment of the multi-phase ISM could be used for direct comparison with observations, coupled with radiative transfer calculations for various observational probes \citep[e.g.][]{wada00,wada05}. Fine structures of molecular gas associated with the spiral arms and in the inter-arm regions in external galaxies, which could be compared with the present numerical model, will be revealed by high resolution observations, e.g. by the Atacama Large Millimeter/Submillimeter Array. | 7 | 10 | 0710.0407 |
0710 | 0710.0588_arXiv.txt | Outflows from active galactic nuclei (AGNs) seem to be common and are thought to be important from a variety of perspectives: as an agent of chemical enhancement of the interstellar and intergalactic media, as an agent of angular momentum removal from the accreting central engine, and as an agent limiting star formation in starbursting systems by blowing out gas and dust from the host galaxy. To understand these processes, we must determine what fraction of AGNs feature outflows and understand what forms they take. We examine recent surveys of quasar absorption lines, reviewing the best means to determine if systems are intrinsic and result from outflowing material, and the limitations of approaches taken to date. The surveys reveal that, while the fraction of specific forms of outflows depends on AGN properties, the overall fraction displaying outflows is fairly constant, approximately 60$\%$, over many orders of magnitude in luminosity. We emphasize some issues concerning classification of outflows driven by data type rather than necessarily the physical nature of outflows, and illustrate how understanding outflows probably requires more a comprehensive approach than has usually been taken in the past. | The role of outflows from quasars and active galactic nuclei (AGN) has recently become an important feature in the overall framework of how galaxies and star formation processes evolve over cosmic time. Mergers and other interactions triggering AGN seem to provide feedback affecting the larger scale environment. Recent efforts to include the effects of this so-called AGN feedback focus on two modes: a ``radio'' mode whereby a relativistic jet heats the surrounding interstellar and intercluster media \citep[e.g.,][]{best07}, and a ``quasar'' mode whereby a lower velocity but higher mass outflow also helps to clear out post-merger shrouding gas and quenches star formation \citep*[e.g.,][]{tdm05}. We focus on this second mode in this paper. For this mode, a number of questions require addressing. How common are outflows? Do all AGN have outflows? What drives outflows? Is there a single all-governing structure of AGN? Answering these questions will help us to understand the role AGN outflows with respect to issue of feedback, and other important issues like chemical enrichment and accretion. In the ensuing sections we aim to achieve several goals: (1) to review the ways in which outflows are detected in AGN over all luminosity scales; (2) to comment on the merits of various catalogs of outflows; and (3) to arrive at the true (possibly property-dependent) observed frequency of outflows. In its most basic interpretation, the observed frequency of outflows can be equated with the fraction of solid angle (from the view point of the central black hole) subtended by outflowing gas. This interpretation assumes that all AGN feature outflows and that not all sight-lines to the emitting regions are occulted by the outflow. Alternatively (and equally simplistic), the frequency can be interpreted as the fraction of the duty cycle over which AGN feature outflows (assuming the outflow subtends 4$\pi$\ steradians). The actual conversion of the fraction of AGN featuring spectroscopic evidence of outflow to the solid angle subtend by such outflows has been treated by \citet{cren99} and \citet*{cren03b}. This computation involves further knowledge of the line-of-sight covering factor (that is, the fraction of lines-of-sight that reach the observer that are occulted by the outflow) as well as an understanding of range of solid angle sampled by the AGN used (e.g., Type 1 versus Type 2 AGN). The true situation is likely in between these two extremes, and may depend also on properties we can not currently constrain, such as the time since the AGN was triggered. Additionally, we strive here to build a case that more effort should be made to consider outflows of all types together. Often data limitations of one sort or another have led to the study of limited parts of parameter space (e.g., outflow velocity or velocity dispersion), creating artificial or at least biased divisions. There appears to be a continuous range in properties of outflows and these should only be regarded as fundamentally different when there is clear evidence to reach such a conclusion. Below we discuss the identification of outflows (\S 2) and the data-driven subcategories (\S 3). We show an illustrative example of how combining the different outflow subclasses may lead to a more unified physical understanding of outflows (\S 4). Finally, we bring together the different survey methodologies to determine an overall fraction of AGN displaying the signatures of outflows (\S 5) and summarize the case for more global studies of the outflow phenomenon. We adopt a cosmology with $\Omega_M$ = 0.3, $\Omega_{\Lambda}$ = 0.7, and H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$. | 7 | 10 | 0710.0588 |
|
0710 | 0710.0893_arXiv.txt | We use two-band imaging data from the Advanced Camera for Surveys on board the Hubble Space Telescope for a detailed study of NGC\,1533, an SB0 galaxy in the Dorado group surrounded by a ring of \HI. NGC\,1533 appears to be completing a transition from late to early type: it is red, but not quite dead. Faint spiral structure becomes visible following galaxy subtraction, and luminous blue stars can be seen in isolated areas of the disk. Dust is visible in the color map in the region around the bar, and there is a linear color gradient throughout the disk. We determine an accurate distance from the surface brightness fluctuations (SBF) method, finding $\mM=31.44\pm0.12$ mag, or $d = 19.4\pm1.1$ Mpc. We then study the globular cluster (GC) colors, sizes, and luminosity function (GCLF). Estimates of the distance from the median of the GC half-light radii and from the peak of the GCLF both agree well with the SBF distance. The GC specific frequency is $S_N=1.3\pm0.2$, typical for an early-type disk galaxy. The color distribution is bimodal, as commonly observed for bright galaxies. There is a suggestion of the redder GCs having smaller sizes, but the trend is not significant. The sizes do increase significantly with galactocentric radius, in a manner more similar to the Milky Way GC system than to those in Virgo. This difference may be an effect of the steeper density gradients in loose groups as compared to galaxy clusters. Additional studies of early-type galaxies in low density regions can help determine if this is indeed a general environmental trend. | The Hubble Space Telescope (\hst) has opened the door to our understanding of extragalactic star cluster systems, revealing numerous globular clusters (GCs) in early-type galaxies (e.g.\ Gebhardt \&\ Kissler-Patig, 1999; Peng \etal\ 2006) as well as ``super-star clusters'' (SSCs) in late-type galaxies (Larsen \&\ Richtler 2000), especially starbursts (e.g. Meurer \etal\ 1995, Maoz \etal\ 1996). In early-type systems the color distribution of the GCs is often bimodal, consisting of a blue metal-poor component and a red metal-rich component (e.g. West \etal\ 2004; Peng \etal\ 2006). This observation gives a hint to the connection between the early- and late-type systems. Mergers often have particularly strong starbursts and rich populations of SSCs, as seen for example in NGC4038/39 - ``the Antennae'' system (Whitmore \&\ Schweizer 1995; Whitmore \etal\ 1999), and hierarchical merging is one possible origin for the redder population of GCs in early-type galaxies (e.g., Ashman \& Zepf 1998; Beasley \etal\ 2002; Kravtsov \& Gnedin 2005). Much of the research on GC systems has concentrated on galaxy clusters which are rich in early-type galaxies (e.g., the ACS Virgo Cluster Survey, C{\^o}t{\'e} \etal\ 2004; the ACS Fornax Cluster Survey, Jord{\'a}n \etal\ 2007). Early-type galaxies in groups and the field are somewhat less studied, particularly with \hst\ and the Wide Field Channel (WFC) of its Advanced Camera for Surveys (ACS). The relatively wide ($3\farcm4$) field of view of the ACS WFC combined with its fine pixel sampling make it an exceptional tool for imaging GCs out to a few tens of Mpc where they have measurable angular sizes (Jord{\'a}n \etal\ 2005). Here we report \hst\ ACS/WFC imaging of the SB0 (barred lenticular) galaxy NGC~1533 in the Dorado group. This group is in the ``Fornax wall'' (Kilborn \etal\ 2005) and hence at a similar distance to the Fornax cluster (e.g., Tonry \etal\ 2001). Dorado is interesting in that it is richer than the Local Group but still dominated by disk galaxies (its brightest members being the spiral NGC~1566 and the S0 NGC~1553), and its members have \HI\ masses similar to non-interacting galaxies with the same morphology (Kilborn \etal\ 2005). While the apparent crossing time of the group is only $\sim 13$\%\ of the age of the universe (Firth \etal\ 2006; see also Ferguson \&\ Sandage 1990), the most recent analyses conclude the group is unvirialized (Kilborn \etal\ 2005; Firth \etal\ 2006), which may explain the richness in spirals and \HI. \begin{figure*} \plottwo{f1a.eps}{f1b.eps} \caption{\textit{Left panel:} \HI\ contours from Ryan-Weber \etal\ (2003) are overlaid on a ground-based $R$-band image of the NGC\,1533 field from the SINGG survey (Meurer \etal\ 2006). The outermost contour (bold) is at a column density of $10^{20}$ cm$^{-2}$, and the contours increase in steps of $0.5{\times}10^{20}$ cm$^{-2}$. The small companion galaxies IC~2038/2039 are in the upper right corner of the image. NGC\,1533 itself is in an \HI\ ``hole'' (the galaxy center is not detected), and this distribution has been described as a ring. \textit{Right panel:} ACS HRC (green) and WFC (blue) fields of view for the two HST roll angles described in the text (labeled 1 and 2). The outlines of the camera fields are overlaid on a $\sim\,$9\arcmin\ portion of the SINGG $R$-band image. North is up and East is to the left in both panels. } \label{fig:ACSfields} \end{figure*} NGC~1533 is the seventh brightest member of the Dorado group, with $M_V{\,\approx\,}{-}20.7$. It lies within the virial radius, but is a $\sim\,$2-$\sigma$ velocity outlier (Kilborn \etal, 2005; Firth \etal\ 2006) so that it is moving at high speed through the intra-group medium. A vast \HI\ arc is seen in the outskirts of NGC~1533 connected to the Sdm galaxy IC~2038 and the small S0 galaxy IC~2039 (Ryan-Weber \etal\ 2004). This suggests that NGC~1533 is ``stealing'' ISM from its companions or has cannibalized another gas-rich satellite (see Figure~\ref{fig:ACSfields}, left panel). As is typically seen in S0 galaxies, star formation is weak in NGC~1533. Observations of this galaxy in spectroscopic surveys note the presence of emission lines (Jorgensen \etal\ 1997; Bernardi \etal\ 2002); the nuclear spectrum available from the 6dF survey (Jones \etal\ 2005) shows [{\sc N ii}]6584 and weak \Halpha. \Halpha\ imaging from the Survey of Ionization in Neutral Gas Galaxies (SINGG, an \Halpha\ imaging survey of \HI\ selected galaxies; Meurer \etal\ 2006) shows a few weak \HII\ regions beyond the end of its bar (the nucleus is too bright to allow faint nuclear \HII\ regions to be detected in the SINGG images), as well as a scattering of very faint ``intergalactic \HII\ regions''. These are discussed in more detail by Ryan-Weber \etal\ (2004) who show that they are so faint that it would only take one to a few O stars to ionize each one. Although its current rate is low, the star formation in NGC~1533 illustrates another possible channel for building up cluster systems in early-type galaxies: slow re-ignited star formation in ISM stripped from companions. The ACS WFC images of NGC~1533 used in the present study were obtained with \hst\ as ``internal parallel images'' while the ACS High Resolution Channel (HRC) was pointed at the intergalactic \HII\ regions (\hst\ GO Program 10438; M. Putman, PI). The HRC observations are discussed elsewhere (Werk \etal\ 2007, in preparation). Here we use the WFC observations to measure the structural properties of the galaxy, characterize its GC population, and use the GC luminosity function (GCLF), GC half-light radii, and surface brightness fluctuations (SBF) to provide accurate distance estimates. The contrast between NGC~1533, a (weakly) star-forming gas-rich barred S0 in a loose group environment, and galaxies in the richer environments of the Virgo and Fornax cluster, provides a useful test of the ubiquity of the various relations found in the denser environments. The following section describes the observations and data reductions in more detail. Sec.~\ref{sec:props} discusses the galaxy morphology, structure, color profile, and isophotal parameters. Sec.~\ref{sec:sbf} presents the SBF analysis and galaxy distance, while Sec.~\ref{sec:gccolors}--\ref{sec:gclf} discuss the GC colors, effective radii, luminosity function, and specific frequency. The final section summarizes our conclusions. \begin{figure*} \epsscale{1.0} \plotone{f2.eps} \caption{\textit{Upper left:} Combined F814W ACS/WFC image of NGC~1533. \textit{Lower left:} Contour map of a 3\farcm6$\times$4\farcm0 portion of the image. Contours are plotted in steps of a factor of two in intensity, with the faintest being at $\mu_I = 20.7$ mag~arcsec$^{-2}$. \textit{Upper right:} The image following galaxy model subtraction, showing the faint spiral structure (the ``plume'' 2\arcmin\ north of the galaxy is a ghost image). The same 3\farcm6$\times$4\farcm0 field is shown; the box marks the central~1\arcmin. \textit{Lower right:} Colormap of the central 1\arcmin\ region of NGC1533, with dark indicating red areas and white indicating blue. The dark spot at center marks the center of the galaxy. Dust can be seen as faint, dark, wispy features. A compact blue star-forming region is visible to the left of the galaxy center, near the center-left of the map. } \label{fig:cmb4} \end{figure*} | We have analyzed deep F606W and F814W images of the galaxy NGC\,1533 and its GC population taken at two roll angles with the ACS/WFC on \hst. Although it is classified as an early-type barred lenticular galaxy, we found faint spiral structure once a smooth fit to the galaxy isophotes was subtracted. The color map shows faint dust features in the area around the bar and inner disk. Previous ground-based \Halpha\ imaging had shown that the galaxy disk contains several faint, compact \HII\ regions. We find that all of these regions have some blue stars associated with them. Four of the \HII\ regions lie within one of the faint spiral arms, and several other blue stars are spread out within the arm. These observations suggest that NGC\,1533 is in the late stages of a transition in morphology from type SBa to~SB0. Transition objects such as NGC\,1533 may be the key to understanding the evolution of the morphology-density relation in galaxy clusters, which is often explained as infalling spirals being transformed into S0s by the harsh cluster environment (e.g., Dressler \etal\ 1997; Postman \etal\ 2005). However, recent evidence at intermediate redshift indicates that the transitions begin outside the clusters in small group environments through galaxy-galaxy interactions (Moran \etal\ 2007). Following infall, the intra-cluster medium then serves to expedite the process by removing the remaining cool gas. With NGC\,1533, we have a close-up view of this transition in the Dorado group, including interaction with the neighboring galaxies IC~2038/2039 (Ryan-Weber \etal\ 2004). From two-dimensional two-component parametric modeling of the galaxy surface brightness, we find a bulge-to-total ratio $B/T\approx0.42$. The half-light radii of the bulge and disk are $\sim\,$7\arcsec\ and $\sim\,$46\arcsec, respectively. We find a best-fitting S\'ersic index $n=2.0$ for the bulge, which can be reasonably approximated by an $r^{1/4}$ law in the 1-D profile. However, the disk has a relatively flat profile over a factor-of-three in radius, from $\sim\,$15\arcsec\ to $\sim\,$45\arcsec, then steepens fairly abruptly beyond $\sim\,$50\arcsec. This gives the disk a very low S\'ersic index of $n\approx0.4$, which might result from past high-speed interactions of NGC\,1533 within the group environment. Overall, the color of NGC\,1533 is that of an evolved, red population, except in the few, small isolated regions where the blue stars occur. The bulge color is $\vi\gta1.22$, similar to cluster ellipticals, and then there is a mild, but significant, linear color gradient throughout the disk. There is a gradual isophotal twist and the isophotes increase in ellipticity out to a semi-major axis distance of 24\arcsec, where $\epsilon$ goes above 0.4 before falling sharply again towards the round outer disk. The peak of the $A_4$ harmonic term, measuring ``diskiness,'' occurs at a smaller semi-major axis of 21\arcsec. This is because the pointed lens-like isophotes occur inside of the bar. We measured the SBF amplitude in four broad radial annuli for each of the two observations at different roll angles. A gradient in the SBF amplitude is clearly detected and follows the color gradient (the bluer outer regions have relatively brighter SBF). By matching our ACS photometry against ground-based \vi\ data for this galaxy, we have accurately calibrated the SBF measurements to obtain distance moduli. We find excellent agreement among the different annuli but with an offset of 0.04 mag in distance between the two observations. However, the distance error is dominated by systematic uncertainty in the color and calibration zero point. We find a final distance modulus $(m{-}M) = 31.44 \pm 0.12$ mag, or $d = 19.4 \pm 1.1$ Mpc. Candidate globular clusters were selected according to color, magnitude, radial position, and FWHM. Analysis of the color distribution of these objects with the KMM algorithm indicates with a very high degree of confidence that the distribution is bimodal with peaks at $\vi\approx0.92$ and 1.22. There is no evidence that the blue GCs become redder at bright magnitudes, the so-called ``blue tilt.'' The absence of this effect in NGC\,1533, an intermediate luminosity galaxy with a small GC population, is consistent with a self-enrichment explanation, since the GCs in such systems do not reach the high masses that they do in richer systems. The sizes of the GC candidates were measured using the Ishape software. By comparing the results from the two different roll angles, we found that the effective (half-light) radii \reff\ have an accuracy of about 13\% down to $\iacs=23$, but are not reliable beyond this. We did not find a significant trend of \reff\ with GC color, although the red-peak GCs have a median \reff\ smaller by $11\pm8$\% than the blue-peak GCs. However, we did find a significant (4\,$\sigma$) trend of \reff\ with galactocentric radius. In this respect, NGC\,1533 is more like the Milky Way than the Virgo early-type galaxies. This may be an effect of the environment: since the sizes of the GCs are limited by the tidal field, and the density gradients will be steeper in small groups such as Dorado or the Local Group, GC sizes should have a stronger dependence on radius in such environments. The dominance of this radial effect may weaken or obscure any relation between size and color. More studies of size and color trends for the GCs of galaxies in loose groups are needed to verify this hypothesis, although this may be difficult because of the low GC populations in such systems. We then used the median half-light GC radius to obtain another estimate of the distance to NGC\,1533. Following Jord\'an \etal\ (2005), we find $d = 18.6\pm2.0$ Mpc, in good agreement with the SBF distance. We modeled the $I_{814}$-band GCLF of NGC\,1533 as a Gaussian using a maximum likelihood fitting routine. The best-fit peak magnitude $m^0_I=22.84^{+0.18}_{-0.24}$ corresponds to $M^0_I \approx -8.6$ for the measured SBF distance, in good agreement with expectations based on other galaxies. The fitted Gaussian dispersion of $\sigma_{\rm LF} = 1.10 \pm 0.15$ mag is in accord with the relation between $\sigma_{\rm LF}$ and galaxy luminosity found recently by Jord\'an \etal\ (2006) for Virgo galaxies. Finally we estimate the GC specific frequency in the analysis region to be $S_N = 1.3\pm0.2$, typical for a disk galaxy. We conclude that the GCs in NGC\,1533 have the same average size, color, and luminosity within the errors as the Virgo early-type galaxies, but the stronger dependence of size on galactocentric distance is more reminiscent of the Milky Way. NGC\,1533 represents an interesting class of transitional objects, both in terms of morphology and environment. A large, multi-band, systematic study of such systems with \hst, similar to the ACS Virgo and Fornax Cluster surveys but focusing on group galaxies, has yet to be undertaken and must await either a revived ACS or Wide Field Camera~3. Such an effort would be extremely valuable in piecing together a more complete picture of the interplay between galaxy structure, globular cluster system properties, and environment. | 7 | 10 | 0710.0893 |
0710 | 0710.1547_arXiv.txt | We present our new advanced model for population synthesis of close-by cooling NSs. Detailed treatment of the initial spatial distribution of NS progenitors and a detailed ISM structure up to 3 kpc give us an opportunity to discuss the strategy to look for new isolated cooling NSs. Our main results in this respect are the following: new candidates are expected to be identified behind the Gould Belt, in directions to rich OB associations, in particular in the Cygnus-Cepheus region; new candidates, on average, are expected to be hotter than the known population of cooling NS. Besides the usual approach (looking for soft X-ray sources), the search in 'empty' $\gamma$-ray error boxes or among run-away OB stars may yield new X-ray thermally emitting NS candidates. | More than 10 years after the discovery of its brightest member RX J1856-3754 \cite{Walter1996}, a group of seven radio-quiet isolated neutron stars (NSs) detected by ROSAT gained an important place in the rich zoo of compact objects. Together with Geminga and several close-by young radio pulsars, these objects form the local population of cooling NSs. Studies of this group of sources already provided a wealth of information on NSs physics (see e.g. \cite{Haberl2007,Page2007,Zane2007} for recent reviews). Since 2001 the number of known close-by radio-quiet NSs has not been growing despite all attempts to identify new candidates. Partly this is due to the fact that all these searches are blind. To advance the identification of new near-by cooling NSs it is necessary to perform a realistic modeling of this population. To investigate the population of close-by young cooling NSs the method of population synthesis is used here. In this short note an advanced population synthesis model is briefly discussed for the population of close-by ($< 3$~kpc) isolated NSs which can be observed via their thermal emission in soft X-rays (the detailed description and full analysis of new results will be presented elsewhere \cite{Posselt2008}). Previously our models were applied to confirm the link between the seven radio quiet NSs (the Magnificent Seven) and the Gould Belt \citep{p03} (Paper I), to study distribution of NSs in the Galaxy and in the solar vicinity \cite{p04} (Paper II), and to test theories of thermal evolution of NSs \citep{pgtb04} (Paper III). The major interest of the present study is to get a hint how to find more objects of this type. | The main aim of this study is to make some advances in the strategy for searching for new isolated cooling NSs. According to our results, new candidates expected to be identified at ROSAT count rates $<0.1$~cts~s$^{-1}$ should be young objects born in rich OB associations behind the Gould Belt. Most of the recent studies \cite{Agueros2006,Chieregato2005,Rutledge2003} looked for new candidates far from the galactic plane. It seems that this is not very promising. Our results indicate that new cooling NSs should be searched in directions of OB associations such as Cyg OB7 and Cep OB3. Considering sky coverage the ROSAT All Sky Survey is currently the best choice to look for new ``cowboys'' in the Cygnus-Cepheus region which is, according to our results, the most promising area. However, the relatively large positional error circle of ROSAT usually includes many possible optical counterparts, especially at these low galactic latitudes. Furthermore one has to exclude variable X-ray sources to find isolated cooling NSs. In this respect the recently published XMM-$Newton$ Slew Survey may become an important database. A major step can be expected from the planned eROSITA all sky survey which - compared with the ROSAT all sky survey -- will provide a factor of $\sim$~10 in soft X-ray sensitivity and factor of $\sim$~4 in energy resolution \cite{Predehl2006}. Some of unidentified $\gamma$-ray sources (already observed by EGRET and forthcoming due to AGILE and GLAST) can be identified as cooling NSs as it was with Geminga and 3EG J1835+5918. In particular, GLAST observations of the 56 EGRET error boxes studied in \cite{Crawford2006} and later cross-correlation with the ROSAT (or/and XMM) data can result in new identification of cooling NSs. Another possibility to find new isolated coolers is to search for (un)bound compact companions of OB runaway stars. More than one hundred OB runaway stars are known in a 1 kpc region around the Sun \citep{Zeeuw1999}. They are characterized by large spatial velocities or/and by large shifts from the galactic plane. Two main origins of these large velocities are currently discussed: dynamical interaction and explosion of a companion in a close binary system. The latter case is interesting for the discussion of search for new close-by cooling NSs. A binary can survive after the first SN explosion in, roughly, 10-20\% of cases. Then one expects to have a runaway system consisting of an OB star and a compact object (most probably a NS). A young NS can appear as a radio pulsar. In \cite{sayer1996} and \cite{philp1996} the authors searched for radio pulsar companions of $\sim 40$ runaway OB stars. Nothing was found. This result is consistent with the assumption that in less than 20\% cases OB stars have radio pulsar companions. Still, it is interesting to speculate that runaway massive stars can have cooling radio quiet NS as companions. Then a companion can be identified as a source of additional X-ray emission. \begin{theacknowledgments} We thank D. Blaschke, H. Grigorian, and D. Voskresensky for data on cooling curves and discussions; A. Mel'nik for discussion of properties of OB associations; R. Lallement for the sodium data; and A. Pires for discussions about the ISM model. S.B.P. was supported by INTAS and Dynasty foundations. \end{theacknowledgments} | 7 | 10 | 0710.1547 |
0710 | 0710.1637_arXiv.txt | A simple speed-up cosmology model is proposed to account for the dark energy puzzle. We condense contributions from dark energy and curvature term into one effective parameter in order to reduce parameter degeneracies and to find any deviation from flat concordance $\Lambda$CDM model, by considering that the discrimination between dynamical and non-dynamical sources of cosmic acceleration as the best starting point for analyzing dark energy data sets both at present and in future. We also combine recent Type Ia Supernova (SNIa), Cosmic Microwave Background (CMB) and Baryon Oscillation (BAO) to constrain model parameter space. Degeneracies between model parameters are discussed by using both degeneracy diagram and data analysis including high redshift information from Gamma Ray Bursts (GRBs) sample. The analysis results show that our model is consistent with cosmological observations. We try to distinct the curvature effects from the specially scaling dark energy component as parameterized. We study the linear growth of large scale structure, and finally show the effective dark energy equation of state in our model and how the matter component coincidences with the dark energy numerically. | It is now well-established that the expansion of our universe is currently in an accelerating phase, supported by the most direct and robust evidence from the redshift - apparent magnitude measurements of the "cosmic lighthouse" type Ia supernova \cite{Perlmutter}, and indirect others such as the observations of Cosmic Microwave Background (CMB) by the WMAP satellite \cite{Spergel,jarosik,hinshaw,page,boomerang,cbi,kuo,sa,planck}, and large-scale galaxy surveys by 2dF and SDSS \cite{cole,sdss1,sdss2,tegmark}. Under the assumption that general relativity is valid on cosmological scale, the combined analysis of different observation data sets indicates a spatially-flat universe with about 70\% of the total energy content of the universe today as so called dark energy with effectively negative pressure responsible for the accelerating expansion (see Ref.~\cite{Peebles} for reviews on this topic). Among multitudinous candidates of dark energy models, the "simplest" and theoretically attractive one might be the so called vacuum energy, i.e. $\rho_{\Lambda} = \Lambda/8 \pi G$, where $\Lambda$ is the cosmological constant, which has been long considered as a leading candidate and works quite well on explaining observations through out the history of our universe at different scales. But the origin or mechanisms responsible for the cosmic accelerating expansion are not very clear. On the other hand, some authors suggest that maybe there does not exist such mysterious dark component, but instead the observed cosmic acceleration is a signal of our first real lack of understanding of gravitational physics \cite{Lue} on cosmic scale. An example is the braneworld theory with the extra dimensions compactified or non-compactified \cite{Dvali,rs,hw,ms2}. Consequently, finding the different cosmological implications to distinguish modified gravity models and dark energy scenario from observations is essentially fundamental to physically understanding of our universe\cite{distinguish}. Along with the matter (mainly cold dark matter) component and possible curvature term, the mysterious dark energy dominates the fate of our universe (we do not consider the radiation component contribution as it is supposed to be very tiny for the current universe evolution, at least for the present discussion interests). Ironically so far we do not know much to either of them, even full of puzzling to some extends. So any progress or reasonable understanding to each of them is undoubtedly valuable. Specifically, the quest to distinguish between dark energy and modified gravity scenario and further to differentiate cosmological constant and dynamical dark energy models from observations has become the focus of cosmology study since it holds the key to new fundamental physics. Although we have built up a successful parametrization to describe the properties and evolution of our universe, and in principle distinct dark energy models live at different sub-space of fully descriptive multi-dimension parameter space, due to serious degeneracies among different parameters, we cannot get tight enough constraints from observations by global fitting various observational data sets. One way to extract useful information from observation data and get hints for fundamental physics from cosmology study is to reduce the dimension of parameter space (thus reduce the parameter degeneracies) with particular purpose in mind without apparently biased input to model parametrization. Since we have not found any evidence of inconsistency of standard $\Lambda$CDM model, including more parameters which describe detailed properties of each component if the "cosmic pie" will complicate the situation to constrain model parameters. In order to find deviation of dark energy equation of state parameter $w$ from $-1$ (e.g. evidence of non-cosmological constant dark energy), the assumption of a flat universe is widely accepted in the literature with claims that curvature is negligible from inflation predictions and with emphasis on combined analysis results \emph{with} prior assumption $w=-1$. On the other hand, typically one looks for evidence of dynamical dark energy in the absence of spatial curvature to get better constraints (for an exception, see \cite{curvature1}). It has been concluded in \cite{curvature} that the non-curvature assumption can induce critically large errors in reconstructing the dark energy equation of state even if the true cosmic curvature is on sub-percent level. These claims motivate us proposing a parameterized dark component term to mimic the effective contributions from either dark energy or curvature term plus the dark energy (It is also possible that the parameterized term we postulate may be from a fundamental theory or reasonably modified gravity model we are seeking), besides the conventional matter term. In the first step, it is reasonable to introduce only one parameter which stands for any kind of deviation from standard cosmology model. In some limit case, it should be reduced to the simple four dimensional (4D) $\Lambda$CDM cosmology. The constraint on this parameter from observations should provide insightful hints to further explore fundamental physics. In the next section, we propose a simple cosmic parametrization for the current universe, a parameterized model for the later evolution of our universe. In section 3 we give various cosmic probes to this model, with comparison to the DGP model Universe\cite{Dvali} and the concordance model with a cosmological constant, i.e, the $\Lambda$CDM model, with the hope to locate new features to this new model. Then in section 4, we discuss the new degeneracies between the parameters we introduced and dark matter content. The possible constraints from high redshift observations are also discussed. The last section devotes discussions and conclusions for the general framework studies to this present model. | We have presented a cosmic model parameterizing the late universe which collapses curvature and dark energy effects into one parameter $B$ that may indicate any deviation from standard flat $\Lambda$CDM model and we find that we can not conclude that the cosmic curvature term is constantly zero, instead it may contribute rich phenomenological effects. In order to show the advantages of our parametrization, we study the degeneracy properties between $B$ and $\Omega_m$, emphasizing the contribution from high redshift distance information from GRB or gravitational waves experiments on-going and up-coming. It is well-known that deducing the number of free parameter without significant physics lost is quite important to constrain cosmology models and to find new physics behind. In this paper we also investigated the DGP cosmological model in the simplest flat geometry case with the extra dimension contribution as an effective "cosmological constant", compared with our parameterized model and the reduction to the power-law $\Lambda$CDM model for the 4D real Universe. We find that the DGP model even in the simplest case is still an interesting candidate for the current cosmic speed-up expansion mechanism at long distances, while we know that in the short ranges the model behaves as 4D conventional gravity. We will exploit the non-compact extra dimension to see its possible existence signatures via cosmic effects in the general DGP model later as a promising model, while we do not intent to discuss the quantum aspects of this model as a basic theory\cite{alu}. As a generalization of the $\Lambda$CDM model with naive cosmological constant as dark energy candidate we has parameterized a curvature like term with new phenomenological features via numerical fittings and show the term explicit physics meanings when we perform the parameter B reduction directly to zero or 2. It may be interesting also to study the general properties of the parameterized term as the matter-energy contents in our Universe continuous equation to see what kind of "matter" it may describe effectively, without specifying the form of the parameter. Besides, the phantom case can be realized too, for example, the equation of state parameter $w=p/\rho<-1$ if we take $B>2$ and quintessence corresponds to $B<2$ with $w=p/\rho>-1$ numerically. We think this picture is in conformity with other popular models and enlarges phenomenological dark energy study possibilities to explain the late-time accelerating expansion of our Universe, thus it is worth of further endeavors. | 7 | 10 | 0710.1637 |
0710 | 0710.5128_arXiv.txt | {} {A number of recent works have suggested that the period-luminosity (PL) relation for the Large Magellanic Cloud (LMC) Cepheids exhibits a controversial nonlinear feature with a break period at 10 days. Therefore, the aim of this Research Note is to test the linearity/nonlinearity of the PL relations for the LMC Cepheids in $BVI_cJHK_s$ band, as well as in the Wesenheit functions.} {We show that simply comparing the long and short period slopes, together with their associated standard deviations, leads to a strictly larger error rate than applying rigorous statistical tests such as the $F$-test. We applied various statistical tests to the current published LMC Cepheid data. These statistical tests include the $F$-test, the testimator test, and the Schwarz information criterion (SIC) method.} {The results from these statistical tests strongly suggest that the LMC PL relation is nonlinear in $BVI_cJH$ band but linear in the $K_s$ band and in the Wesenheit functions. Using the properties of period-color relations at maximum light and multi-phase relations, we believe that the nonlinear PL relation is not caused by extinction errors. } {} | Recently, Fouqu\'{e} et al. (\cite{fou07}) have derived the Galactic Cepheid period-luminosity (PL) relation with several different techniques, including parallax measurements (from {\it Hipparcos} and {\it HST}), variants of the Baade-Wesselink method, and distances inferred from open clusters. We point out that such an approach has been applied before in Ngeow \& Kanbur (\cite{nge04}) and Groenewegen et al. (\cite{gre04}). In addition, Fouqu\'{e} et al. (\cite{fou07}) also derive Large Magellanic Cloud (LMC) PL relations in the $BVR_cI_cJHK_s$ band, and refer to the work of Sandage et al. (\cite{san04}), which suggests a possible change of slope for the LMC PL relation at 10 days. In fact, there are several other papers on the topic of nonlinear\footnote{By nonlinearity we mean that the PL relation can be broken into two relations, with a break period adopted at 10 days.} LMC PL relations (see Kanbur \& Ngeow \cite{kan04,kan06}; Kanbur et al. \cite{kan07}; Ngeow et al. \cite{nge05}; Ngeow \& Kanbur \cite{nge06a,nge06b}; Koen et al. \cite{koe07}). These previous works concentrate on the $VI_c$ band (Kanbur \& Ngeow \cite{kan04,kan06}; Ngeow \& Kanbur \cite{nge06b}) or $V$ band only (Ngeow \& Kanbur \cite{nge06a}; Kanbur et al. \cite{kan07}), with data mostly from the OGLE (Optical Gravitational Lensing Survey, Udalski et al. \cite{uda99}) database. For the $JHK_s$ band PL relations, Ngeow et al. (\cite{nge05}) investigated possible nonlinearities using the 2MASS data from Nikolaev et al. (\cite{nik04}) that cross-correlated with the MACHO LMC Cepheids, and a random-phase correction to derive the mean magnitudes of these 2MASS data. Our motivation for this Research Note is to extend the previous work in $BVI_cJHK_s$ band, using the LMC Cepheid data from Fouqu\'{e} et al. (\cite{fou07}), with various rigorous statistical tests. The $JHK_s$ band data used in Fouqu\'{e} et al. (\cite{fou07}) are the 2MASS data matched to the OGLE Cepheids, and the mean magnitudes are derived using the method presented in Soszy\'{n}ski et al. (\cite{sos05}), which is different from the data used in Ngeow et al. (\cite{nge05}). It is important to test the nonlinearity results in $JHK_s$ band results with different Cepheid samples and different methods deriving the $JHK_s$ mean magnitudes. As emphasized in Ngeow \& Kanbur (\cite{nge06a}), statistical tests are needed to test and detect the existence of the nonlinear PL relation. We also point out that in searching for nonlinearity or a change of slope at 10 days, the method of comparing the short and long period slope with their associated standard deviations is more prone to error than applying a statistical test, such as the $F$-test as indicated by the following, purely analytical example. A statement such as the ``the slope is $x\pm \delta x$'' means that the probability that the slope is in the interval $(x-\delta x,x+\delta x)$ is $ 1 - \alpha$, where $\alpha$ is the desired significance level. Then if $A$ is the event that the calculated short period slope is wrong and $B$ is the event that the calculated long period slope is wrong, we have $P(A)=\alpha$ and $P(B) = \alpha$. Then in comparing the short and long period slopes using just their calculated standard deviations, the probability of at least one mistake is $P(A\cup B) = P(A) + P(B) - P(A\cap B) = 2{\alpha} - {\alpha}^2$. If $1 > {\alpha} > 0$, then $2{\alpha} - {\alpha}^2 > {\alpha}$. If the $F$-test or any other statistical test is carried out to the level of significance ${\alpha}$, then this states that the probability of making an error in just comparing long and short period slopes through their standard deviations is greater than the probability that the $F$-test makes a mistake. In essence the $F$-test compares both short and long period slopes (as well as the zero-points) of the nonlinear PL relations simultaneously. \begin{table*} \caption{$F$-test Results of the LMC PL Relations.} \label{tab1} \begin{center} \begin{tabular}{lcccccccccc} \hline\hline Band & $a_S$ & $b_S$ & $\sigma_S$ & $N_S$ & $a_L$ & $b_L$ & $\sigma_L$ & $N_L$ & $F$ & $p(F)$ \\ \hline $B$ & $-2.628\pm0.072$ & $17.493\pm0.046$ & 0.262 & 618 & $-2.402\pm0.192$ & $17.419\pm0.238$ & 0.316 & 96 & 7.10 & 0.001 \\ $V$ & $-2.899\pm0.052$ & $17.148\pm0.033$ & 0.191 & 621 & $-2.763\pm0.141$ & $17.127\pm0.176$ & 0.233 & 95 & 6.83 & 0.001 \\ $I_c$& $-3.073\pm0.035$ & $16.657\pm0.022$ & 0.126 & 604 & $-2.951\pm0.104$ & $16.609\pm0.129$ & 0.162 & 88 & 7.15 & 0.001 \\ $J$ & $-3.237\pm0.040$ & $16.330\pm0.025$ & 0.126 & 481 & $-3.035\pm0.151$ & $16.184\pm0.179$ & 0.134 & 48 & 5.00 & 0.007 \\ $H$ & $-3.347\pm0.036$ & $16.116\pm0.023$ & 0.114 & 481 & $-3.099\pm0.137$ & $15.925\pm0.162$ & 0.122 & 48 & 7.69 & 0.001 \\ $K_s$& $-3.294\pm0.043$ & $16.027\pm0.028$ & 0.137 & 481 & $-3.211\pm0.144$ & $15.992\pm0.171$ & 0.128 & 18 & 1.98 & 0.140 \\ \hline $W_{bi}$ & $-3.463\pm0.021$ & $15.933\pm0.013$ & 0.074 & 598 & $-3.507\pm0.055$ & $15.999\pm0.068$ & 0.086 & 88 & 0.886 & 0.413 \\ $W_{vi}$ & $-3.349\pm0.019$ & $15.897\pm0.012$ & 0.069 & 601 & $-3.316\pm0.050$ & $15.883\pm0.062$ & 0.078 & 87 & 1.631 & 0.196 \\ \hline \end{tabular} \\ \end{center} The subscripts $_S$ and $_L$ are for the short ($\log P<1.0$) and long period Cepheids, respectively, while $a$, $b$ and $\sigma$ are the slope, zero-point and dispersion of the fitted PL relations. \end{table*} | Combining the results from the three statistical tests presented in the previous sections, we find that there is strong statistical evidence to suggest the LMC PL relation is nonlinear in the $BVI_cJH$ band but linear in the $K_sW_{bi}W_{vi}$ band. Including additional data from Persson et al. (\cite{per04}) for the $JHK_s$ band does not alter the results as well. We have to emphasize that both of the testimator and SIC methods are applied to the $BI_cJHK_s$ band and the Wesenheit functions for the first time, in contrast to the $V$ band data that has been studied in Kanbur et al. (\cite{kan07}). The nonlinear LMC PL relation has been found from a Cepheid sample that consists of OGLE Cepheids only (Kanbur \& Ngeow \cite{kan04}). To extend the OGLE sample, mostly at the long period end, Ngeow \& Kanbur (\cite{nge06a}) included various additional data from literature (see table 1 of Ngeow \& Kanbur \cite{nge06a}), and again found strong evidence of nonlinearity of the LMC PL relation. In this Research Note, results using the Fouqu\'{e} et al. (\cite{fou07}) data alone, and with additional data from Persson et al. (\cite{per04}), further supports the conclusion given in Ngeow \& Kanbur (\cite{nge06a}): that the sample selection does not play an important role in detecting the nonlinear LMC PL relation. However, Fouqu\'{e} et al. (\cite{fou07}) have suggested that the mixture of data used in previous work may lead to the nonlinearity seen in the statistical tests. This may certainly be the case and the analysis of a homogeneous sample, such as that provided by the ``LMC shallow survey'' (Fouqu\'{e} 2007; Gieren 2007 -- private communication) is desirable. The nonlinearity of the PL relation that is seen in the optical and $JH$ band but not in the reddening insensitive $K_s$ band and the Wesenheit function may suggest that extinction is the cause of the nonlinearity. However, extinction is not the only explanation and there is some evidence against the hypothesis of extinction errors as a cause for the apparent nonlinearity. The linearity of the $K_s$ band PL relation, as compared to other shorter wavelength PL relations, is expected from black-body arguments (Ngeow \& Kanbur \cite{nge06a}). Simply speaking, the temperature variation dominates the luminosity variation in the optical, and extends to $JH$ band for Cepheid-like temperatures. But in the $K_s$ band the luminosity variation is dominated by the radius variation of Cepheid variables. The proposed mechanism that may cause the nonlinear PL relation, the interaction between the hydrogen ionization front and the stellar photosphere (Kanbur \& Ngeow \cite{kan06}), will only affect the temperature variation and not the radius variation. The linearity of the Wesenheit functions is also not a surprise, and has been studied and discussed in Ngeow \& Kanbur (\cite{nge05w}) and in Koen et al. (\cite{koe07}), and will not be repeated here. Since the additional data used is mainly at the long period end, the possibility remains of systematic errors in reddening as a function of period. However we note that a reddening error as a function of long period LMC Cepheids would also change the observed properties of LMC Cepheids at other phases. A reddening error as a function of period such that LMC Cepheids obey a linear PL relation at mean light would force the LMC Cepheids to have a period-color relation such that they get bluer or hotter at maximum light as the period increases (see Figure \ref{pcmax} for a schematic illustration). This is in stark contrast to the behavior of Galactic Cepheids and long period LMC Cepheids, which are known to have a flat period-color relation at maximum light (Code \cite{cod47}; Simon et al \cite{sim93}; Kanbur \& Ngeow \cite{kan04,kan06}; Kanbur et al \cite{kan04a}). Moreover, it is difficult to explain, theoretically, how a Cepheid could get hotter at maximum light as the period increases. The PL relation at phase $\sim0.8$ described in Ngeow \& Kanbur (\cite{nge06b}) presents clearly the dramatic nature of the nonlinearity at 10 days and the dynamic nature of the PL relation as a function of phase. It is difficult to reconcile this behavior as being due to sampling errors and/or reddening errors. It is worth to point out that the mean light PL relation used in the literature is an average of the PL relations in all phases (Kanbur \& Ngeow \cite{kan04}; Kanbur et al. \cite{kan04a}; Ngeow \& Kanbur \cite{nge06b}). nonlinearity of the PL relations at certain phases will certainly affect the linearity/nonlinearity of the mean light PL relation. \begin{figure} \resizebox{\hsize}{!}{\includegraphics{8812fig5.eps}} \caption{Schematic illustration for the argument with PC(max) relation. Top panels show the observed PC relations (after corrected for extinction) at mean (top-left panel) and maximum (top-right panel) light. Bottom panels show that if additional extinction as function of period to make the mean light PC relation linear, then the same extinction will cause the colors at maximum light get bluer as period increases, which are against observation and theoretical expectation.} \label{pcmax} \end{figure} We have to remind that the data used in this study (and in most of our previous work) were published data that have been corrected for extinction using the ``state-of-the-art'' and well-developed methodology. If extinction error is believed to be the cause of nonlinear PL relations, it would imply that the extinction correction done previously in the literature is incorrect and/or incomplete. This would affect the previous work that using these extinction corrections, and those results need to be revised in future work. | 7 | 10 | 0710.5128 |
0710 | 0710.5544_arXiv.txt | We use current theoretical estimates for the density of long cosmic strings to predict the number of strong gravitational lensing events in astronomical imaging surveys as a function of angular resolution and survey area. We show that angular resolution is the most important factor, and that interesting limits on the dimensionless string tension $\G$ can be obtained by existing and planned surveys. At the resolution of the Hubble Space Telescope ($0\farcs14$), it is sufficient to survey of order a few square degrees -- well within reach of the current HST archive -- to probe the regime $\G\sim10^{-7}$. If lensing by cosmic strings is not detected, such a survey would improve the limit on the string tension by a factor of two over that available from the cosmic microwave background. Future high resolution imaging surveys, covering a few hundred square degrees or more, either from space in the optical or from large-format radio telescopes on the ground, would be able to further lower this limit to $\G\lesssim10^{-8}$. | Superstrings of cosmic size \citep[introduced by ][]{Kib76} are a generic prediction of a number of string theory models \citep[see, e.g., ][and references therein]{Pol04,D+K05}. Given their macroscopic nature, they are in principle detectable through astronomical observations. Therefore, they provide a perhaps unique opportunity for direct empirical tests of the physics of the very early Universe. Considering strings whose only interactions are gravitational, all of their effects are controlled by the global constant dimensionless string tension~$\G$. Current limits on this parameter are given mainly by the properties of the cosmic microwave background (CMB) and by studies of pulsar timing. As far as the former is concerned, cosmic strings produce a smooth component in the CMB power spectrum and a non-gaussian signature in the CMB anisotropy map \citep[e.g.,][]{L+W05,J+S07}. As far as the latter is concerned, strings produce a stochastic gravitational wave background, detectable in the time series of pulsars \citep[e.g.,][]{KTR94,D+V05}. Current limits are $\G \lesssim 3 \times 10^{-7}$ from the CMB~\citep{PWW06}, and perhaps one or two orders of magnitude more stringent from pulsar timing depending on the details of the statistical analysis and the assumed string loop size distribution. An up-to-date review of current observational limits is given by~\citet{Pol07}. The idea of detecting cosmic strings by observing their gravitational lensing effect dates back to~\citet{Vil84}. Briefly, cosmic strings produce a conical space time resulting in a very clear strong lensing signature, i.e.\ they produce identical (neither parity-flipped, nor magnified or sheared) offset replica images of background objects, separated by an angle proportional to the string tension. Thus, strong gravitational lensing provides an opportunity for the {\it direct} detection of cosmic strings, and even a single detected event would provide a measurement of the string tension, independent on the overall demographics of cosmic strings. A number of past studies have identified cosmic string lens candidates in optical surveys, but unfortunately none so far has withstood the test of higher resolution imaging \citep{AHP06,Saz++07}. In this paper we use current theoretical knowledge about the abundance of cosmic strings to predict the number of lensing events as a function of~$\G$ for a realistic set of current and future imaging surveys. Our calculations show that, for sufficiently high angular resolution and sufficiently high (yet currently allowable) string tension, imaging surveys will either be able to detect cosmic string lenses or to at least set interesting limits on the dimensionless string tension parameter. For simplicity, we restrict our analysis to long strings, i.e. strings that are the size of the cosmic volume, and in particular straight with respect to the typical image separation, neglecting the contribution from string loops. \citep[For a discussion of the lens statistics in future radio surveys from the loop population we refer the reader to the recent paper by ][]{MWK07} This assumption simplifies significantly the treatment and, since they do not depend on the detailed topology of the string network, nor on the timescales for gravitational decay, makes our predictions quite robust. | \begin{enumerate} \item Present-day high-resolution imaging surveys are capable of probing the putative cosmic string tension parameter to a factor of two lower than the current CMB limit, making lensing both competitive with, and complementary to, the pulsar timing methods. As has been noted before, in the event of a detection, gravitational lensing would provide a direct measurement of the tension, and perhaps the velocity, of this string. \item The main practical considerations in detecting string lensing events are two-fold: firstly, the expected faint pairs of images must first be carefully deblended and then understood in the context of neighbouring events; secondly, the survey geometry should be such that the fields are large enough to contain the characteristic multiple neighbouring events, but sparsely distributed to ensure that the global, not local, lensing rate is being probed. \item The upper bound on the tension, from the failure to detect a single string lensing event in a given survey, is principally determined by the available angular resolution. Relatively little is gained from studying an area of sky greater than some critical value: for typical ground-based optical image resolution, this critical survey size is a few tens of square degrees. Likewise, string tensions of $\G\sim10^{-8}$ should be able to be investigated in future surveys with image resolutions of 10-100~milliarcsec covering a few hundred square degrees. \end{enumerate} | 7 | 10 | 0710.5544 |
0710 | 0710.2963_arXiv.txt | {The identification of increasingly smaller signal from objects observed with a non-perfect instrument in a noisy environment poses a challenge for a statistically clean data analysis.} {We want to compute the probability of frequencies determined in various data sets to be related or not, which cannot be answered with a simple comparison of amplitudes. Our method provides a statistical estimator for a given signal with different strengths in a set of observations to be of instrumental origin or to be intrinsic.} {Based on the spectral significance as an unbiased statistical quantity in frequency analysis, Discrete Fourier Transforms (DFTs) of target and background light curves are comparatively examined. The individual False-Alarm Probabilities are used to deduce \em conditional \rm probabilities for a peak in a target spectrum to be real in spite of a corresponding peak in the spectrum of a background or of comparison stars. Alternatively, we can compute joint probabilities of frequencies to occur in the DFT spectra of several data sets simultaneously but with different amplitude, which leads to \em composed \rm spectral significances. These are useful to investigate a star observed in different filters or during several observing runs. The composed spectral significance is a measure for the probability that none of coinciding peaks in the DFT spectra under consideration are due to noise.} {\sc Cinderella \rm is a mathematical approach to a general statistical problem. Its potential reaches beyond photometry from ground or space: to all cases where a quantitative statistical comparison of periodicities in different data sets is desired. Examples for the composed and the conditional \sc Cinderella \rm mode for different observation setups are presented.} {} | \label{s1} % The micromag precision, achieved by the MOST\footnote{MOST is a Canadian Space Agency mission, jointly operated by Dynacon Inc., the University of Toronto Institute of Aerospace Studies, the University of British Columbia, and with the assistance of the University of Vienna, Austria.} (Microvariability \& Oscillations of STars) mission (Walker et al.~2003; Matthews 2004), does not only provide exciting new results in asteroseismology, but reveals instrumental problems which challenge our data reduction techniques (see Sect.\,\ref{s1.1}). Cosmic ray impacts on the detector, stray light, positioning errors of the satellite, and thermal stability problems introduce periodic and, in the worst case, pseudo-periodic effects into photometric measurements. All this calls for new techniques in data reduction and analysis (see Sect.\,\ref{s1.2}). Space observations in general can provide an unprecedented amount of measurements, requiring an enhanced degree of automatic data analysis without sacrificing accuracy and reliability. In this context, {\sc SigSpec} (Reegen~2007) was developed to combine the Discrete Fourier Transform (DFT) -- a standard method to determine stellar pulsation frequencies -- with a clean statistical quantity: the spectral significance of a peak in an amplitude or power spectrum by comparison to white noise. The basic idea of {\sc Cinderella} is to use target and comparison data sets simultaneously for a cross-identification of artifacts in the frequency domain. It is the first technique permitting a statistically unbiased and quantitative comparison of different (not necessarily photometric) time series in the {\em frequency} domain. Being applicable to practically all measurements of physical quantities over time, {\sc Cinderella} has the potential to become a valuable tool beyond the scope of micromag space photometry. \subsection{The MOST mission}\label{s1.1} The first space telescope designed and built for photometric stellar seismology was EVRIS (Vuillemin et al.~1998), a 10-cm photoelectric telescope aboard the MARS-96 probe, but it unfortunately did not achieve the transfer orbit. An instrument providing photometric information on a large scale useful for asteroseismology was NASA's WIRE satellite, whose primary scientific goal of infrared mapping failed, but a 5-cm star tracker telescope with a CCD detector turned out to permit stellar photometry of remarkable quality (e.\,g., Buzasi et al.~2000). The MOST satellite launched in June, 2003, assumed the role as a precursor to the CNES-led mission COROT (Baglin et al.~2004), which was successfully launched on December 27, 2006, and which is producing extremely useful space photometric data of hitherto unprecedented accuracy and volume. MOST, WIRE and COROT are low-Earth-orbit (LEO) missions with comparable environmental effects (e.g., cosmic radiation, stray light scattered from the Earth's surface). A further commonality of all three missions is the requirement to extract asteroseismic information from a series of up to hundreds of thousands of CCD frames (or sub-rasters, respectively), each of which may consist of a few hundred to several million pixels. Hence, the present work may apply to other LEO space photometry missions and to ground-based multi-object photometry. The MOST telescope is a 15-cm Maksutov optical telescope, supplied with a single broadband filter and initially with two identical CCD detectors: one used for science data acquisition, the other for the {\em Attitude Control System} (ACS). Thanks to the low mass of $54$\,kg and the ACS developed by Dynacon, Inc. (Groccott, Zee \& Matthews~2003; Carroll, Rucinski \& Zee~2004), a pointing stability to approximately $\pm 1\arcsec$ rms is achieved. In {\em Fabry Imaging} mode the telescope entrance pupil is imaged onto the CCD via a Fabry microlens as is shown by Figs.\,7 and 8 of Walker et al.~(2003). Each Fabry Image is an annulus with an outer diameter of 44 pixels. The pixels in a square subraster outside the annulus are used to estimate the background. MOST also obtains {\em Direct Imaging} photometry of typically $1-6$ stars, based on defocussed images (FWHM $\sim$ 2.2 pixels; Rowe et al.~2006; Huber \& Reegen~2008), and {\em Guide Star} photometry of about $20-30$ stars (Aerts et al.~2006; Saio et al.~2006). \subsection{Data reduction}\label{s1.2} The data reduction described by Reegen et al.~(2006) applies linear correlations between pairs of target and background pixels for stray light correction. This so-called {\em decorrelation technique} is also applicable to simultaneous photometry of several stars, in this case correlating variable vs.~constant stars. Fig.\,\ref{fig1} illustrates the performance of the Fabry imaging photometry with MOST data of $\beta$ CMi (Saio et al.~2007). The blue graph refers to the raw data and the red graph to the reduced light curve. The overall noise level decreased by an order of 10, and so did the harmonics of the orbital frequency of the spacecraft, ($\approx 14.2$\,d$^{-1}$ for $101.4$\,min; Walker et al.~2003). However, instrumental peaks (dotted green lines) persisted on a lower level and their amplitudes still exceeded the stellar signal (main frequencies: $3.257$\,d$^{-1}$ \& $3.282$\,d$^{-1}$; dotted black line). \begin{figure}\includegraphics[width=256pt]{f1.eps} \caption{The raw light curve ({\em blue}) of the MOST Fabry target $\beta$\,CMi and after data reduction ({\em red}). Harmonics of the satellite's orbital frequency ($\approx 14.2$\,d$^{-1}$; {\em dotted green}), the detected stellar signal ($3.257$\,d$^{-1}$ \& $3.282$\,d$^{-1}$; {\em dotted black}) are indicated.}\label{fig1} \end{figure} \subsection{\sc SigSpec}\label{s1.3} {\sc SigSpec} (Reegen~2007), is based on DFT amplitude spectra and consecutive prewhitening of dominant peaks. But instead of considering the peak with the highest amplitude to be significant and estimating the reliability roughly in terms of signal-to-noise ratio, the {\em Probability Density Function} (PDF) is employed. The PDF depends on the frequency and phase of the examined peak using white noise as a reference. The mean photometric magnitude in a time series is usually reduced to zero before evaluating the DFT. {\sc SigSpec} the resulting statistical consequences into account, and is furthermore not restricted to Gaussian distributed residuals. The {\em False-Alarm Probability} is a frequently used statistical quantity in time series analysis. It is the probability of a peak at a given amplitude level to be generated by noise. Formally it is obtained through integration of the PDF. To avoid problems in computing extremely low numerical values, {\sc SigSpec} returns a quantity called {\em spectral significance} (hereafter abbreviated by ``sig''), which is the negative logarithm of the False-Alarm Probability. It gives the number of uncorrelated data sets needed, containing pure noise, so that a peak in the Fourier domain appears which is comparable in amplitude and phase to the peak under consideration in the observed data. Although {\sc SigSpec} prevailed as a powerful tool for analyzing MOST photometry, it occasionally suffered from the weakness of having to refer to uncorrelated (i.e. white) noise. \subsection{The virtue of {\sc Cinderella}}\label{s1.4} Frequencies with individual amplitudes and phases (``peaks'') in the DFT spectra of a target and comparison data sets are examined by {\sc Cinderella} for compatibility. In other words, {\sc Cinderella} allows us to investigate whether these data sets are related by any physical (deterministic) process. The procedure is the same if the comparison data represent sky background or a star with a different frequency spectrum as the target star, which -- in the best case -- is a constant star. Subsequently, the terms ``target star'' and ``comparison star'' will be used, keeping in mind that everything discussed here readily applies to sky readings instead of comparison stars as well. Obviously, all compared data sets have to be observed under similar circumstances. An extension of the method to handle more than one comparison data set is useful for multi-object environments, such as photometry in a field. In conditional mode, {\sc Cinderella} establishes a quantitative comparison of significant frequencies occurring at the same time in at least two different data sets. It returns a statistically robust value, called {\em conditional sig}, for the probability that a peak in the spectrum of one data set is not (deterministically) related to a peak in the other data set(s) within a given frequency resolution. The alternative composed mode is dedicated to testing whether peaks in different DFT spectra with similar frequencies are ``real'', in the sense of not due to noise. The corresponding quantity, the {\em composed sig} is a measure for the probability that none of the examined peaks is due to noise. \subsection{Frequency resolution}\label{s1.6} The question how to set the frequency difference acceptable for the consideration of peaks as coincidental is crucial to the examination of corresponding peaks in different DFT spectra. In this context, an alternative definition to the Rayleigh resolution, \begin{equation}\label{eqRayleigh} \delta f_R := \frac{1}{\Delta t}\: , \end{equation} with $\Delta t$ denoting the total time interval width of the time series is introduced by Kallinger, Reegen \& Weiss~(2007). They suggest to additionally employ the sig for a peak amplitude according to \begin{equation}\label{eqKallinger} \delta f_K := \frac{1}{\Delta t\sqrt{\mathrm{sig}\left( A\right)}}\: , \end{equation} for obtaining a more realistic criterion for matching peaks ({\em frequency resolution}) than provided by Eq.\,(\ref{eqRayleigh}). Their numerical simulations show an excellent compatibility of this quantity, subsequently termed {\em Kallinger resolution}, to the {\em frequency error} derived by Montgomery \& O'Donoghue~(1999). For practical applications, it is useful to enhance the flexibility of {\sc Cinderella} by introducing an exponent $z$ and to re-define the frequency resolution according to \begin{equation}\label{eqFRes} \delta f := \frac{1}{\Delta t\left[\sqrt{\mathrm{sig}\left( A\right)}\right] ^z}\: , \end{equation} where $z$ usually attains values in the range $\left[ 0,1\right]$. The Rayleigh resolution is obtained for $z=0$, whereas $z=1$ yields the Kallinger resolution. | \label{s5} % This paper introduces a technique to interpret periodicities in an ensemble of data of common origin. {\sc Cinderella} relies on {\sc SigSpec} (Reegen~2007), thus benefitting from a correct employment of the complex phase information in Fourier Space on the one hand and a clean statistical description of interrelation of datasets on the other. The conditional {\sc Cinderella} mode is based on a quantitative comparison between one target and one or more comparison datasets and returns a measure of the probability (conditional sig) for periodicities identified in the target data to be deterministically related (to be `unique') to the target. The composed {\sc Cinderella} analysis returns a measure of the joint probability (composed sig) that a given periodicity observed in individual datasets -- but with different signal strengths -- is not due to noise. Such datasets could contain, e.g., measurements of the same target in different observing runs or with different instruments (e.g., different filters or simultaneous spectroscopy and photometry). Our experience (as outlined in our examples in Sect.\,\ref{s3}) confirms that {\sc Cinderella} reliably identifies residual instrumental signal in the MOST data even after a fairly sophisticated data reduction in the time-domain and also provides quantitative arguments to distinguish intrinsic from instrumental signal. {\sc Cinderella} is a statistically correct technique replacing what experienced observers achieve based on their ``good feelings'' when evaluating, for example, differential photometry, but, of course, the method is not limited to photometry. It quantitatively determines conditional and composed probabilities for matching peaks in DFT spectra of any kind of datasets containing periodicities. \subsection*{Acknowledgements} PR received financial support from the Fonds zur F\"orderung der wis\-sen\-schaft\-li\-chen Forschung (FWF, projects P14546-PHY, P14984-PHY) Furthermore, it is a pleasure to thank D.\,B.~Guenther (St.~Mary's Univ., Halifax), M.~Hareter, D.~Huber, T.~Kallinger (Univ.~of Vienna), R.~Kusch\-nig, J.\,M. Matthews (UBC, Vancouver), A.\,F.\,J.~Moffat (Univ.~de Montreal), D.~Punz (Univ.~of Vienna), S.\,M. Rucinski (D.~Dunlap Obs., Toronto), D.~Sasselov (Harvard-Smithsonian Center, Cambridge, MA), G.\,A.\,H.~Walker (UBC, Vancouver), and K.~Zwintz (Univ.~of Vienna) for valuable discussion and support with extensive software tests. Finally, we address our very special thanks to S.\,M.~Mochnacki (University of Toronto) for his careful revision and valuable comments that substantially improved the presentation of this work. | 7 | 10 | 0710.2963 |
0710 | 0710.4068_arXiv.txt | To optimise the science results of the asteroseismic part of the CoRoT satellite mission a complementary simultaneous ground-based observational campaign is organised for selected CoRoT targets. The observations include both high-resolution spectroscopic and multi-colour photometric data. We present the preliminary results of the analysis of the ground-based observations of three targets. A line-profile analysis of 216 high-resolution FEROS spectra of the $\delta$ Sct star HD~50844 reveals more than ten pulsation frequencies in the frequency range 5--18 d$^{-1}$, including possibly one radial fundamental mode (6.92 d$^{-1}$). Based on more than 600 multi-colour photometric datapoints of the $\beta$ Cep star HD~180642, spanning about three years and obtained with different telescopes and different instruments, we confirm the presence of a dominant radial mode $\nu_1=5.48695\,$d$^{-1}$, and detect also its first two harmonics. We find evidence for a second mode $\nu_2=0.3017$\,d$^{-1}$, possibly a g-mode, and indications for two more frequencies in the 7--8 d$^{-1}$ domain. From Str\"omgren photometry we find evidence for the hybrid $\delta$ Sct/$\gamma$ Dor character of the F0 star HD~44195, as frequencies near 3 d$^{-1}$ and 21 d$^{-1}$ are detected simultaneously in the different filters. | The asteroseismic window of the CoRoT satellite mission aims at the monitoring of several types of pulsators along the Main Sequence. To optimise the science results, its targets are carefully chosen and selected. Preparatory observations from the ground have been a key stone in the selection process \cite{ref1}. With the CoRoT satellite successfully launched (December 2006), simultaneous ground-based observations are very important and are complementary to the space data. Multi-colour photometry provides colour information, which allows identification of the degree $\ell$ by means of amplitude ratios and phase shifts, while high-resolution spectroscopy allows the detection of high-degree modes and the identification of both the degree $\ell$ and the azimuthal order $m$. In this framework a simultaneous ground-based observing campaign was organised by the CoRoT/SWG Ground-based Observations Working Group. Large Programme proposals (i.e. guaranteed observing time during four observing seasons) were submitted and accepted at ESO La Silla (FEROS/2.2m; 10+5 nights per semester), Observatoire de Haute Provence (SOPHIE/1.92m; 10+5 nights per semester) and Calar Alto Astronomical Observatory (FOCES/2.2m; 10+10 nights per semester), with the aim to obtain multi-site time-series of high-resolution spectra of a selection of $\delta$ Sct, $\gamma$ Dor, $\beta$ Cep and Be CoRoT primary and secondary targets. The first two observing seasons (winter 2006-2007 and summer 2007), (nearly) coinciding with CoRoT's IR01 (first Initial Run) and LRc1 (first Long Run in the center direction), have been succesfully completed. \Tref{logbook} gives an overview of the targets and the amount of spectra obtained. In addition to the Large Programmes, and in continuation of a project started three years ago, observing time has been awarded at smaller telescopes with multi-colour photometric instruments (Str\"omgren photometry: 90cm@Sierra Nevada Observatory (SNO), 1.5m@San Pedro M\'{a}rtir Observatory (SPMO); Geneva photometry: 1.2m Mercator@Observatorio Roque de los Muchachos (ORM); Johnson photometry: Konkoly Observatory (KO)). In particular, 18 and 14 consecutive nights have been awarded with the $uvby\beta$ photometers at SPMO and SNO, respectively, in Nov-Dec 2006 and 2007. \begin{center} \begin{table}[h] \caption{\label{logbook} Logbook of the spectroscopic observations, obtained in Jan-Feb 2007 and May-Jul 2007 with the FEROS, SOPHIE and FOCES instruments, dedicated to a selection of targets of the CoRoT IR01 (Feb-Apr 2007), LRc1 (summer 2007) and LRa1 (winter 2007-2008) runs. The different columns give the target name, its V magnitude, spectral type, variable type, name of the CoRoT run, amount of spectra obtained with FEROS, SOPHIE and FOCES, respectively. The targets indicated in boldface are discussed in this paper. \\} \centering \begin{tabular}{llllllll} \br target & V & Sp.T. & type & CoRoT run & FEROS & SOPHIE & FOCES \\ \mr {\bf HD~50844} & 9.09 & A2 & $\delta$ Sct & IR01 & 216 & & \\ HD~50747 & 5.45 & A4 & SB2 & IR01 & 17 & 14 & 6\\ HD~51106 & 7.35 & A3m & SB2 & IR01 & 15 & 14 & 4\\ HD~50846 & 8.43 & B5 & EB & IR01 & 16 & 12 & 4\\ \hline HD~49434 & 5.74 & F1V & $\gamma$ Dor & LRa1 & 71 & 444 & 75 \\ HD~50209 & 8.36 & B9Ve & Be star & LRa1 & 68 & & \\ \hline HD~181555 & 7.52 & A5 & $\delta$ Sct & LRc1 & 343 & 66 & 285 \\ HD~174966 & 7.72 & A3 & $\delta$ Sct & LRc1 & & & 119 \\ {\bf HD~180642} & 8.27 & B1.5III & $\beta$ Cep & LRc1 & 213 & 35 & \\ HD~181231 & 9.69 & B9.0V & Be star & LRc1 & 72 & & \\ \br \end{tabular} \small{SB2=double-lined spectroscopic binary; EB=eclipsing binary} \end{table} \end{center} In the next sections we report on the preliminary results of the analysis of the ground-based datasets of the $\delta$ Sct star HD~50844 and the $\beta$ Cep star HD~180642, and give an outlook towards an interesting candidate CoRoT target, the hybrid $\delta$ Sct/$\gamma$ Dor star HD~44195. | We obviously are in a challenging era of asteroseismology. The preparatory work and the simultaneous ground-based observations of the CoRoT mission require a huge effort of the asteroseismic community. Challenging objects are chosen (e.g. faint targets, fast rotating stars) and in more than one way are we pushing the limits. The ambitious choice of targets asks for new methodologies and analysis techniques, and requires a close collaboration between theoreticians and observers. With joined forces, with shared expertise and with complementary data from space and from the ground we have excellent prospects to probe the interiors of stars. \ack KU acknowledges financial support from a \emph{European Community Marie Curie Intra-European Fellowship}, contract number MEIF-CT-2006-024476. This work is supported by the italian ESS project, contract ASI/INAF I/015/07/0, WP 03170, and by the Research Council of Leuven University under grant GOA/2003/04. | 7 | 10 | 0710.4068 |
0710 | 0710.4859_arXiv.txt | {Many binary stellar systems in which the primary star is beyond the asymptotic giant branch (AGB) evolutionary phase show significant orbital eccentricities whereas current binary interaction models predict their orbits to be circularised.} {In the search for a mechanism to counteract the circularising effect of tidal interaction we analyse how the orbital parameters in a system are modified under mass loss and mass exchange among its binary components.} {We propose a model for enhanced mass-loss from the AGB star due to tidal interaction with its companion, which allows a smooth transition between the wind and Roche-lobe overflow mass-loss regimes. We explicitly follow its effect along the orbit on the change of eccentricity and orbital semi-major axis, as well as the effect of accretion by the companion. We calculate timescales for the variation of these orbital parameters and compare them to the tidal circularisation timescale.} {We find that in many cases, due to the enhanced mass loss of the AGB component at orbital phases closer to the periastron, the net eccentricity growth rate in one orbit is comparable to the rate of tidal circularisation. We show that with this eccentricity enhancing mechanism it is possible to reproduce the orbital period and eccentricity of Sirius system, which under the standard assumptions of binary interaction is expected to be circularised. We also show that this mechanism may provide an explanation for the eccentricities of most barium star systems, which are expected to be circularised due to tidal dissipation. } {By proposing a tidally enhanced model of mass loss from AGB stars we find a mechanism which efficiently works against the tidal circularisation of the orbit. This mechanism can explain the significant eccentricities observed in binary systems containing a white dwarf and a less evolved companion, which are predicted to be circularised due to their proximity, such as Sirius and systems with barium stars.} | Detached binary systems containing a white dwarf and a relatively unevolved companion, i.e., a main-sequence star or a (sub)giant, are a useful tool to understand the binary evolution of systems with an asymptotic giant branch (AGB) star, given that their orbital and chemical properties do not change significantly from the moment that the primary finished its AGB evolution and became the current white dwarf. An example of such a system is Sirius, a 2.1 M$_{\odot}$ main sequence star with a 1.05 M$_{\odot}$ white dwarf companion in a 50-year orbit \citep[e.g., ][]{1960JO.....43..145V,1978ApJ...225..191G}. Under the standard picture of binary evolution with tidal interaction this system should have circularised when the primary became an AGB star, as we show in $\S$\ref{sec:sirius} of this paper. However, this system has an eccentricity $e=0.59$ and is not an exception. A similar problem is faced when considering the barium-star systems, which are red giants with over-abundances of $s$-process elements (prominently barium) with white dwarf companions. The best current explanation for the enhancement of $s$-process elements in these stars is that they accreted mass from their companion when it was an AGB star. Due to the large size of AGB stars these systems are expected to be circularised by tidal interaction for periods smaller than about 3500-4000 days \citep{2003ASPC..303..290P}. However, barium-star systems with period as short as 600 days are observed to be significantly eccentric \citep{1998A&A...332..877J}. The problems stated above are indications that a mechanism must exist that counteracts the circularising effect of the tides by the time the primary star is on the AGB. \citet{1995A&A...293L..25V} propose that enhanced mass loss at periastron could enhance the eccentricity and \citet{2000A&A...357..557S} shows that this works when the AGB star fills its Roche lobe during periastron passages. It has also been proposed that the tidal interaction between a binary system and a circumbinary disk could account for an eccentricity enhancing mechanism \citep{1996A&A...314L..17W,1998Natur.391..868W}. However, \citet{2000A&A...357..557S} argues that the masses of observed circumbinary disks are too small to counteract the circularisation. Mass loss from stars in binary systems is commonly treated as single-star wind mass-loss as long as the stars are inside their Roche lobes, while when one star fills its Roche lobe the mass-loss rate increases abruptly to the high values that correspond to Roche-lobe overflow. This approximation is fairly accurate for stars with a steep density gradient in their atmospheres, but it is not appropriate for the case of AGB stars, given their large atmospheric pressure scale height and their weakly bound envelope. The fact that AGB stars undergo dynamical pulsations further reduces density gradient in the layers above the photosphere \citep[e.g.,][]{1988ApJ...329..299B}. We propose a prescription for enhanced mass loss which smoothly grows from the single star wind mass loss rate to the Roche lobe overflow mass loss rate as the radius of the star approaches its Roche lobe radius. This gives a variable mass loss rate along eccentric orbits which is higher at orbital phases closer to the periastron, even if the star does not fill its Roche lobe. This effect works in the same way as discussed above, but no filling of the Roche lobe is needed, so that (during the AGB phase) this effect is permanently competing against tidal circularisation. \citet{1988A&A...205..155B} carried out calculations of the variation of orbital parameters considering instantaneous mass transfer and only linear momentum conservation. Angular momentum conservation was not taken into account in this exploratory study. An extension including angular momentum conservation was made by \citet{2000A&A...363..660L}, but they make use of the orbital average of the distance between the components and orbital angular velocity to carry out their calculations, which in the case of a variable mass-loss along the orbit does not give the correct results. In $\S$\ref{sec1} we revise the variation of orbital parameters due to stellar mass loss and mass transfer by taking into account the conservation laws of linear and angular momentum with respect to the centre of mass of the actual binary system. We do not make use of the orbital averages a priori, but give the variations as a function of the orbital phase, allowing different rates of mass loss and mass transfer along the orbit. In $\S$\ref{sec2} we calculate the rate of change of the eccentricity according to different assumptions of mass loss and accretion. In $\S$\ref{sec3} we evaluate the competition between the tidal circularisation and the eccentricity pumping due to our proposed mass loss, and show that the latter can prove effective in counteracting the former for systems such as Sirius and barium stars in which their AGB component did not fill its Roche lobe. A summary is presented in $\S$\ref{sec4}. | \label{sec4} We have revised the description of the evolution of orbital parameters due to mass loss and mass transfer in a binary system where the primary is an AGB star. We also propose a tidally enhanced mass-loss rate for AGB stars in binary systems which allows a smooth transition between wind mass loss and mass loss due to Roche lobe overflow. With our revised prescription for the evolution of the orbital eccentricity and the fact that our proposed mass loss is not constant along an eccentric orbit, we find an eccentricity enhancing mechanism which % counteracts the circularising effect of tidal dissipation. We have shown that the standard scenario of binary interaction cannot explain the observed orbital parameters of the eccentric binary system Sirius unless the strength of tidal dissipation in AGB stars is at least 3 orders of magnitude smaller than in normal red giants. On the other hand, our models with the eccentricity-enhancing mechanism can reproduce the orbital properties of the Sirius system with a reasonable choice of the tidal dissipation parameter ($f'=0.25$). Our eccentricity enhancing mechanism also allows binary systems containing barium stars to remain eccentric with periods as short as about 1000 days under reasonable parameter assumptions, while under the same assumptions in the standard scenario of tidal circularisation only systems with periods longer than about 2500 days are expected to show a significant eccentricity. We also show that we can explain the eccentricities of almost all barium star systems which do not fill their Roche lobes if the uncertain convective tidal dissipation strength in AGB stars is reduced by a factor of 10, compared to what is measured for red giant stars. Moreover, this assumption allows Roche-lobe overflow to occur in significantly eccentric systems, which may be the key to explain shorter-period barium-star systems which are still eccentric. Whether our eccentricity-enhancing mechanism is a viable solution to the problem of explaining short-period eccentric binaries with (former) AGB stars, in particular barium stars, needs to be tested with a binary population synthesis model. Results of such a test will be given in a separate paper \citep{BMP07}. Preliminary barium star population synthesis calculations indicate that when the full evolutionary history of a binary population is taken into account, a reduction of the tidal dissipation strength of only $f'=0.25$ is needed to reproduce the observed systems. This is consistent with our findings about the values of $f'$ needed to reproduce the Sirius system. | 7 | 10 | 0710.4859 |
0710 | 0710.1198_arXiv.txt | We investigate the influence of an interaction between dark energy and dark matter upon the dynamics of galaxy clusters. We obtain the general Layser-Irvine equation in the presence of interactions, and find how, in that case, the virial theorem stands corrected. Using optical, X-ray and weak lensing data from 33 relaxed galaxy clusters, we put constraints on the strength of the coupling between the dark sectors. Available data suggests that this coupling is small but positive, indicating that dark energy might be decaying into dark matter. Systematic effects between the several mass estimates, however, should be better known, before definitive conclusions on the magnitude and significance of this coupling could be established. | Cosmological accelerated expansion is by now a well-established observational fact \cite{cosmoaccel,wmap,bao}, leading either to an asymptotically de Sitter cosmology, plagued with an astonishingly small cosmological constant, or else to a universe filled up to 80\% with a strange dynamical component with negative pressure -- dark energy \cite{darkenergy}. If dark energy contributes a significant fraction of the content of the Universe, it is natural, in the framework of field theory, to consider its interactions with the remaining fields of the Standard Model and well-motivated extensions thereof. For lack of evidence to the contrary, interactions of dark energy or dark matter with baryonic matter and radiation must be either inexistent or negligible. Nevertheless, some level of interaction between dark energy and the dark matter sector, which is present in most extensions of the Standard Model, is still allowed by observations. The possibility that dark energy and dark matter can interact has been studied in \cite{amendola}-\cite{[14]}, among others. It has been shown that the coupling between a dark energy (or �quintessence�) field and the dark matter can provide a mechanism to alleviate the coincidence problem \cite{amendola,[11]}. A suitable choice of the coupling, motivated by holographic arguments, can also lead to the crossing of the �phantom barrier� which separates models with equations of state $w = p/\rho > -1$ from models with $w < -1$ \cite{[12]} -- see also \cite{[14],[15]}. In addition, it has been argued that an appropriate interaction between dark energy and dark matter can influence the perturbation dynamics and affect the lowest multipoles of the CMB spectrum, accounting for the observed suppression of the quadrupole \cite{[13],[16]}. The strength of the coupling could be as large as the fine structure constant \cite{[13],[17]}. Recently, it was shown that such an interaction could be inferred from the expansion history of the universe, as manifested in, e.g., the supernova data together with CMB and large-scale structure\cite{[18]}. Nevertheless, the observational limits on the strength of such an interaction remain weak \cite{[19]}. A complementary and fundamentally different way in which the coupling between dark energy and dark matter can be checked against the observations is through its impact on large-scale structure. If dark energy is not a cosmological constant, it must fluctuate in space and in time -- and, in particular, if dark energy couples to dark matter, then it must surely be dynamical. If that is the case, dark energy affects not only the expansion rate, but the process of structure formation as well, through density fluctuations, both in the linear \cite{[11]}, \cite{[20]}-\cite{[23]} and the non-linear \cite{[24],[25]} regimes. The growth of dark matter perturbations can in fact be enhanced due to the coupling between these two components \cite{[13],[14],[26]}. Recently, it was suggested that the dynamical equilibrium of collapsed structures would be affected by the coupling of dark energy to dark matter, in a way that could be observed in the galaxy cluster Abell A586 \cite{[27]}. The basic idea is that the virial theorem is distorted by the non-conservation of mass caused by the coupling. In this paper we show precisely how the Layser-Irvine equation, which describes the flow to virialization \cite{[28]}, is changed by the presence of the coupling, in such a way that the final state of equilibrium violates the usual virial condition, $2K + U= 0$, where $K$ and $U$ are respectively the kinetic and the potential energies of the matter constituents in an isolated system. We show that this violation leads to a systematic bias in the estimation of masses of clusters if the usual virial conditions are employed. Although it is still possible that systematic errors from observations smear the results, the fact that some shift in the mean value of the coupling for two independent sets of observations (compared to the third set) may signalize some new physics. Even though the uncertainties associated with any individual galaxy cluster are very large, by comparing the naive virial masses of a large sample of clusters with their masses estimated by X-ray and by weak lensing data, we may be able to constrain such a bias and to impose tighter limits on the strength of the coupling than has been achieved before. | We have estimated the effective coupling between dark energy and dark matter through the internal dynamics of galaxy clusters. In the presence of coupling, the flow of mass and energy between the components changes the virial condition in a way that can be tested by comparing different estimators for the mass of clusters. We searched for this signature in 33 galaxy clusters for which reliable X-ray, weak lensing and optical data were available. Our results indicate a weak preference for a small but positive effective coupling constant $\bar\zeta$ -- in line with predictions made by some of us \cite{[12],[13]}. Since the statistical significance is still low ($\Delta \chi^2/{\rm d.o.f.} \sim 0.2$), it is paramount that more clusters (with homogeneous mass determinations and good control of systematics) be tested. If a significant indication of such coupling is still found, this would open a tantalizing new window on the nature of the dark sector. | 7 | 10 | 0710.1198 |
0710 | 0710.3648_arXiv.txt | {} {We present a qualitative analysis of key (but yet unappreciated) linear phenomena in stratified hydrodynamic Keplerian flows: \emph{(i)} the occurrence of a vortex mode, as a consequence of strato-rotational balance, with its transient dynamics; \emph{(ii)} the generation of spiral-density waves (also called inertia-gravity or $g\Omega$ waves) by the vortex mode through linear mode coupling in shear flows.} {Non-modal analysis of linearized Boussinesq equations were written in the shearing sheet approximation of accretion disk flows.} {It is shown that the combined action of rotation and stratification introduces a new degree of freedom, vortex mode perturbation, which is in turn linearly coupled with the spiral-density waves. These two modes are jointly able to extract energy from the background flow, and they govern the disk dynamics in the small-scale range. The transient behavior of these modes is determined by the non-normality of the Keplerian shear flow. Tightly leading vortex mode perturbations undergo substantial transient growth, then, becoming trailing, inevitably generate trailing spiral-density waves by linear mode coupling. This course of events -- transient growth plus coupling -- is particularly pronounced for perturbation harmonics with comparable azimuthal and vertical scales, and it renders the energy dynamics similar to the 3D unbounded plane Couette flow case.} {Our investigation strongly suggests that the so-called bypass concept of turbulence, which has been recently developed by the hydrodynamic community for spectrally stable shear flows, can also be applied to Keplerian disks. This conjecture may be confirmed by appropriate numerical simulations that take the vertical stratification and consequent mode coupling into account in the high Reynolds number regime.} | According to classical fluid dynamics, unmagnetized disk flows in Keplerian rotation (more generally: with angular momentum increasing outward and with no extremum of vorticity) are spectrally stable; however, there is irrefutable observational evidence that such disks have to be turbulent. Due to this apparent contradiction, disk turbulence is often considered as some sort of mystery. However, an analogous dilemma that existed in laboratory/engineering flows has been solved by the hydrodynamic community in the 90s of the last century, where a breakthrough was accomplished in the comprehension of turbulence in spectrally/asymptotically stable shear flows (e.g. in the plane Couette flow). Let us briefly recall the essence of that breakthrough (see Chagelishvili et al. 2003 for details). Traditional stability theory followed the approach of Rayleigh (1880) where the instability is determined by the presence of exponentially growing modes that are solutions of the linearized dynamic equations. Only recently has one become aware that operators involved in the the modal analysis of plane shear flows are not normal, hence that the corresponding eigenfunctions are non-orthogonal and would strongly interfere (Reddy et al. 1993). For this reason, the emphasis was shifted in the 90s from the analysis of long time asymptotic flow stability to the study of short time behavior. It was established that \emph{asymptotically/Rayleigh stable flows allow for linear transient growth of vortex and/or wave mode perturbations} (cf. Gustavsson 1991; Butler and Farrell 1992; Reddy and Henningson 1993; Trefethen et al. 1993). This fact incited a number of fluid dynamicists to examine the possibility of a subcritical transition to turbulence, with the linear stable flow finding a way to bypass the usual route to turbulence (via linear classical/exponential instability). On closer examination, the perturbations reveal rich and complex behavior in the early transient phase, which leads to the expectation that they may become self-sustaining when there is nonlinear positive feedback. Based on the interplay of linear transient growth and nonlinear positive feedback, a new concept emerged in the hydrodynamic community for the onset of turbulence in spectrally stable shear flows and was named \emph{bypass transition} (cf. Boberg \& Brosa 1988; Butler \& Farrell 1992; Farrell \& Ioannou 1993; Reddy \& Henningson 1993; Gebhardt \& Grossmann 1994; Henningson \& Reddy 1994; Baggett et al. 1995; Grossmann 2000; Reshotko 2001; Chagelishvili et al. 2002; Chapman 2002). The bypass scenario differs fundamentally from the classical scenario of turbulence. In the classical model, exponentially growing perturbations permanently supply energy to the turbulence and they do not need any nonlinear feedback for their self-sustenance, so the role of nonlinear interaction is just to reduce the scale of perturbations to that of viscous dissipation. In the bypass model, nonlinearity plays a key role. The nonlinear processes are conservative, but in the case of positive feedback, they ensure the repopulation of perturbations that are able to extract energy transiently from the mean flow. The self-sustenance of turbulence is then the result of a subtle and balanced interplay of linear transient growth and nonlinear positive feedback. Consequently, thorough examination of the nonlinear interaction between perturbations is a problem of primary importance, and the first step is to search and to describe the linear perturbation modes that will participate in the nonlinear interactions. Such linear transient growth is also at work in rotating hydrodynamic disk flows; however, the Coriolis force causes a quantitative reduction of the growth rate there which delays the onset of turbulence. Keplerian flows are therefore expected to become turbulent for Reynolds numbers a few order of magnitudes higher than for plane subcritical flows (see: Longaretti 2002; Tevzadze et al. 2003). The possibility of an alternate route to turbulence gave new impetus to the research on the dynamics of astrophysical disks (Lominadze et al. 1988; Richard \& Zahn 1999; Richard 2001; Ioannou \& Kakouris 2001; Tagger 2001; Longaretti 2002; Chagelishvili et al. 2003; Tevzadze et al. 2003; Klahr \& Bodenheimer 2003; Yecko 2004; Afshordi et al. 2004 Umurhan \& Regev 2004; Umurhan \& Shaviv 2005; Klahr 2004; Bodo et al. 2005; Mukhopadhyay et al. 2005; Barraco \& Marcus 2005; Johnson \& Gammie 2005a, 2005b; Umurhan 2006). By adapting the progress of the hydrodynamic community to the disks flow, this research is promising for solving the disks' hydrodynamic turbulence problem. But it remains to be seen whether this route to turbulence actually applies to astrophysical disks. Compared to plane shear flows, these possess two additional properties: differential rotation and vertical stratification. Separate studies of these factors show that each exerts a stabilizing effect on the flow: these include numerical calculation of the stability of unstratified flows by Shen et al. (2006), experiments on Keplerian rotation without stratification by Ji et al. (2006), estimates of the growth rates with stratification by Brandenburg \& Dintrans (2006). However, it appears that the combined action of differential rotation and stratification introduces a new degree of freedom that may influence the flow stability and lead to turbulence at a high enough Reynolds number. Indeed, it has been shown that strato-rotational flows may exhibit global instability in bounded domains (Dubrulle et al. 2005). However, it is probable that local disk dynamics will also lead to hydrodynamic turbulence. The study of the linear perturbations in strato-rotational flow in the local limit can be found in Tevzadze et al. (2003; hereafter T03); it is shown there that the combined action of rotation and stratification generates an aperiodic vortex mode, which undergoes nonmodal transient growth. We conjecture that this transient growth may be the main energy source for the turbulence in the bypass scenario. Although the importance of the transient exchange of energy between perturbations and mean flow has now been realized by most working in the field, only a few seem aware that another linear process may play also a crucial role, namely the linear coupling of modes, which allows for transient exchange of energy between them. As shown by Chagelishvili et al. (1997a,b) and Gogoberidze et al. (2004), the energy exchange between modes is inherent to shear flows (as the transient exchange of energy between the mean flow and perturbations), and it determines in many respects the diversity of perturbation modes and, therefore, of nonlinear processes. Once one fully realizes the role of nonlinear processes in the bypass concept discussed above, it becomes evident that the neglect of linear mode coupling may lead to an incorrect picture of nonlinear (and, consequently, turbulent) phenomena. There are signs of such linear mode coupling in the simulations performed by Klahr (2004), Barraco \& Marcus (2005), and Brandenburg \& Dintrans (2006), but apparently they were not identified as such. Compressive waves are present in the simulation by Johnson \& Gammie (2005b), along with vortical perturbations, but their origin is not recognized, namely linear mode coupling. On this mode coupling, attention is focused in T03 and Bodo et al. (2005). The latter paper studies the linear dynamics of an imposed two-dimensional pure vortex mode perturbation, in a compressible Keplerian disk with constant mean pressure and density. (Two-dimensionality, i.e. the neglect of cross-disk variation, is only correct for perturbations with characteristic scales close to or larger than the vertical stratification scale of the disk.) Two modes -- a vortex mode and a spiral-density wave mode -- exist in the system, and they are strongly coupled. This investigation points out the importance of mode coupling and the necessity of considering compressibility for dynamic processes with characteristic scales close to or larger than the disk thickness. In T03 we studied the linear dynamics of three-dimensional small-scale perturbations (with characteristic length scales much shorter than the disk thickness) in compressible, vertically (stably) stratified Keplerian disks. \emph{The first novelty} presented in that paper is the occurrence of an interplay between the disk rotation and stratification. Separately, each of these factors is stabilizing. However, their interplay gives rise to a new vortex/aperiodic mode that is able to extract the basic flow energy transiently. \emph{The second novelty} is the existence of a linear coupling of that vortex mode with spiral-density wave (SDW) modes, that makes SDWs valuable participants of the dynamical processes. Furthermore, we compared the linear dynamics of the small-scale perturbations in steady stratified disks with that of the perturbations in unbounded plane Couette flow. We showed that just the linear coupling of the disk modes makes the dynamics in the disk and plane flows similar to each other (provided the Reynolds number of the disk flow is chosen about three orders of magnitudes higher than in the plane flow). This similarity suggests that the bypass concept can also be applied to disk flow and it further motivates investigation in this direction. In this respect, the present paper is a sequel of T03: we focus again on the mathematical and physical aspects of mode coupling, while introducing a significant simplification by neglecting the rotational-acoustic waves. This is justified by the fact that the characteristic timescale of these modes is much shorter than for the two other modes, when the characteristic lengthscale of the perturbations is much shorter than the disk thickness. Then the rotational-acoustic waves are not coupled to the vortex and SDW modes, and they play a negligible role in the slow, small-scale dynamics. That is why we can cut out the rotational-acoustic waves (i.e., the flow compressibility) without detriment to the dynamic picture and confine ourselves to the Boussinesq approximation. Keeping just vortex and SDW modes, this approximation simplifies mathematical aspects of the problem and allows an advance in the analytical description and comprehension of mode coupling. The paper is organized as follows. In Sect. 2 we introduce the physical approximations and the mathematical formalism, and describe the linear strato-rotational balance and the perturbation modes. In Sect. 3 we present the qualitative and quantitative analysis of the linear dynamics of perturbations. We summarize and discuss the results in Sect. 4. | Our motivation in studying mode coupling in the presence of vertical stratification is quite clear: in unstratified rotating flows the strato-rotational balance is absent (see Eq. \ref{vort1}), hence there will be no vortex mode, whose role is so powerful in extracting the shear flow energy. Moreover, in stratified rotating flows mode coupling causes the generation of SDW, which are able to conserve the extracted energy. Thus one expects astrophysical disks, which are vertically stratified, to demonstrate intrinsically different behavior compared to unstratified rotating flows. Therefore it is not possible to draw conclusions about the stability of astrophysical disks by considering unstratified rotating flows, as has been done often in the past and again in recent investigations: analytical (Balbus 2006), numerical (Shen et al. 2006), and experimental (Ji et al. 2006). Our analysis shows that, in the local limit, spectrally stable stratified Keplerian disks allow for two modes of perturbations -- vortex and SDW -- that are jointly able to extract the background flow energy and determine the disk dynamical activity in the small-scale range. These modes are linearly coupled due to the non-normality of Keplerian/shear flow. The coupling is asymmetric: the vortex mode is able to generate the related SDW, but the inverse is not true. This mode coupling is transient (like the energy exchange between perturbations and the basic flow): the SFH of the vortex mode generates the wave SFHs during the brief time interval where it switches from leading to trailing, thus rendering the dynamics non-adiabatic. At first sight, the mode coupling described here seems similar to the phenomenon of geostrophic adjustment, since both lead to the generation of the spiral density waves. However, these processes are intrinsically different. Geostrophic adjustment is an initial value problem that describes the transition from an initially unbalanced state to that of geostrophic balance. In the general case, part of the initial conditions containing no potential vorticity (the wave component) will radiate away as inertial-gravity waves (see, e.g., Pedlosky 2003) during the geostrophic adjustment, whereas our initial conditions do not include zero potential vorticity corresponding to the wave component. Hence, we have eliminated the wave generation due to the process of the initial geostropic adjustment. The waves generated in our case stem from linear mode coupling induced by the velocity shear. Moreover, geostrophic adjustment is mainly a nonlinear process that leads to equilibrium. In contrast, wave generation due to mode coupling is a linear process and does not describe the relaxation of the system, but hopefully the opposite: it promotes the transition to turbulence. The linear dynamics of each leading SFH of vortex mode proceeds in the following way: initially, the SFH extracts energy from the basic flow and it grows. At the same time $~k_x(t)/k_y \to -0$. Becoming trailing, the vortex SFH generates the related SFH of SDW. In what follows, while tilting (i.e., increasing $~k_x(t)/k_y$), the wave SFH keeps the energy, whereas the vortex SFH gives its energy back to the basic flow, so the energy gained by the leading vortex SFH is conserved by the SDW. This course of events -- the transient growth plus coupling -- is strongly pronounced for SFHs with $~k_z/k_y \sim 1$ and makes their energy dynamics similar to that of a 3D unbounded plane Couette flow. We are aware that there are differences between them: in the Couette case, only the vortex mode participates in the dynamical process, whereas in the disk case, this role is played (in the local limit) by the symbiosis between vortex and SDW perturbations. However, given the similarities that we have discussed, we conjecture that the bypass concept, which has been developed for the transition to turbulence in the Couette flow, is also applicable to rotating stratified disks. One question remains, of course: will the nonlinear interactions provide positive feedback that is efficient enough for that bypass transition? To answer it, more numerical simulations are needed, which must include the vertical stratification and grasp the mode coupling. | 7 | 10 | 0710.3648 |
0710 | 0710.3224_arXiv.txt | We report the detection of OH~1667~MHz and wide \hi~21cm absorption at $z \sim 0.67$ towards the red quasar 1504+377, with the Green Bank Telescope and the Giant Metrewave Radio Telescope. The \hi~21cm absorption extends over a velocity range of $\sim 600$~km/s blueward of the quasar redshift ($z=0.674$), with the new OH~1667~MHz absorption component at $\sim -430$~\kms, nearly coincident with earlier detections of mm-wave absorption at $z \sim 0.6715$. The atomic and molecular absorption appear to arise from a fast gas outflow from the quasar, with a mass outflow rate ${\dot M} \sim 12 M_\odot$~yr$^{-1}$ and a molecular hydrogen fraction $f_{\rm H_2} \equiv (N_{\rm H_2}/N_{\rm HI}) \sim 0.2$. The radio structure of 1504+377 is consistent with the outflow arising due to a jet-cloud interaction, followed by rapid cooling of the cloud material. The observed ratio of HCO$^+$ to OH column densities is $\sim 20$~times higher than typical values in Galactic and high-$z$ absorbers. This could arise due to small-scale structure in the outflowing gas on sub-parsec scales, which would also explain the observed variability in the \hi~21cm line. | \label{intro} The quasar 1504+377 is a rare case of a radio-loud active galactic nucleus (AGN) hosted by a disk galaxy (e.g. \citealt{perlman96,carilli97}). The flat-spectrum radio emission arises from a compact core and a one-sided jet to the southwest, with the jet axis aligned (within $\sim 15^\circ$) with the major axis of the host galaxy \citep{polatidis95,fomalont00}. The AGN is heavily reddened ($r$-K~=~5.1) and was not detected in a deep R-band image, suggesting a high level of dust obscuration \citep{stickel96}. Consistent with this, strong redshifted mm-wave molecular absorption has been detected towards the radio source \citep{wiklind96}, with two absorption complexes at $z \sim 0.6734$ (system~A) and $z \sim 0.6715$ (system~B), close to the redshift of the host galaxy ($z = 0.674 \pm 0.001$; \citealt{stickel94}). Besides the mm-wave transitions, \hi~21cm, OH~1665~MHz and OH~1667~MHz absorption have all been detected from system~A, with strong, wide profiles extending over a velocity range of $\gtrsim 100$~\kms\ \citep{wiklind96,carilli97,kanekar02}. In contrast, the mm-wave absorption in system~B is quite narrow [full-width-at-half-maximum (FWHM)$ \sim 15$~\kms] and neither \hi~21cm nor OH absorption have been detected at this redshift \citep{carilli97,carilli98}. This is the only $z > 0.1$ mm-wave absorber that has not hitherto been detected in OH or \hi~21cm absorption \citep{wiklind94,wiklind95,wiklind96,wiklind96b,chengalur99,kanekar02,kanekar03c} and is thus an excellent candidate for a deep search in these transitions. Beside studying physical conditions in the interstellar medium (ISM) of the QSO host, the detection of these lines would, in principle, also allow one to test for changes in the fundamental constants from $z \sim 0.67$ to the present epoch \citep{darling03,chengalur03,kanekar04b}. Unfortunately, the OH~1665~MHz line from $z \sim 0.6715$ lies at the same frequency as the known 1667~MHz absorption from $z \sim 0.6734$ \citep{kanekar02}, implying that it (and the latter 1667~MHz line) cannot be used to probe fundamental constant evolution. We report here a search for the other three redshifted OH ground-state lines (at rest frequencies of 1667.3590, 1612.2310 and 1720.5299~MHz) and the \hi~21cm line towards 1504+377 with the Giant Metrewave Radio Telescope (GMRT) and the Green Bank Telescope (GBT), resulting in the detection of OH~1667~MHz and \hi~21cm absorption at $z \sim 0.6715$. | \label{sec:discuss} \begin{figure} \centering \epsfig{file=fig3.eps,height=3.3truein,width=3.3truein} \caption{Final difference spectrum between the \hi~21cm optical depth spectra towards 1504+377 in 2003 and 2006, with optical depth difference (in units of $10 \times \tau_{diff}$) plotted against heliocentric velocity in \kms\, relative to $z = 0.674$. The original difference spectrum had a resolution of $\sim 0.54$~\kms; this was boxcar-smoothed to, and resampled at, a resolution of $\sim 4.8$~\kms\ to produce this spectrum.} \label{fig:variability} \vskip -0.1in \end{figure} \subsection{Variability in the \hi~21cm profile} \label{sec:21cm-change} Fig.~\ref{fig:variability} shows a plot of the difference between the \hi~21cm optical depths in 2003 and 2006 versus heliocentric velocity, in \kms, relative to $z = 0.674$. The strong features in the difference spectrum between $\sim -150$~\kms\ and $\sim -70$~\kms\ indicate significant changes ($\sim 10\%$ of the line depth) in the \hi~21cm profile between 2003 and 2006. Note that the difference cannot be due to a simple scaling of one or both of the spectra, as different spectral components show changes of opposite sign. While the possibility that the observed change might be due to RFI cannot be ruled out, no evidence was seen for RFI at these frequencies, in either these or our other 850~MHz GBT datasets. The profile ``variability'' is coincident with the strongest spectral components, with the rest of the profile showing no evidence for changes within the noise. Variability in redshifted \hi~21cm profiles has been seen earlier in two damped Lyman-$\alpha$ systems, at $z \sim 0.524$ towards 0235+164 \citep{wolfe82} and $z \sim 0.3127$ towards 1127$-$145 \citep{kanekar01c}. While changes in the latter two profiles have been detected on far shorter timescales (a few days) than in 1504+377, it is interesting that all three sources contain highly compact ($\sim $~mas-scale) components. Possible explanations for the observed changes towards 1504+377 include refractive scintillation in the Galactic interstellar medium (for which the background source need not be compact; \citealt{macquart05}), or transverse motion of a source component on VLBI scales \citep{briggs83c}. Both models require small-scale structure in the 21cm optical depth of the absorbing gas. \subsection{Physical conditions in the absorbing gas} \label{sec:1504} The radio core of 1504+377 and the nucleus of the host galaxy are coincident within the errors ($\sim 1''$) in the R-band image of \citet{stickel94}. At mm-wave frequencies, the core dominates the quasar flux density, with very little emission coming from the steep-spectrum jet \citep{wiklind96}. The core is also likely to be extremely compact at these frequencies, implying that both mm-wave absorbers arise along a single line of sight, which must also pass extremely close to the centre of the host galaxy. \citet{wiklind96} noted that it is impossible to produce two absorption components at very different velocities in such circumstances if the absorbing gas is in pure rotational motion. The large separation ($\sim 330$~\kms) between the two observed absorption velocities is thus suggestive of the presence of strong non-circular orbits; these authors argued in favour of a scenario in which the broad absorption from system~A originates close to the nucleus (in a nuclear ring or a bar), while the narrow absorption of system~A arises in a more-distant cloud in the disk of the host galaxy. In this picture, the systemic redshift is $z \sim 0.6715$. On the other hand, \citet{carilli97} used the fact that the optical emission redshift of the host galaxy ($z = 0.674 \pm 0.001$) is in excellent agreement with that of the higher-redshift complex to argue that the latter is the systemic redshift. They also pointed out that the outflow velocity of system~A in this model ($\sim 330$~\kms) is too large for a cloud in the outer disk of the parent galaxy and suggested the possibility that it might arise in a high-velocity cloud, due to tidal interactions between the host galaxy and a nearby object seen in the R-band image of \citet{stickel94}. Our new GBT \hi~21cm spectrum of Fig.~\ref{fig:21cm} clearly shows that the two absorption systems are, in fact, part of a continuous absorption complex, spanning $\sim 600$~\kms\ and extending from the optical redshift of $z \sim 0.674$ out to $z \sim 0.6706$. The 21cm absorption lies entirely blueward of the optical redshift, implying that it must arise in gas that is outflowing from the quasar. The large velocity spread of the \hi\ outflow in 1504+377 is similar to that seen in a number of low-redshift AGNs \citep{morganti05}. These authors note that all known fast \hi\ outflows have been detected in radio galaxies in early or re-started phases of their radio activity. There is also evidence that the most likely mechanism to explain such fast \hi\ outflows is interaction between the radio jets and the surrounding interstellar medium (e.g. \citealt{morganti05b}), with rapid cooling taking place in the gas after a jet-cloud interaction, as expected from numerical simulations (e.g. \citealt{fragile04}). The fact that 1504+377 shows no extended radio structure (the outer jet extends to only $\sim 55$~mas, i.e. $\sim 387$~pc, from the nucleus; \citealt{polatidis95}) suggests that it too is in a early phase of its radio activity. Recent 5~GHz VLBI observations \citep{bolton06} have found a new north-eastern extension, which was not seen in earlier (deeper) images (e.g. \citealt{fomalont00}), demonstrating that the source is currently in an active phase. Finally, the fact that the radio structure in 1504+377 is strongly one-sided (e.g. \citealt{fomalont00}) indicates that the jet lies close to the line of sight towards the core. The above suggestion that jet-cloud interactions are responsible for local gas cooling is consistent with the fact that mm-wave absorption (which takes place in cold gas and, as noted earlier, must arise towards the core) is seen at multiple velocities along the line of sight. It thus appears that the wide \hi~21cm and molecular absorption towards 1504+377 arise in outflowing gas from the AGN that is cooling rapidly after an interaction with the south-western radio jet. This is the highest redshift at which such a high-velocity outflow has been observed (e.g. \citealt{morganti05}) and, perhaps more interesting, the first case where molecular gas has been detected in the outflow. The ${\rm H_2}$ fraction is $f_{\rm H_2} = \left[N_{\rm H_2}/N_{\rm HI}\right] \le 2 \times (\ts/100) (\tx/10) (f_{\rm OH}/f_{\rm 21})$. \citet{morganti05} assume $\ts \sim 1000$~K to estimate \hi\ column densities for sources in their sample due to the proximity of the gas to the AGN and the likely presence of shocks. Using this value for consistency gives a molecular fraction of $f_{\rm H_2} \sim 0.2$ in the outflowing gas. We estimate the mass outflow rate ${\dot M}$ using the model of \citet{heckman00}, in which a constant-velocity, mass-conserving wind flows into a solid angle $\Omega$ from a minimum radius $r_*$, viz. \begin{equation} {\dot M} = 30 \left[ \frac{\Omega}{4\Pi}\right] \left[ \frac{r_*}{1 \: {\rm kpc}} \right] \left[ \frac{N_{\rm H}}{10^{21} \: {\rm cm^{-2}}} \right] \left[ \frac{v}{300 \:{\rm km s^{-1}}} \right] \: M_\odot \: {\rm yr}^{-1}, \end{equation} \noi where $v$ is the outflow velocity and $N_{\rm H}$, the total hydrogen column density of the outflowing gas. We will assume that the minimum radius $r_*$ is $\sim 10$~pc, the size of the radio core, and, following \citet{morganti05}, that $\Omega = \Pi$~steradians and $v = \:{\rm FWBN}/2 \sim 300$~\kms. The total hydrogen column density at $z \sim 0.67$ is $N_{\rm H} = \left[ N_{\rm HI} + 2 \times N_{\rm H_2} \right] \sim 1.6 \times 10^{23}$~\cm, again assuming $\ts \sim 1000$~K. This leads to an estimated mass outflow rate of ${\dot M} \sim 12 M_\odot$~yr$^{-1}$, comparable to estimates in nearby fast \hi\ outflows \citep{morganti05}. \citet{wiklind96} noted that HCO$^+$ is highly over-abundant in system~B, enhanced by at least an order of magnitude relative to expected abundances in chemical models. The ratios of HCO$^+$ to CO and HCO$^+$ to HCN column densities here are $3 - 5$ times larger than in system~A. While such large differences in relative abundances between HCO$^+$ and species such as CO, HCN, etc, have been observed in Galactic clouds \citep{lucas98}, the ratio of OH to HCO$^+$ column densities in both the Galaxy and a sample of four redshifted HCO$^+$ and OH absorbers has been found to be fairly constant, with $N_{\rm HCO^+}/N_{\rm OH} \sim 0.03$ \citep{liszt96,kanekar02} over more than two orders of magnitude in HCO$^+$ column density. \citet{liszt04} found this ratio to show a large spread (by a factor of $\sim 4$) in the clouds towards Cen.A and NGC1052, with the HCO$^+$ and OH lines also showing very different kinematics, but argued that this could be explained by differing source structure and foreground free-free opacity at the OH and HCO$^+$ frequencies, source variability between observing epochs, and excitation effects at high OH column densities ($\gtrsim 10^{15}$~\cm; \citealt{langevelde95}). Conversely, system~B has $N_{\rm HCO^+}/N_{\rm OH} \sim 0.5 \times (\tx/10) (0.46/f_{\rm OH})$, discrepant by more than an order of magnitude from the expected value. However, 1504+377 is highly compact at both mm-wave and cm-wave frequencies (with a cm-wave core-fraction of $\sim 46$\%; \citealt{carilli97}) and the HCO$^+$ and OH lines have very similar FWHMs [$\sim 16.5 \pm 2.2$~\kms\ (OH) and $\sim 15.2 \pm 0.9$~\kms\ (HCO$^+$)], making it likely that they arise from similar gas. Increasing $\tx$ by an order of magnitude could resolve this problem but such high $\tx$ values have never been seen in the Galaxy (e.g. \citealt{liszt96}). Similarly, the ratio of peak optical depths in the HCO$^+$ and \hi~21cm lines in system~A is $R \equiv \tau_{HCO^+}/\tau_{\rm 21} \sim 30$, far larger than that seen in Galactic clouds ($0.1 \le R \le 6$; \citealt{lucas96,liszt96}). \citet{carilli97} point out that high values of $R$ could result from either far warmer \hi\ or low molecular dissociation, but this would not explain the discrepancy in the ratio of OH and HCO$^+$ column densities. If the latter is not due to real chemical differences between OH and HCO$^+$ (which seems unlikely; \citealt{liszt00}), a plausible explanation is extreme small-scale structure in the opacity of the absorbing gas on sub-parsec scales, smaller than the size of the radio core at cm wavelengths. This could arise due to internal shocks or turbulence in the rapidly outflowing gas. As noted earlier, the observed variability in the \hi~21cm absorption at $z \sim 0.674$ suggests similar small-scale structure at a different location in the outflow, which could also account for the large velocity difference ($\sim 15$~\kms) in peak OH and HCO$^+$ absorption in system~A \citep{kanekar02}. Monitoring the mm-wave lines for variability would be one way of testing this hypothesis. Finally, comparisons between the OH, \hi~21cm and HCO$^+$ redshifts from an absorber can be used to test the evolution of fundamental constants \citep{darling03,chengalur03}. However, the above possibility of small-scale structure in the absorbing gas makes it likely that any such comparisons in the absorbing gas towards 1504+377 will be dominated by local systematic velocity offsets. We conclude that this absorber is unlikely to be useful for the purpose of probing fundamental constant evolution. | 7 | 10 | 0710.3224 |
0710 | 0710.3189_arXiv.txt | We present a simple and efficient method to set up spherical structure models for $N$-body simulations with a multimass technique. This technique reduces by a substantial factor the computer run time needed in order to resolve a given scale as compared to single-mass models. It therefore allows to resolve smaller scales in $N$-body simulations for a given computer run time. Here, we present several models with an effective resolution of up to $1.68 \times 10^9$ particles within their virial radius which are stable over cosmologically relevant time-scales. As an application, we confirm the theoretical prediction by \citet{2005MNRAS.360..892D} that in mergers of collisonless structures like dark matter haloes always the cusp of the steepest progenitor is preserved. We model each merger progenitor with an effective number of particles of approximately $10^{8}$ particles. We also find that in a core-core merger the central density approximately doubles whereas in the cusp-cusp case the central density only increases by approximately 50\%. This may suggest that the central region of flat structures are better protected and get less energy input through the merger process. | Resolution is a key issue in $N$-body simulations. In state-of-the-art cosmological $N$-body simulations today, structures can be fully resolved down to the scale of a fraction of a per cent of the virial radius \citep{2007ApJ...657..262D}. But there are many problems where even higher resolution in central regions of structures is needed. Also, one often would like to study a certain problem without the cosmological context, i.e. one needs a possibility to set up isolated structures with high resolution in order to study dynamical effects in detail and with precise control of the initial condition, which can be very difficult within the framework of a cosmological $N$-body simulation. For example, the question of whether the central dark matter density profile of haloes that form in cosmological $N$-body simulations is cuspy or cored needs at least a resolution of $\approx 10^{-3}~\rvir$ to be answered. Another example are centrally flat profiles: in order to resolve isolated structures with flat central profiles, a lot of particles are needed since in flat profiles the resolution scaling with the number of particles is the slowest (see below for more details about the scaling of resolution with the number of particles). Another problem we had in mind when developing the multimass technique presented in this paper, was the dynamics of super-massive black hole binaries in the centre of remnants of galaxy mergers. There, the relevant scales are of order of a few pc $\approx 10^{-6}-10^{-5}~\rvir$ for Milky Way size dark matter haloes. \begin{figure*} \centering \includegraphics[width=0.49\textwidth]{g10.res.eps} \includegraphics[width=0.49\textwidth]{g15.res.eps} \caption{Plot of $r_1$, $r_{100}$, $\rrel(\tdyn(\rvir)/10)$, $\rrel(\tdyn(\rvir))$ as a function of $\Nvir$ for a spherical structure with $\cvir = 10$, $\alpha = 1$, $\beta = 3$ but different central profiles: $\gamma = 1$ (left) and $\gamma = 3/2$ (right). One can nicely see that in general the constraint on $\rres$ is set by the scale where relaxation becomes important since $\rres = \max(r_{100}(\Nvir),\rrel(t_0,\Nvir))$ for a given simulation time $t_0$. The asymptotic scaling for $r_N$ and $\rrel$ is given in the plots; the discrepancy between the sampled scale $r_{100}$ and the relaxation scale $\rrel$ as a function of $\Nvir$ is bigger for larger $\gamma$.} \label{fig:resolution} \end{figure*} We illustrate the resolution problem in more detail with a commonly used family of spherically symmetric density profiles: the so called $\alpha\beta\gamma$-models family \citep{1996MNRAS.278..488Z}. An $\alpha\beta\gamma$-model density profile is given by \begin{equation} \label{eq:rho} \rho(r) = \frac{\rho_0}{\left(\frac{r}{\rs}\right)^\gamma \left(1+\left(\frac{r}{\rs}\right)^\alpha \right)^{\left(\frac{\beta-\gamma}{\alpha}\right)}}~, \end{equation} where $\gamma$ determines the inner slope and $\beta$ the outer slope of the density profile, whereas $\alpha$ controls the transition between the inner and outer region. The normalisation is given by $\rho_0$ and $\rs$ is the scale radius defined by $\rs \equiv \rvir/\cvir$ ($\cvir$ is the virial concentration). Many models used by observers as well as theorists to describe structures in the universe belong to this family, e.g. the Hernquist profile \citep{1990ApJ...356..359H}, the NFW profile \citep{1996ApJ...462..563N} or the Moore profile \citep{1998ApJ...499L...5M} are just special cases of the general form (\ref{eq:rho}). An obvious minimal criterion for a given length scale to be resolved is, that it has to be populated with enough particles by the sampling procedure. For example, if one would like to set up an isolated structure given by the density profile (\ref{eq:rho}), then the radius $r_{N}$ that contains $N$ particles is simply given by the solution of \begin{equation} \label{eq:rNdef} M(r_{N}) = N m~, \end{equation} where $M(r)$ is the enclosed mass given by \begin{equation} \label{eq:Menc} M(r) \equiv \int_0^r 4 \pi \rho(x) x^2 \d x~, \end{equation} and $m$ is the mass of the particles which can be expressed in virial quantities for single-mass models as $m = \Mvir / \Nvir$. In the central asymptotic regime (i.e. $r_{N} \ll \rs$), equation (\ref{eq:rNdef}) can be solved for $r_{N}$ to yield \begin{equation} \label{eq:rN} r_{N} = \left(\frac{3-\gamma}{4 \pi} \frac{N m}{\rho_0 \rs^{\gamma}}\right)^{\frac{1}{3-\gamma}}~. \end{equation} For $N = 1$ one obtains the scale of the location of the innermost particle, $\rimp = r_1$.\footnote{An alternative definition for the location of the innermost particle can be made with the mean particle separation given by $h(r) \equiv \sqrt[3]{\frac{m}{\rho(r)}}$. The two definitions are essentially equivalent and only differ by the geometric factor $\left(\frac{4 \pi}{3-\gamma}\right)^{\frac{1}{3-\gamma}}$ which is of order unity for most useful profiles.} But of course one particle is not enough to resolve that scale. If we say at least 100 particles are needed in the innermost bin so that the bin is well resolved, then $r_{100}$ provides a better estimate of the resolved scale. Another constraint to resolution comes from the fact that often the structures that one would like to simulate can be very well approximated as collisionless systems, i.e. the local relaxation time-scale is much longer than the age of the system. This property must be preserved in $N$-body simulations of such collisionless systems. Otherwise, purely artificial relaxation processes due to under-resolving the structures will have a dominant effect on the dynamics of the system \citep{2003MNRAS.338...14P,2004MNRAS.348..977D,2004MNRAS.349.1117B}. This is not a numerical artefact since a real astrophysical system with a small particle number will also be subject to relaxation processes. The problem is, that in most cases, we can not simulate the astrophysical system under study with the number of particles the system would have in nature, e.g. a galaxy size dark matter halo would have of the order of $O(10^{67})$ dark matter particles if we assume that the dark matter particle has a mass of 100 GeV/c$^2$. Therefore, this effect due to under-resolving the system with not enough particles will always be a limitation of collisionless $N$-body simulations. In principle, there are only few known dynamically stable systems in the universe, e.g. the Keplerian 2-body system or the figure-8 rotation of 3 bodies \citep{1993PhRvL..70.3675M, 2000AM....152..881C}.\footnote{This is only true in the Newtonian regime. In General Relativity, not even a 2-body orbit is dynamically stable and the orbit decays due to the emission of gravitational radiation \citep{1979RvMP...51..447D, 1977QJRAS..18....3I, 1997RvMP...69..337A}.} Dynamical effects like relaxation or evaporation will sooner or later lead to the disruption of any system. The issue is always on what time-scale the system is actually stable. The time-scale of interest is normally of the order of the age of the universe and for most cases collisionless astrophysical systems can be regarded as perfectly stable within that time-scale. This is different in $N$-body simulations. Here, one just tries to generate models that show the desired stability behaviour during the time-scale of interest with the least amount of particles needed. Artificial $N$-body models will be subject to such disruption effects much sooner than the real astrophysical system one wants to study. We illustrate this in more details now. The local relaxation time is defined by \begin{equation} \label{eq:trelax} \trel(r) \equiv \frac{N(r)}{\ln(N(r))}~\tdyn(r)~, \end{equation} where \begin{equation} \label{eq:tdyn} \tdyn(r) \equiv 2 \pi \sqrt{\frac{r^3}{G M(r)}} \end{equation} is the dynamical or orbital time at radius $r$ and $N(r) \equiv M(r)/m$ denotes the number of particles within $r$. Here, we have a slightly different normalisation than in the usual expression for the local relaxation time since we dropped the factor of 8 in front of the logarithmic term in the denominator. We found better agreement with results from $N$-body simulations with this normalisation (see section \ref{chap:tests} below for more details). Had we kept the factor of 8, then our definition (\ref{eq:trelax}) for spherically symmetric structures would be equivalent to the empirical expression found by \citet{2003MNRAS.338...14P}. Normally, the Coulomb logarithm is given by $\ln(\Lambda) = \ln(b_{\mathrm{max}}/b_{\mathrm{min}})$, where $b_{\mathrm{max}}$ and $b_{\mathrm{min}}$ are the maximum and minimum impact parameters of the particles under consideration. Since the minimum impact parameter is related to the softening length and the latter scales with the number of particles, we prefer the direct formulation of the Coulomb logarithm as a function of the number of particles, $\ln(N(r))$. In a simulation run for a time $t_0$ relaxation processes become important on a scale $\rrel$ given by the solution of \begin{equation} \label{eq:rrelaxdef} \trel(\rrel) = t_0~. \end{equation} In the central asymptotic regime (i.e. $\rrel \ll \rs$), equation (\ref{eq:rrelaxdef}) can be inverted for $\rrel$ to yield \begin{equation} \label{eq:rrelax} \rrel(t_0) = \left( \frac{W_{-1}(X)~\frac{(3 - \gamma) m}{4 \pi \rho_0 \rs^{\gamma}}} {-\frac{1}{2} \frac{6-\gamma}{3-\gamma} \frac{\pi}{t_0} \left( \frac{3-\gamma}{G \pi \rho_0 \rs^{\gamma}} \right)^{\frac{1}{2}}} \right)^{\frac{2}{6-\gamma}} \end{equation} where $W_{-1}$ denotes the $k = -1$ branch of the Lambert W function \citep{1758AH......3..128L,1996ACM.....5..329C} and \begin{equation} X \equiv -\frac{1}{2} \frac{6-\gamma}{3-\gamma} \frac{\pi}{t_0} \left(\frac{3-\gamma}{G \pi \rho_0 \rs^{\gamma}}\right)^{\frac{1}{2}} \left( \frac{(3 - \gamma) m}{4 \pi \rho_0 \rs^{\gamma}} \right)^{\frac{\gamma}{2 (3-\gamma)}}~. \end{equation} In Fig. \ref{fig:resolution} we plot $r_1$, $r_{100}$, $\rrel(\tdyn(\rvir)/10)$ and $\rrel(\tdyn(\rvir))$ as a function of $\Nvir$ by setting $m = \Mvir / \Nvir$. The virial radius $\rvir$ is defined so that the enclosed average density within $\rvir$ is given by \begin{equation} \frac{\Mvir}{4 \pi \rvir^3/3} = \rhovir = \Deltavir \rho_{\mathrm{crit},0} = 1.41\times10^{4}~\Mo~\kpc^{-3} \end{equation} where $\Deltavir = 178~\Omega_{\mathrm{M},0}^{0.45} \approx 104$ \citep{1998ApJ...503..569E} for our choice of cosmology with $\Omega_{\mathrm{M},0} = 0.3$, $\Omega_{\Lambda,0} = 0.7$, $\rho_{\mathrm{crit},0} = {3 H_0^2}/{8 \pi G}$ and $H_0 = 70~\km~\s^{-1}~\Mpc^{-1}~(h_0 = 0.7)$, which we use throughout this paper. The dynamical time at the virial radius is then \begin{equation} \tdyn(\rvir) = \sqrt{\frac{3 \pi}{G \Deltavir \rho_{\mathrm{crit},0}}} = 12.2~\Gyr~. \end{equation} For example for a galaxy size dark matter halo with $\Mvir = 10^{12}/h_0~\Mo = 1.43 \times 10^{12}~\Mo$ one would obtain $\rvir = 289~\kpc$ and the values for systems of different virial mass $\Mvir$ can be obtained by the simple scaling relation \begin{equation} \rvir = 289~\kpc \sqrt[3]{\frac{\Mvir}{1.43 \times 10^{12}~\Mo}}~. \end{equation} Although the virial radius is a rather artificial definition of the size of a dark matter halo and the virialised region of haloes in cosmological $N$-body simulations is generally much larger \citep{2006ApJ...645.1001P,2007ApJ...667..859D} it is a convenient normalisation and cut-off scale for isolated models since we are anyway mainly interested in the central dynamics of the structure. We set $\cvir = 10$, $\alpha = 1$, $\beta = 3$ in both cases and present plots for $\gamma = 1$ (left) respectively $\gamma = 3/2$ (right). We only plot these quantities for $\Nvir \geq 10^{6}$ since the expressions (\ref{eq:rN}) and (\ref{eq:rrelax}) are only valid in the central asymptotic regime. One can see that $r_N \propto \Nvir^{-\frac{1}{3-\gamma}}$ and $\rrel \sim \Nvir^{-\frac{2}{6-\gamma}}$ since the Lambert W function only varies very slowly as a function of $\Nvir$ which reflects the weak dependence of the logarithmic term of the local relaxation time. For a given particle resolution $\Nvir$ and simulation time $t_0$, the radius $\rres$ that we can still resolve with correct collisionless physics (i.e. this scale does not suffer from too much artificial relaxation) is given by \begin{equation} \rres = \max(r_{100}(\Nvir),\rrel(t_0,\Nvir))~. \end{equation} As one can see from Fig. \ref{fig:resolution}, this resolution scale $\rres$ is in general set by $\rrel(t_0,\Nvir)$ for isolated high resolution structures. It is worth remarking here, that for structures assembled hierarchically in a cosmological $N$-body simulation, the amount of relaxation is significantly larger and $\rrel$ scales slower as a function of $\Nvir$ \citep{2004MNRAS.348..977D}. Although a structure might be sampled with enough particles at a certain scale at the final time, the relaxation time at that scale would have been much smaller in the past since particles were in lower mass structures during the hierarchical growth. By inspecting Fig. \ref{fig:resolution}, we see that more than approximately of order $O(10^{12})$ particles in the centre of a structure with $\gamma = 1$ are needed in order to resolve scales of $\approx 10^{-5}~\rvir$. It is worth remarking, that with the same number of particles $\Nvir$ much smaller scales are populated in the $\gamma = 3/2$ profile than in the $\gamma = 1$ profile. Generally, the steeper the central profile, the more the particles are concentrated. But unfortunately, the relaxation scale $\rrel$ does not scale equally fast so that the discrepancy as a function of $\Nvir$ becomes bigger for larger values of $\gamma$. Nonetheless, such an enormous amount of particles per structure is hardly doable today - even with large supercomputers. But since we only need this high resolution at the very centre of the structure, our solution to this problem is to use models where we only populate regions of the phase space that are in the centre or will reach the centre in the future with high resolution particles. In section \ref{chap:method} we present the simple idea behind the multimass models and present stability tests in section \ref{chap:tests}. In section \ref{chap:preservation}, we test Dehnen's prediction with high resolution mergers and we summarize our results in section \ref{chap:conclusions}. | \label{chap:conclusions} The multimass technique for modelling haloes is a simple method to perform high resolution $N$-body simulations. An early version of this technique without orbit dependent refinement has already been used successfully in several applications which also include cosmological structure formation simulations \citep{2005MNRAS.364..665D,2006MNRAS.368.1073G}. With a careful choice of parameters it is possible to gain over an order of magnitude in computer run time for a given resolution scale. As an application of this technique we confirm the earlier findings that core-core mergers lead to a cored merger remnant while cups-cusp mergers lead to a cuspy merger remnant with high resolution multimass $N$-body simulations. In cusp-core merges, the merger remnant has a final profile corresponding to the steepest progenitor which is in excellent agreement with the theoretical predictions. We find that in the core-core case the central density approximately doubles whereas in the cusp-cusp case the central density in the merger remnant only increases by approximately 50\%. This may suggest that the central region of flat structures are better protected and get less energy input from the merger. A software tool called \textsc{halogen} (HALO GENerator) for generating multimass initial conditions is available from the author upon request. | 7 | 10 | 0710.3189 |
0710 | 0710.4556_arXiv.txt | It has been known for over 30 years that Galactic globular clusters (GCs) are overabundant by orders of magnitude in bright X-ray sources per unit mass relative to the disk population. Recently a quantitative understanding of this phenomenon has developed, with a clear correlation between the number of X-ray sources in a cluster, $N_X$, and the cluster's encounter frequency, $\Gamma$, becoming apparent. We derive a refined version of $\Gamma$ that incorporates the finite lifetime of X-ray sources and the dynamical evolution of clusters. With it we find we are able to explain the few clusters that lie off the $N_X$--$\Gamma$ correlation, and resolve the discrepancy between observed GC core radii and the values predicted by theory. Our results suggest that most GCs are still in the process of core contraction and have not yet reached the thermal equilibrium phase driven by binary scattering interactions. | \label{sec:dyn} It was realized more than 30 years ago that Galactic globular clusters (GCs) are overabundant by orders of magnitude in bright X-ray sources per unit mass relative to the disk population \citep{1975ApJ...199L.143C,1975Natur.253..698K}. It was quickly understood that strong dynamical scattering interactions of binaries in the dense cluster cores should be responsible for this overabundance \citep{1987IAUS..125..187V}. With the advances in X-ray astronomy made possible by observatories such as {\em Chandra}, the relationship between X-ray sources and core cluster dynamics has recently been quantitatively studied. \citet{2003ApJ...591L.131P} performed {\em Chandra} observations of many Galactic GCs down to a limiting luminosity of $4 \times 10^{30}\,{\rm erg}/{\rm s}$ in the 0.5--6 keV range (which includes low-mass X-ray binaries [LMXBs] in outburst and quiescence, cataclysmic variables [CVs], millisecond pulsars [MSPs], and magnetically active main sequence binaries [ABs]), and looked for correlations between the number of X-ray sources in each cluster and properties of the cluster itself. They found the strongest correlation with the ``encounter frequency'' $\Gamma$, a rough estimate of the current dynamical encounter rate in the cluster. More recently, \citet{2003ApJ...598..501H} and \citet{2006ApJ...646L.143P} have isolated the quiescent LMXBs (qLMXBs) and CVs, respectively, from the X-ray source populations, and have shown that their numbers are indeed consistent with dynamical formation. These results represent quantitative, empirical evidence that dynamical encounters are responsible for the formation of X-ray sources in clusters. However, they suffer from at least a few drawbacks. First, the correlation between the number of X-ray sources, $N_X$, and the encounter frequency appears to be sub-linear, with $N_X \propto \Gamma^{0.74 \pm 0.36}$, although for LMXBs the exponent is $0.97 \pm 0.5$ \citep{2003ApJ...591L.131P}. Second, there are three clusters for which $N_X$ is significantly larger than predicted by $\Gamma$. In the original work of \citet{2003ApJ...591L.131P} it was already clear that NGC 6397 has an $N_X$ that is $\sim 5$ times larger than predicted by the $N_X$--$\Gamma$ correlation. Recent observations show that $N_X$ is factor of $\sim 2$ times that predicted by $\Gamma$ for NGC 7099 \citep{2007ApJ...657..286L}, and a factor of $\sim 20$ for Ter 1 \citep{2006MNRAS.369..407C}. The common thread among these three clusters is that they are observationally ``core-collapsed,'' while all others in the \citet{2003ApJ...591L.131P} sample are not \citep{2006AdSpR..38.2923G}. (A possible exception is NGC 6752, whose collapsed core status is debated \citep{2003ApJ...595..179F,1995ApJ...439..191L}.) A cluster is observationally termed core-collapsed if its surface brightness profile is consistent with a cusp at the limit of resolution, making it more difficult to measure the core radius. As described below, the collapsed core status of a cluster is linked to its dynamical state, implying that cluster evolution complicates the $N_X$--$\Gamma$ correlation. \begin{figure} \begin{center} \includegraphics[width=0.9\columnwidth]{f1.eps} \caption{Number of observed cluster X-ray sources with $L_X\gtrsim 4 \times 10^{30}\,{\rm erg}/{\rm s}$ for several Galactic GCs versus the encounter rate $\Gamma$. The power-law fit and data points are from \citet{2003ApJ...591L.131P} with the exception of NGC 7099 \citep{2007ApJ...657..286L} and Ter 1 \citep{2006MNRAS.369..407C}. The $N_X$ error bars for Ter 1, NGC 6397, and NGC 7099 represent source counting noise and background source uncertainty, but for the remaining clusters represent only background uncertainty. \label{fig:gamma}} \end{center} \end{figure} | \label{sec:disc} This {\em letter} presents the confluence of three suggestive observational and theoretical results into a self-consistent picture. The first is that of the clusters that have been observed sufficiently to determine their XRB population, the three that are core-collapsed are the same three that have a significant X-ray source overabundance (of a factor of $\sim 2$ to $\sim 20$). The second is the semi-analytical result derived in this {\em letter} that a cluster in the binary burning phase for the last few Gyr should have $\sim 5$ times more dynamically formed X-ray sources than if it were in the core contraction phase for the same time. The third is the recently confirmed discrepancy between observations and theory for the core radii of Galactic GCs, which suggests that only the observationally core-collapsed clusters are in the binary-burning phase. In light of these facts, the conclusion that seems strongly suggested is that most Galactic GCs are currently still in the core contraction phase, while only the $\sim 20\%$ of clusters that are core-collapsed are in the binary burning phase. This goes counter to the widely held belief that most clusters are currently in the binary burning phase, and complicates the many existing studies that have assumed cluster core properties that are constant with time. The implications of this result are manifold. There are many studies of the dynamical production of interesting source populations in clusters which assume core properties that are constant with time. These include predictions of the formation of blue stragglers \citep[e.g.,][]{2004ApJ...605L..29M}, the evolution of the core binary fraction \citep{2005MNRAS.358..572I}, and tidal-capture binaries \citep[e.g.,][]{1994ApJ...423..274D}, among others. Revising the results may be as simple as scaling predicted source numbers, but may not be so simple for other quantities. Studies of GC evolution have shown that clusters starting from very different initial conditions evolve toward a common range of values in many observable structural parameters in the binary-burning phase, including the concentration and ratio of core to half-mass radius \citep{2003ApJ...593..772F,2006MNRAS.368..677H,2007ApJ...658.1047F}. Since most clusters may not be in the binary-burning phase after all, their observed properties are likely to be more strongly correlated with their initial conditions. This makes modeling of clusters a bit more complicated, but on the other hand allows one to more readily deduce something about the initial properties of clusters. Perhaps anticlimactically, our results suggest that the alternative energy sources recently proposed for supporting GC cores are not required. These include the suggestion of IMBHs in {\em many} Galactic GCs \citep{2006astro.ph.12040T}, enhanced stellar mass loss from stellar evolution of physical collision products \citep{chatterjeeposter}, mass segregation of compact remnants in young clusters \citep{2004ApJ...608L..25M}, or evaporation of the stellar-mass black hole subsystem in young clusters \citep{2007MNRAS.379L..40M}. Although the picture painted in this {\em letter} is a suggestive one, there are still several caveats and limitations to our analysis. In the derivation of our refined $\Gamma$ we assume that the core binary fraction and abundance of compact objects are constant over the time of integration. Neither is strictly true, although we expect they will not vary enough to significantly change the overabundance value we derive. We have used only one possible expression for the evolution of the core radius in core contraction. While we expect the general behavior to be very similar to what we have assumed here, more work needs to be done to determine if it is universal. Additionally, our analysis ignores the effect of Galactic tidal stripping on cluster mass, which would make a cluster appear overabundant in X-ray sources, and may be relevant for NGC 6397 \citep{2003ApJ...591L.131P}. On the observational side, there are some uncertainties in evaluating $\Gamma$, which is dependent on quantities that are somewhat difficult to measure for core-collapsed clusters. | 7 | 10 | 0710.4556 |
0710 | 0710.0375_arXiv.txt | \spider\ is a long-duration, balloon-borne polarimeter designed to measure large scale Cosmic Microwave Background (CMB) polarization with very high sensitivity and control of systematics. The instrument will map over half the sky with degree angular resolution in I, Q and U Stokes parameters, in four frequency bands from $96$ to $275$ GHz. \spider's ultimate goal is to detect the primordial gravity wave signal imprinted on the CMB B-mode polarization. One of the challenges in achieving this goal is the minimization of the contamination of B-modes by systematic effects. This paper explores a number of instrument systematics and observing strategies in order to optimize B-mode sensitivity. This is done by injecting realistic-amplitude, time-varying systematics in a set of simulated time-streams. Tests of the impact of detector noise characteristics, pointing jitter, payload pendulations, polarization angle offsets, beam systematics and receiver gain drifts are shown. \spider's default observing strategy is to spin continuously in azimuth, with polarization modulation achieved by either a rapidly spinning half-wave plate or a rapidly spinning gondola and a slowly stepped half-wave plate. Although the latter is more susceptible to systematics, results shown here indicate that either mode of operation can be used by \spider. | \label{sec:intro} In the past decade, a wealth of data have pointed to a ``standard model'' of the Universe, composed of $\sim 5\%$ ordinary matter, $\sim 22\%$ dark matter and $\sim 73\%$ dark energy in a flat geometry. The flatness of the Universe, the near isotropy of the CMB, and the nearly-scale-invariant nature of the primordial scalar perturbations from which structure grew support the existence of an early accelerating phase dubbed ``inflation''. A necessary by-product of inflation is tensor perturbations from quantum fluctuations in gravity waves. A detection of this Cosmological Gravity-Wave Background (CGB) would give strong evidence of an inflationary period and determine its energy scale, while a powerful upper limit would point to more radical inflationary scenarios, e.g., involving string theory, or some non-inflationary explanation of the observations. The CGB imprints a unique signal in the curl-like, or B-mode, component of the polarization of the CMB; detection of a B-mode signal can be used to infer the presence of a CGB at the time of decoupling. Direct detection of the gravity waves is many decades off; an advanced Big Bang Observer successor to LISA has been suggested as a way to achieve this \citep{BBO1,BBO2}. Thus a measurement of the primordial B-modes is the only feasible near-term way to detect the CGB and have a new window to the physics of the early Universe \citep{Bock:2006}. A CGB with a potentially measurable amplitude is a by-product of the simplest models of single field inflation which can reproduce the scalar spectral tilt observed in current combined CMB data~\citep{spergel2007, MacTavish:2005yk}. Examples are chaotic inflation from power law inflaton potentials~\citep{linde:1983,linde:2005} or natural inflation from cosine inflaton potentials involving angular (axionic) degrees of freedom~\citep{adams93}. The amplitude is usually parameterized in terms of the ratio of the tensor power spectrum to the scalar power spectrum, $r={\cal P}_t (k_p)/{\cal P}_s (k_p)$, evaluated at a comoving wavenumber pivot $k_p$, typically taken to be $0.002 \ {\rm Mpc}^{-1}$. Chaotic inflation predicts $r\approx 0.13$ for a $\phi^2$ potential and $r \approx 0.26$ for a $\phi^4$ potential, and natural inflation predicts $r \approx 0.02-0.05$. The potential energy $V$ driving inflation is related to $r$ by $V \approx (10^{16} \ {\rm Gev})^4 r/0.1$. Low energy inflation models have low or negligible amplitudes for the CGB. To get the observed scalar slope and yet small $r$ requires special tuning of the potential. These are often more complicated, multiple field models, e.g., \citep{linde:1993}, or string-inspired brane or moduli models~\citep{kallosh2007}. Given the collection of models it is difficult to predict a precise range for the expected tensor level and the prior probability for $r$ should be considered as wide open. Recent CMB data have reached the sensitivity level required to constrain the amplitude of and possibly characterize the gradient-like, or E-mode, component of the polarization \citep{Kovac:2002fg,hedman2002,Readhead:2004xg,Montroy:2006,page2007,quad2007}. A significant complication of the measurements is that the E-mode amplitude is an order of magnitude lower than the total intensity. In addition, galactic foregrounds such as synchrotron and dust are expected to be significant at these amplitudes~\citep{kogut2007}. Furthermore the polarization properties of foregrounds are largely unknown. Constraining B-modes presents an even greater challenge as it is a near certainty that polarized foregrounds will dominate the signal. The next generation of CMB experiments will benefit from a revolution in detector fabrication in the form of arrays of antenna-coupled bolometers~\citep{Goldin:2002,kuo2006}. The antenna-coupled design is entirely photo-lithographically fabricated, greatly simplifying detector production. In addition, the densely populated antennas allow a very efficient use of the focal plane area. \spider\ will make use of this technological advance, in the form of 2624 polarization sensitive detectors observing in four frequency bands from $96$ to $275$ GHz. A multi-frequency observing strategy is a necessary requirement to allow for a subtraction of the foreground signal. \spider\ will observe over a large fraction of the sky at degree scale resolution producing high signal-to-noise polarization maps of the foregrounds at each frequency. Extraordinarily precise understanding of systematic effects within the telescope will be required to measure the tiny B-mode signal. This {\it paper} presents a detailed investigation of experimental effects which may impact \spider's measurement of B-modes. The strategy is to simulate a \spider\ flight time-stream injecting systematic effects in the time domain. The aim is to determine the level of B-mode contamination at subsequent stages of the analysis. With these results one can set stringent requirements on experimental design criteria, in addition to optimizing the telescope's observing strategy. The outline of this paper is as follows. Section~\ref{sec:instrument} gives an overview of the instrument, flight and observing strategy. Section~\ref{sec:method} describes the details of the simulation methodology. Results for several systematic effects are presented in Section~\ref{sec:results}. Section~\ref{sec:conclusions} concludes with a summary and discussion of the results. | \label{sec:conclusions} The results from Section~\ref{sec:results} are summarized in terms of experimental specifications in Table~\ref{tab:summary}. While results are \spider-specific the order of magnitude of various effects can be translated to other CMB polarization experiments. The RMS B-mode signal for $r = 0.1$ is roughly 10000 $\rm nK^2$ ($\sim$1436 $\rm nK^2$ at $\ell = 8$ and $\sim$7379 $\rm nK^2$ at $\ell = 80$), and scales linearly with r. Experimental specifications are set by limiting the allowed systematic residual level to a factor of $\sim10$ smaller than the B-mode signal for $r = 0.01$. While the simulations were signal-only the impact of large low frequency detector noise (1/f noise) is reflected in the large scale degradation of the B-mode signal for the stepped half-wave plate mode of operation. Rapid, continuous half-wave plate modulation mitigates this effect entirely. Rapid, continuous gondola rotation also works but only with iterative map-making which accurately recovers the larger scale signal. It is important to note that the effects studied in Sections~\ref{sec:modes} and Sections~\ref{sec:noise} (naive versus iterated maps, spinning more slowly, stepping half-wave plate versus spinning half-wave plate) will degrade the signal-to-noise achieved on the bandpowers. These effects differ from the systematics studied in Section~\ref{sec:point} to Section~\ref{sec:drift} (pointing reconstruction errors, polarization angle uncertainty, uncorrected ghosting, uncorrected gain drifts) which will ultimately bias the final result. The requirements on the biasing effects are more difficult treat than the signal-to-noise issues. The impact of systematics on B-mode polarimeter experiments is also discussed in~\cite{hhz:2003} and more recently ~\cite{odea:2007}, where analytical methods are used for calculating the B-mode spectrum bias. The results are useful for setting experimental ``benchmark parameters'' at the very earliest phases of instrument design. This work goes a step further by considering the impact of systematics in the map/time-domain; A necessary step in the evolution of an experiment which aims to measure the tiny primordial, gravity wave signal. \begin{table*} \centering \begin{tabular}{|c|c|c|} \hline\hline \space {\bf Systematic} & {\bf Experimental Spec.} & {\bf Comments} \\ \hline\hline Receiver & & for 110dps\\ 1/f knee & $< 200$ mHz & gondola spin \\ \hline Receiver & & for 36dps\\ 1/f knee & $< 100$ mHz & gondola spin \\ \hline Pointing & & sufficient for\\ Jitter & $< 10'$ & $\ell < 50$ \\ \hline Absolute& & \\ Pol. Angle Offset & $< 0.25$ deg. & \\ \hline Relative& & \\ Pol. Angle Offset & $< 1$ deg. & \\ \hline Knowledge of & & sufficient for\\ Beam Centroids & $< 1'$ & $\ell < 30$ \\ \hline Optical & & \% TOD\\ Ghosting & $< 2\%$ & contamination\\ \hline Calibration & & \\ Drift & $< 3.0\%$ & in phase \\ \hline Calibration & & \\ Drift & $< 0.1\%$ & out of phase \\ \hline\hline \end{tabular} \caption{\rm Summary of experimental specifications based on simulation results. Realistic-amplitude, time-varying systematics are injected in the simulated time streams. Maps are reconstructed without any attempt to correct for the systematic errors. Experimental specifications are set by limiting the allowed systematic residual level to a factor of $\sim10$ smaller than the B-mode signal for $r = 0.01$. The nominal operating mode is a 36 dps gondola spin rate, with the half-wave plate stepping $22.5^{\circ}$ once per day, with 10 iterations (sufficient to recover the residual levels of the continuously-rotating half-wave plate case) of the map-maker, a Jacobi iterative solver~\citep{Jones:2007}. }\label{tab:summary} \end{table*} | 7 | 10 | 0710.0375 |
0710 | 0710.5285_arXiv.txt | Gamma-ray burst afterglow observations in the \emph{Swift} era have a perceived lack of achromatic jet breaks compared to the \emph{BeppoSAX} or pre-\emph{Swift} era. Specifically, relatively few breaks, consistent with jet breaks, are observed in the X-ray light curves of these bursts. If these breaks are truly missing, it has serious consequences for the interpretation of GRB jet collimation and energy requirements, and the use of GRBs as cosmological tools. Here we address the issue of X-ray breaks that are possibly `hidden' and hence the light curves are misinterpreted as being single power laws. We do so by synthesising XRT light curves and fitting both single and broken power laws, and comparing the relative goodness of each fit via Monte Carlo analysis. Even with the well sampled light curves of the \emph{Swift} era, these breaks may be left misidentified, hence caution is required when making definite statements on the absence of achromatic breaks. | \label{section:introduction} The afterglow emission of Gamma-Ray Bursts (GRBs) is well described by the blast wave, or fireball, model \citep{rees1992:mnras258,meszaros1998:apj499}. This model details the temporal and spectral behaviour of the emission that is created by external shocks when a collimated ultra-relativistic jet ploughs into the circumburst medium, driving a blast wave ahead of it. The level of collimation, or jet opening angle, has important implications for the energetics of the underlying physical process and the possible use of GRBs as standard candles. The signature of the collimation, according to simple analytical models, is an achromatic temporal steepening or `jet break' at $\sim 1$ day in an otherwise decaying, power law light curve; from the time of this break, the jet opening angle can be estimated \citep{rhoads1997:ApJ487}. Since the launch of the \emph{Swift} satellite \citep{gehrels2004:ApJ611}, this standard picture of afterglows has been called into question by the lack of observed achromatic temporal breaks, up to weeks in a few bursts (e.g., \citealt{panaitescu2006:MNRAS369,burrows2007:astro.ph2633}). In some afterglows, a break is unobserved in both the X-ray and optical light curves, while in other bursts a break is observed in one regime but not the other (e.g., \citealt{liang2007:arXiv0708}). In the \emph{BeppoSAX} era, most well sampled light curves were in the optical regime, while in the \emph{Swift} era most well sampled light curves are in the X-ray regime. Our expectations of the observable signature of a jet break are hence based on the breaks observed pre-\emph{Swift}, predominately by optical telescopes, and the models which explained them, notably those of \cite{rhoads1997:ApJ487,rhoads1999:ApJ525} and \cite{sari1999:ApJ519}. It is not clear that the breaks will be identical in both the X-ray and optical regimes. In the cases of GRB\,990510 and GRB\,060206, both have clear breaks in the optical but only marginal breaks in the X-ray (\citealt{kuulkers2000:ApJ538,curran2007:MNRAS381} respectively). Regardless of this, GRB\,990510 is taken as a prototypical achromatic jet break, while the achromatic nature of the GRB\,060206 break is only evident when supported by the broad-band spectral indices. In this paper we address the issue of X-ray breaks which are possibly `hidden' and hence the light curve misinterpreted as being a single power law. We do so by synthesising X-ray light curves and fitting both single and broken power laws, and comparing goodness of each fit via the F-test. In \S\ref{section:method} we introduce our method while in \S\ref{section:results} we present the results of our Monte Carlo analysis. In \S\ref{section:discussion} we discuss the implications of these results in the overall context of GRB observations and we summarise our findings in \S\ref{section:conclusion}. | \label{section:conclusion} As underlying smoothly broken power laws may be \emph{hidden} as single power laws, in XRT light curves, we should exercise caution in ruling out breaks based solely on a comparison of the nominal fitted slopes. We hence need to be cautious in implying chromatic breaks from optical and X-ray light curves where an X-ray break is not ruled out with a high degree of certainty. Multi-wavelength temporal and spectral data are required to make a confident statement on the absence or presence of an achromatic break. There may be a bias towards detecting breaks from bursts with a constant density circumburst medium, as these are most obvious and easily detectable. We have shown that the fitted temporal slopes of smoothly broken power laws show significant variation about central values and that, especially in the case of a smooth break, the underlying values are difficult to extract via a fit. More accurate and wavelength specific descriptions of breaks, likely involving numerical simulations of the jet dynamics, are necessary to better understand the observable signature of breaks and to clarify whether the breaks could vary somewhat between wavebands. | 7 | 10 | 0710.5285 |
0710 | 0710.2693_arXiv.txt | {Recently, we have developed a method useful for mapping large-scale horizontal velocity fields in the solar photosphere. The method was developed, tuned and calibrated using the synthetic data. Now, we applied the method to the series of Michelson Doppler Imager (MDI) dopplergrams covering almost one solar cycle in order to get the information about the long-term behaviour of surface flows. We have found that our method clearly reproduces the widely accepted properties of mean flow field components, such as torsional oscillations and a pattern of meridional circulation. We also performed a periodic analysis, however due to the data series length and large gaps we did not detect any significant periods. The relation between the magnetic activity influencing the mean zonal motion is studied. We found an evidence that the emergence of compact magnetic regions locally accelerates the rotation of supergranular pattern in their vicinity and that the presence of magnetic fields generally decelerates the rotation in the equatorial region. Our results show that active regions in the equatorial region emerge exhibiting a constant velocity (faster by $60 \pm 9$~\mps{} than Carrington rate) suggesting that they emerge from the base of the surface radial shear at $0.95\ R_\odot$, disconnect from their magnetic roots, and slow down during their evolution. | The largest scale velocity fields in the solar photosphere consist of rotation profile and a meridional flow pattern. Basically, the differential rotation is described as an integral of the zonal component $v_\varphi$ of the studied flow field. The integrated flow field may be obtained using spectroscopic method, using tracer-type measurements, or using helioseismic inversions. The latest case allows to measure the solar rotation not only as a function of heliographic latitude, but also as a function of depth. From the helioseismic inversion we know that throughout the convective envelope, the rotation rate decreases monotonically toward the poles by about 30~\%. Angular velocity contours at mid-latitudes are nearly radial. Near the surface at the top of the convection zone there is a layer of a large radial shear in the angular velocity. At low and mid-latitudes there is an increase in the rotation rate immediately below the photosphere which persists down to $r \sim 0.95~R_\odot$. The angular velocity variation across this layer is roughly 3~\% of the mean rotation rate and according to the helioseismic analysis of \cite{2002SoPh..205..211C} the angular velocity $\omega$ decreases within this layer approximately as $r^{-1}$, where $r$ is a radial coordinate. At higher latitudes, the situation is less clear. For the overview of solar differential rotation measurements see \cite{1985SoPh..100..141S} or a more recent review by \cite{2000SoPh..191...47B}. The \emph{torsional oscillations}, in which narrow bands of faster than average rotation, interpreted as zonal flows, migrate towards the solar equator during the sunspot cycle, were discovered by \cite{1980ApJ...239L..33H}. Later research \citep{2001ApJ...559L..67A} found that there exist two different branches of torsional oscillations. At latitudes below about 40\,$^\circ$, the bands propagate equatorward, but at higher latitudes they propagate poleward. The low-latitude bands are about 15\,$^\circ$ wide in latitude. The flows were studied in surface Doppler measurements \citep{2001ApJ...560..466U}, and also using local helioseismology \citep{1997ApJ...482L.207K}. The surface pattern of torsional oscillations penetrate deeply in the convection zone, possibly to its base, as suggested by \cite{2002Sci...296..101V}. The amplitude of the angular velocity variation is about 2--5~nHz, which is roughly 1~\% of the mean rotational rate (5--10~\mps). The direct comparison between different techniques inferring the surface zonal flow pattern \citep{2006SoPh..235....1H} showed that the results are sufficiently coherent. The surface magnetic activity corresponds well with the torsional oscillation pattern -- the magnetic activity belt tends to lie on the poleward side of the faster-rotating low-latitude bands. The magnetic activity migrates towards the equator with the low-latitude bands of the torsional oscillations as the sunspot cycle progresses \citep{2004ApJ...603..776Z}. Some studies \citep[e.\,g.][]{2002ApJ...575L..47B} suggest that meridional flows may diverge out from the activity belts, with the equatorward and poleward flows well correlated with the faster and slower bands of torsional oscillations. In a recent theoretical study by \cite{2007ApJ...655..651R} it was suggested that the poleward-propagating high-latitude branch of the torsional oscillations can be explained as a response of the coupled differential rotation/meridional flow system to periodic forcing in midlatitudes of either mechanical (Lorentz force) or thermal nature. The equatorward-propagating low-latitude branch is most likely not a consequence of the mechanical forcing alone, but rather of thermal origin. The axisymmetric flow in the meridional plane is generally known as the \emph{meridional circulation}. The meridional circulation in the solar envelope is much weaker than the differential rotation, making it relatively difficult to measure. Two principal methods are widely used to measure the meridional flow: feature tracking and direct Doppler measurement. There are several difficulties complicating the measurements of the meridional flow using tracers. Sunspots and filaments do not provide sufficient temporal and spatial resolution for such studies. Sunspots also cover just low latitudinal belts and do not provide any informations about the flow in higher latitudes. Doppler measurements do not suffer from the problem associated with the tracer-type measurements, however they introduce another type of noisy phenomena. It is difficult to separate the meridional flow signal from the variation of the Doppler velocity from the disc centre to the limb. Using different techniques, the parameters of the meridional flow show large discrepancies. It is generally assumed that the solar meridional flow in the close subphotospherical layers is directed poleward with one cell per hemisphere. Such flow is also produced by early global hydrodynamical simulation such as \cite{1982ApJ...256..316G}. As reviewed by \cite{1996ApJ...460.1027H}, the surface or near sub-surface velocities of the meridional flow are generally in the range 1--100~\mps, the most often measured values lie within the range of 10--20~\mps. The flow has often a complex latitudinal structure with both poleward and equatorward flows, multiple cells, and large asymmetries with respect to the equator. \cite{2004ApJ...603..776Z} used the time-distance helioseismology to infer the properties of the meridional flow in years 1996--2002. They found the meridional flows of an order of 20~\mps, which remained poleward during the whole period of observations. In addition to the poleward meridional flows observed at the solar minimum, extra meridional circulation cells of flows converging toward the activity belts are found in both hemispheres, which may imply plasma downdrafts in the activity belts. These converging flow cells migrate toward the solar equator together with the activity belts as the solar cycle evolves. \cite{2002ApJ...575L..47B} measured the meridional flow (and torsional oscillations) using the time-distance helioseismology and found the residual meridional flow showing divergent flow patterns around the solar activity belts below a depth of 18~Mm. The most complete maps of the torsional oscillations and the meridional flow available at present have been constructed on the basis of Mt.~Wilson daily magnetograms \citep[see][]{ulrich90}. The measurements cover more than 20 years (since 1986) and the results obtained using this very homogenous material agree well with the properties described above. The modern dynamo flux-transport models use the meridional flow and the differential rotation as the observational input. In the models by Dikpati et al. (\citeauthor{2006ApJ...638..564D} \citeyear{2006ApJ...638..564D} or \citeauthor{2006ApJ...649..498D} \citeyear{2006ApJ...649..498D}) the return meridional flow at the base of the convection zone is calculated from the continuity equation. They found the turnover time of the single meridional cell of 17--21~years. The meridional flow is assumed to be essential for the dynamo action, global magnetic field reversal and forecast of the future solar cycles. There are known many relations of the differential rotation profile to the phase of the progressing solar cycle -- see e.\,g. \cite{2003SoPh..212...23J} or \cite{2005ApJ...626..579J} -- showing for example different properties of the differential rotation profile in the odd and even solar cycles. The rotation of the sunspots in relation to their morphological type was studied e.\,g. by \cite{1986AA...155...87B} who found that more evolved types of sunspots (E, F, G and H type) rotate slower than less evolved types. \cite{2004SoPh..221..225R} investigated Greenwich Photoheliographic Results for the years 1874--1976 and found a clear evidence for the deceleration of the sunspots in the photosphere with their evolution. \cite{2002aprm.conf..427H} found that the leading part of a complex sunspot group rotate about 3~\% faster than the following part. The dependence of the rotation of sunspot on their size and position in the bipolar region was investigated by \cite{1994SoPh..151..213D}. They explained the observed behaviour through a subtle interplay between the forces of magnetic buoyancy and drag, coupled with the role of the Coriolis force acting on rising flux tubes. This dynamics of rising flux tubes also explains the faster rotation of smaller sunspots. In average, sunspots rotate about 5~\% faster than the surrounding plasma. In the theoretical study \citep{2004SoPh..220..333B} based on 3-D numerical simulations of compressible convection under the influence of rotation and magnetic fields in spherical shells, the author stated that in the presence of magnetic field the Maxwell stresses may oppose the Reynolds stresses and therefore the angular momentum is propagated more to the poles than without the presence of magnetic fields. As a consequence the rotation profile is more differential in the periods of lowered magnetic activity and it leads to the increase of the rotation in low latitudes. This behaviour was observed in many studies, e.\,g. \cite{1990ApJ...357..271H}. The subject of this work is a verification of the performance of the method described in \cite{svanda06} (hereafter Paper~I) on the real data and the investigation of long-term properties of the flows at largest scales obtained with this method. We shall also discuss the influence of magnetic fields on the measured zonal flow in the equatorial region. This topic will be studied more in detail in one of the next papers in the series. \begin{figure*} \centering \resizebox{\textwidth}{!}{\includegraphics{fig01.ps}} \caption{Mean meridional flow in time and heliographic latitude. It can be clearly seen that for almost all the processed measurements a simple model of one meridional cell per hemisphere would be sufficient. However, some local corruptions of this simple idea can be noticed on both hemispheres.} \label{fig:meridional} \end{figure*} \begin{figure*} \centering \resizebox{\textwidth}{!}{\includegraphics{fig02.ps}} \caption{Torsional oscillations. The residua of the mean zonal flow with respect to its parabolic fit displayed in time and heliographic latitude. In the period of weak magnetic activity the pattern of belts propagating towards the equator is very clear. In the periods of stronger magnetic activity the flow field is influenced by the local motions in active regions and therefore the pattern of torsional oscillations is not clearly seen.} \label{fig:torsional} \end{figure*} | We have verified that the method developed and tested using the synthetic data (Paper~I) is suitable for application to real data obtained by the MDI on-board SoHO and maybe also to the data that will be produced by its successor Helioseismic Michelson Imager (HMI) on-board the Solar Dynamic Observatory (SDO). HMI will have a greater resolution and will cover larger time span than two months each year. We verified that the long-term evolution of the horizontal velocity fields measured using our method is in agreement with generally accepted properties. During the periodic analysis of the equatorial area we found two suspicious periods in the real data, which are not present in the control data set containing the inclination of the solar axis towards the observer, the quantity that can bias systematically and periodically the results by a few \mps. The periods of 1.8~year and 4.7~years need to be confirmed using a more homogenous data set. We also found that the presence of the local magnetic field generally speeds-up the region occupied by the magnetic field. However, we cannot conclude that there exists a dependence of this behaviour for different types of sunspots. We can generally say that the more evolved types of active regions rotate slower than the young ones, however the variance of the typical rotation rate is much larger than the differences between the rates for each type. We have found that the distribution of active regions rotation is bimodal. The faster-rotating cases correspond to new and growing active regions. Their almost constant rotation speed suggests that they emerge from the base of the surface radial shear at $0.95\ R_\odot$. The decaying and recurrent regions rotate slower with a wider scatter in their velocities. This behaviour suggests that during the sunspot evolution, sunspots loose the connection to their magnetic roots. Both regimes alternate with a period of approximately one Carrington rotation in years 2001 and 2002, which suggests that new active regions emerge in groups and may have a linked evolution. | 7 | 10 | 0710.2693 |
0710 | 0710.2370_arXiv.txt | We present an investigation of quasar colour-redshift parameter space in order to search for radio-quiet red quasars and to test the ability of a variant of the KX quasar selection method to detect quasars over a full range of colour without bias. This is achieved by combining IRIS2 imaging with the complete Fornax Cluster Spectroscopic Survey to probe parameter space unavailable to other surveys. We construct a new sample of 69 quasars with measured $b_J - K$ colours. We show that the colour distribution of these quasars is significantly different from that of the Large Bright Quasar Survey's quasars at a 99.9\% confidence level. We find 11 of our sample of 69 quasars have signifcantly red colours ($b_J - K \geq 3.5$) and from this, we estimate the red quasar fraction of the $K \leq 18.4$ quasar population to be 31\%, and robustly constrain it to be at least 22\%. We show that the KX method variant used here is more effective than the UVX selection method, and has less colour bias than optical colour-colour selection methods. | \label{introduction} \citet{1995Natur.375..469W} compared the $B - K$ colours of Parkes Half-Jansky Flat-Spectrum Sample (PHJFS) \citep{1997MNRAS.284...85D} quasars to those of Large Bright Quasar Survey (LBQS) \citep{1995AJ....109.1498H} quasars. The PHJFS quasars were found to have a significantly broader and redder distribution of $B - K$ colours than LBQS quasars. This disparity should not exist according to the unified model of Active Galactic Nuclei (AGN) galaxies \citep{1993ARA&A..31..473A}. \citet{1995Natur.375..469W} concluded that there exists a population of red quasars that at that time had not been discovered. \citet{1995Natur.375..469W} investigated this possibility using a simple model of LBQS selection effects, and demonstrated that the LBQS was biased against the detection of red quasars. The bias was caused by the blue magnitude limit of the LBQS, because in a blue-flux limited survey, red quasars need to be intrinsically brighter than blue quasars to be detected \citep{1995Natur.375..469W}. The bias causes fewer red quasar detections, because the luminosity function of quasars is steep at the bright end \citep{1991MNRAS.251..482B}. Other surveys with a blue-flux limit are similarly biassed, independent of the quasar selection method. Using their simple model, \citet{1995Natur.375..469W} estimated that the red quasars missed by previous surveys constitute as much as 80\% of all quasars. Red quasars are defined by a colour criterion. The best colour to use for this criterion is that made from filters separated by the largest wavelength range possible, because anomalously redder quasar colours stand out more. For the datasets used in this paper, the colour covering the largest wavelength range is $U - K$. We used the next largest colour $b_J - K$, as far more objects, $\sim$ 3 times as many, have a measured $b_J - K$ colour compared to the number with a measured $U - K$ colour. The additional advantage of having used $b_J - K$ is that it is possible to compare the colours of the quasar sample presented here with the $B - K$ colour used in \citet{2004ApJ...607...60G} and the $g - K$ colour of \citet{2004AJ....128.1112H}. Throughout this paper we define a red quasar by the criterion, $b_J - K \geq$ 3.5 \citep{2002ApJ...564..133G}, because when describing the optical-infrared region of quasar spectra with a power-law, $S(\lambda) \propto \lambda^{\alpha}$, this colour corresponds to an effective optical-infrared spectral index of $\alpha \geq 1$. The existence of red quasars, other than in radio-selected quasar samples, over most of the $b_J - K$ colour range of the PHJFS quasars, satisfying the above criterion, has already been demonstrated \citep{2004AJ....128.1112H, 2003AJ....126.1131R, 2004ApJ...607...60G}. The \citet{2004ApJ...607...60G} red quasar sample contains the reddest quasars found, with $B - K$ colours of $5 \leq B - K \leq 8$. \citet{2004AJ....128.1112H, 2003AJ....126.1131R} found red quasars in the Sloan Digital Sky Survey (SDSS)\citep{2000AJ....120.1579Y}, and \citet{2004ApJ...607...60G} found them in a combination of the 2-Micron All Sky Survey (2MASS) \citep{2006AJ....131.1163S} and an Automatic Plate Measuring CATalogue (APMCAT) of POSS/UKST sky survey plates \citep{APMCAT}. Despite confirming the existence of red quasars, current observations have not explained the mechanism or mechanisms responsible for their existence. Dust obscuration by dust at the quasar redshift, in either the host galaxy or the quasar accretion disk \citep{1995Natur.375..469W}, was the first mechanism proposed for quasar reddening. \citet{1995Natur.375..469W} proposed this mechanism as a correlation between $B - K$ colour and redshift was not observed. The result of reddening a composite optically selected quasar spectral energy distribution (SED) with dust at the quasar redshift, was found to be consistent with the dust mechanism, because the reddening resulted in the SED of a red PHJFS quasar \citep{1995Natur.375..469W}. As an alternative, \citet{1998MNRAS.295..451B} suggested that the $B - K$ spread in the \citet{1995Natur.375..469W} sample is partly or entirely caused by host galaxy starlight. Quasars with flat-radio-spectra are typically hosted by later galaxy types \citep{1998MNRAS.295..451B}, and the old stellar population of late type galaxies is a stronger emitter in K than B, reddening quasars \citep{1998MNRAS.295..451B}. An analysis in \citet{1998MNRAS.301..975M} of the host galaxy contribution to the red $B - K$ colours of the PHJFS quasars in \citet{1995Natur.375..469W} was inconclusive, but suggested host galaxy contribution was insufficient to cause the red $B - K$ colours observed. Further work in \citet{2006MNRAS.367..717M} confirmed that host galaxy starlight can definitely redden quasars, particularly low luminosity and low redshift resolved quasars. In accordance with \citet{2006MNRAS.367..717M}, in this paper we use the criteria that luminous unresolved quasars, which are not located at low redshift, are unlikely to be reddened by host galaxy contribution. Alternatively, \citet{1996Natur.379..304S} proposed that a synchrotron component of quasar emission created a K-band excess, reddening quasars. An analysis of 157 PHJFS quasars and 12 LBQS quasars found the shape of the SEDs was consistent with both line-of-sight dust and synchrotron emission mechanisms \citep{2000PASA...17...56F}. In \citet{2001AJ....121.2843B} it was suggested that because any synchrotron component to quasar emission is weaker in radio-quiet quasars than in radio-loud quasars, any synchrotron emission will only contribute to a small fraction of any reddening. Based on the quasar sample contained therein, \citet{2001AJ....121.2843B} suggested that any radio-quiet quasar redder than $b_J - K \geq$ 3.7 is too red to be the result of synchrotron emission, and is most likely the result of dust obscuration. Throughout this paper we use the criterion that any red quasar that is also radio-quiet, is unlikely to be reddened by synchrotron emission. Finally, \citet{1997MNRAS.288..138M} suggested that a small subset of red quasars was caused by the lensing of a quasar by a dusty galaxy. Not only does the dust in the lens redden the quasar but the dusty galaxy potentially magnifies the host galaxy starlight contribution \citep{2002ApJ...564..133G}. This last mechanism was proposed after \citet{1997MNRAS.288..138M} observed a similar phenomenon to \citet{1995Natur.375..469W}; with radio-selected lenses having redder colours than optically selected lenses. Whichever mechanism creates red quasars, they are possibly an evolutionary stage. After observing low-ionisation Broad Absorption Line (BAL) quasars, \citet{2002AJ....123.2925L} suggested a model of quasar evolution where a quasar occupies a thick dusty torus after forming, eventually the torus dissipates leaving behind an optically blue quasar. Alternately, quasars may emerge from dusty starburst galaxies, residing in the dust stirred up by mergers \citep{2005AAS...20717401U}. Either evolutionary sequence would naturally create a situation where quasars are reddened by dust at the quasar redshift. A third evolutionary path has been proposed where red quasars are an evolutionary link between Ultra-Luminous Infrared Galaxies (ULIRGs) and optically selected quasars \citep{2006Ap&SS.302...17C}. Determining the fraction of red quasars and any dependence on redshift, luminosity, radio flux, etc. will indicate if red quasars are an evolutionary stage, and which evolutionary path red quasars belong to. Depending on the mechanism, red quasars might explain the quasar contribution to the X-Ray Background (XRB). Using a theoretical model, \citet{1994MNRAS.270L..17M} showed that a ratio of two to three times as many dust obscured to unobscured active galactic nuclei (AGN) galaxies can adequately account for the existence of the XRB. The model also accounts for the observed spectrum and the source counts in the soft and hard X-ray bands. The latest theoretical models of the XRB require a smaller ratio, lying somewhere between 0.6 to 1.5 \citep{2007A&A...463...79G}. If it can be shown that red quasars are the result of dust obscuration, and that the red quasar fraction is within 37.5\% to 60\%, then red quasars are the most probable source of the quasar contribution to the XRB. The existence of red quasars has various implications, the main one is that it can no longer be assumed that existing quasar samples are representative of the entire quasar population. The statistics of a sample reflects the properties of the sub-set of a population from which the sample is drawn, previous quasar samples, such as the LBQS, are biassed to the inclusion of blue quasars; therefore, the statistics of such samples predominantly reflects the blue quasar population. Use of existing quasar datasets requiring a sample that is representative of the entire quasar population is therefore subject to review. Direct examples of this are that the luminosity function of quasars and the change in number density with redshift need to be re-evaluated while including red quasars. In order to learn more about red quasars, we need to spectroscopically observe quasars over their entire colour range. To achieve this, future surveys will have to select quasar candidates for spectroscopic follow-up using a technique that is not biased against either red or blue quasars. One candidate for a suitable selection technique is the KX method, developed in \citet{2000MNRAS.312..827W}. The KX method utilises the power-law nature of quasar SEDs at long wavelengths combined with morphological classification. Stars experience a turnover in their SED in H-band while quasars follow a power-law, allowing discrimination between stars and quasars. The stellar morphology of quasars discriminates between quasars and galaxies. Observationally, this requires constructing a colour-colour plot of all objects with a stellar morphology, using an optical colour and a colour that straddles the H-band. In this paper we explore regions of the quasar colour-redshift space unavailable in previous work, and test a KX method variant. We aim to find the boundaries of the quasar colour-redshift parameter space that is occupied. In doing so, we will improve the estimate of the red quasar fraction and possibly shed light on the mechanism that creates red quasars. In testing a KX method variant, we intend to verify if it is a viable method of selecting quasars, without bias, from the entire quasar colour-redshift parameter space. We intend to apply a KX method variant to our quasar sample, and then contrast it with established methods of selecting potential quasars for spectroscopic follow-up. To achieve these goals we use a sub-sample of the quasars identified in the Fornax Cluster Spectroscopic Survey (FCSS) \citep{2000A&A...355..900D}, a spectroscopic survey of all extended objects to $b_J = 19.8$ and point sources to $b_J = 21.5$. The sub-sample consists of all quasars that match to an object in Ks-band imaging to a depth of K = 18.4, obtained using the Infrared Imager and Spectrograph 2 (IRIS2) instrument on the 3.92-m Anglo-Australian Telescope (AAT). The Ks-band imaging is deeper than 2MASS allowing us to explore parts of the $b_J - K$ colour space unavailable in \citet{2004ApJ...607...60G}, which was limited to $b_J - K \geq$ 5 by the K = 14.5 magnitude limit of 2MASS. Furthermore, the quasar sample used covers a larger range of redshift than was examined in \citet{2004AJ....128.1112H,2003AJ....126.1131R}, which were limited to z $\leq$ 2.2. Because the FCSS used no selection criteria, our quasar sample is bias free, and ideal for testing the ability of the KX method to select quasars from a K magnitude limited survey. Additionally, most of the FCSS objects have a U magnitude as well as a $b_J$ magnitude, allowing the KX method to be directly contrasted with traditional methods for optically selecting quasars. The structure of this paper is as follows. In Section \ref{data} the datasets used in this paper are described. Section \ref{redquasars} details the detection of red quasars in the quasar sample and an analysis of the quasar sample. This analysis involves a comparison of the $b_J - K$ colours of the sample quasars to LBQS quasars and a determination of the effect of our selection on the quasar sample constructed here. Section \ref{redquasars} concludes with the calculation of the fraction of red quasars. The effectiveness of the KX method is assessed in Section \ref{testKX} and following this are our conclusions in Section \ref{conclusion}. \section[]{Photometric and Spectroscopic Data} \label{data} The sample of quasars used here and their photometric information were obtained by combining three datasets. Astrometry and optical photometry were obtained from the Automatic Plate Measuring CATalogue (APMCAT) of POSS/UKST sky survey plates \citep{APMCAT}. IRIS2 imaging was used to obtain the K-band photometry used here. The necessary spectroscopic identifications were provided by the Fornax Cluster Spectroscopic Survey (FCSS) catalogue of \citet{2000A&A...355..900D}. The datasets were combined by independently matching the astrometry of K-band catalogue objects and FCSS objects to that of APMCAT catalogue objects. In the rest of this section, we describe the datasets in more detail. The FCSS catalogue is a blind spectroscopic survey of Fornax objects in the APMCAT catalogue. The FCSS is an ideal source of spectroscopic identifications because the lack of target selection precludes biasing. In the FCSS, spectroscopy was carried out on all extended sources to $b_J \leq 19.8$ and all point sources to $b_J \leq 21.5$. Not every source was spectroscopically observed, and the fraction of observed sources is referred to as the spectroscopic completeness. The FCSS obtained a redshift and spectroscopic identification for more than 90\% of the objects observed, referred to as the redshift completeness. In most cases where a spectroscopic identification could not be obtained, neither was a redshift because of the poor spectrum quality. Because the redshift completeness was consistently high for the entire FCSS, we have combined the redshift completeness with the spectroscopic completeness, and throughout the paper the term `completeness' refers to this combination. The completeness of the FCSS is 96\% for $b_J < 20.5$, 82\% for $20.5 \leq b_J < \leq 21$ and 36\% for $21 < b_J \leq 21.5$ sources. We account for the completeness by appropriately weighting quasars during our analysis in later sections. IRIS2 was used in December 2001 to obtain Ks-band imaging using 60s exposures. The raw data were processed using the Observatory Reduction and Acquisition Control project - Data Reduction pipeline (ORAC-DR) outlined in \citet{1999ASPC..172...11E,1999ASPC..172..171J}. The reduced IRIS2 imaging dataset was then analysed with SExtractor \citep{1996A&AS..117..393B} to obtain a photometric catalogue of corrected isophotal magnitudes. Corrected isophotal magnitudes were used instead of PSF or small aperture magntiudes, because PSF and small aperture magnitudes minimise any possible extended contribution to the flux of sources in the IRIS2 imaging. The Ks-band catalogue magnitude zero-point was then calibrated, by matching point sources in the Ks-band catalogue to the 2MASS point source catalogue. Calibrating also facilitated the conversion from Ks magnitudes to the K-band magnitudes used throughout this paper. After calibration, the magnitude limit of the K-band photometry was found to be 18.4 magnitudes. In the FCSS, quasars were spectroscopically identified as objects with broad permitted lines with a measured full width half-maximum $>$ 1100 km $s^{-1}$ \citep{2001MNRAS.324..343M}. After matching the IRIS2 imaging catalogues and the FCSS quasar identifications to the APMCAT object positions using blind matching, to a radius of 10"; a catalogue of all the necessary photometry was created, and all quasars within both the IRIS2 imaging K magnitude limit, and the FCSS spectroscopic identifications $b_J$ magnitude limit were identified. This resulted in a sample of 69 spectroscopically identified quasars with U, $b_J$ and K magnitudes. All 69 of these quasars are APMCAT point sources; therefore, our quasar sample is limited to $K \leq 18.4$ and $b_J \leq 21.5$. Of these 69 quasars, 62 have an R magnitude as required by our KX method variant. The missing R magnitudes are caused by the different magnitude limits in the $b_J$ and R bands of the APMCAT catalogue, 22.5 and 21 magnitudes respectively. | \label{conclusion} In this paper, we constructed a sample of 69 quasars with measured $b_J - K$ colours, using a combination of APMCAT, IRIS2 imaging and FCSS spectroscopic identifications. 11 of these quasars are red, satisfying $ b_J - K \geq$ 3.5, and all of these red quasars are radio-quiet according to the All Sky Optical Catalogue of Radio/X-Ray Sources. In accordance with \citet{2001AJ....121.2843B} and \citet{2006MNRAS.367..717M}, as the 11 red quasars found here are unresolved, luminous and radio-quiet sources predominantly located at $z > 1$, this strongly indicates that the main cause of the red $b_J - K$ colours is dust at the quasar redshift. Comparing our quasar sample to LBQS quasars, we found that the two $b_J - K$ colour distributions are different at the 99.9\% confidence level, with our sample having a significantly broader, redder distribution. Analysis of the uncertainties in the $b_J - K$ colours demonstrated that neither our red quasar detections or our comparison to the LBQS quasars is affected by them. A second analysis, on the effects of our datasets magnitude and completeness limits, revealed that they limited our measured $b_J - K$ distribution to $ b_J - K <$ 5. As red quasars are observed up to this colour limit, we concluded that the true $b_J - K$ distribution of quasars is even broader and redder than observed here. From the observed $b_J - K$ colour distribution, we robustly constrained the red quasar fraction of the $K \leq 18.4$ quasar population to be greater than 22\%, and from a model estimated it to be 31\%. Using the quasar sample constructed here, the viability of a KX method variant was tested. Using a $b_J - R$ vs $R - K$ plot, the KX method variant was capable of separating the quasar sample out from the other objects in the APMCAT catalogue. Comparing the KX method variant to the UVX and 2QZ multicolour methods, the KX method variant was found to be as effective in the number of quasars it selected; however, it was superior at selecting quasars independent of colour. For those reasons, the KX method variant used here is an ideal technique for future large surveys to use in selecting potential quasars, whether red or blue, for spectroscopic follow-up. | 7 | 10 | 0710.2370 |
0710 | 0710.2000_arXiv.txt | We present a comprehensive spectroscopic imaging survey of the distribution and kinematics of atomic hydrogen (HI) in 16 nearby spiral galaxies hosting low luminosity AGN, observed with high spectral and spatial resolution (resolution: $\sim$20\arcsec, $\sim$5~km~s$^{-1}$) using the NRAO Very Large Array (VLA). The sample contains a range of nuclear types, ranging from Seyfert to star-forming nuclei and was originally selected for the NUclei of GAlaxies project (NUGA) - a spectrally and spatially resolved interferometric survey of gas dynamics in nearby galaxies designed to identify the fueling mechanisms of AGN and the relation to host galaxy evolution. Here we investigate the relationship between the HI properties of these galaxies, their environment, their stellar distribution and their AGN type. The large-scale HI morphology of each galaxy is classified as ringed, spiral, or centrally concentrated; comparison of the resulting morphological classification with AGN type reveals that ring structures are significantly more common in LINER than in Seyfert host galaxies, suggesting a time evolution of the AGN activity together with the redistribution of the neutral gas. Dynamically disturbed HI disks are also more prevalent in LINER host galaxies than in Seyfert host galaxies. While several galaxies are surrounded by companions (some with associated HI emission), there is no correlation between the presence of companions and the AGN type (Seyfert/LINER). | \label{sec:intro} Active Galactic Nuclei (AGN) represent some of the most extreme conditions and include the most powerful individual objects to be found throughout the Universe. Nowadays, the phenomena of nuclear activity is generally understood to be the result of accretion of material onto a SuperMassive Black Hole (SMBH); very likely through the infall of gas from its host galaxy \citep[e.g.][]{Rees}. Although SMBHs exist in most galaxies, only a small fraction of all galaxies exhibit nuclear activity \citep{Huchra,Ho97b, Miller}. The majority of AGNs show low-ionization narrow emission-line regions (LINERs), which have luminosities lower than Seyfert galaxies and quasars. Originally, LINERs and Seyferts were considered to form a continuous distribution in the usual line-ratio classification diagrams. Recent investigations of the host properties of emission-line galaxies from the Sloan Digital Sky Survey \citep{Kew06} confirmed that their nuclear activities are due to accretion by a SMBH. But interestingly, LINERs and Seyferts are also clearly separable in emission line ratio diagrams \citep{Gro06}, and that these two classes have distinct host properties \citep{Gro06}, implying that their distinction is in fact not simply an arbitrary division. An open question is therefore what mechanisms give rise to the different AGN types observed. \par A coherent picture of AGN feeding, e.g. how to remove the angular momentum from host galaxy gas so that it can reach the central parsec, is still missing. Therefore it is expected that a hierarchy of mechanisms might combine to transport gas from large kpc scales down to the inner pc scales \citep{Shl90, Com03, Jog04}. Additionally, trigger mechanisms driven by the galactic environment, like minor mergers and tidal forces, might play a role as well \citep{Barnes}. A larger gas fraction for Seyferts than for normal galaxies was suggested by \cite{Hunt99a} and a prevalence of stellar rings (RC3 classification) in active galaxies (Seyfert and LINERs) compared to non-active galaxies was suggested \citep{Hunt99b}. NICMOS imaging of the centers of 250 nearby galaxies \citep{Hunt04} revealed systematic morphological differences between HII/starburst (most disturbed), Seyfert (intermediate disturbed), LINER and normal galaxies (most regular). The study of the circumnuclear regions of 24 Seyfert2 galaxies using HST images \citep{Mar99} revealed a dominance of nuclear spirals, suggesting that spiral dust lanes may be responsible for feeding gas to the central engines. \par Theoretical simulations have also made progress in addressing the nature of nuclear fueling with respect to different types of gravitational instabilities and their feeding efficiency: nested bars \citep[e.g.][]{Shl89, Friedli, Mac00, Eng04}, gaseous spiral density waves \citep[e.g.][]{Eng00, Mac02, Mac04a, Mac04b}, m = 1 perturbations \citep[e.g.][]{Shu, Jun96, Gar00} and nuclear warps \citep{Sch00} have all been suggested as possible transport mechanisms. \par In order to distinguish models for nuclear fueling, observations of neutral gas with high angular and velocity resolution are required. Therefore, the IRAM key project NUclei of GAlaxies \citep[ NUGA; see][]{Gar03} was established - a spectroscopic imaging survey of gas in the centers of nearby low luminosity AGN. As most of the gas in galaxy nuclei is in the molecular phase, the survey used millimeter CO lines to conduct a detailed mapping of molecular gas dynamics at {\em high-resolution} (0.5$\arcsec$) in the central kiloparsec of AGN hosts. The CO-NUGA survey (using the IRAM Plateau de Bure Interferometer) reveals a wealth of nuclear gas distribution and kinematics which are studied in detail: (a) m=1 gravitational instabilities (one-arm spirals and lopsided disks; NGC~4826: \cite{Gar03b}), (b) m=2 instabilities (two-arm spiral wave) expected to form in stellar bar potentials show only small amounts of molecular gas coincident with the AGN (e.g. NGC~4569: \cite{Boo07}, NGC~4579: \cite{Gar05} NGC~6951: Schinnerer et al., in prep.), (c) stochastic spirals that are related to non self-gravitating instabilities, and rings \citep[NGC~7217:][]{Com04}, and (d) large scale warps that might extend into the central kiloparsec \citep[NGC~3718:][]{Kri05}. Besides the CO studies also extended radio continuum components resembling jet emission from AGN were detected \citep{Kri07a} and a molecular gas disk/torus of dense gas was observed in the HCN line emission \citep[NGC~6951:][]{Kri07b}. Additionally, Garc{\'{\i}}a-Burillo et al. (\citeyear{Gar05}) derived in a pilot study the gas inflow rates via measurements of the gravitational torque onto the molecular gas in the central kiloparsec for 4 galaxies. \par To complement these CO data and provide a more complete census of the ISM and gas flows from the outskirts of the galaxy disks to the very center, the HI-NUGA project was initiated. HI-NUGA provides observations of the HI gas distributions and kinematics for 16 galaxies (including archival data for 2 galaxies) of the NUGA sample using NRAO's Very Large Array (VLA). Determining the HI distribution and kinematics is necessary to establish whether characteristics of the nuclear dynamics, and hence the nuclear activity, depend on overall host galaxy properties. Possible dependencies might also exist between nuclear modes and large scale drivers (i.e. tidal forces exerted by companions), or the overall content/distribution of the atomic gas. Since the HI gas extends in most disk galaxies much further than the optical disk, it is more loosely bound in the outer region because of the decrease of the gravitational potential. Hence, it provides an ideal tool to identify interactions and tidal features \citep[e.g.,][]{Sim87, Mun95}. Furthermore, due to its dissipative nature the gas is more sensitive to dynamical disturbances, both internal ones such as non-axisymmetric potentials or external ones (e.g. tidal interactions). Thus, one could expect, for example, to find a prevalence of Seyfert activity with asymmetries in the gas distribution and kinematics.\par In this paper we present for 16 nearby active galaxies data obtained in the 21 cm emission line of neutral hydrogen using the VLA. The relation between HI properties, such as the large scale environment (HI/optical companions, HI disk asymmetries) and the AGN type (Seyfert, LINER, HII) are analyzed in detail. The outline of the paper is as follows: We describe the HI-NUGA sample, the observation and data reduction in \S \ref{sec:obs}. The results and derived HI properties are presented in \S \ref{sec:res}, along with a discussion of the presence of companions and tidal disturbances of the HI disks and any correlation with AGN type. A comparison of AGN types and optical (stellar) light distribution is also examined. The results are discussed in the context of correlations with the HI environment (e.g. tidal forces) and the AGN type (\S \ref{sec:dis}). A summary is presented in \S \ref{sec:sum}. | \label{sec:dis} The relationship between large scale environment and nuclear activity has been long debated in the literature. Several studies have found indications for correlations between the environment and nuclear activity \citep{Sto01, Cha02, Mar03}. In particular it was suggested that interacting galaxies or galaxies with companions exhibit a significant excess of nuclear activity compared to isolated galaxies \citep{Dah84, Kee85, Raf95}. On the other hand no relation was seen by other studies \citep{Vir00, Sch01, Fue88, Mac89, Lau95}. Hence the issue of possible relations between the environment and nuclear activity appears to be still controversial. Furthermore, results from Keel et al. (\citeyear{Kee85}) suggested that nuclear phenomena might likely be triggered by a tidally induced inflow of gas from the disk to the nuclear regions, rather than gas transfer between the interacting galaxies themselves. Keel (\citeyear{Kee96}) also found that Seyfert galaxies in pairs actually display smaller kinematic disturbances than non-Seyfert galaxies in pairs, which is obviously in disagreement with the hypothesis that tidal interactions are necessary for the transport of angular momentum and the fueling of the SMBH. Since most of these studies are based on optical/IR imaging, they are in principle less sensitive to distortions than our study of the atomic gas that reacts most readily to tidal disturbances. It should be pointed out that a detailed study of HI gas properties for active versus non-active galaxies (Mundell et al., in prep.) is under-way. \par In the context of different AGN types, the analysis of a sample of 451 active galaxies (Sy, LINER, Transition, HII, and absorption-line galaxies) from the Palomar survey \citep{Sch01} showed no correlation between AGN-type and the percentage of galaxies with nearby companions after taking morphological differences of the host galaxies into account. We also see no evidence for a correlation between the fraction of companions and the AGN type present (Seyfert, LINER galaxies), neither from our HI study nor for cataloged optical companions listed in NED (see \S \ref{subsec:res_env}). Note that our sample has a limited number of LINERs (7 galaxies) and Seyfert galaxies (7) and hence the existence of possible relations can not be completely excluded. \par \subsection{HI environment: tidal forces and their correlation with disturbed HI disks} \label{subsec:dis_env} Companion galaxies can possibly disturb the HI gas in a galactic disk via tidal forces and hence affect the fueling of the center with gas \citep[e.g.][]{Kee85}. The sensitivity and endurance of HI to trace the strength and prevalence of tidal interactions among Seyfert galaxies is discussed in detail by Greene et al. (\citeyear{Gre03}). In our analysis we found companions for about half of our sample. In almost all cases the produced tidal forces are smaller than the binding forces of the affected host galaxy (indicated as $\textbf{Q}<0$, see \S \ref{subsec:res_env} and Tab.~\ref{tab_sat}). Only the system NGC5953/54 exhibits signs for very strong gravitational interaction as mentioned in \S \ref{subsec:res_env}. Therefore, most of the disturbances in our sample (6/7 galaxies) can not be explained simply by tidal forces presently at work. The most probable explanations for these disturbances are described in the following, listed in decreasing relevance: \begin{itemize} \item Interaction with a companion now further or far away. If we assume a relaxation timescale of the disturbance of 3$\cdot 10^8$yr (typical time for one rotation of a galaxy) and a maximum fly-by velocity of $\sim$ 500 km s$^{-1}$, the involved companion is expected to be found within a radial distance of $\sim$ 150 kpc from the disturbed galaxy. Since 71\% of the disturbed disks show companions in a reasonable distance (projected distance is less than 150 kpc), tidal interactions in the past seem to be primarily responsible of the disturbances identified in the HI disks. \item Ram pressure stripping \citep{Cay90, Vol00, Vol01}. In particular for NGC~1961 stripping by intergalactic material was suggested \citep{Sho82}. But since none of our galaxies with disturbed HI disk lie in a massive group or cluster environment, gas stripping due to ram pressure is not very likely to explain the presence of disturbances in our sample where no nearby companion is found. Furthermore, the outer HI disk, where disturbances are usually seen, has been removed in clusters due to ram pressure stripping, indicated also by a small HI radius (e.g. NGC~4569, NGC~4579 as part of the Virgo cluster; see \S\ref{subsec:dis_stellar}) \item Minor merging, whereby the companion has now fully merged and has left no optical trace. This might be the case for NGC~4736 and NGC~2782 where no companions are found in a reasonable distance for tidal interaction. \item Large gas accretion from cosmic filaments: Asymmetries in the gas accretion may cause disturbances \citep[for effects of gas accretion on spiral disk dynamics see][]{Bou02}. \end{itemize} \subsection{HI morphology and comparison with the stellar distribution} \label{subsec:dis_stellar} The comparison of the radial density profiles between the HI gas distribution and the stellar distribution revealed significant deviations as expected: 1.) The extent of the HI disk is on average 1.7 times the optical radius, indicated by the Holmberg radius. Only NGC~3627, NGC~4569 and NGC~4579 show a smaller HI radius than optical radius. This can be explained by the fact that they are all members of interacting groups and/or by their rapid motion through an intracluster medium. NGC~3627 is part of the Leo Triplet group and the past encounter with NGC~3628 could explain the spatial conincidence of both the stars and the gas \citep{Zha93}. The truncated disks of NGC~4569 (and NGC~4579) are most probably a signature of strong ram pressure stripping in the past by the intracluster medium which pervades the Virgo Cluster \citep{Cay94}. However, most of the galaxies in our sample show a larger HI disk than their optical one. 2.) The radial density profiles exhibit for most of our galaxies a deficiency of HI gas in the inner part of the galaxy disk. This is in general explained by the phase transition from atomic to molecular gas in neutral ISM \citep{Young}. \subsection{Correlations between HI gas properties and AGN activity type} \label{subsec:dis_correl} As described in \S \ref{subsec:res_agn} our analysis revealed that the number of galaxies with disturbed HI disks is higher for LINER galaxies than for Seyfert galaxies. But since several mechanisms which can not easily be distinguished can cause these asymmetries (see \S \ref{subsec:dis_env}), it is not possible to draw any strong conclusions explaining the higher fraction for LINERs. \par Our study of the HI morphology revealed a significantly higher percentage of galaxies with a HI gas ring for LINER than for Seyfert galaxies. Interestingly, only the study of the Extended 12$\mu$m Galaxy Sample \citep{Hunt99b} indicated a prevalence of stellar rings in active galaxies (Seyfert and LINERs), where LINERs have elevated rates of inner rings, while the Seyfert host galaxies have outer ring fractions several times those in normal galaxies. However, we found no HI gas rings in Seyfert host galaxies. Note that stellar rings are not preferentially found in Seyfert, LINERs, or non-AGNs for our sample by using the optical classification from RC3 listed in NED (see classification in Tab.~\ref{table_intro}; indicated as R or r). To summarize, HI rings are more often found in LINERs while no strong correlation with stellar rings is present. \par One possible explanation for an abundance of HI gas rings in LINERs may be a common evolution of the gas distribution in the disk together with the nuclear activity where both are subject to the influence of a present bar or previous one which has now dissolved. This becomes important since rings are linked observationally and phenomenologically to barred galaxy dynamics \citep[for a general review see][]{But96}, and hence, expected to be seen after the bar had enough time to redistribute the gas toward the end of its life-time. Thus, a time evolution of AGN types seems to be possible, where Seyfert and LINERs represent different phases of the galaxy activity cycle: Seyfert galaxies are the ones where the fueling process has just been triggered (e.g. through disturbances and/or bar dynamics) while LINERs are the ones where the triggering mechanism has already distributed the gas in a more stable new configuration (rings) that does no longer support the massive inflow of gas. Interestingly, disturbances, which are assumed to be a sign for a recent trigger of gas inflow, are preferentially found in LINERs in our sample. That would suggest that LINERS are the first stage of activity, just after the triggering: The gas is still distributed in rings, and the fueling of the AGN has just started by the redistribution of the gas. However, as most of the disturbances in our sample are probably due to tidal interactions in the past, as explained in \S \ref{subsec:dis_env}, the currently observed nuclear activity must not be identical to the one during the ongoing tidal interaction. \par A general scenario for self-regulated activity in low-luminosity AGNs was developed by \cite{Gar05}, in which the onset of nuclear activity is explained as a recurrent phase during the typical lifetime of any galaxy. In this scenario the activity in galaxies is related to that of bar instabilities, expecting that the active phases are not necessarily coincident with the phase where the bar has its maximum strength. Since the infall of gas driven by a bar is self-destructive \citep{Bou02}, i.e. it weakens and destroys the bar, the potential returns to axisymmetry and the gas piles up in a stable configuration (i.e. in rings at the resonances). At this stage, torques exerted by the gravitational potential are negligible and other competing mechanisms, e.g. viscous torques, must transport the gas in the center of the galaxy. The periods of Seyfert/LINER activity (each lasting $\sim 10^8$ years) are suggested to appear during different evolutionary stages of a bar episode (typically characterized by a lifetime $\sim 10^9$ years), depending on the competition between viscosity and gravity torques. The prevalence of HI rings in LINERS, derived in this work, could be explained by the scenario proposed by \cite{Gar05}, where AGN activity is linked to the evolutionary state of the bar-induced gas flow. However, to further substantiate this link requires deriving the complete gravity torque budget based on the HI distribution (see Haan et al. in prep.). Regarding the relative HI radius ($R_{HI}$/$R_{optical}$) we found a slightly larger HI extent for Seyfert galaxies than for LINERs, which increases even more when neglecting galaxies which lie in galaxy clusters or groups. In contrast, our study of the large environment, the HI gas content (by using the ratio $M_{HI}$/$M_{dyn}$), and the relative HI mass ($M_{HI}$/$M_{dyn}$) revealed no significant correlation with the AGN-type. This might be due to the limited number of galaxies in our sample. Hence we conclude that a detailed HI study with a larger sample may reveal more additional information on the interplay between the gaseous component and the AGN type present. | 7 | 10 | 0710.2000 |
0710 | 0710.0005_arXiv.txt | {} {UVES spectra of the very young ( $\sim$\,10$^{7}$ years) peculiar B-type star HR\,6000 were analyzed in the near-UV and visual spectral regions (3050-9460\,\AA{}) with the aim to extend to other spectral ranges the study made previously in the UV using IUE spectra.} { Stellar parameters \teff=12850\,K, \logg=4.10, and $\xi$=0\,km\,s$^{-1}$, as determined from H$_{\beta}$, H$_{\gamma}$, H$_{\delta}$ Balmer profiles and from the \ion{Fe}{i}, \ion{Fe}{ii} ionization equilibrium, were used to compute an individual abundances ATLAS12 model. We identified spectral peculiarities and obtained final stellar abundances by comparing observed and computed equivalent widths and line profiles. } { The adopted model fails to reproduce the (b-y) and c color indices. The spectral analysis has revealed: the presence of emission lines for \ion{Mn}{ii}, \ion{Cr}{ii}, and \ion{Fe}{ii}; isotopic anomalies for \ion{Hg}, \ion{Ca}; the presence of interstellar lines of \ion{Na}{i} at $\lambda\lambda$ 3302.3, 3302.9, 5890, 5896\,\AA{}, and of \ion{K}{i} at 7665, 7699\,\AA{}; the presence of a huge quantity of unidentified lines, which we presume to be mostly due to \ion{Fe}{ii} transitions owing to the large Fe overabundance amounting to [+0.7]. The main chemical peculiarities are an extreme overabundance of Xe, followed by those of Hg, P, Y, Mn, Fe, Be, and Ti. The most underabundant element is Si, followed by C, N, Al, S, Mg, V, Sr, Co, Cl, Sc, and Ni. The silicon underabundance [$-$2.9] is the lowest value for Si ever observed in any HgMn star. The observed lines of \ion{He}{i} can not be reproduced by a single value of the He abundance, but they require values ranging from [$-$0.8] to [$-$1.6]. Furthermore, when the observed and computed wings of \ion{He}{i} lines are fitted, the observed line cores are much weaker than the computed ones. From the present analysis we infer the presence of vertical abundance stratification for He, Mn, and possibly also P. } {} | HR\,6000 (HD\,144667) is one of the most remarkable chemically peculiar (CP) stars. It does not fit any of the CP subclasses, but it seems to combine abundance anomalies from a variety of Bp sub-types. It forms with the star HR\,5999 (HD\,144668) the common proper motion visual binary system $\Delta$199 or Dunlop\,199 (Bessell \& Eggen, 1972). The angular separation between the HR\,6000, which is the brighter more massive component, and the secondary component, the well-known Herbig Ae star HR\,5999, is about 45\arcsec{}. The main interest to study the chemical composition of HR\,6000 comes from the generally poor understanding of the occurrence of abundance anomalies in such a young object with an estimated age of the order of 10$^{7}$ years. In fact, the $\Delta$\,199 double system is located close to the center of the Lupus\,3 molecular cloud which is populated by numerous T\,Tauri stars. This has led to an assumption that this system has the same age as the cloud, which is estimated to be (9.1$\pm$2.1)$\times$10$^{6}$ years (James et al., 2006). However, after the Hipparcos mission, the membership of HR\,5999 and HR\,6000 to the cloud became rather questionable due to the uncertainties in the distance determination for the Lupus cloud. The distance of HR\,5999 and HR\,6000 measured by Hipparcos are 208$\pm$38\,pc and 240$\pm$48\,pc, respectively, while the Lupus cloud distance is 150$\pm$10\,pc according to Crawford (2000). Comer\`on et al.\ (2003) assigned a distance of about 200\,pc to the Lupus cloud, but this determination was made by assuming a priori that $\Delta$199 belongs to the cloud. A ROSAT survey of Herbig Ae/Be stars presented by Zinnecker \& Preibisch (1994) has detected strong X-ray emission coming from the direction of HR\,6000. To explain the X-ray origin, a possible T\,Tauri companion for HR\,6000 was postulated by van den Ancker et al.\ (1996) on the basis of an infrared excess in the energy distribution relative to predictions made for an effective temperature of \teff=14000\,K. Siebenmorgen et al.\ (2000) fitted very short spectral regions of observed flux in the mid IR to a black body of 13000\,K and noticed that the observed spectrum was featureless. Neither spectrum nor radial velocity variability was found for HR\,6000 by Andersen \& Jaschek (1984), who studied optical spectra taken at different epochs. However, van den Ancker et al.\ (1996) discovered long-term variations in the u, v, and b Str\"omgren magnitudes with approximate amplitude of 0\fm{}03$\pm$0\fm{}01 in u, 0\fm{}02$\pm$0\fm{}01 in v, and 0\fm{}01 in b. No variations in y were found. Kurtz \& Marang (1995) discovered variations of 0.008 mag in V with a period near to 2\,d and suggested that this period could be possibly caused by rotational variation in a spotted magnetic star. This would imply a pole on orientation for HR\,6000 as they measured $v \sin i$$\le$5~km~s$^{-1}$. Catanzaro et al.\ (2004) were the first ones who provided abundances from the optical range. Previously, optical spectra were studied only by Andersen \& Jaschek (1984) and Andersen et al.\ (1984) who identified the spectra from 3323\,\AA{} to 5317\,\AA{} and discussed the abundances on the basis of the line intensities. It is remarkable that the stellar parameters derived by Catanzaro et al.\ (2004) from the Balmer profiles (\teff=12950$\pm$50~K, \logg=4.05$\pm$0.01) are 1000\,K lower than those previously adopted and deduced from the photometry by Castelli et al.\ (1985) (\teff=14000\,K, \logg=4.0), by Smith (1997) (\teff=13990\,K, \logg=4.29) and by van den Ancker et al.\ (1996) (\teff=14000\,K, \logg=4.3). In this paper we examine the whole optical spectrum of HR\,6000 from 3050\,\AA{} to 9460\,\AA{} with the aim to extend the analysis performed on IUE spectra by Castelli et al.\ (1985) to the visible region and to investigate the possible contamination of the stellar spectrum by a close T\,Tau companion. We re-examine the parameter determination, the line identification, and the abundances. We also searched for the presence of peculiarities like emission lines and isotopic anomalies which were recently discovered in HgMn stars. Owing to the availability of spectra taken at different epochs we also investigated possible spectral variabilities. The comparison of the observed and computed spectra as described in this paper is available at the web-address given in the footnote\footnote{http://wwwuser.oat.ts.astro.it/castelli/hr6000/hr6000.html}. | The present study of HR\,6000 has led to the following results: the measured radial velocity indicates that HR\,6000 most likely belongs to the Lupus cloud so that its age is of the order of 10$^{7}$ years as it was found for the Lupus\,3 cloud (James et al., 2006). Interstellar lines from the cloud can be observed in the spectrum from the ultraviolet to the visible, with displacements of the order of $-$0.05 $\div$ $-$0.15\,\AA{} from the stellar lines. UVES spectra observed with a six months time interval do not show spectrum variability for both stellar and interstellar lines, but a low variability of 1.2\,km\,s$^{-1}$ in the radial velocity and marginal variabilities of the weak spectral features which are however at the level of the noise. Different methods for the parameter determination have led to large differences in \teff{}. While the gravity values are found in the range between 4.1 and 4.4\,dex with errors of the order of 0.1$\div$0.2\,dex, \teff{} changes from 13900$\pm$300\,K, (if obtained from the UBV$\beta$ and uvby$\beta$ photometry), down to 12850$\pm$50\,K when determined from the Balmer profiles, and can become as low as 11200$\pm$300\,K when derived from the VI$_{c}$J photometry. Furthermore, as far as the Balmer profiles are concerned, the parameters which well reproduce H$_{\beta}$, H$_{\gamma}$, and H$_{\delta}$ predict too narrow H$_{\alpha}$ wings. All these discrepancies in the parameter determinations could be accounted for by the presence of the Lupus cloud with could affect the colors in such a way to invalidate the standard reddening relations. Also a spectral contamination by a weak-emission T\,Tauri (WTT) companion, as suggested by van den Ancker et al.\ (1996), can not be excluded in an absolute way, although the only spectroscopic sign of its possible presence is a weak variable broad absorption at the position of \ion{Li}{i} 6707\,\AA\ and a somewhat better agreement between observed and computed spectra longward of 6000\,\AA{} for a combined synthetic spectrum obtained from that HR\,6000 and one computed for \teff{}=3500\,K, \logg=4.0, [M/H]=0.0. This companion could be either physically related with HR\,6000 to form a close binary system or could be located in the foreground of HR\,6000. Finally, the different parameters derived from the different determinations could be due to the extremely peculiar nature of HR\,6000 which may make the classical model atmospheres inadequate to reproduce all the observed stellar quantities. In fact, the abundance stratification inferred for some elements, helium in particular, would require more refined models in which the hypothesis of constant abundances throughout the atmosphere is dropped. By using an ATLAS12 model with parameters \teff=12850\,K, \logg=4.1, $\xi$=0.0\,km\,s$^{-1}$ we have computed a synthetic spectrum for HR\,6000 from 3050\,\AA\ to 9460\,\AA\ and have compared it with the observed spectrum. This comparison has shown that most of the elements lighter than Ca are significantly underabundant, except for Be, Na, and P. The Si underabundance ([-2.9]) is remarkable because it is even lower than that of 46 Aql ([$-$1.0]) which was claimed by Sadakane et al. (2001) to be the lowest one found in HgMn stars. A striking peculiarity of HR\,6000 is the lack of any overabundance for heavy elements with Z$>$40, except for Xe and Hg. The line spectrum of \ion{Xe}{ii} is similar to that we observed in a preliminary analysis of UVES spectra of 46 Aql, another \ion{Xe} rich star: while most of the lines lie at the laboratory wavelength, some other \ion{Xe}{ii} lines seem to be shifted by about $-$0.1\,\AA\ from the predicted position. The same behaviour can be observed also in HD\,175640, although to a less extent, owing to its lower \ion{Xe} overabundance. Unfortunately, the small number of known transition probabilities and the ignorance of the \ion{Xe}{ii} isotopic structure due to the lack of atomic data put strong limitations in the study of this element in the spectra of the CP stars. The large overabundance of [+0.7] for iron generates a very rich line spectrum in which numerous, still unclassified \ion{Fe}{ii} lines are observed. A very large number of lines probably due to \ion{Fe}{ii} remain unidentified. The similar iron overabundance of 46 Aql ([+0.65]) gives rise to an impressive close resemblance between the spectra of the two stars. Abundances in the ultraviolet obtained from a re-analysis of the IUE spectra studied by Castelli et al. (1985) agree with abundances derived from the UVES spectra for all the elements, except for carbon and phosphorus. While the carbon identification in the optical spectrum is rather questionable owing to the weakness of the observed lines, the difference in the phosphorus abundance is similar to that yielded by the individual \ion{P}{ii} and \ion{P}{iii} lines. There are several signs of vertical abundance stratification in HR\,6000 for He, P, Mn, and Fe. The peculiar shape of the \ion{He}{i} profiles is discussed for the first time in this paper. The profiles have cores too intense as compared to the wings, a fact indicating strong vertical He abundance stratification (Dworetsky, 2004; Bohlender, 2005). Furthermore, the inferred He abundance decreases with increasing wavelength ranging from an underabundance of [-0.8] at 4000\,\AA\ to [-1.6] at 6000\,\AA. For \ion{Mn}{ii}, the abundance shortward of the Balmer discontinuity is larger by about 0.6\,dex than that longward of the Balmer discontinuity; for phosphorus, the \ion{P}{ii} abundance is larger by 0.2\,dex than the \ion{P}{iii} abundance, but this discrepancy lies within the mean square error of the average abundances; for iron, the abundance from \ion{Fe}{ii} high excitation lines is generally higher by about 0.2\,dex than that from \ion{Fe}{ii} low excitation lines. Similar to other studied HgMn stars, also HR\,6000 shows emission lines. The most numerous ones belong to \ion{Mn}{ii}, followed by those of \ion{Fe}{ii} and \ion{Cr}{ii}. While the first two elements are overabundant in HR\,6000, chromium has solar abundance. The star exhibits also isotopic anomalies for Hg and Ca. In both cases the most heavy isotope is the predominant one. | 7 | 10 | 0710.0005 |
0710 | 0710.3626_arXiv.txt | Globular clusters produce orders of magnitude more millisecond pulsars per unit mass than the Galactic disk. Since the first cluster pulsar was uncovered twenty years ago, at least 138 have been identified -- most of which are binary millisecond pulsars. Because of their origins involving stellar encounters, many of these systems are exotic objects that would never be observed in the Galactic disk. Examples include pulsar---main sequence binaries, extremely rapid rotators (including the current record holder), and millisecond pulsars in highly eccentric orbits. These systems are allowing new probes of the interstellar medium, the equation of state of material at supra-nuclear density, the mass distribution of neutron stars, and the dynamics of globular clusters. | The first globular cluster (GC) pulsar was identified 20 years ago in the cluster M28 after intense efforts by an international team \cite{lbm+87}. Since then at least 138 GC pulsars\footnote{For an up-to-date catalog of known GC pulsars, see Paulo Freire's website at \url{http://www.naic.edu/~pfreire/GCpsr.html}}, the vast majority of which are millisecond pulsars (MSPs), have been found. Finding these GC pulsars has required high-performance computing, sophisticated algorithms, state-of-the-art instrumentation, and deep observations with some of the largest radio telescopes in the world, primarily Parkes, Arecibo, and the Green Bank Telescope (GBT). The payoff has been an extraordinarily wide variety of science. Low-Mass X-ray Binaries (LMXBs) have been known to be orders-of-magnitude more numerous per unit mass in GCs as compared to the Galactic disk since the mid-1970s \cite{kat75,cla75}. This overabundance is due to the production of compact binary systems containing primordially-produced neutron stars via stellar interactions within the high-density cluster cores. Since LMXBs are the progenitors of MSPs, this dynamics-driven production mechanism also applies to them, and it has made GCs (particularly the massive, dense, and nearby ones) lucrative targets for deep pulsar searches. Camilo \& Rasio \cite{cr05} produced an excellent review of the first 100 GC pulsars in 2005. This current review provides a significant update to Camilo \& Rasio as it concentrates on the advances made (primarily with the GBT) within the past several years, including almost 40 additional pulsars and over 50 new timing solutions. | 7 | 10 | 0710.3626 |
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0710 | 0710.2882_arXiv.txt | We study the initial mass function (IMF) of one of the most massive Galactic star-forming regions NGC 3603 to answer a fundamental question in current astrophysics: is the IMF universal, or does it vary? Using our very deep, high angular resolution $JHK_{S}L'$ images obtained with NAOS-CONICA at the VLT at ESO, we have successfully revealed the stellar population down to the subsolar mass range in the core of the starburst cluster. The derived IMF of NGC 3603 is reasonably fitted by a single power law with index $\Gamma \sim -0.74$ within a mass range of $0.4 - 20$ \msun, substantially flatter than the Salpeter-like IMF. A strong radial steepening of the IMF is observed mainly in the inner $r \lesssim 30''$ field, indicating mass segregation in the cluster center. We estimate the total mass of NGC 3603 to be about $1.0 - 1.6 \times 10^4$ \msun. The derived core density is $\geq 6 \times 10^4$ \msun pc$^{-3}$, an order of magnitude larger than e.g., the Orion Nebula Cluster. The estimate of the half-mass relaxation time for solar-mass stars is about $10 - 40$ Myr, suggesting that the intermediate- and low-mass stars have not yet been affected significantly by the dynamical relaxation in the cluster. The relaxation time for the high-mass stars can be comparable to the age of the cluster. We estimate that the stars residing outside the observed field cannot steepen the IMF significantly, indicating our IMF adequately describes the whole cluster. Analyzing thoroughly the systematic uncertainties in our IMF determination, we conclude that the power law index of the IMF of NGC 3603 is $\Gamma = -0.74^{+0.62}_{-0.47}$. Our result thus supports the hypothesis of a potential top-heavy IMF in massive star-forming clusters and starbursts. | \label{introduction} One of the most interesting properties of massive star-forming regions is the stellar initial mass function (IMF). Since the pioneering work by \citet{sal55}, which led to a standard picture of the IMF, the so-called Salpeter IMF ($dN/d\log \mathcal{M} \propto \mathcal{M}^\Gamma$ with $\Gamma = -1.35$, where $\mathcal{M}$ is stellar mass), numerous efforts have been made to understand the IMF of many types of objects such as field populations, stellar associations, and open and globular clusters, covering the whole stellar mass range from massive OB stars down to substellar brown dwarfs. It is currently accepted that the IMF follows a single Salpeter-like power law in the high- to intermediate-mass range ($\mathcal{M} \gtrsim 1$ \msun). It becomes flatter towards subsolar masses and peaks at a characteristic mass of several tenths of a solar mass. It then declines towards the brown dwarf mass range. Therefore, several analytical expressions have been proposed for the standard IMF, for example, a lognormal distribution \citep{mil79,sca86}, a segmented power law distribution \citep{sca98, kro01}, and a combination of both \citep{cha03}. As these representations of the IMF fit reasonably well the various types of stellar populations, the idea of the \textit{universality} of the IMF was born. A universal IMF basically suggests a universal star formation mechanism. However, it is somewhat intuitive that different physical conditions such as the density, velocity fields, chemical composition, and tidal forces in the natal molecular clouds can lead to different star formation processes and, consequently, some variability in the IMF. Indeed, there is growing evidence for the variable IMF. There are significant variations in the power law index and the characteristic masses at which the distribution shows the peak or at which the power law breaks. A typical example of an IMF that does not follow the Salpeter-like distribution is the so-called top-heavy IMF in starburst galaxies and young, massive star-forming regions. Does the IMF really vary among stellar populations? To answer this question, we need to clarify, for at least some stellar populations, if the observed IMF variations are without any doubt true deviations from the Salpeter-like IMF or if they can simply be accommodated by the combined effects of observational, theoretical, and statistical uncertainties? Many studies have so far addressed this question. In order to address the question and to obtain new insights into the IMF of an intense starburst environment, we study the IMF of the massive star-forming HII region \objectname{NGC 3603} based on unprecedented spacial resolution observations of its central starburst cluster obtained by NACO at the Very Large Telescope (VLT), as well as a wider field with the Infrared Spectrometer and Array Camera (ISAAC) at VLT at the European Southern Observatory (ESO). Compared to extragalactic starbursts (e.g. \objectname{M82}), young star-forming regions in the Milky Way and in the Magellanic Clouds -- so-called nearby starburst templates -- are close enough to resolve the individual stars with current powerful adaptive optics (AO)-assisted ground-based telescopes and space-based facilities. These objects include, for example, the \objectname{Arches} cluster, the \objectname{Quintuplet} cluster, \objectname{R136}, and \objectname{NGC 3576}. More nearby but less massive star-forming regions are the \objectname{Trapezium} cluster in the Orion Nebula Cluster (ONC), the \objectname{Pleiades} cluster, the \objectname{Taurus} cluster, and \objectname{IC 348}. The IMF of NGC 3603 has been presented in several studies, and recent works have derived somewhat flat IMFs with a power law index of $\Gamma = -0.73$ for $1 - 30$ \msun\ in \citet{eis98}, $\Gamma = -0.9$ for $2.5 - 100$ \msun\ in \citet{sun04}, and $\Gamma = -0.91 \pm 0.15$ for $0.4 - 20$ \msun\ in \citet{sto06}. As another example of the IMF of massive Galactic star clusters, the Arches cluster has shown a slightly flat IMF with $\Gamma = -0.6$ to $-1.1$ for intermediate- and high-mass stars \citep{fig99,sto05,kim06}. Studies of the IMF of the Arches and other clusters are presented in \S~\ref{c-discussion}. \subsection{NGC 3603} NGC 3603 is one of the most luminous, optically visible HII regions in the Milky Way \citep{gos69} with its global properties such as $L$(H$_{\alpha}$) $\sim 1.5 \times 10^{39}$ ergs s$^{-1}$ \citep{ken84}, and the total mass of molecular clouds $\sim4 \times 10^{5}$ \msun\ \citep{gra88}. Its total bolometric luminosity is as large as 10$^{7}$ $L_{\odot}$, 2 orders of magnitude larger than that of the ONC, and just an order of magnitude smaller than the other well known starburst template R136 in 30 Dor in the Large Magellanic Cloud (LMC). In particular, the starburst cluster located in the northern part of the gigantic HII complex including \objectname{HD 97950} -- a very compact Trapezium-like system with plenty of high-mass components \citep{wal73} -- at its center, is one of the most massive and the densest star-forming regions known in the Galaxy. The starburst cluster consists of three WNL stars, six O3 stars, and many other late O- and B-type stars in a volume of less than a cubic light year, providing most of the ionizing radiation in the giant HII region \citep{mof83, cla86, mof94, dri95,hof95, cro98}. It thus shows remarkable similarities with R136. The distance of NGC 3603 from the Sun has been estimated in many earlier works by means of both photometric and kinematic analysis \citep[e.g.][]{vdb78,mel82}, and the currently accepted value is about 7 $\pm$ 1 kpc. Owing to its intrinsic properties such as its proximity, relatively low visual extinction of only $A_{V} = 4 - 5$ mag, and extreme compactness and brightness, NGC 3603 is one of the most suitable Galactic templates of starburst phenomena in distant galaxies. | \label{c-summary} Our study is aiming at a fundamental question of current star formation research. Is the IMF universal, or does it vary with environment? To answer this question, we measured the IMF of NGC 3603 -- one of the most massive Galactic star-forming regions -- from our NIR observations with the AO system NACO at the VLT/ESO. In the following we summarize the main results from our study. \begin{enumerate} \item \textbf{NIR Photometry}\\ Our very deep, high angular resolution $JHK_{S}L'$-band images obtained by NACO show unprecedented details of the core of the starburst cluster in NGC 3603. Together with the wider field ISAAC $JHK_{S}$ images, we could successfully derive magnitudes and positions of almost 10,000 stars in the dense cluster up to $r \sim 110''$ covering the mass range from the most massive stars down to $\sim0.4$ \msun. The brightest 256 stars in the NACO images of the inner $r \leq 13''$ are listed and cross-identified with potential counterparts in previous studies. \item \textbf{Age, extinction, and disk fraction}\\ Based on the fitting of the stellar evolution models in the CMDs and CCDs, we derived the age of the PMS stars of $0.5 - 1.0$ Myr and the upper limit for the age of the MS stars of $\sim2.5$ Myr, suggesting a slight age spread in the cluster. The derived average foreground extinction is $A_{V} = 4.5 \pm 0.5$ mag, and the foreground extinction increases by $\Delta{A_{V}} \sim 2.0$ mag towards larger radii ($r \gtrsim 55''$). Using the $K_{S} - L'$ versus $J - H$ CCD, we derived a circumstellar disk fraction of $\sim25 \pm 10$\% for stars with a mass of $\geq 0.9$ \msun\ in the central cluster ($r \leq 10''$). \item \textbf{IMF and its radial variation}\\ Applying the field star rejection and the incompleteness correction, the \textit{K}LF for 7514 stars simultaneously detected in the $JHK_{S}$ bands follows a power law with no obvious turnover or truncation within the detection limit of $m_{K_{S}} \sim 17.4$ mag (based on the $J$-band 50\% completeness of $\sim19.4$ and the typical color of $J - K_{S} \sim 2$). Within the mass range of $0.4 - 20$ \msun\, the IMF is well described by a single power law with a power law index $\Gamma \sim -0.74$. We found the power law index decreasing from $\Gamma \sim -0.3$ at $r \leq 5''$ to $\Gamma \sim -0.8$ at $r \sim 30''$. The strong steepening occurs in the inner $r \lesssim 13''$, pointing towards mass segregation in the very center of the cluster. No significant variation of the IMF is found for larger radii ($r \gtrsim 30''$). \item \textbf{Size, mass, and dynamical status}\\ Fitting a King model to the radial density profile of stars with a mass of $0.5 - 2.5$ \msun, we derived a core radius of $\sim4''.8$ ($\sim0.14$ pc at $d \sim 6$ kpc). As the radial density decreases even at the limits of our field of view, we can give a firm lower limit of $r = 110''$ ($\sim3.2$ pc) for the cluster size. We also derive an upper limit of $r = 1260''$ ($\sim37$ pc) for the tidal radius of the cluster. The de-projected King model allowed us to extrapolate to the total mass of NGC 3603. Assuming a single power law IMF with index $\Gamma = -0.74$ within the mass range of $0.1 - 100$ \msun, we found a total mass of about $1.0 - 1.6 \times 10^4$ \msun. The half-mass radius is found to be within $25'' - 50''$ ($0.7 - 1.5$ pc). The derived core mass density of the cluster is $\geq 6 \times 10^4$ \msun\ pc$^{-3}$. We estimate a half-mass relaxation time of approximately $10 - 40$ Myr for stars with a typical mass of 1 \msun, an order of magnitude larger than the age of the PMS population in the cluster ($\lesssim1$ Myr). This implies that the intermediate- and low-mass stars have not yet experienced significant dynamical relaxation. However, the relaxation time of the high-mass stars is expected to be an order of magnitude shorter and is comparable to the cluster age. We could thus not conclude if the observed mass segregation of the high-mass stars is caused by dynamical evolution or if it is primordial. Indeed it can be due to the combination of them. We compute that our images with a maximum radius of $r \sim 110''$ cover at least $\sim67$\% of intermediate- and low-mass stars of NGC 3603. The stars outside the observed field cannot steepen the IMF by more than $\Delta\Gamma \lesssim 0.16$. Considering also the fact that the IMF does not significantly change beyond $r \gtrsim 30''$, we conclude that the observed IMF is representative for the whole NGC 3603 stellar cluster, irrespective of the mass segregation in the very center. \item \textbf{Systematic uncertainties of the IMF}\\ We thoroughly analyzed the systematic errors in the IMF determination. In particular we derived the errors from uncertainties in the age, distance, foreground extinction, stellar evolution models, metallicity, individual extinction, incompleteness correction, cluster membership, unresolved binaries, and the stellar mass outside the observed field. Combining all errors we derived the power law index of $\Gamma = -0.74^{+0.62}_{-0.47}$. Assuming a Gaussian probability distribution, we conclude that the probability that the IMF is as steep as the Salpeter IMF ($\Gamma = -1.35$) is less than $\sim10$\%. \end{enumerate} Our result thus supports the hypothesis of a top-heavy IMF in massive star-forming clusters and starburst galaxies. One potentially common property among such clusters showing flat IMFs would be the high stellar density in the core of the cluster. | 7 | 10 | 0710.2882 |
0710 | 0710.0763_arXiv.txt | Magnetic fluctuations in the solar wind are distributed according to Kolmogorov's power law $f^{-5/3}$ below the ion cyclotron frequency $f_{ci}$. Above this frequency, the observed steeper power law is usually interpreted in two different ways: a dissipative range of the solar wind turbulence or another turbulent cascade, the nature of which is still an open question. Using the Cluster magnetic data we show that after the spectral break the intermittency increases toward higher frequencies, indicating the presence of non-linear interactions inherent to a new inertial range and not to the dissipative range. At the same time the level of compressible fluctuations raises. We show that the energy transfer rate and intermittency are sensitive to the level of compressibility of the magnetic fluctuations within the small scale inertial range. We conjecture that the time needed to establish this inertial range is shorter than the eddy-turnover time, and is related to dispersive effects. A simple phenomenological model, based on the compressible Hall MHD, predicts the magnetic spectrum $\sim k^{-7/3+2\alpha}$, which depends on the degree of plasma compression $\alpha$. | \label{Intro} Solar wind, which is highly turbulent, represents a unique opportunity to investigate turbulence in natural plasmas via \emph{in situ} measurements \cite{sw2,noi}. In non-magnetized fluids, where the energy injection scale is far from the dissipation one, the intermediate scales (inertial range) are described by the universal power law Kolmogorov spectrum $k^{-s}$ with $s=5/3$. This law depends neither on the energy injection nor on the energy dissipation processes. In the solar wind, and in the interplanetary space in general, the mean free path roughly corresponds to the Sun-Earth distance and the usual dissipation via collisions is negligible. At the same time, in a magnetized plasma there is a number of characteristic space and temporal scales. Investigating solar wind turbulence at these scales is then a challenging topic from the point of view of basic plasma physics. Here we will focus our discussion on the ion scales, namely the ion inertial length $\lambda_i = c/\omega_{pi}$ and the ion cyclotron frequency $f_{ci}= eB/m_i$. At these scales the fluid-like approximation of plasma dynamics, the usual Magnetohydrodynamic (MHD) description, breaks down in favor of a more complex description of plasma. Solar wind turbulent spectrum of magnetic field fluctuations follows a $\sim f^{-5/3}$ power law below the ion cyclotron frequency $f_{ci}$. For $f>f_{ci}$ the spectrum steepens significantly, but is still described by a power law $f^{-s}$, with $s \in (2-4)$ \citep{Leamon98,smith06}. \emph{In situ} solar wind measurements provide time series data, i.e. information on frequency in Fourier space. How to get any information on wave vectors? For fluctuations with velocities much smaller than the plasma bulk velocity $V$, the Taylor hypothesis is valid: the observed variations on a time scale $\delta t$ correspond to variations on the spatial scale $\delta r=V \delta t$. Therefore there is a direct correspondence between $f$ and $k$ spectra. % If the solar wind turbulence was a mixture of linear wave modes, Taylor hypothesis would be well verified only below the spectral break: the bulk velocity is superalfv\'enic ($V>V_A$, where $V_A$ being the Alfv\'en speed) and so the low frequency fluctuations can be considered as frozen in plasma. However, whistler waves (with $f > f_{ci}$ and phase speed $V_{\varphi} > V_A$) do not satisfy the Taylor hypothesis assumption. In this study we assume that in the solar wind there are no whistler waves above the spectral break frequency. This last assumption is supported by results recently obtained in the Earth's magnetosheath \citep{Mangeney2006}. Using the Taylor hypothesis the observed solar wind spectrum below the break is usually attributed to the Kolmogorov's spectrum $\sim k^{-5/3}$. Above the spectral break, the spectral steepening, $\sim k^{-3}$, can be interpreted in two different ways. Some authors associate it to the dissipation range \citep{Leamon98,Leamon99,Leamon2000,Bale2005,smith06}. Others suggest that after the spectral break another turbulent cascade takes place \citep{Biskamp1996,Ghosh96,stawicki,Li01,Galtier06}. In ordinary fluid flows the dissipation range is described by an exponential function \citep{frisch}. While in the solar wind, a well-defined power law is observed after the break point and not an exponential. Note that power spectra, i.e. the second order statistics, completely describe Gaussian, or statistically independent, fluctuations. However, as is well known, fluctuations cannot be described by a Gaussian statistics in the low-frequency part of the solar wind turbulence \citep{noi}. Deviations from Gaussianity, i.e. intermittency \citep{frisch}, may be quantified by the flatness, the forth-order moment of fluctuations. In this paper we investigate the nature of magnetic fluctuations in the high frequency range of the solar wind turbulence, that is usually called dissipation range. We find that in this range the flatness increase with frequency. This is similar to what is going on in the low-frequency range. The presence of intermittency together with the well defined power law in the high-frequency part of the spectrum suggests another turbulent cascade rather than a dissipation range. This small scale cascade is observed to be much more compressible than the Kolmogorov-like inertial range, in agreement with the previous observations by the Wind spacecraft \citep{Leamon98}. We show that the energy transfer rate and intermittency are sensitive to the level of compressibility of the turbulent fluctuations in this range. Finally, we propose a simple phenomenological model, based on the compressible Hall MHD, which allows to explain the observed range of the spectral indices in the high frequency part of the solar wind spectrum. | In this paper we investigate small scales turbulent fluctuations in the solar wind (i.e. at frequencies above the spectral break at the vicinity of $f_{ci}$). Taking into account the Taylor hypothesis ($\ell=V/f$) the frequency domain above the spectral break covers $\sim (40- 2000)$~km space scales that corresponds to $(0.3- 15)\lambda _i$. In this range we found evidences for strong departure from Gaussian statistics and the presence of intermittency while the spectrum presents a well-defined power law. Both the presence of a power-law spectrum and the absence of global self-similarity, seems to be quite in contrast with the role of ``dissipative range". In usual fluid turbulence, the dissipative range \citep{frisch} starts with a rough exponential cutoff; in the near dissipation range the intermittency increases as far as the Gaussian fluctuations dissipate faster than the coherent structures \citep{near-diss}; then the fluctuations become self-similar, the singularities being smoothed by dissipation. In the solar wind turbulence we observe a completely different picture. After the spectral break in the vicinity of $f_{ci}$ the flatness increases as a power law indicating that non-linear interactions are at work to build up a new inertial range. This small scale cascade is much more compressible than the lower frequency Alfv\'enic cascade. This sudden change in nature of the turbulent fluctuations can happen due to a partial dissipation of magnetic fluctuations at the spectral break: the left-hand Alfv\'enic fluctuations with $k_{\|}\gg k_{\perp}$ are damped by the ion-cyclotron damping \citep{Ghosh96}. Above the break, however, a new \emph{magnetosonic cascade} takes place up to the electron characteristic scales. This energy cascade is seems to be dominated by the fluctuations with $k_{\perp}\gg k_{\|}$ \citep{fouad06,Mangeney2006}. We found, as well, that the plasma compressibility controls the statistics of the magnetic field fluctuations. Preliminary results show that the increase of the level of the plasma compressibility leads to the spectrum steepening and increase of the intermittency. To explain this small scale compressible cascade we introduce here for the first time a simple phenomenological model based on the compressible Hall MHD: we find that the magnetic energy spectrum follows a $E(k) \sim k^{-7/3 + 2\alpha}$ law, depending on the degree of plasma compressibility $\alpha$. While far from a complete model of small scale turbulence, this simple model can explain observed variations of the spectral index within the high frequency part of the solar wind turbulent spectrum \citep{Leamon98,smith06} by different degree of plasma compressibility. | 7 | 10 | 0710.0763 |
0710 | 0710.5831_arXiv.txt | {} {Microquasars are ideal natural laboratories for understanding accretion/ejection processes, studying the physics of relativistic jets, and testing gravitational phenomena. Nevertheless, these objects are difficult to find in our Galaxy. The main goal of this work is to increase the number of known systems of this kind, which should allow better testing of high-energy phenomena and more realistic statistical studies of this galactic population to be made.} {We have developed an improved search strategy based on positional cross-identification with very restrictive selection criteria to find new MQs, taking advantage of more sensitive modern X-ray data. To do this, we made combined use of the radio, infrared, and X-ray properties of the sources, using different available catalogs.} {We find 86 sources with positional coincidence in the NVSS/XMM catalogs at galactic latitudes $|$b$|$ $\leq$ 10$^{\circ}$. Among them, 24 are well-known objects and the remaining 62 sources are unidentified. Out of 86 sources, 31 have one or two possible infrared counterparts in the 2MASS catalog. For the fully coincident sources, whenever possible, we analyzed color-color and hardness ratio diagrams and found that at least 3 of them display high-mass X-ray binary characteristics, making them potential microquasar candidates.} {} | Some of the most attractive objects in the Galaxy are the enigmatic microquasars (MQs). On a smaller scale, they copy the characteristics exhibited by distant quasars \cite{mirabel99}. These systems are X-ray binaries (XRBs) containing compact objects like stellar black holes or neutron stars that accrete matter from a companion star. They are known to emit from radio to X-ray energies \cite{mirabel94} and possibly up to TeV gamma-ray energies, as in the case of Cygnus X-1 \cite{albert07}. Evidence that jets of accreting X-ray binaries can accelerate particles to TeV energies has nevertheless been presented by Corbel et al. (2002). MQs combine two important aspects of relativistic astrophysics: accreting black holes or neutron stars identified by the production of hard X-rays around accreting disks and relativistic radio jets detected by means of their synchrotron emission. These binary systems are ideal natural laboratories for understanding accretion/ejection process and other gravitational phenomena. However, they seem to be rare objects in our Galaxy. In order to enable more robust statistical studies, it is necessary to increase the number of known MQs. Of the 15 currently confirmed MQs in the galaxy, 6 belong to the high-mass X-ray binary (HMXB) class and 9 are of the low-mass X-ray binary (LMXB) kind \cite{paredes05}. Finding new MQs candidates is not an easy task. Considerable effort in the past has been put in to increasing their number. A few searches for radio-emitting XRB systems and therefore new MQ candidates, based on cross-identifications between radio, infrared, and X-ray catalogs, have been carried out in the past, but without enough success \citep[e.g.][]{paredes02,ribo04}. The method of looking for such objects in the Galaxy usually includes a number of steps with very restrictive selection criteria for the sources being investigated. Mainly, a number of competing emission mechanisms and several physical parameters should be associated to the same source or system, and they basically include the properties of jet emission and XRB behavior. The detection of emission at radio wavelengths could be the signature that such relativistic jets, mainly emitting via incoherent synchrotron emission from very high-energy electrons spiraling in magnetic fields, are present in the object. Ultraviolet and infrared emissions are characteristics displayed by the normal star companions, and X-ray radiation is most efficient at revealing accretion-powered sources, such as binary stars. As an example, a previous list of MQ candidates obtained from cross-correlation of the radio NVSS catalog \cite{condon98} and the ROSAT Bright Source Catalog (RBSC) sources of the Galactic plane was presented by Paredes et al. (2002). They found 35 possible MQ candidates based on a hardness ratio criterion i.e. $HR1+\sigma(HR1)\geq 0.9$, where $HR1=([0.5-2.0 {\rm ~keV}]-[0.1-0.4 {\rm ~keV}])/([0.5-2.0 {\rm ~keV}]+[0.1-0.4 {\rm ~keV}])$. However, since ROSAT soft X-rays are strongly absorbed by the interstellar medium, their resulting list of MQ candidates could be fairly reduced. Given the impossibility that ROSAT will detect X-ray photons at energies above 2.4 keV, and because X-rays in MQs are expected to be highly energetic, such a criterion could not be efficient enough to select possible MQ candidates. Here, we develop an improved search strategy that is also based on very restrictive, but improved, selection criteria aimed at finding new MQs in the Galaxy. Moreover, we take advantage of modern and more sensitive multiwavelength data. In Sect.~\ref{search} we describe the search strategy and sample definition. The general analysis and statistical results are presented in Sect.~\ref{results}. Finally, we summarize our main conclusions in Sect.~\ref{summary}. | \label{summary} In this work we have presented an improved search strategy for finding new MQ candidates based on positional cross-identifications. By analyzing radio and modern more sensitive X-ray data, we found 86 positional-coincidence sources in the NVSS and XMM catalog, with galactic latitudes $|b| \leq 10^{\circ}$. Among them, 24 are well-known objects and the rest of the 62 sources are unidentified. Out of the 86 sources 31 have one or two infrared counterpart candidates in the 2MASS catalog. For all the sources, when possible, the hardness ratio was used to assess the low-mass or high-mass XRB properties of the objects. At least 29 unidentified objects fulfill these characteristics. Using the 2MASS counterpart of well-known MQs, we found that MQs with a high-mass nature follow a lineal behavior in the infrared color-color diagram, while LMXB MQs display a more random distribution. Based on this information, we suggest a possible region in the infrared diagram for new MQ candidates. As a result, we found 3 objects that we propose as likely galactic HMXBs and MQ candidates. A follow-up study of these three sources will be reported on a future paper (Combi et al. 2007). Finally, it is important to mention that the whole analysis carried out here only covers a minute fraction of the expected MQ candidates. The NVSS survey does not cover all of the sky. It is complete up to $-40^{\circ}$ in declination. In addition, the current XMM catalog only represents $\sim$ 2\% of the galactic plane ($|$b$|$ $\leq$ 10$^{\circ}$). Thus, a naive extrapolation of the 3 MQ candidates to the whole Galactic plane suggests that about 120 new candidates could be expected in the future. | 7 | 10 | 0710.5831 |
0710 | 0710.0139_arXiv.txt | Recent observations of the Galactic center revealed a nuclear disk of young OB stars, in addition to many similar outlying stars with higher eccentricities and/or high inclinations relative to the disk (some of them possibly belonging to a second disk). Binaries in such nuclear disks, if they exist in non-negligible fractions, could have a major role in the evolution of the disks through binary heating of this stellar system. We suggest that interactions with/in binaries may explain some (or all) of the observed outlying young stars in the Galactic center. Such stars could have been formed in a disk, and later on kicked out from it through binary related interactions, similar to ejection of high velocity runaway OB stars in young clusters throughout the galaxy. | 7 | 10 | 0710.0139 |
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0710 | 0710.0625_arXiv.txt | In a previous paper, we modeled the oscillations of a thermally-supported (Bonnor-Ebert) sphere as non-radial, linear perturbations following a standard analysis developed for stellar pulsations. The predicted column density variations and molecular spectral line profiles are similar to those observed in the Bok globule B68 suggesting that the motions in some starless cores may be oscillating perturbations on a thermally supported equilibrium structure. However, the linear analysis is unable to address several questions, among them the stability, and lifetime of the perturbations. In this paper we simulate the oscillations using a three-dimensional numerical hydrodynamic code. We find that the oscillations are damped predominantly by non-linear mode-coupling, and the damping time scale is typically many oscillation periods, corresponding to a few million years, and persisting over the inferred lifetime of gobules. | Recent observations suggest that the internal structure of some of the small isolated molecular clouds known as starless cores or Bok globules \citep{Bok1948} is well approximated by a balance of thermal and gravitational forces on which is superposed a pattern of non-radial oscillations \citep{Lada2003}. Previously, the internal oscillations of these starless cores have been noted in simulations of collapse \citep{Hennebelle2003}, magnetized clouds \citep{Galli2005}, modeled as one-dimensional radial oscillations by means of a numerical simulation that includes radiative energy losses \citep{KetoField2005} and as three-dimensional non-radial perturbations assuming isothermal conditions \citep{Keto2006}. Analysis of these models and comparison of simulated molecular spectral line profiles with those observed in the dark clouds L1544 \citep{Caselli2002a} and B68 \citep{Lada2003} indicate that the motions in some dark clouds are consistent with large-scale hydrodynamic oscillations, i.e., sound waves. Nevertheless, a number of questions remain as yet unanswered. Is the underlying assumption of the oscillations as perturbations valid for the amplitudes required to produce the observed velocities? Based on the observations of B68, the perturbation amplitudes in our previous three-dimensional model were $\sim25$\%. Does the linear model correctly predict the velocity field for such large amplitudes or would significant differences appear in the non-linear regime? What is the lifetime of the perturbations and thus the velocities within the clouds? For hydrodynamic oscillations to remain a viable explanation of the velocity field observed in B68 they must have lifetimes in excess of a crossing time. Supersonic oscillations are strongly damped via shocks. However, in the subsonic regime the dominant damping mechanism is less clear. In the one-dimensional model \citep{KetoField2005}, the dissipation of the oscillations occurs via radiative losses with an estimated timescale of $10\,\Myr$. In the non-radial perturbative model \citep{Keto2006} damping is explicitly ignored. For mode amplitudes larger than linear perturbations, non-linear mode-mode coupling provides an additional potentially significant damping mechanism. Are the oscillations stable or growing? Are there preferred modes? In general we expect the oscillations to damp over time, and we expect that the longest-lived modes of oscillation would be the long wavelength modes. Neither the one-dimensional non-linear analysis nor the three-dimensional linear analysis is able to adequately address this question. In this paper we address these questions by modeling the oscillations of dark clouds with a three-dimensional, non-linear, numerical hydrodynamic simulation. We find that the dominant damping mechanism is indeed non-linear mode coupling, the longest wavelength mode lifetimes are on the order of a few $\Myr$ and produce velocity and density fields qualitatively similar to those observed. Section \ref{CM} summarizes the numerical methods employed, sections \ref{QP} and \ref{DP} discuss the viability of quadrupole and dipole oscillation models, and the conclusions are contained in section \ref{C}. | \label{C} Despite their strongly non-linear nature, large-amplitude oscillations of pressure-supported molecular clouds reproduce many of the qualitative features of observations of starless cores. These include the complex velocity structures and highly distorted column densities found in B68 and L1544, as well as mimicking rotation, expansion, and contraction depending upon the mode structure and viewing angle. This may complicate efforts to discern the global dynamical state of these objects from observations of molecular line shapes alone. However, it may be possible to distinguish oscillations from collapse by comparing the spectral line profiles and velocities of volatile molecules such as N2H+ and NH3 against those of more refractory species such as CO, HCO+, or CS \citep{KetoField2005}. In a simple breathing mode oscillation, the gas velocities are highest towards the outer regions of the cloud where the abundance of the carbon species is not depleted by freeze-out. In contrast, in the gravitational collapse of unstable BE spheres, the gas velocities are highest in the center of the cloud where the emission from the non-depleting nitrogen molecules is highest. We note that we have restricted our attention to hydrodynamic oscillations of pressure-bounded, thermally-supported spheres. This description seems appropriate for globules such as B68 that are surrounded by hotter gas and have density profiles that closely match those of Bonnor-Ebert spheres. Whether the Bonnor-Ebert description is appropriate for embedded cores remains to be determined. For example, if magnetic stresses contribute significantly to cloud support, then additional classes of oscillations exist, associated with the magnetic field \citep[see,\eg,][]{Hennebelle2003,Galli2005}. The lifetimes of hydrodynamic oscillations of starless cores like B68 are sufficiently long (few $10^6$ yr) to be a significant fraction of the expected lifetime of a molecular cloud ($10^7$ yr). The oscillation lifetimes are dominantly limited by non-linear mode-mode coupling. Nevertheless, large amplitude oscillations persist for many periods (and thus many sound crossing times), and are therefore not expected to be particularly rare. There are no signs of cloud fragmentation, and thus amplitudes comparable to those discussed here are insufficient to produce binaries without the inclusion of additional physical processes. | 7 | 10 | 0710.0625 |
0710 | 0710.4094_arXiv.txt | A subgroup of dwarf galaxies have characteristics of a possible evolutionary transition between star-forming systems and dwarf ellipticals. These systems host significant starbursts in combination with smooth, elliptical outer envelopes and small HI content; they are low on gas and unlikely to sustain high star formation rates over significant cosmic time spans. We explore possible origins of such starburst ``transition" dwarfs using moderately deep optical images. While galaxy-galaxy interactions could produce these galaxies, no optical evidence exists for tidal debris or other outer disturbances, and they also lack nearby giant neighbors which could supply recent perturbations. Colors of the outer regions indicate that star formation ceased $>$1~Gyr in the past, a longer time span than can be reasonably associated with the current starbursts. We consider mechanisms where the starbursts are tied either to interactions with other dwarfs or to the state of the interstellar medium, and discuss the possibility of episodic star formation events associated with gas heating and cooling in low specific angular momentum galaxies. | The formation and evolution of dwarf galaxies remains poorly understood. How these processes produce both rotationally supported (dIrr) as well as dynamically warm (dE/dS0s) dwarf systems remain puzzles, despite considerable effort. An outstanding question in dwarf galaxy evolution is whether there is an evolutionary connection between the various morphological classes of dwarf galaxies (e.g. van den Bergh 1977, van Zee et al. 1998, van Zee et al. 2001, Lisker et al. 2007). Specifically, can dwarfs of one morphological type evolve by some process into a different type? For example, do some dIrrs evolve into dE/dS0s or did the dE/dS0s we see in the Universe originally form as that type? These questions regarding the formation and evolution of dwarf galaxies are of particular importance in light of hierarchical cold dark matter (CDM) galaxy formation models which predict large numbers of dwarf galaxies and suggests that larger galaxies are built up from the accretion of many low mass halos (e.g. White \& Frenk, 1991). A related issue concerns the impact of star formation on dwarf galaxy evolution and particularly the role of episodic star formation (e.g. Searle, Sargent \& Bagnuollo 1973, Lee et al. 2002). How does star formation relate to morphology in low mass galaxies? For example, where do blue compact dwarfs (BCDs) fit in the dwarf galaxy zoo as starbursting objects and are they related to either dIrrs or dE/dS0s (e.g. Sung et al. 1998, Gil de Paz \& Madore 2005, Noeske et al. 2005)? In particular, it has been suggested that BCD galaxies may evolve into dE/dS0 galaxies after they lose their gas, either through supernovae-driven winds during episodes of intense star formation (Marlowe et al. 1999, Tajiri \& Kamaya 2002), or through stripping processes induced by galaxy-galaxy interactions (Drinkwater \& Hardy 1991). Or, are some of these objects dEs where star formation has been renewed through gas capture as suggested by Silk, Wyse, \& Shields (1987)? Previous studies, which have concentrated mainly on optical morphological differences between dwarf galaxies and to some extent on their neutral gas content, have not firmly established a direct evolutionary connection between dIs and dEs in galaxy groups and other low density environments (e.g. van Zee, Salzer \& Skillman 2001, Noeske et al. 2005). Even in the Local Group the situation is unclear. Many Local Group dwarfs have complex star formation histories and kinematic peculiarities, and at least five of the faint dwarfs possess optical morphologies which place them in a ``transition" or ``mixed-type" structural category between dI and dE/dS0 (e.g. Mateo 1998, Grebel, Gallagher, \& Harbeck 2003). While nearby dE galaxies, such as NGC~185 and NGC~205 in the local group, support star formation, the rates are extremely low and consistent with these objects' small gas supplies. More distant galaxy samples also contain dwarfs with early-type structures that also host HI gas and often also some young stars. As in the Local Group, these transition dwarfs frequently are small, low luminosity objects containing $\lesssim$ 10$^6$ M$_{\odot}$ of HI and having correspondingly low star formation rates (e.g. Bouchard et al. 2005). On the other hand, the Virgo Cluster of galaxies appears to contain a complete sequence of more luminous dwarfs with declining and increasingly concentrated star formation, extending from typical BCDs through to blue core dEs (Gallagher \& Hunter 1989, Lisker et al. 2006,2007). While moderate luminosity galaxies with early-type outer structures and young stellar populations in their centers also exist in less dense environments (e.g. NGC~404, del Rio et al. 2004; NGC~5102; Deharveng et al. 1997), they are sufficiently rare that it is difficult to determine if evolutionary sequences exist. In this paper we explore possible histories for a sample of actively star forming luminous dwarf galaxies with transitional properties between BCDs and dEs residing in loose group environments. These systems are fairly isolated and show little indication of a recent interaction or merger. They have an intriguing combination of characteristics which do not allow them to easily be categorized as either dI/BCD or dE systems, but rather show a mixture of the two. The key properties which indicate these low luminosity galaxies {\it may} be in the midst of an evolutionary transition from dI/BCD to dE include: \begin{itemize} \item They are actively forming stars with star formation rates between 0.1 and 1 M$_{\odot}$/yr. \item The star formation is currently centrally concentrated, with the outer regions composed of an older stellar population. \item Although actively forming stars, the starburst is fueled by very little HI gas, with M$_{HI}/L_B \lesssim$\ 0.1. This compares with M$_{HI}/L_B >$ 0.2 for typical dIs (Roberts \& Haynes 1994) and BCDs (Pustilnik et al. 2002) and M$_{HI}/L_B <$ 0.1 for dEs. Star formation can continue in our sample galaxies only for about another 10$^9$ years at their current rate, based on their HI content; in a few cases, including molecular gas will at most double this time (e.g. Jackson et al. 1989, Gordon 1991). \item Unlike many other star forming dwarf galaxies including BCDs they have high oxygen abundance ratios, with 12 + log(O/H) $>$ 8.4 as we found in Dellenbusch et al. (2007). \item Their outer optical colors are similar to those of typical BCDs with (B-R)$_0$ of order 1 (Gil de Paz et al. 2005), but bluer than the (B-R)$_0$ $\approx$ 1.3 - 1.4 expected for slightly fainter dEs (Conselice et al. 2003). In this regard our sample resembles the Gil de Paz et al. "E-type" BCDs. \item These objects have smooth outer isophotes which are much more indicative of early- rather than late-type galaxy structures. \end{itemize} Table~\ref{properties} quantifies several of these properties. In this paper we examine the significance of these final two characteristics. In \S 2 we discuss the observations and data reduction. In \S 3 we describe our analysis and modeling processes used to examine the data for evidence of tidal debris and fine structures. In this section we also describe the results of this analysis. We discuss two possible evolutionary pictures to explain the galaxies' star formation histories as a means of transition in \S 4. \S 5 contains a brief summary and our conclusions. | In summary, we have presented results of deep R-band imaging for five starbursting transition dwarf galaxies. Our observational results show: \begin{enumerate} \item All five galaxies exhibit smooth elliptical outer isophotes. \item We find no evidence of extended tidal debris or other indications of a recent major interaction (e.g. tails, shells or ripples). \item Fine structure exists in the central regions that are associated with the ISM and star formation, with dust structures and HII regions being prominent in several of the galaxies. \item Outer envelope colors are consistent with having no star formation in the outer regions for the past several Gyr. This timescale could be longer if star formation slowly declined rather than stopping suddenly. \end{enumerate} We consider two possible classes of evolutionary histories which could be consistent with the observed properties and locations of transition dwarf galaxies: an internal ISM instability and a past interaction or merger event. The interaction or dwarf-dwarf merger scenario remains a possibility, although we find no strong evidence to support it. After several billion years it is not clear if we should still find evidence of tidal debris in the outer regions of the galaxies. However, long term products of dwarf interactions and dwarf-dwarf mergers in particular have not been well modeled (e.g. the effects of large dark matter contents on merger products). The other possibility of an unstable ISM seems promising to explain the central star formation in these objects. In this case the eventual evolution to an inactive state will be slower, possibly requiring several starburst cycles. Observationally the model can be tested by looking for the young stellar populations in fading starburst phases in dE-like dwarfs. A remaining question, which is particularly puzzling, is the issue of timescales. If these objects last experienced significant star formation in their outer regions $\gtrsim$ 1 Gyr ago, why are they currently undergoing bursts of centrally concentrated star formation? What mechanisms control the structure of the ISM such that these galaxies now support rapid star formation, and will these circumstances lead to galaxies with little or no star formation, i.e. objects resembling dEs, within the next $\sim$ 1 Gyr? | 7 | 10 | 0710.4094 |
0710 | 0710.3161_arXiv.txt | For a density that is not too sharply peaked towards the center, the local tidal field becomes compressive in all three directions. Available gas can then collapse and form a cluster of stars in the center, including or even being dominated by a central black hole. We show that for a wide range of (deprojected) S\'ersic profiles in a spherical potential, the tidal forces are compressive within a region which encloses most of the corresponding light of observed nuclear clusters in both late-type and early-type galaxies. In such models, tidal forces become disruptive nearly everywhere for relatively large S\'ersic indices $n \ga 3.5$. We also show that the mass of a central massive object (CMO) required to remove all radial compressive tidal forces scales linearly with the mass of the host galaxy. If CMOs formed in (progenitor) galaxies with $n \sim 1$, we predict a mass fraction of $\sim 0.1-0.5$\,\%, consistent with observations of nuclear clusters and super-massive black holes. While we find that tidal compression possibly drives the formation of CMOs in galaxies, beyond the central regions and on larger scales in clusters disruptive tidal forces might contribute to prevent gas from cooling. | \label{sec:intro} It is now well-known that the masses of supermassive black holes (SMBHs) in the centres of galaxies and bulges correlate with the stellar velocity dispersion, $\Mbh \propto \sigma_\star^{\alpha}$ with $\alpha \sim 4-5$ \citep[e.g.][]{Ferrarese+01, Gebhardt+01, Tremaine+02}, as well as nearly linearly with the mass of these spheroids, $\Mbh \propto \Msph^{1.12\,\pm\,0.06}$ \citep[e.g.][]{MF01, HaringRix04}. \cite{Ferrarese+06}, \cite{WH06} and \cite{Rossa+06} also found that the masses of nuclear (star) clusters (NCs), which are present in many both late and early-type galaxies \citep[see, e.g.][]{Boeker+02, Cote+06}, are similarly related to the properties of the host galaxy \citep[see also][]{GD07}. Recently, \cite{McLaughlin+06} proposed momentum feedback, from accretion onto SMBHs or from stellar and supernovae winds in the case of NCs, as an explanation for the observed relations. We investigate if central massive objects (CMOs), in the form of NCs, may have formed from gas being tidally compressed in the centers of galaxies. This effect, resembling compressive shocking of globular clusters by the Galactic disk \citep{Ostriker+72}, has been studied by \cite{Valluri93} in the context of tidal compression of a (disk) galaxy in the core of galaxy cluster. At the scale of galaxies, \cite{DasJog99} argue for tidal compression of molecular clouds in the centers of flat-core early-type galaxies and ultraluminous galaxies as an explanation for the presence of observed dense gas. Very recently, an independent study by \cite{Masi07} emphasised the potential importance of compressive tidal forces. Low luminosity early-type galaxies and late-type galaxies share an overall luminosity profile with relatively low central power-law slopes. The fact that NCs are predominantly found in such galaxies \citep[see e.g.][]{Cote+06} may provide an interesting link between the presence of CMOs in galaxies and the properties of the host galaxy. In the present study, we investigate whether tidal forces may help explaining this link. We first derive the radial component of the tidal force associated with a (deprojected) S\'ersic profile in \S~\ref{sec:tidalcompression}. We then examine in \S~\ref{sec:scaling} how this applies to simple models of CMO hosts, including early-type and late-type galaxies. The corresponding results are then briefly discussed in \S~\ref{sec:disc}, and conclusions are drawn in \S~\ref{sec:concl}. | \label{sec:concl} We have built simple spherical models following deprojected S\'ersic profiles to show that compressive tidal forces are naturally present in the central region when the S\'ersic index $n \la 3.5$. For $n \ga 3.5$, the radial component of the tidal forces is disruptive almost everywhere (i.e., for $r / R_e > 10^{-3}$). Observed nuclear (star) clusters in early-type and late-type galaxies have extents and/or apparent locations which are within the tidally compressed region. If we assume that most NCs form when their host galaxies have density profiles corresponding to rather low S\'ersic indices $n \sim 1$, we have shown that the masses of the NCs are consistent with $M_+$, the mass above which the radial tidal force becomes disruptive due the presence of the central massive object. In this picture, the predicted $M_+$ scales linearly with the host galaxy mass (or the mass of the spheroidal component) with $M_+/\Msph \sim 0.1-0.5$\,\% for $n \sim 1$, in agreement with what is observed for both NCs and super-massive black holes in the centers of (more luminous) galaxies. If indeed compressive tidal forces are a key actor in the formation of CMOs, only late-type galaxies have, today, the required gas content and density profiles ($n \sim 1$), which allow the recurrent and common formation of CMOs (in the form of NCs). This is consistent with the fact that young NCs are predominantly found in late-type spiral galaxies. Finally, while we find that tidal compression possibly drives the formation of CMOs in galaxies, beyond the central regions and on larger scales in clusters disruptive tidal forces might contribute to prevent gas from cooling. Such a simple scenario must be tested and extended to accomodate galaxies with e.g., core-Sersic surface brightness profiles \citep[see e.g.][]{Ferrarese+06b} as well as to allow more realistic (non-spherical, multi-component) galaxy models. Moreover, using specific (stellar) mass-to-light ratios for the NCs and virial mass estimates for the host galaxy enables a direct comparison in mass instead of luminosity. Finally, hydrodynamical simulations are needed to examine the role of compressive tidal forces in the evolution of galaxies (and cluster). | 7 | 10 | 0710.3161 |
0710 | 0710.3999_arXiv.txt | {G39$-$27/289 is a common proper motion pair formed by a white dwarf (WD0433$+$270) and a main-sequence star (BD$+$26~730) that apparently has been classified as a member of the Hyades open cluster. Previous studies of the white dwarf component yielded a cooling time of $\sim4$ Gyr. Although it has not been pointed out before explicitly, this result is 6 times larger than the age of the Hyades cluster, giving rise to an apparent conflict between the physics of white dwarfs and cluster main-sequence fitting.} {We investigate whether this system belongs to the Hyades cluster and, accordingly, give a plausible explanation to the nature of the white dwarf member.} {We have performed and analyzed spectroscopic observations to better characterize these objects, and used their kinematic properties to evaluate their membership to the Hyades. Then, different mass-radius relations and cooling sequences for different core compositions (He, C/O, O/Ne and Fe) have been employed to infer the mass and cooling time of the white dwarf.} {From kinematic and chemical composition considerations we believe that the system was a former member of the Hyades cluster and therefore has an evolutionary link with it. However, the evidence is not conclusive. With regards to the nature of the white dwarf component, we find that two core compositions --- C/O and Fe --- are compatible with the observed effective temperature and radius. These compositions yield very different cooling times of $\sim$4 Gyr and $\sim$1 Gyr, respectively.} {We distinguish two posssible scenarios. If the pair does not belong to the Hyades cluster but only to the Hyades stream, this would indicate that such stream contains rather old objects and definitely not coeval with the cluster. This has interesting consequences for Galactic dynamics. However, our favoured scenario is that of a white dwarf with a rather exotic Fe core, having a cooling time compatible with the Hyades age. This is a tantalizing result that would have implications for the thermonuclear explosion of white dwarfs and explosion theories of degenerate nuclei.} | The determination of reliable ages is of obvious importance to both astrophysics and cosmology, but not exempt of many complications --- see, e.g., \cite{von01}. The age of a star is perhaps the most difficult property to estimate and, further, it nearly always depends on rather strong model assumptions (Mamajek et al.~2007). An exception is the method that uses radioactive decay to directly estimate stellar ages (Cayrel et al.~2001; Frebel et al.~2007), but this has rather restricted applicability. Another emerging technique is asteroseismology, which has the potential to provide accurate ages (Miglio \& Montalb\'an 2005), but this will require high-precision photometric data from upcoming space missions and also stars with well-constrained physical properties. In the meantime, the use of open clusters and main sequence stellar evolutionary models continues to be the most widely used method to infer the ages of stars in the Galaxy. In spite of the still-present uncertainties, such as convective overshoot, chemical composition anomalies,\ldots~the field has reached sufficient maturity to provide ages that are reliable to better than $\sim$10\% --- see, for instance, \cite{pau06}. The analogous method of using eclipsing binaries (Ribas et al.~2000) yields similar (model-dependent) accuracy. The study of white dwarfs has made very valuable contributions to numerous areas of astrophysics, and estimating individual ages of stars is no exception. The main advantage of white dwarfs is the conceptual simplicity of their evolution, which can be described as a simple cooling process. Modelling of the cooling sequences makes white dwarfs powerful stellar chronometers, accurate to about 25\% (Silvestri et al.~2005). Among the major uncertainties of this method are the pre-white dwarf evolution time and the effects of the chemical composition of the core. \begin{table*}[t] \begin{center} \caption{Photometric data and stellar parameters derived for BD$+$26~730.} \small{ \begin{tabular}{lccccccc} \hline \hline \noalign{\smallskip} $V$ & $J$ & $H$ & $K$ & $T_{\rm eff}$ (K) & [Fe/H] & $Z$ & $\log(L/L_{\sun})$ \\ \noalign{\smallskip} \hline \noalign{\smallskip} 8.42$\pm$0.02 & 5.945$\pm$0.023 & 5.400$\pm$0.018 & 5.240$\pm$0.023 & $4,595\pm30$ & $0.03\pm0.09$ & $0.021\pm0.004$ & $-0.527\pm0.021$ \\ \noalign{\smallskip} \hline \hline \end{tabular} \label{tab:spar}} \end{center} \end{table*} In this paper we discuss a case that presents an interesting puzzle in which the age estimate stemming from the white dwarf cooling sequence is in apparent conflict with the age estimate coming from cluster membership. The object is G39$-$27/289, a common proper motion pair formed by a DA white dwarf (WD0433$+$270) and a K type star --- BD$+$26~730 (Holberg et al.~2002) --- which is a well studied variable star (V833 Tau) and also a single-lined spectroscopic binary with a very low-mass companion (Tokovinin et al.~2006). The members of a common proper motion pair were likely born simultaneously and with the same chemical composition (Wegner 1973; Oswalt et al.~1988). Since the components are well separated ($\sim 2,200$ AU in this case), no mass exchange has taken place and they have evolved as isolated stars. Thus, it is logical to assume that both components of G39$-$27/289 have the same age and the same original metallicity. The K-star component, and, by extension, its common proper motion companion, were classified as Hyades members by \cite{per98} and therefore would have an age ranging from about 0.6 to 0.7 Gyr. However, this stands in obvious conflict with the cooling time of WD0433$+$270, which has been estimated to be about 4 Gyr (Bergeron et al.~2001). In this work, we carry out a detailed study of these objects, using both information present in the literature and also our own spectroscopic observations, with the objective of unveiling their nature and their possible membership to the Hyades open cluster. We evaluate the different scenarios and discuss their respective implications to white dwarf physics and Galactic dynamics. | In this work we have presented new spectroscopic observations of the members of the common proper motion pair G39$-$27/289. Using these observations and the available photometry we have better characterized these objects, which has helped us to understand their nature. Considering the kinematic properties of the members of the pair and the lithium detection in BD$+$26~730, we favour the scenario in which the common proper motion pair is indeed a former member of the Hyades cluster, and thus its members have a coeval age of $\sim$0.6--0.7 Gyr. Having evaluated different compositions for WD0433$+$270, the young age inferred from cluster membership is only compatible with the case of an Fe core, which would have an associated cooling time of $\sim$1 Gyr. The agreement is not perfect, but the modelling of the cooling evolution of Fe-core white dwarfs could still have some associated uncertainties. The existence of white dwarfs with an Fe-rich core has important consequences for the models of thermonuclear explosion of stellar degenerate cores. According to the theory of stellar evolution, all stars that develop an Fe core experience a collapse to a neutron star or a black hole, regardless of the mass loss rate assumed. However, there is still an alternative possibility to avoid this final fate that relies in the failure of the thermonuclear explosion of a degenerate white dwarf near the Chandrasekhar limit (Isern et al.~1991). Our current view of a thermonuclear explosion is as follows. Once the thermonuclear runaway starts in the central regions of a white dwarf, the ignition front propagates outwards and injects energy at a rate that depends on the velocity at which matter is effectively burned leading to the expansion of the star. At the same time, electron captures on the incinerated matter efficiently remove energy at a rate determined by the density causing the contraction of the star. Therefore, depending on the ignition density and the velocity of the burning front, the outcome can be different. It is known that He-degenerate cores always explode, that O/Ne and Fe degenerate cores always collapse and that C/O-degenerate cores can explode or collapse depending on the ignition density. The existence of Fe-rich white dwarfs would imply the possibility of an intermediate behaviour between those discussed above in which an Fe remnant is left after a mild explosion. A detailed theory explaining the formation of Fe-core white dwarfs is still to be developed and up to now the possibility of their existence has been suggested mainly from observational evidences. Other examples of possible Fe-core white dwarfs have been found in the past. \cite{pro98} singled out two objects whose radii and masses were too small to fit the typical C/O mass-radius relation. These authors analyzed a sample of white dwarfs, calculated their masses using either orbital parameters, gravitational redshifts or spectroscopic fits, and determined their radii independently from the knowledge of effective temperatures and distances. \cite{pro98} indicated that in those two cases, a core made by Fe was the only plausible explanation that could account for their small radii. Besides the Fe-core hypothesis, there is an alternative scenario that we cannot rule out. Although we think it is unlikely, the common proper motion pair might not be related in any way to the Hyades cluster. This would eliminate the age constraint and thus relax the requirement of an exotic composition for the white dwarf core. In this case, the age of the objects (both WD0433$+$270 and BD$+$26~730) could be of $\sim$4 Gyr. This would have some interesting consequences to our current view of star streams or moving groups in th Galaxy. As pointed out by \cite{che99}, the Hyades stream contains at least three distinct age groups with ages up to 2 Gyr. The upper limit was probably a consequence of the use of A--F stars as kinematic tracers, which would naturally have an age cutoff at about 2--3 Gyr. With an age of 4 Gyr, WD0433$+$270 and BD$+$26~730 would indicate the existence of an even older population in the Hyades stream thus completely ruling out any coevality within the members. This would lend strong support to the model of a resonant origin for the Hyades stream (Famaey et al.~2007). Another implication of such an old age is the conflict with lithium detection, which would imply much lower destruction rate than expected. Definitive proof supporting one of these scenarios will have to await an unambiguous determination of the mass of this object or, alternatively, a more conclusive study on its evolutionary link with the Hyades cluster. | 7 | 10 | 0710.3999 |
0710 | 0710.1678_arXiv.txt | Numerous cosmological hydrodynamic studies have addressed the formation of galaxies. Here we choose to study the first stages of galaxy formation, including non-equilibrium atomic primordial gas cooling, gravity and hydrodynamics. Using initial conditions appropriate for the concordance cosmological model of structure formation, we perform two adaptive mesh refinement simulations of $\sim$$10^8 \Ms$ galaxies at high redshift. The calculations resolve the Jeans length at all times with more than 16 cells and capture over 14 orders of magnitude in length scales. In both cases, the dense, $10^5$ solar mass, one parsec central regions are found to contract rapidly and have turbulent Mach numbers up to 4. Despite the ever decreasing Jeans length of the isothermal gas, we only find one site of fragmentation during the collapse. However, rotational secular bar instabilities transport angular momentum outwards in the central parsec as the gas continues to collapse and lead to multiple nested unstable fragments with decreasing masses down to sub-Jupiter mass scales. Although these numerical experiments neglect star formation and feedback, they clearly highlight the physics of turbulence in gravitationally collapsing gas. The angular momentum segregation seen in our calculations plays an important role in theories that form supermassive black holes from gaseous collapse. | Since the first investigations of galaxy interactions \citep{Holmberg41} using light bulbs, the use of numerical simulations in galaxy formation has developed dramatically. Not only gravity but also hydrodynamics and cooling are standard ingredients in the sophisticated computer models studying galaxy formation and interactions. In hierarchical structure formation, dark matter (DM) halos merge to form larger halos while the gas infalls into these potential wells \citep{Peebles68, White78}. \citeauthor{White78} provided the basis for modern galaxy formation, in which small galaxies form early and continuously merge into larger systems. As more high redshift galaxies were observed in the following 10 years, \citet{White91} refined the theory to address the observed characteristics in these galaxies. In their model, the halo accumulates mass until the gas cools faster than a Hubble time, \tH, which usually occurs when atomic hydrogen line, specifically \lya, cooling is efficient. This happens when the halo has \tvir~$>$~10$^4$ K, where the cooling function sharply rises by several orders of magnitude because the number of free electrons able to excite hydrogen greatly increases at this temperature \citep{Spitzer78}. One can define a cooling radius, \rcool, in which the interior material is able to cool within a Hubble time. Once the halo reaches this first milestone, \rcool~ increases through additional accretion and cooling. A rapid baryonic collapse ensues when \tcool~$\lsim$~\tdyn~ \citep{Rees77}. The material accelerates towards the center, and its density quickly increases. In the model discussed in White \& Frenk, this collapse will halt when one of the following circumstances occurs. First, angular momentum can prevent the gas from collapsing further, and the system becomes rotationally supported. Afterwards, this disc fragments and star formation follows. Alternatively, star formation does not necessarily develop in a disc component, but the energy released by stars during their main sequence and associated supernovae (SNe) terminates the collapse. These concepts have been applied also to the earliest galaxies in the universe \citep{Mo98, Oh02, Begelman06, Lodato06}. Many studies \citep[e.g.][]{Ostriker96, Haiman97b, Cen03, Somerville03, Wise05} demonstrated that OB-stars within protogalaxies at $z > 6$ can produce the majority of photons required for reionization. These protogalaxies contain an ample gas reservoir for widespread star formation, and the accompanying radiation propagates into and ionizes the surrounding neutral intergalactic medium. Several high redshift starburst galaxies have been observed that host ubiquitous star formation at $z > 6$ \citep{Stanway03, Mobasher05, Bouwens06}. Additionally, supermassive black holes (SMBH) more massive than 10$^8 \Ms$ are present at these redshifts \citep[e.g.][]{Becker01, Fan02, Fan06}. Finally, a reionization signature in the polarization of the cosmic microwave background (CMB) at z $\sim$ 10 \citep{Page07} further supports and constrains stellar and SMBH activity at high redshifts. The distinction between SMBH formation and a starburst galaxy should depend on the initial ingredients (i.e. seed BHs, metallicity, merger histories) of the host halo, but the evolution of various initial states is debatable. It is essential to study the hydrodynamics of high redshift halo collapses because the initial luminous object(s) that emerges will dynamically and thermally alter its surroundings. For example, as the object emits ultraviolet radiation, the nearby gas heats and thus the characteristic Jeans mass increases, which may inhibit the accretion of new gas for future star formation \citep{Efstathiou92, Thoul96}. The following work will attempt to clarify early galaxy formation by focusing on protogalactic (\tvir~$>10^4$ K) halos and following their initial gaseous collapse. \citet[][hereafter Paper I]{Wise07a} studied the virialization of protogalactic halos and the virial generation of supersonic turbulence. In this paper, we address the gas dynamics of the continued, turbulent collapse of a halo and study the evolution and characteristics of the central regions. In later studies, we will introduce the effects from primordial star formation and feedback and \hh~cooling. The progressive introduction of new processes is essential to understand the relevance of each mechanism. We argue that our results may be relevant for scenarios that envisage SMBH formation from gaseous collapses. \citet{Loeb94} and \citet{Bromm03} conducted smoothed particle hydrodynamics (SPH) simulations that focused on the collapse of idealized, isolated protogalactic halos. The former group concluded that a central $10^6 \Ms$ SMBH must exist to stabilize the thin gaseous disc that forms in their calculations. \citeauthor{Bromm03} considered cases with and without \hh~chemistry and a background UV radiation field. They observed the formation of a dense object with a mass $M \sim 10^6 \Ms$, or $\gsim 10\%$ of the baryonic matter, in simulations with no or strongly suppressed \hh~formation. These calculations without metal cooling and stellar feedback are useful to explore the hydrodynamics of the collapse under simplified conditions. \citet{Spaans06} analytically studied the collapse of 10$^4$ K halos with an atomic equation of state. They find that $\sim$0.1\% of the baryonic mass results in a pre-galactic BH with a mass $\sim$$10^5 \Ms$. \citet{Lodato06} also found that $\sim$5\% of the gas mass in $M = 10^7 \Ms$ halos at $z \sim 10$ becomes unstable in a gaseous disc and forms a SMBH. Recently, \citet{Clark07} studied the effects of metal and dust cooling on the fragmentation of a collapsing protogalactic core with varying metallicities ($Z = 0, 10^{-6}, 10^{-5} Z_\odot$) and found the gas fragmenting ten times as much in the $10^{-5} Z_\odot$ case than the primordial case. In addition, the fragments in the primordial case are biased toward larger masses. A runaway gaseous collapse requires angular momentum transport so material can inflow to small scales and form a central object. The stability of rotating gaseous clouds have been subject of much interest over the last four decades and was thoroughly detailed by the work of \citet[][hereafter EFE]{Chandra69}. In the 1960's and 1970's, studies utilizing virial tensor techniques \citep[EFE;][]{Lebovitz67, Ostriker69, Ostriker73a}, variational techniques \citep{LyndenBell67, Bardeen77}, and N-body simulations \citep{Ostriker73b} all focused on criteria in which a stellar or gaseous system becomes secularly or dynamically unstable. The first instability encountered is an $m = 2$ bar-like instability that is conducive for angular momentum transport in order to form a dense, central object. \citet{Begelman06} investigated the conditions where a gaseous disc in a pre-galactic halo would become rotationally unstable to bar formation \citep[see][]{Christodoulou95a, Christodoulou95b}. They adapt the ``bars within bars'' scenario \citep{Shlosman89, Shlosman90}, which was originally formulated to drive SMBH accretion from a gaseous bar that forms within a stellar galactic bar, to the scenario of pre-galactic BH formation. Here a cascade of bars form and transport angular momentum outwards, and the system can collapse to small scales to form a quasistar with runaway neutrino cooling, resulting in a central SMBH. The simulations detailed below show how many central bar-like instabilities form. In \S2, we describe our simulations and their cosmological context. In the following section, we present our analysis of the halo collapse simulations and investigate the structural and hydrodynamical evolution, the initial halo collapse, rotational instabilities, and the importance of turbulence. In \S4, we discuss the relevance of angular momentum transport and rotational instabilities in early galaxy and SMBH formation. There we also examine the applicability and limitations of our results and desired improvements for future simulations. Finally we conclude in the last section. \begin{figure*} \resizebox{\textwidth}{!}{\rotatebox{0}{\includegraphics*{f1a}}} \resizebox{\textwidth}{!}{\rotatebox{0}{\includegraphics*{f1b}}} \caption{An overview of the final state of the collapsing protogalactic gas cloud. Slices of log gas density in cm$^{-3}$ are shown through the densest point in the halo. The \textit{first} and \textit{three} rows show simulation A, and the \textit{second} and \textit{fourth} rows show simulation B. The columns in the top two rows from left to right are slices with a field of view of 10 kpc, 1 kpc, 100 pc, and 1 pc. For the bottom two rows, the fields of view are 0.01pc, 20AU, 0.2AU, and 4 R$_\odot$. Note that each color scale is logarithmic, spans 5 orders of magnitude, and is unique for every length scale.} \label{fig:slices} \end{figure*} | We have simulated the hydrodynamics and collapse of a protogalactic gas cloud in two cosmology AMR realizations. Our focus on the hydrodynamics presents a basis for future studies that consider stellar and BH feedback. In the idealized case presented, we find a central dense object forms on the order of 10$^5 \Ms$ and $r \lsim 5$ pc. This central object is not rotationally supported and does not fragment in our simulations. However our results do not dismiss disc formation in protogalaxies because rotationally supported disc formation may begin after the initial central collapse. Disc formation may be sensitively affected by feedback from the central object. These simulations highlight the relevance of secular bar-like instabilities in galaxy formation and turbulent collapses. Similar bar structures are witnessed in primordial star formation simulations. As low angular momentum infalls, it gains rotational energy as it conserves angular momentum. This induces an $m$ = 2, bar-like instability that transports angular momentum outwards, and the self-similar collapse can proceed without becoming rotationally supported and exhibits a density profile $\rho \propto r^{-12/5}$. This process repeats itself as material infalls to small scales that is indicative of the ``bars within bars'' scenario. We see three and four occurrences of embedded secular instabilities in the two realizations studied here. We also find that supersonic turbulence influences the collapse by providing a channel for the gas to preferentially segregate according to its specific angular momentum. The low angular momentum material sinks to the center and provides the material necessary for a central collapse. Here the possibilities of a central object include a direct collapse into a SMBH \citep[e.g.][]{Bromm03}, a starburst \citep[e.g.][]{Clark07}, or a combination of both \citep[e.g.][]{Silk98}. All of these cases are viable in the early universe, and the occurrence of these cases depends on the merger history, local abundances in the halo, and the existence of a seed BH. Moreover, star formation should occur whether a central BH exists or not. Perhaps the frequency of these different protogalactic outcomes may be traced with either 3D numerical simulations that consider star and SMBH formation and feedback along with metal transport or Monte Carlo merger trees that trace Pop III star formation, metallicities, and BHs. We will attempt the former approach in future studies to investigate protogalactic formation in more realistic detail. | 7 | 10 | 0710.1678 |
0710 | 0710.1352_arXiv.txt | Stellar feedback, expanding {\sc H~ii} regions, wind-blown bubbles, and supernovae are thought to be important triggering mechanisms of star formation. Stellar associations, being hosts of significant numbers of early-type stars, are the loci where these mechanisms act. In this part of our photometric study of the star-forming region NGC~346/N~66 in the Small Magellanic Cloud, we present evidence based on previous and recent detailed studies, that it hosts at least two different events of triggered star formation and we reveal the complexity of its recent star formation history. In our earlier studies of this region (Papers I, III) we find that besides the central part of N~66, where the bright OB stellar content of the association NGC~346 is concentrated, an arc-like nebular feature, north of the association, hosts recent star formation. This feature is characterized by a high concentration of emission-line stars and Young Stellar Objects, as well as embedded sources seen as IR-emission peaks that coincide with young compact clusters of low-mass pre-main sequence stars. All these objects indicate that the northern arc of N~66 encompasses the most current star formation event in the region. We present evidence that this star formation is the product of a different mechanism than that in the general area of the association, and that it is triggered by a wind-driven expanding {\sc H~ii} region (or bubble) blown by a massive supernova progenitor, and possibly other bright stars, a few Myr ago. We propose a scenario according to which this mechanism triggered star formation away from the bar of N~66, while in the bar of N~66 star formation is introduced by the photo-ionizing OB stars of the association itself. | Massive OB stars, not having an optically visible pre-main sequence (PMS) contraction phase, appear almost immediately after their birth on the main sequence (e.g. Stahler \& Palla 2005). They are mostly grouped in stellar associations, loose concentrations of stars, which host also significant numbers of intermediate- and low-mass PMS stars (see review by Brice\~{n}o et al. 2007). When the far-UV radiation of the bright OB stars reaches the surface of the parental molecular cloud, a photo-dissociated region (PDR) develops. Nebula LHA~115-N~66 or in short N~66 (Henize 1956), the brightest {\sc H~ii} region in the Small Magellanic Cloud (SMC), being very rich in early-type stars, is certainly an excellent example of an extragalactic PDR. The stellar association NGC~346, located in the central part of the nebula, hosts the largest sample of spectroscopically confirmed OB stars in the SMC (Massey et al. 1989), which have been the subject of several previous investigations (Niemela et al. 1986; Walborn et al. 2000; Evans et al. 2006). These massive stars evolve rapidly, and immediately start to ionize the cloud, blowing its material away. The degree of ionization decreases outwards, and a thin barrier develops that segregates the ionized from the atomic gas. Molecular hydrogen is not fully ionized behind this front, but partly dissociated, while CO molecules which are located a bit deeper into the cloud are more easily dissociated than H$_{2}$ by absorbing UV photons. A correlation between H$_{2}$ infrared emission and CO lines, characteristic of a PDR, has been found for the region of NGC~346/N~66 by Contursi et al. (2000) and Rubio et al. (2000). These authors define the ``bar'' of N~66 as the oblique bright emission region, extending from southeast to northwest centered on NGC~346. They suggest that {\sl star formation has taken place as a sequential process in the bar of N~66}, resulting in several embedded sources, seen as IR-emission peaks in the 2.14~\micron\ H$_{2}$ line and the ISOCAM LW2 band (5 - 8 \micron). These peaks are alphabetically numbered from ``A'' to ``I'' (Contursi et al. 2000), with the association NGC~346 itself coinciding with peak ``C''. Recent studies reconstruct the star formation history in nearby Galactic OB associations, and provide observational evidence for sequential and triggered star formation in their vicinity. Preibisch \& Zinnecker (2007) conclude from their study of the Scorpius-Centaurus OB association that the formation of whole OB subgroups (each consisting of several thousand stars) requires large-scale triggering mechanisms such as shocks from expanding wind and supernova driven super-bubbles surrounding older subgroups\footnote{According to these authors, other triggering mechanisms, like radiatively driven implosion of globules, also operate, but seem to be secondary processes, forming only small stellar groups rather than whole OB subgroups with thousands of stars.}. Since the low-mass stellar members of associations remain in their PMS phase for \lsim~30~Myr, they play a key r\^{o}le in the understanding of star formation in the vicinity of these systems (Brice\~{n}o et al. 2007). Consequently, the recent discovery of a plethora of low-mass PMS stars in the region of NGC~346/N~66 with photometry from the {\sl Advanced Camera for Surveys} (ACS) onboard HST (Nota et al. 2006; Gouliermis et al. 2006, hereafter Paper~I) can contribute significantly to the clarification of the mechanisms that may act in this extraordinary star-forming region. The subsequent investigation of these PMS stars in NGC~346/N~66 (Sabbi et al. 2007, hereafter SSN07; Hennekemper et al. 2008, hereafter Paper~III) demonstrated that indeed they are clustered in several compact concentrations, some of them coinciding with the IR-emission peaks of Contursi et al. (2000; detected also with {\sl Spitzer}, see \S 2.1), verifying the existence of stellar subgroups in the region, similarly to galactic associations. Not being able to resolve differences in age smaller than 1 - 2 Myr in the individual CMDs of these sub-clusters, SSN07 suggest that ``all sub-clusters appear coeval with each other''. According to these authors, this coevality is a signature of the star formation conditions predicted by the hierarchical fragmentation of a turbulent molecular cloud model (Klessen \& Burkert 2000; Bonnell \& Bate 2002; Bonnell et al. 2003). However, our analysis on the clustering properties of the PMS stars in NGC~346/N~66 (Paper~III) showed that {\sl there are significant age differences between some sub-clusters and the association itself}. Specifically, three compact PMS clusters located to the north of the bar of N~66 are found to be not older than 2.5~Myr, while the CMD of the main body of the association NGC~346 show indications of an underlying older PMS population. This clearly suggests that these northern clusters are probably formed {\sl after} the central stellar association. As we discuss later, triggering mechanisms for star formation such as those described by models of ionization shock fronts from OB stars (e.g. Kessel-Deynet \& Burkert 2003) or wind-driven shock waves (e.g. Vanhala \& Cameron 1998) may explain better their formation. Indeed, while a sequential star formation mechanism has been suggested to take place {\sl in} the bar of N~66, around the association (Rubio et al. 2000), the existence of young compact PMS clusters {\sl away} from it to the north of the association, does not quite fit in the hypothesis that this mechanism propagated from the center of N~66 along the bar. The triggering agent of these clusters may well be located outside of the bar. In this paper we consider the findings from earlier and recent comprehensive investigations of the region NGC~346/N~66 to attempt a clearer understanding of the mechanisms that shape the recent star formation of this outstanding extragalactic star-forming region. Our aim is to provide answers to two important questions: i) {\sl Is the star formation away from the association NGC~346 the product of triggered fragmentation of the cloud alone?} ii) {\sl Is the photo-dissociation by the early-type stars of NGC~346 in the center of N~66 the only triggering mechanism in the region?} In section~2 we present evidence that this region has a far more complicated recent star formation history than what was previously considered, and that the most recent star formation event could not have been triggered by the central association. In section~3 we propose a scenario for the recent star formation in the northern part of the region NGC~346/N~66 and we indicate a nearby massive supernova progenitor, located at the northeast of the bar of N~66, as the triggering agent. We also discuss the possibility that other nearby massive stars away from the association may have also contributed to this mechanism. We present supporting evidence to this scenario using analyses of massive stellar content and gas kinematics of this region. Concluding remarks are given in section~4. | 7 | 10 | 0710.1352 |
|
0710 | 0710.2398_arXiv.txt | {% Differential rotation can be detected in single line profiles of stars rotating more rapidly than about $v\,\sin{i} = 10$\,km\,s$^{-1}$ with the Fourier transform technique. This allows to search for differential rotation in large samples to look for correlations between differential rotation and other stellar parameters. I analyze the fraction of differentially rotating stars as a function of color, rotation, and activity in a large sample of F-type stars. Color and rotation exhibit a correlation with differential rotation in the sense that more stars are rotating differentially in the cooler, less rapidly rotating stars. Effects of rotation and color, however, cannot be disentangled in the underlying sample. No trend with activity is found. } | Stars are born from molecular clouds carrying net angular momentum that makes every star rotate more or less rapidly. During their evolution stars are being braked more or less efficiently, but even after several Gyrs of efficient braking stars as strongly braked as the Sun still show substantial angular velocity. This rotation in general must be expected to be differential, i.e. angular velocity changes with depth and latitude. In a rotating star, it is already the temperature difference due to rotationally induced surface gravity gradients which leads to meridional flow and differential rotation. More severe effects like magnetic forces and inhomogeneities in the convective structure may lead to even more severe effects of differential rotation. The star with the best studied rotation law is the Sun, its equatorial angular velocity is roughly 20\,\% higher than the polar one, which has been observed centuries ago by following the rotation of spot groups on the solar surface. Today it is believed that these spots are due to magnetic activity of which at least the cyclic part is generated by a dynamo mechanism requiring the shear that is caused by differential rotation. Detecting differential rotation on stars other than the Sun unfortunately is not as straightforward as on our host star. Other stars cannot (yet) be spatially resolved so that we cannot just follow the latitude-dependent motion of spot groups -- on some stars such spots might even be absent. Instead, one has to apply indirect techniques. One way to detect differential rotation is the reconstruction of the stellar surface using Doppler Imaging (DI) and comparing the spot's migration at different epochs. This method was applied successfully to many objects, and a summary of the state of the art is presented by Collier Cameron in this volume. DI correlates spot configurations between different epochs that can in principle be temporally separated as long as the lifetime of the spots, i.e. on the order of month in the Sun. Hence DI allows the detection of very small angular velocity differences. This method, however, requires that for each epoch the star is observed for a full rotation period in order to reconstruct a good picture of its surface. Depending on the rotation period and the brightness of the target that requires large amounts of telescope time, which in case of faint targets have to be quite big to ensure the small exposure times necessary for a good resolution of the surface. A different approach to detect differential rotation is to scrutinize the shape of the rotationally broadened line profiles, which yields characteristic differences between rigid and differential rotation. This method was pioneered in the seventies (Gray, 1977). The problem of the degeneracy of differential rotation, limb darkening and inclination was investigated by several authors, (see Reiners \& Schmitt 2002a and references therein). These authors also provide a recipe how to detect differential rotation in stellar line profiles. Since then, differential rotation was searched for in more than a thousand spectra of stars from spectral types A--G. This resulted in many detections of differential rotation. The results of this project are published in Reiners \& Schmitt (2003a, 2003b), Reiners \& Royer (2004), and Reiners (2006). In this paper, I summarize the results concentrating on the dependency of differential rotation on temperature, rotation, and activity. I focus on the F-stars, i.e. I will take into account data from different samples for stars of colors $0.2 < B-V < 0.6$. | I have presented the results from the analysis of several hundred high resolution spectra looking for differential rotation in F-stars. The fraction of stars with signatures of differential rotation in the overall sample of 200 stars is shown as a function of $B-V$ color, projected rotation velocity $v\,\sin{i}$, and normalized X-ray activity $\log{L_\mathrm{X}/L_\mathrm{bol}}$. The fraction of differential rotators shows a clear trend in color and rotation velocity, two parameters which unfortunately are not uncorrelated in the sample. So far the conclusion is that among the cooler, less rapidly rotating objects the fraction of differentially rotating stars is larger than among the hotter, more rapidly rotating stars. No trend is seen in the fraction of differential rotators as a function of activity. From the current sample, it cannot be decided whether slow rotation, low temperature, or both conditions together drive the high fraction of differentially rotating stars. To decide this, observations in a well defined sample containing more slowly rotating early F-stars as well as more rapidly rotating late F-stars are required. | 7 | 10 | 0710.2398 |
0710 | 0710.2351_arXiv.txt | We use a statistical sample of $\sim$500 rich clusters taken from 72 square degrees of the Red-Sequence Cluster Survey (RCS-1) to study the evolution of $\sim$30,000 red-sequence galaxies in clusters over the redshift range 0.35$<$z$<$0.95. We construct red-sequence luminosity functions (RSLFs) for a well-defined, homogeneously selected, richness limited sample. The RSLF at higher redshifts shows a deficit of faint red galaxies (to $M_V\ge$ -19.7) with their numbers increasing towards the present epoch. This is consistent with the `down-sizing` picture in which star-formation ended at earlier times for the most massive (luminous) galaxies and more recently for less massive (fainter) galaxies. We observe a richness dependence to the down-sizing effect in the sense that, at a given redshift, the drop-off of faint red galaxies is greater for poorer (less massive) clusters, suggesting that star-formation ended earlier for galaxies in more massive clusters. The decrease in faint red-sequence galaxies is accompanied by an increase in faint blue galaxies, implying that the process responsible for this evolution of faint galaxies is the termination of star-formation, possibly with little or no need for merging. At the bright end, we also see an increase in the number of blue galaxies with increasing redshift, suggesting that termination of star-formation in higher mass galaxies may also be an important formation mechanism for higher mass ellipticals. By comparing with a low-redshift Abell Cluster sample, we find that the down-sizing trend seen within RCS-1 has continued to the local universe. | Clusters of galaxies are ideal laboratories for studying galaxy evolution since they contain many galaxies seen at the same epoch in close proximity. Their cores are dominated by early-type galaxies, which are the major component of the high mass end of the galaxy stellar mass function locally. There is now a good deal of evidence that cluster early-type galaxies formed the bulk of their stars at high redshift and thereafter simply evolved passively with little or no residual star-formation. One such line of evidence is the tight sequence they form in color--magnitude space \citep[e.g.,][]{visv,bow92}, `the red-sequence'. A similar red-sequence is also seen for early-type galaxies in the field out to at least z$\sim$1 \citep{Bell:2004lb}. Furthermore, {\it all} galaxies appear to be divided into two distinct populations: the passively-evolving red-sequence and the actively star-forming `blue cloud'. Only a small amount of residual star-formation (less than $\sim$10\% of the galaxies' past averaged star-formation rate) is necessary to move a galaxy from the red-sequence to the blue cloud. Therefore, early-type galaxies can provide unique insight into the history of star-formation, as traced by objects in which star-formation has already been terminated. In the local universe, the probability of a galaxy belonging to the red-sequence or blue cloud depends on its stellar mass and its environment \citep{Baldry:2006bq}. It is likely that the other fundamental parameter governing the properties of a galaxy is the epoch at which it is observed. Thus, in order to build a complete picture of galaxy evolution, we need to study the colors of galaxies as a function of mass (or luminosity), environment and redshifts. The classical picture for the formation of galaxies proposes a single `monolithic collapse` \citep{Eggen:1962yu}, with stars in elliptical galaxies being formed in a single burst, thereafter evolving passively \citep{Partridge:1967nh,Sandage:1970lj}. This very simple model predicts remarkably well many of the properties and scaling relations of elliptical galaxies. In the current hierarchical paradigm, structure forms in a `bottom-up' sense, as galaxies and clusters are built from the merging of smaller units. Recently, there is growing evidence that star-formation has evolved in a `top-down` sense with more massive galaxies being most actively star-forming in the past and the bulk of the star-formation activity moving toward less massive galaxies as the universe ages. Although this seems intuitively at odds with hierarchical models, scenarios have been proposed in which star-formation progresses in this anti-hierarchical manner \citep{De-Lucia:2006sa}. Whereas previous generations of semi-analytic models in this hierarchical framework suggested that the most massive early-type galaxies should be younger than less massive ones \citep{Baugh:1996vs,Kauffmann:1998mh}, in order to reconcile with the observed down-sizing trend, the prediction is now that although the most massive early-types assembled their {\it mass} later than lower mass early-types, the stellar mass has been built up through a series of gas-poor mergers which do not result in additional star-formation. Hence the earlier formation times of the stellar populations in more massive galaxies is recovered. Despite numerous signs of merging in early-type galaxies \citep[e.g.,][]{van-Dokkum:2005dl,Tran:2005xc}, it remains an open question how important mergers are in their formation and evolution. The problem of disentangling how a galaxy assembled its mass from how it assembled its stars is a difficult one. Several studies of field galaxies have reported this down-sizing or anti-hieararchical trend in star-formation \citep[e.g.,][]{Bell:2004lb,Juneau:2005ft,Faber:2005yw,Bundy:2005bc,Scarlata:2007oz}. Initial results suggested that the comoving number density of massive early-type galaxies had evolved more than could be accounted for by passive evolution alone, and that `dry merging` of massive galaxies was required \citep{Bell:2004lb,Faber:2005yw}. More recently, it has been suggested that most, if not all, of the evolution can be attributed to the termination of star formation and pure passive evolution; and that a significant contribution from merging is not required \citep{Cimatti:2006vz,Scarlata:2007oz}. In galaxy clusters, down-sizing appears to be supported by the spectroscopic ages of red-sequence galaxies as a function of mass \citep{Nelan:2005ml}. In distant clusters (z$\sim$0.8), a deficit of faint red-sequence galaxies relative to local clusters has been claimed, in accordance with this picture \citep{De-Lucia:2004xa,Tanaka:2005mk,de-Lucia:2007li}. However, all of these high-redshift works have found a large cluster-to-cluster scatter in their samples of sizes of approximately 1-10 clusters, indicating that a large, statistical sample is crucial to such studies. The relative contributions of passive evolution versus dry merging to explain the evolution of the number density of early-type galaxies both in the field and in clusters is still an open question. In this paper we present results using the first statistical sample of galaxy clusters drawn from a well-defined, wide area, homogenous survey covering a large redshift range, 0.35$<$z$<$0.95. We present the survey data in \S2 and detail our method for constructing composite clusters in \S3. In \S4 we examine the Red-Sequence Luminosity Function and use the ratio of luminous-to-faint red-sequence cluster galaxies to trace its evolution with redshift and dependence on cluster mass. In \S5 we discuss our results and compare with other studies of early-type galaxy evolution both in clusters and the field, and in \S6 we present our conclusions. Throughout we assume a cosmology of $H_0=$70 km s$^{-1}$Mpc$^{-1}$ (and $h=H_0/$100 km s$^{-1}$Mpc$^{-1}$), $\Omega_{\rm M}=$0.3 and $\Omega_\Lambda=$0.7. | We have studied the properties of red-sequence galaxies in a well-defined statistical sample of galaxy clusters over the redshift range 0.35$<$z$\le$0.95. Each redshift bin of width $\Delta$z$=$0.1 contains $\approx$50 clusters and $\approx$5000 red-sequence galaxies. Our main results are: 1) The faint end of the red-sequence, as measured by the ratio of luminous-to-faint galaxies, declines with increasing redshift. This implies that star formation has not yet ended in the faintest cluster galaxies at high redshift. The red-sequence is built up at the faint end as star-formation proceeds to progressively less luminous (less massive) galaxies, consistent with the down-sizing scenario \citep{Cowie:1996xw}. 2) The turnover of the faint end of the RSLF is dependent on the cluster richness (mass) in the sense that for more massive clusters, the deficit of faint red-sequence galaxies is less than that for less massive clusters. This is an indication that star formation ended earlier for faint galaxies in richer clusters than in poorer clusters. 3) The decline in the faint end of the red-sequence toward higher redshift is accompanied by an increase in the total (i.e., including blue galaxies) cluster LF. This suggests that the build up of faint, red galaxies may be driven largely by the termination of star-formation in low mass galaxies. A similar increase of blue galaxies is also seen at the brighter end of the LF, suggesting that (at least some of) the build up of high mass early-type galaxies may also be attributed to the termination of star-formation. Future work will add the $B$- and $V$-band imaging of RCS-1 fields \citep{Hsieh:2005fq} and the accompanying photometric redshift catalog to examine the luminosity functions of RCS clusters. The $\sim$1000 square degree next generation survey, RCS-2 \citep{Yee:2007qb}, will provide an order of magnitude larger sample to improve upon the statistics of the current work. | 7 | 10 | 0710.2351 |
0710 | 0710.0953_arXiv.txt | {}{To investigate variability and to model the pulsational behaviour of AGB variables in the intermediate-age LMC cluster NGC 1846.}{Our own photometric monitoring has been combined with data from the MACHO archive to detect 22 variables among the cluster's AGB stars and to derive pulsation periods. According to the global parameters of the cluster we construct pulsation models taking into account the effect of the C/O ratio on the atmospheric structure. In particular, we have used opacities appropriate for both O-rich stars and carbon stars in the pulsation calculations.} {The observed P-L-diagram of NGC 1846 can be fitted using a mass of the AGB stars of about 1.8\,M$_{\sun}$. We show that the period of pulsation is increased when an AGB star turns into a carbon star. Using the mass on the AGB defined by the pulsational behaviour of our sample we derive a cluster age of $1.4\times10^{9}$ years. This is the first time the age of a cluster has been derived from the variability of its AGB stars. The carbon stars are shown to be a mixture of fundamental and first overtone radial pulsators.}{} | The cluster NGC 1846 belongs to an intermediate age population in the Large Magellanic Cloud (LMC). In contrast to the Milky Way the LMC harbours many clusters with an age around 1 to 5 Gyr (e.g. Girardi et al.\,\cite{GCBB95}). These clusters thus provide an interesting possibility to study the late stages of stellar evolution for stars around 1.5 to 2.5 M$_{\sun}$, while the globular clusters belonging to the Milky Way trace the evolution of stars of only up to about 0.9 M$_{\sun}$. While there is a general agreement in the literature that NGC 1846 is indeed an intermediate age cluster, the published values on global parameters of this system show some scatter. Derived metallicities can be roughly divided into two groups, one around [Fe/H]$=$-1.5 (Dottori et al.\,\cite{DPB83}, Leonardi \& Rose \cite{LR03}) and one around [Fe/H]$=$-0.7 (Bessell et al.\,\cite{BWL83}, Bica et al.\,\cite{BDP86}, Olszewski et al.\,\cite{OSSH91}, Beasley et al.\,\cite{BHS02}). Probably the first determination of the metallicity was done by Cohen (\cite{C82}) giving a value in between the two discussed metallicity levels, namely [Fe/H]$=$-1.1. The most recent determination of the cluster's metallicity is given by Grocholski et al.\,(\cite{GCS06}), who derive a value of [Fe/H]$=-$0.49$\pm0$ based on the Ca triplet strength of 17 individual cluster members measured with the VLT. Age values for NGC 1846 scatter between less than 1 Gyr (Frantsman \cite{F88}) and 4.3 Gyr (Bica et al.\,\cite{BDP86}). Mackey \& Broby Nielsen (\cite{MB07}) give cluster ages from 1.5 to $2.5 \times 10^{9}$ years based on a comparison of the colour magnitude diagram and theoretical isochrones. The range in age thereby results from the usage of two different sets of isochrones. These ages correspond to masses on the AGB of about 1.3 to 1.8 M$_{\odot}$. Mackey \& Broby Nielsen (\cite{MB07}) also report on the probable existence of two populations in NGC 1846 with similar metallicity but separated in age by about 300 Myr. A reddening of $A_{V}$$=$0.45\,mag has been determined by Goudfrooij et al.\,(\cite{GGKM06}). Grocholski et al.\,(\cite{GCS06}) note that NGC 1846 is probably suffering from differential reddening without giving any detailed numbers. Keller \& Wood (\cite{KW06}) give a somewhat lower reddening of $E(B-V)$$=$0.08, i.e. an $A_{V}$ value of 0.25\,mag. The first studies of the stellar content of NGC 1846 were published by Hodge (\cite{H60}) and Hesser et al.\,(\cite{HHU76}). A number of luminous stars on the Asymptotic Giant Branch (AGB) were identified and published with finding charts by Lloyd-Evans (\cite{LE80}). Throughout this paper we will use the naming given in Lloyd-Evans' publication (LE{\it XX}) except for H39, which follows the numbering from Hodge (\cite{H60}). Details on individual AGB stars have been published in a number of papers. Frogel et al. (\cite{FPC80}, \cite{FMB90}) and Aaronson \& Mould (\cite{AM85}) gave near infrared photometry and $m_{\rm bol}$ values for most of Lloyd-Evans' AGB stars and a few more cluster objects, unfortunately without any finding chart. Tanab\'{e} et al.\,(\cite{T98}) searched this cluster for extreme infrared stars, i.e.~stars with a high circumstellar absorption in the visual range, but found none. No investigations on the variability of the AGB stars in NGC 1846 has been published up to now. However, as most AGB stars are pulsating (forming the group of long period variables or LPVs) it seemed very likely that some light variability would be found in these stars. In a previous paper (Lebzelter \& Wood \cite{LW05}) we showed that the variability analysis of LPVs in a single stellar population like the globular cluster 47 Tuc allows one to investigate various fundamental aspects of AGB stars like the evolution of the pulsation mode or mass loss. With NGC 1846 we have chosen an interesting alternative target since its AGB stars are more massive and include a number of carbon stars. | \label{discus} Having matched the O-rich models to the observed values of $T_{\rm eff}$, we ensure that the models matched the observed pulsation periods of the O-rich stars. If not, the mass was altered, and new calculations of the giant branch track and the pulsation periods were made. It was found that a mass of 1.8\,M$_{\odot}$ gave the best fit between observations and theory for the O-rich stars. This corresponds to a cluster age of $1.4\times10^{9}$ years using the isochrones of Girardi et al. (\cite{GBBC00})\footnote{Using the BaSTI isochrones (Bedin et al. \cite{basti05}) we find an age of $1.9\times10^{9}$ years for a 1.8\,M$_{\odot}$ star to reach the AGB (using the same parameters as Mackey \& Broby Nielsen \cite{MB07}).}. The fit is shown in Figure~\ref{pl_fig} and the model details are given in Table~\ref{puls_mods}. It is clear that all the O-rich stars, with the exception of LE\,17 near $\log P = 1.8$ and $\log L$/L$_{\odot} = 3.49$, are first or second overtone pulsators (the stars with long secondary periods are not considered here). It is clear from the models that as the C/O ratio increases beyond 1.0, the periods of all modes increase rather rapidly due to the increase in radius and decrease in $T_{\rm eff}$ caused by the molecules in the C-rich atmospheres. Furthermore, it can be seen that the C stars fall quite nicely on two sequences corresponding to pulsation in the first overtone and fundamental modes. In the absence of these C-rich models, it would have been natural to explain the C stars as fundamental mode pulsators scattering broadly around an extension of the fundamental mode O-rich sequence. The fit of the models to the observed C star points is not perfect. This is not surprising given the rather elementary method used to compute the C-rich opacities, combined with the fact that we do not really know the C/O ratios in these stars. Furthermore, as mentioned above, the stars move up and down in luminosity during a He-shell flash cycle by factors of about 2 in luminosity (Vassiliadis \& Wood \cite{VW93}). The dashed lines in Figure~\ref{pl_fig} show how the fundamental and first overtone periods would vary over a shell flash cycle for a star with C/O = 2. It seems most likely that the two C stars with $\log L$/L$_{\odot} < 3.6$ are near the luminosity minimum of a shell flash cycle (where AGB stars spend about 20\% of their time). In general, the theoretical and observed periods for the C-stars agree reasonably well. This indicates that the radii and $T_{\rm eff}$ values of the models agree well with the real values for these stars and the assumption that the temperatures indicated by the near infrared colours are misleading. It is somewhat strange that the most luminous of the C stars appear to have shorter pulsation periods than the less luminous C stars. One way for this to occur would be to decrease the C/O ratio in the most luminous C stars by hot-bottom burning. However, at 1.8 M$_{\odot}$, the mass is too low for hot-bottom burning, since AGB masses greater than about 4 M$_{\odot}$ are required for hot-bottom burning to be significant. Another explanation may be that the nonlinear pulsation period differs from the linear pulsation period when the pulsation amplitude increases, as noted in Lebzelter \& Wood (\cite{LW05}). The visual light amplitudes give no indication for this as the amplitudes of the carbon stars are all very similar (see Fig.\,\ref{tanabe}). In the infrared, as shown in Fig.\,\ref{lcir}, the star with the largest amplitude in our data set is LE5 which is also the most luminous star of our sample. This may indeed indicate that the shorter periods at the tip of the AGB can be explained by nonlinear effects. However, as mentioned above, our near infrared time series are too short for a definite conclusion on this point. A reduction in stellar mass due to significant amounts of mass loss can not explain the shorter periods of the most luminous C stars since this would {\em increase} the pulsation period. In general, we have not found it necessary to invoke any significant reduction in mass with luminosity up the AGB in these calculations. We note that this contrasts with the stars in 47\,Tuc studied by Lebzelter \& Wood (\cite{LW05}) where a significant amount of mass loss was found on the AGB. As noted briefly in Section 4, this is as expected. The stars being studied here and in 47\,Tuc appear to have relatively low mass loss rates where a Reimers' mass loss law may apply. In such a case, the mass loss rate is proportional to $LR/M$. For a given $L$, this functional dependence suggests that the intermediate mass ($\sim$1.8 M$_{\odot}$) stars in NGC\,1846 will have lower mass loss rates than the low mass ($\sim$0.9 M$_{\odot}$) 47\,Tuc stars because of the higher mass. In addition, the giant branch temperature increases with stellar mass so the radius will be smaller at a given $L$ in NGC\,1846. The consequence of the expected lower mass loss rate in NGC\,1846 is that a smaller amount of mass is lost on the giant branches up to a given $L$. The fractional change in mass in NGC\,1846 is even smaller because of the higher mass. An independent check of the mass loss can be obtained from mid-IR photometry looking for indications of an infrared excess due to dust emission. NGC\,1846 was covered by the SAGE mid-IR survey (Meixner et al.\,\cite{sage06}). Blum et al. (\cite{Blum07}) find [8]-[24] as an indicator for mass loss. For five of the carbon stars in our sample (LE1, LE2, LE3, LE6 and LE11) the corresponding photometry can be found in the SAGE point source data base. All of them show an [8]-[24] colour typical for low mass loss stars. We recall that Tanab\'{e} et al. (\cite{T98}, \cite{T04}) did not find any dust enshrouded objects in this cluster. For two sources Tanab\'{e} et al. (\cite{T04}) list no identification from the literature. Their object 16 is obviously the variable LW2. Object 8, the weakest object in their list of mid infrared sources, corresponds to LW4. Using the SAGE list of point sources we checked for bright objects at 24\,$\mu$m with no or a very weak optical counterpart. Indeed we found three sources fulfilling this criterion, but their membership to the cluster is not very likely. In the appendix we give a more detailed description of these three sources. Thus we can derive from our findings that the AGB stars of NGC 1846 lose their mass in a quite short time at the end of the AGB phase. The location of LE17 in the P-L-diagram remains somehow a mystery. While the two carbon stars found at $\log L$/L$_{\odot}$ = 3.57 nicely agree with the expected location of a TP-AGB star close to its luminosity minimum, it is unlikely that this explanation also holds for LE17, which is the O-rich star slightly below in the P-L-diagram (Fig.\,\ref{pl_fig}). If it would be in the minimum of its TP cycle, it would be in a region exclusively occupied by C-rich stars during its maximum. A possibility may be that the star is located at the "knee" of the fundamental mode sequence at its maximum. Alternatively, the derived temperature and luminosity of this star may be wrong due to higher reddening of this object. Circumstellar reddening probably plays only a minor effect as there is no clear indication for this from mid-IR photometry (SAGE; Tanab\'{e} et al.\,\cite{T04}). We further note that the star nicely fits onto the giant branch of NGC 1846 (Fig.\,\ref{hrd}), thus there is no indication for a significant reddening. For the same reason we may safely assume that the star is indeed a member of the cluster. Further investigations are required to understand the behaviour of this star. | 7 | 10 | 0710.0953 |
0710 | 0710.2484_arXiv.txt | \medskip \noindent We assume the validity of the Standard Model up to an arbitrary high-energy scale and discuss what information on the early stages of the Universe can be extracted from a measurement of the Higgs mass. For $M_h \simlt 130 \gev$, the Higgs potential can develop an instability at large field values. From the absence of excessive thermal Higgs field fluctuations we derive a bound on the reheat temperature after inflation as a function of the Higgs and top masses. Then we discuss the interplay between the quantum Higgs fluctuations generated during the primordial stage of inflation and the cosmological perturbations, in the context of landscape scenarios in which the inflationary parameters scan. We show that, within the large-field models of inflation, it is highly improbable to obtain the observed cosmological perturbations in a Universe with a light Higgs. Moreover, independently of the inflationary model, the detection of primordial tensor perturbations through the $B$-mode of CMB polarization and the discovery of a light Higgs can simultaneously occur only with exponentially small probability, unless there is new physics beyond the Standard Model. | The search for the Higgs boson and the measurement of its properties are one of the primary goals of the LHC. The mere discovery of a Higgs can be viewed as a possible indication of additional new physics not far from the electroweak scale, because of the high sensitivity to short-distance quantum corrections of the mass term associated to a fundamental scalar. However, even in the absence of any new-physics discovery at the LHC, a measurement of the Higgs mass can give us useful hints on the structure of the theory at very high energies. This is because, at large field values, the Higgs potential can develop an instability or become non-perturbative, depending on the precise value of the Higgs quartic coupling $\lambda$ or, ultimately, on the Higgs mass $M_h$. Because of the logarithmic dependence of $\lambda$ on the energy, such considerations~\cite{con,Str,CEQ,thermal1,Hambye} can test the properties of the theory up to extremely small distances, which are otherwise totally unaccessible to any imaginable collider experiment. In this paper, we want to use the same considerations for a different purpose. Rather than trying to infer new particle-physics properties at small distances, we will assume the validity of the Standard Model (SM) up to an arbitrary high-energy scale, and find what information on the early stages of the Universe can be extracted from a measurement of the Higgs mass. We will first obtain that in the Higgs mass range $114 \gev \simlt M_h\simlt 130 \gev$, where the electroweak vacuum is potentially metastable, the absence of excessive thermal Higgs field fluctuations in the early Universe imposes a bound on the reheat temperature after inflation $T_{RH}$. Then we will discuss the interplay between the quantum Higgs fluctuations generated during the primordial stage of inflation and the cosmological perturbations which are either currently observed in the form of Cosmic Microwave Background (CMB) anisotropies, or might be detected in the near future in the form of tensor (gravity waves) perturbations. The key ingredient is that all these fluctuations depend upon the Hubble rate during inflation and therefore, under certain assumptions, it is possible to relate the amount of Higgs fluctuations to observable properties of the CMB. However, excessive fluctuations of the Higgs field during inflation can pose a threat to the stability of the present vacuum, if the Higgs mass lies in the metastability window. We will assume that the various initial inflationary patches of the Universe are characterized by different microphysical parameters. In this sense we take the point of view that the underlying theory has many vacua, realized in different patches of the Universe. This picture, usually referred to as the landscape \cite{landscape}, has been put forward especially in the context of string theory. Under this assumption, from CMB observations and from a measurement of the Higgs mass, we can derive probabilistic conclusions on the properties of our Universe. In particular, within the class of large-field models of inflation, we will compute the probability to have a Universe at the end of inflation which both survived the quantum Higgs fluctuations and has the right amount of observed cosmological perturbations. Such probability is extremely (exponentially) small, if the Higgs mass is below 124~GeV (for the present central value of the top mass). Moreover, we find that the discovery of a light Higgs together with the detection of primordial tensor perturbations through the $B$-mode of CMB polarization (at a level quantitatively described in sect.~6) would imply that we live in a very atypical Universe, whose probability decreases exponentially when the Higgs mass decreases. Such discovery could be interpreted as an indirect evidence for the existence of new physics beyond the Standard Model at some intermediate energy scale. The paper is organized as follows. In sect.~2 we briefly review the Higgs mass instability window. In sect.~3 the bounds on the reheating temperature after inflation are discussed. In sect.~4 we compute the survival probability of the electroweak vacuum during inflation and in sects.~5 and~6 we relate it to the curvature and tensor perturbations, respectively. Section~7 states our conclusions. The appendix contains technical details relevant to the calculation of the Higgs mass instability window. | In this paper we have investigated some possible cosmological implications of the Higgs mass measurement. If the LHC discovers a Higgs with mass in the metastability window shown in fig.~\ref{fig:window}, and does not find direct evidence for other new physics, then there is a concrete possibility that we live in a metastable state. If we assume that the pure SM is valid up to very large energy scales, then stability against field fluctuations tests properties of the early Universe. We have revisited the known considerations about thermal fluctuations, interpreting the result as an upper limit on the reheating temperature $T_{RH}$ after inflation. The bound is summarized in fig. \ref{mhTRH}. We have also discussed the possibility that the inflationary vacuum fluctuations destabilize the electroweak vacuum. If the Hubble rate is large enough during inflation and the Higgs mass is light, then the danger exists that the classical value of the Higgs field is pushed above its instability point causing the collapse of the corresponding inflating domain. This does not necessarily pose a problem, since all the surviving domains will be characterized by the correct electroweak vacuum. However, interesting probabilistic conclusions can be reached under the assumption of a ``landscape" scenario in which the inflationary parameters take different values in different patches of the Universe. In the context of large-field models of inflation we have argued that, among the surviving regions, only an exponentially tiny fraction of them will be characterized by the correct amount of cosmological perturbations. If the Higgs mass is found below $(120-130)$~GeV, either we live in a very special, and exponentially unlikely, domain or new physics must exist below the scale $\Lambda$. Moreover, our considerations can be directly related to the amount of primordial gravity waves which can be measured through their imprint on the CMB. The discoveries of a light Higgs boson and of tensor modes would be mutually conflicting (at least in a probabilistic sense) if the condition in \eq{trul} is not satisfied. Again, this evidence could be interpreted as an indication of new physics modifying the extrapolation of the Higgs potential to large field values. This result is valid, independently of the particular inflationary model considered. Finally, a measurement of the Higgs mass below about 130~GeV provides a direct bound, shown in fig.~\ref{conformal}, on the quantity $\xi H^2$, where $\xi$ is the (negative) coupling between the Higgs bilinear and the curvature, and $H$ is the maximal Hubble rate during inflation. This bound becomes particularly interesting in case of detection of primordial gravity waves, which provide a measurement of $H$. Note that the bound in fig.~\ref{conformal} does not depend on any statistical consideration of parameter scanning. | 7 | 10 | 0710.2484 |
0710 | 0710.5437_arXiv.txt | We have performed a series of three-dimensional simulations of a starburst-driven wind in an inhomogeneous interstellar medium. The introduction of an inhomogeneous disk leads to differences in the formation of a wind, most noticeably the absence of the ``blow-out'' effect seen in homogeneous models. A wind forms from a series of small bubbles that propagate into the tenuous gas between dense clouds in the disk. These bubbles merge and follow the path of least resistance out of the disk, before flowing freely into the halo. Filaments are formed from disk gas that is broken up and accelerated into the outflow. These filaments are distributed throughout a biconical structure within a more spherically distributed hot wind. The distribution of the inhomogeneous interstellar medium in the disk is important in determining the morphology of this wind, as well as the distribution of the filaments. While higher resolution simulations are required in order to ascertain the importance of mixing processes, we find that soft X-ray emission arises from gas that has been mass-loaded from clouds in the disk, as well as from bow shocks upstream of clouds, driven into the flow by the ram pressure of the wind, and the interaction between these shocks. | In 1963, \citeauthor{LS1963} first detected an outflow of gas along the minor axis of M82. \citet{CC1985} proposed a model in which a galactic scale outflow could be powered by the combined kinetic energy from supernovae. Starburst galaxies, with their characteristically high star formation rates, provide the perfect environments for these winds to develop. Indeed, galactic winds are ubiquitous in starburst galaxies, having been observed in many nearby galaxies and inferred in galaxies at high-redshifts \citep[see][~and references therein]{VCB2005} The best studied galactic wind is the outflow in M82, which is clearly visible in the light of H$\alpha$, displaying a vast filamentary system extending several kpc along the minor axis of the galaxy \citep{SB1998}. These filaments lie on the surface of a mostly hollow structure and rotate in the same direction as the disk \citep{Greve2004}. As with other galactic winds \citep[e.g. NGC 253:][]{SDW2003}, the wind in M82 is asymmetric, with the northern outflow more chaotic than the southern outflow. The filaments can be traced to the nuclear region and display both shell and loop-like structures \citep{Oetal2002}. The formation of these filaments is currently not well understood, but they are thought to be either disk or halo gas that has been entrained into the outflow. The morphology of galactic winds can vary. Outflows often display asymmetries, varying degrees of collimation and may be tilted with respect to the minor axis. While many outflows are limb-brightened \citep[e.g. NGC 3079;][]{Vetal1994}, the optical filaments can also fill the volume rather than remain confined to the surface of the biconical outflow \citep{VB1997}. The host galaxy itself plays an important role in determining the morphology of a wind, with its size and structure affecting the degree of collimation \citep{SS2000} and expansion of the outflow \citep{SetalB2004,Getal2005,Martin2005}. Recent Chandra observations have revealed increasing detail in the X-ray emission from galactic winds. One of the most striking results of these observations is the close spatial relationship with the H$\alpha$ emitting gas \citep[e.g.][]{Setal2000,Setal2002,SetalA2004,CBV2002,MKH2002,Getal2005,OWB2005b}, suggesting a close physical connection. Thus, a successful model of a galactic wind needs to explain this relationship. \citet{Setal2002} provide a summary of several theories for the origin of the X-ray emission that could explain this correlation. These mechanisms involve shocked disk or halo gas that has been swept up into the wind, in the form of dense clouds or shells. Over the past few decades, numerous simulations have been made of starburst-driven winds \citep{TI1988,TB1993,Setal1994,Setal1996,DB1999,TM1998,SS2000,TSM2003}. \citet{Setal1994} performed two-dimensional, axisymmetric simulations of a galactic wind in an isothermal ISM, with varying densities and temperatures. They concluded that the H$\alpha$ filaments form from disk gas that has been entrained into the flow and that the X-ray emission most likely arises from shocked disk and halo gas. More recently, \citet{SS2000} performed a series of simulations, focusing on the energetics and X-ray emission from the wind. As with \citet{Setal1994}, their simulations were two-dimensional and axisymmetric with an isothermal ISM. They found that a large fraction of the soft X-ray emission in their model comes from shock-heated ambient gas and from the interfaces between cool dense and hot tenuous gas. While these simulations provide some insight into the origin of the X-ray and H$\alpha$ emission, the homogeneous nature of these models and their symmetry renders them incapable of forming significant filamentary structures, limiting their ability to constrain the emission processes. In order to improve upon previous models and to gain a better understanding of the origin of the H$\alpha$ filaments and X-ray emission, we have performed a series of three-dimensional simulations of a galactic wind in an inhomogeneously distributed interstellar medium (ISM). The introduction of inhomogeneity is important as the interstellar medium in a galaxy disk is highly complex in all its phases \citep[see for example][~and references therein]{EE2001}. Inhomogeneity is also crucial in the development of a wind, as energy from massive stars formed in dense molecular clouds in the starburst region may be radiated away before a wind could form. A wind is more likely to develop from the kinetic energy from stellar winds adjacent to the diffuse gas surrounding the clouds. The inhomogeneous structure of the ISM is also likely to affect the distribution of filaments throughout the wind, producing asymmetric and tilted outflows. It is likely that the size and strength of the starburst itself plays an important role in determining the morphology. Many starburst galaxies, such as M82 and NGC 3079, contain circumnuclear starbursts, with their resultant outflows being strong and violent \citep{SB1998,Vetal1994}. Other starbursts are weaker and have less prominent outflows. An example is NGC 4631, which is currently undergoing a disk-wide starburst \citep{SetalA2004}. \citet{TSM2003} investigated the formation of the emission line filaments by modeling the formation of a wind from several super star clusters. They proposed that kiloparsec long filaments are formed from stationary and oblique shocks. In this paper we present a different model, which follows a similar approach to that of \citet{SS2000}, but introduces an inhomogeneous disk. We follow the evolution of a starburst-driven wind in different ISM conditions and discuss the effect of the inhomogeneity of the disk on the morphology of the wind. We consider the morphology of the H$\alpha$ emitting filaments separately and investigate their origin. Finally, the luminosity of the soft and hard X-ray emitting gas is calculated and we suggest an origin for the soft X-ray emission. | We have performed a series of three-dimensional simulations of a starburst-driven galactic wind designed to test the evolution of the wind in different ISM conditions. By conducting three-dimensional simulations we are able, for the first time, to study the morphology and dynamics of the entire outflow. The introduction of an inhomogeneous disk enables us to study the development of asymmetries and the interaction of the wind with clouds in the disk. The results of these simulations are as follows- \begin{enumerate} \item The interstellar medium plays a pivotal role in the evolution of a galactic wind. The interaction of the wind with clouds in the disk results in asymmetries and tilted outflows on the small scale. Nevertheless, it is likely that inhomogeneities in the halo are the cause of the large-scale asymmetries in an outflow. \item The distribution of gas surrounding the starburst region assists in collimating the outflow. The thickness of the disk and the location of the starburst are important factors in determining the degree of collimation, with the degree of collimation increasing with the amount of gas surrounding the starburst region. \item The base of the outflow is well confined within a radius of 200 pc over the 2 Myr time frame of the simulation as a result of the high density of the disk gas. \item The H$\alpha$ filaments form from the breakup of clouds in the starburst region, the fragments of which are then accelerated by the ram pressure of the wind. Filaments are also formed from gas that has been stripped from the sides of the starburst region. The distribution and mass of the filaments is affected by the distribution of clouds in the vicinity of the starburst region. \item The H$\alpha$ filaments appear as strings of disk gas that form a filled biconical structure inside of a more spherical hot wind. The filaments are distributed throughout this structure, but do not trace the true extent to the wind defined by the hot gas. \item The calculated soft X-ray luminosities up until 2.0 Myr are of the order of 10$^{38}$ - 10$^{39}$ erg s$^{-1}$ and the hard X-ray luminosities of 10$^{38}$ erg s$^{-1}$. These luminosities are dependent on the volume of the wind and would be larger for a more evolved outflow. \item Interior to the swept-up shell of halo gas, soft X-ray emission originates in the same region as the H$\alpha$ emitting gas. While higher resolution simulations are needed to confirm X-ray emission from mixing processes, we find 4 mechanisms that give rise to Soft X-rays: (i) The mass-loaded wind, (ii) the intermediate temperature interface between the hot wind and cool filaments, (iii) bow shocks, and (iv) interactions between bow shocks. The shell is also a major contributor to the soft X-ray emission, but has no associated H$\alpha$ emission. \item The hard X-ray emission originates from gas in the starburst region. \end{enumerate} The results of these simulations indicate that the host galaxy itself and the environment in which it is situated is a major determinant in the morphology of the outflow. The emission processes that contribute to the H$\alpha$ and soft X-ray emission may vary from one galaxy to the next. Whether the H$\alpha$ emission originates from photoionization or from shock-heating (or both) cannot be determined from these simulations. However, we do find an abundance of filamentary T $\sim$ 10$^4$ K gas that has been accelerated into the outflow, forming a biconical shaped region that is commonly observed in starburst winds. The source of the soft X-ray emission is also likely to depend upon the environment of the host galaxy. In the case of M82 it is plausible that the interaction of the wind with the surrounding HI clouds is also a contributor to the soft X-ray emission, in addition to the processes mentioned above. The observed spatial relationship between the H$\alpha$ and soft X-ray emitting gas can be explained when considering emission processes interior to the wind, such as bow-shocks and the mass-loaded component of the wind. In addition, the presence of the strong X-ray emitting shell with no associated H$\alpha$ emission is interesting. While the ultimate fate of the shell is unknown at present, this result argues for the presence of X-ray emission more extended than the filamentary H$\alpha$ gas. \citet{Setal2002} suggest that this emission may be detectable in more distant starburst galaxies. In future work we shall investigate the evolution of a wind over a larger time frame and spatial extent than our current study, and look at the total energy budget of the outflow. It is also important to further test the effect of resolution on the filaments, and in particular the associated soft X-ray emission that arises through mixing of the hot wind and the cooler disk gas, and will be the subject of a subsequent paper. | 7 | 10 | 0710.5437 |
0710 | 0710.1434_arXiv.txt | We discuss the possibility of accurately estimating the source number density of ultra-high-energy cosmic rays (UHECRs) using small-scale anisotropy in their arrival distribution. The arrival distribution has information on their source and source distribution. We calculate the propagation of UHE protons in a structured extragalactic magnetic field (EGMF) and simulate their arrival distribution at the Earth using our previously developed method. The source number density that can best reproduce observational results by Akeno Giant Air Shower Array is estimated at about $10^{-5}~{\rm Mpc}^{-3}$ in a simple source model. Despite having large uncertainties of about one order of magnitude, due to small number of observed events in current status, we find that more detection of UHECRs in the Auger era can sufficiently decrease this so that the source number density can be more robustly estimated. 200 event observation above $4 \times 10^{19}~{\rm eV}$ in a hemisphere can discriminate between $10^{-5}$ and $10^{-6}~{\rm Mpc}^{-3}$. Number of events to discriminate between $10^{-4}$ and $10^{-5}~{\rm Mpc}^{-3}$ is dependent on EGMF strength. We also discuss the same in another source model in this paper. | \label{introduction} The nature of the sources of ultra-high-energy cosmic rays (UHECRs) is still poorly known though many efforts in detection of these particles with UHE energies. It is one of challenging problems in astroparticle physics. About 10 years ago, Akeno Giant Air Shower Array (AGASA) reported small-scale anisotropy of the UHECR arrival distribution within its angular resolution while it did large-scale isotropy with a harmonic analysis \cite{takeda99,takeda01}. It found five doublets and one triplet as event clusterings. It is enough evidence that origin of UHECRs is astrophysically point-like source. However, no obvious astronomical counterparts to the observed UHECRs have been found. One of this reasons is magnetic fields in the universe. Extragalactic magnetic field (EGMF) is poorly known theoretically and observationally. As an upper limit for its strength, $B$, and correlation length, $l_c$, $B~{l_c}^{1/2} < (1~{\rm nG})(1~{\rm Mpc})^{1/2}$, by Faraday rotation measurements of radio signals from distant quasars, is often adopted \cite{kronberg94}. Based on its upper limit, UHECR propagation has been discussed in simply uniform turbulent magnetic field \cite{yoshiguchi03,aloisio04,berezinsky06}. It is also observationally known that clusters of galaxies have strong magnetic fields with 0.1 - a few $\mu$G at its center \cite{vallee04}. This is shown that there are not only uniform EGMF but also relatively strong EGMF incidental to the local structure. Recently, several groups have performed simulations of large-scale structure formation with magnetic fields \cite{sigl03,sigl04,dolag05}. They find that magnetic field traces the local density distribution. Such magnetic field plays an important role in UHECR propagation since it is stronger than the uniform one. In our previous study \cite{takami06}, we constructed a structured EGMF model that reflects the local structures and discussed predicted arrival distribution of UHE protons at the Earth in consideration with their propagation processes. We also considered Galactic magnetic field (GMF) in the propagation. As a result, the number density of UHECR sources was constrained to $\sim 10^{-4}~{\rm Mpc}^{-3}$ under our luminosity-weighted source model (explained in section \ref{model}) from observed arrival distribution. The EGMF strength was normalized only to $0.4~\mu{\rm G}$ at the center of the Virgo cluster. However, observational measurements of magnetic field in a galaxy cluster result in 0.1 - a few $\mu$G, so that the strength of magnetic field has large uncertainty of about one order of magnitude. Thus, it is important that the propagation process and the arrival distribution are investigated in several strengths of EGMF. In this study, we calculate propagation of UHE protons in a structured EGMF with several strengths for normalization and simulate the arrival distributions at the Earth. At first, comparing them to observational results, we constrain the source number density. We use AGASA data for this purpose because AGASA has observed the most number of events. Such constraint is dependent on UHECR source model. We discuss about two simple source models. We find that the number density has large uncertainty of about an order of magnitude due to small number of observed events at present. So, we also discuss the possibility of a decrease in the uncertainty with future observations. In section \ref{model}, our models of UHECR source distribution, structured EGMF, and GMF are briefly explained. In section \ref{method}, a numerical method of UHE proton propagation, construction of the arrival distribution and statistical analysis are explained. We present our results in section \ref{result}, and summarize in section \ref{summary}. | \label{summary} In this study, we discuss the possibility of accurately estimating the source number density of UHECRs with small-scale anisotropy. Comparison between simulated arrival distribution and the observational results enables us to estimate the source number density. In order to construct the arrival distribution, we calculate the propagation of UHE protons in a structured EGMF with several strengths consistent with measurements of magnetic field in clusters of galaxies. The GMF is also considered. We find that the source number density of $10^{-5}~{\rm Mpc}^{-3}$ in the normal source model and $10^{-4}~{\rm Mpc}^{-3}$ in the luminosity-weighted source model can best reproduce the AGASA results, which are weakly dependent on strength of our structured EGMF. However, these have large uncertainty of about one order of magnitude due to the small number of observed events. So, we discuss the possibility that future observations decrease the uncertainty. In the normal source model, Auger can distinguish $10^{-5}~{\rm Mpc}^{-3}$ and $10^{-6}~{\rm Mpc}^{-3}$ sufficiently by our method in the near future. If the structured EGMF is zero or very weak, $10^{-4}~{\rm Mpc}^{-3}$ is also discriminated from the less number density in 500 event observation above $4 \times 10^{19}~{\rm eV}$. In stronger EGMF, more observations are requested because cosmic rays are diffused more strongly. Number of events that needed for the distinction depends on EGMF strength. In the luminosity-weighted source model, $10^{-3}~{\rm Mpc}^{-3}$ can be distinguished from the less number density by Auger. The distinction between the less number densities is difficult due to large uncertainty which originates from different injection powers of the sources. \begin{figure}[t] \begin{center} \includegraphics[width=0.48\linewidth]{fig8a.eps} \hfill \includegraphics[width=0.48\linewidth]{fig8b.eps} \caption{Distribution of $\chi^2$s calculated from arrival protons above $4 \times 10^{19}~{\rm eV}$, in the normal source model. The GMF is not considered The strength of the structured EGMF is normalized to 0.1 $\mu$G and that of an uniform turbulent field is 0 ({\it upper panel}), 1 nG ({\it middle panel}), and 10 nG ({\it lower panel}). The numbers of events are set to be 200 ({\it left}) and 500 events ({\it right}) within the southern hemisphere.} \label{fig_chi_offset} \end{center} \end{figure} We exclusively adopt the AGASA results in this study. However, High Resolution Fly's Eye (HiRes) claims no significant small-scale anisotropy, contrary to AGASA \cite{abbasi04}. This discrepancy is one of well-known problems in UHECR experiments. At present, this is not statistically significant due to the small number of observed events \cite{yoshiguchi04}. It will be able to solved by new experiments with large aperture, such as Auger, Telescope Array \cite{tahp}, and Extreme Universe Space Observatory \cite{eusohp}. Hence, it is possible that these experiments do not observe sufficient small-scale anisotropy in the future. If Auger does not observe small-scale clusterings during 2007 (maybe it detects about 200 events above $4 \times 10^{19}~{\rm eV}$), the source number density is estimated at about $10^{-4}~{\rm Mpc}^{-3}$ or more in the normal source model by definition of $\chi^2$ in figure \ref{fig_chidisL0}. It is comparable with number density of active galactic nuclei\cite{loveday92}. Our EGMF model within 100 Mpc has about 95\% of volume without magnetic field. In our model, uniform magnetic field is not considered. According to an upper limit mentioned in section \ref{introduction}, deflection angle of UHE protons with energy of $E$ during propagation of distance, $d$, is estimated as \begin{equation} \theta < 3^{\circ} \left( \frac{E}{10^{20}~{\rm eV}} \right)^{-1} \left( \frac{d}{100~{\rm Mpc}} \right)^{1/2} \left( \frac{l_c}{1~{\rm Mpc}} \right)^{1/2} \left( \frac{B}{1~{\rm nG}} \right). \end{equation} This deflection is expected to generate more isotropic arrival distribution of UHECRs. Therefore, uniform turbulent magnetic field affects determination of the source number density. As a demonstration, we show $\chi^2$ distribution for $B_{\rm EG} = 0.1~\mu{\rm G}$, including the uniform field, in figure \ref{fig_chi_offset}. The normal source model is adopted. The number of events is set to be 200 ({\it left panel}), and 500 ({\it right panel}). The strengths of the uniform turbulent field are 0 ({\it upper panel}), 1 ({\it middle panel}), and 10 nG ({\it lower panel}). The distributions are shifted to lower value in stronger uniform turbulent field. The shifts are larger in more source number density. In stronger turbulent uniform magnetic field, number of events needed for the distinction between $10^{-5}$ and $10^{-6}~{\rm Mpc}^{-3}$ is smaller while it becomes difficult to discriminate $10^{-4}$ from $10^{-5}~{\rm Mpc}^{-3}$ due to the diffusion of cosmic rays. Strength of uniform EGMF is also important for estimating the source number density. In this work, we adopt isotropic arrival distribution of UHECRs as a template of a future result since results of new experiments are still unpublished. Auger should detect a few times more number of UHE events than that of AGASA since it started observation about three years ago. Its result will provide us beneficial information on the nature of UHECR sources. \subsubsection*{Acknowledgements:} The work of H. T. is supported by Grants-in-Aid from JSPS Fellows. The work of K. S. is supported by Grants-in-Aid for Scientific Research provided by the Ministry of Education, Science and Culture of Japan through Research Grants S19104006. | 7 | 10 | 0710.1434 |
0710 | 0710.4482_arXiv.txt | % We present recent results from the \htmladdnormallinkfoot{Laboratory for Cosmological Data Mining}{http://lcdm.astro.uiuc.edu} at the National Center for Supercomputing Applications (NCSA) to provide robust classifications and photometric redshifts for objects in the terascale-class Sloan Digital Sky Survey (SDSS). Through a combination of machine learning in the form of decision trees, k-nearest neighbor, and genetic algorithms, the use of supercomputing resources at NCSA, and the cyberenvironment Data-to-Knowledge, we are able to provide improved classifications for over 100 million objects in the SDSS, improved photometric redshifts, and a full exploitation of the powerful k-nearest neighbor algorithm. This work is the first to apply the full power of these algorithms to contemporary terascale astronomical datasets, and the improvement over existing results is demonstrable. We discuss issues that we have encountered in dealing with data on the terascale, and possible solutions that can be implemented to deal with upcoming petascale datasets. | We summarize work carried out as part of the Laboratory for Cosmological Data Mining, a partnership between the Department of Astronomy and the National Center for Supercomputing Applications (NCSA) at the University of Illinois at Urbana-Champaign (UIUC), in collaboration with the Automated Learning Group (ALG) at NCSA, and the Illinois Genetic Algorithms Laboratory at UIUC. This combination of expertise allows us to apply the full power of machine learning to contemporary terascale astronomical datasets. | Given the petascale datasets planned for the next decade, it is vitally important that contemporary data mining can be carried out successfully on this scale. In turn, this requires robust techniques on the terascale. We encountered numerous issues that were relevant to realizing this goal: \begin{itemize} \item Because we are using tens to hundreds of parallel nodes and streaming many GB of data, D2K must be invoked via batch script, negating the advantages of its GUI interface. The resulting lack of an integrated cyberenvironment results in batch scripts that contain many tens of settings, manually set file locations and commands, making them prone to error. \item Job submission is inflexible, subject to fixed wallclock times and numbers of nodes, unpredictable queuing times and no recourse if a job fails due a bug or hardware problem. \item The large datasets must be stored on the Unitree mass storage system, which is occasionally subject to outages in access or significant wait times. In combination with the queuing system for batch jobs described above, this can make new scripts time-consuming to debug. \item There is no way in which to fully explore the huge parameter space (more than $10^{15}$ combinations of settings for decision trees) of the machine learning algorithms. Genetic algorithms were used to optimize the training features, and could be used similarly to optimize the algorithm. \item The present lack of fainter training data forces us to extrapolate in order to classify the whole SDSS. While the data and results we obtain are well-behaved, it will always be the case in astronomy that some form of extrapolation is ultimately required. This result is simply due to the fact that photometry will always be available several orders of magnitude fainter than spectroscopy, due to the physical difficulties in obtaining spectra of faint sources. Thus while our supervised learning represents a vital proof-of-concept over a whole terascale survey, ideally it should be extended with semi-supervised or unsupervised algorithms to fully explore the regions of parameter space that lie beyond the available training spectra. It is worth noting, however, that our training features, the object colors, are largely consistent beyond the limit of the spectroscopic training set. \item The data size is such that integrating the SQL database with D2K via JDBC is impractical, and the data must be stored as flat files. As database engines become more sophisticated, however, it could in the future become possible to offload partial or entire classification rules to a database engine. Doing so, however, would require supercomputing resources for the database engine, which results in an entirely new class of problems. \end{itemize} In moving to the petascale, further issues include: \begin{itemize} \item Conventional hardware, in the form of large clusters of multicore compute nodes, is scaleable to the petascale, however, field-programmable gate arrays (FPGAs), graphical processing units, and Cell processors may be more suited to many data mining tasks, due to their embarrassingly parallel nature. The LCDM group, in collaboration with the Innovative Systems Laboratory at NCSA, has demonstrated results on FPGAs using an SRC-6 MAPE system (Brunner, Kindratenko, \& Myers 2007), which include running a kNN algorithm, although the implementation is not trivial because the algorithm must be rewritten. \item For many applications on the petascale, the performance becomes I/O limited. This is quantified by Bell, Gray, \& Szalay (2006), who apply Amdahl's law (Amdahl 1967) that one byte of memory and one bit per second of I/O are required for each instruction per second, to predict that a petaflop-scale system will require one million disks at a bandwidth of $100~{\mathrm{MB~s^{-1}}}$ per disk. They also state that data should be stored locally (i.e., not transferred over the internet), if the task requires less than 100,000 CPU cycles per byte of data. Many contemporary scientific applications are such that local storage is favored by over an order of magnitude. \end{itemize} | 7 | 10 | 0710.4482 |
0710 | 0710.5167_arXiv.txt | We present results from numerical relativity simulations of equal mass, non-spinning binary black hole inspirals and mergers with initial eccentricities $e\le 0.8$ and coordinate separations $D \ge 12\,M$ of up to 9 orbits (18 gravitational wave cycles). We extract the mass $M_\mathrm{f}$ and spin $a_\mathrm{f}$ of the final black hole and find, for eccentricities $e\lesssim 0.4$, that $a_\mathrm{f}/M_\mathrm{f} \approx 0.69$ and $M_\mathrm{f}/M_\mathrm{adm} \approx 0.96$ are {\em independent} of the initial eccentricity, suggesting that the binary has circularized by the merger time. For $e \gtsim 0.5$, the black holes plunge rather than orbit, and we obtain a maximum spin parameter $a_\mathrm{f}/M_\mathrm{f} \approx 0.72$ around $e = 0.5$. | 7 | 10 | 0710.5167 |
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0710 | 0710.2437_arXiv.txt | % We report the discovery of a 16--20\,$M_\mathrm{Jup}$ radial velocity companion around the very young ($\sim$3\,Myr) brown dwarf candidate Cha\,H$\alpha$\,8 (M5.75--M6.5). Based on high-resolution echelle spectra of Cha\,H$\alpha$\,8 taken between 2000 and 2007 with UVES at the VLT, a companion was detected through RV variability with a semi-amplitude of 1.6\,km\,s$^{-1}$. A Kepler fit to the data yields an orbital period of the companion of 1590~days and an eccentricity of $e$=0.49. A companion minimum mass $M_2\sin i$ between 16 and 20\,$M_\mathrm{Jup}$ is derived when using model-dependent mass estimates for the primary. The mass ratio $q\equiv M_2/M_1$ might be as small as 0.2 and, with a probability of 87\%, it is less than 0.4. Cha\,H$\alpha$\,8 harbors most certainly the lowest mass companion detected so far in a close ($\sim$ 1\,AU) orbit around a brown dwarf or very low-mass star. From the uncertainty in the orbit solution, it cannot completely be ruled out that the companion has a mass in the planetary regime. Its discovery is in any case an important step towards RV planet detections around BDs. Further, Cha\,H$\alpha$\,8 is the fourth known spectroscopic brown dwarf or very low-mass binary system with an RV orbit solution and the second known very young one. | Search for planetary or brown dwarf (BD) companions to BDs are of primary interest for understanding planet and BD formation. There exists no widely accepted model for the formation of BDs (e.g. Luhman et al.~2007). The frequency of BDs in multiple systems is a fundamental parameter in these models. However, it is poorly constrained for close separations: Most of the current surveys for companions to BDs are done by direct (adaptive optics or HST) imaging and are not sensitive to close binaries ($a\la1$\,AU and $a\la10$\,AU for the field and clusters, respectively), and found preferentially close to equal mass systems (e.g. Bouy et al.~2003). Spectroscopic monitoring for radial velocity (RV) variations provides a means to detect close systems. The detection of the first spectroscopic BD binary in the Pleiades, PPl\,15 (Basri \& Mart\'\i n 1999), raised hope to find many more of these systems in the following years. However, the number of confirmed close companions to BDs and very low-mass stars (VLMS, $M \leq 0.1\,M_{\odot}$) is still small. To date, there are three spectroscopic BD binaries known, i.e. for which a spectroscopic orbital solution has been derived: PPl\,15, the very young eclipsing system 2M0535-05 (Stassun, Mathieu \& Valenti 2006), and a binary within the quadruple GJ\,569 (Zapatero Osorio et al.~2004; Simon, Bender \& Prato 2006). They all have a mass ratio close to unity. In particular, no RV planet of a BD/VLMS has been found yet. If BDs can harbor planets at a few AU distance is still unknown. Among the more than 200 extrasolar planets that have been detected around stars by the RV technique, 6 orbit stellar M-dwarfs showing that planets can form also around primaries of substantially lower mass than our Sun. Observations hint that basic ingredients for planet formation are present also for BDs (e.g. Apai et al.~2005). However, the only planet detection around a BD is a very wide 55\,AU system (2M1207, Chauvin et al.~2005), which is presumably formed very differently from the Solar System and RV planets. RV surveys for planets around such faint objects, as BD/VLMS are, require monitoring with high spectral dispersion at 8--10\,m class telescopes. While being expensive in terms of telescope time, this is, nevertheless, extremely important for our understanding of planet and BD formation. We report here on the recent discovery of a very low-mass companion orbiting the BD candidate Cha\,H$\alpha$\,8, which was detected within the course of an RV survey for (planetary and BD) companions to very young BD/VLMS in the Chamaeleon\,I star forming region (Joergens \& Guenther 2001; Joergens 2006; Joergens \& M\"uller 2007). | The companion of Cha\,H$\alpha$\,8 has most certainly a much smaller mass than that of any previously detected close companion of a BD/VLMS. For comparison, all other known spectroscopic BD binaries have mass ratios $>$ 0.6 and the lowest mass BD in these systems has 54\,$M_\mathrm{Jup}$ (Stassun et al.~2006). The discovery of the RV companion of Cha\,H$\alpha$\,8 with its RV semi-amplitude of only 1.6\,km\,s$^{-1}$ is an important step towards RV detections of planets around BD/VLMS. In fact, from the uncertainty of the orbit solution, it cannot be excluded that the companion of Cha\,H$\alpha$\,8 has a mass in the planetary regime ($<13\,M_\mathrm{Jup}$). Follow-up RV measurements at the next phase of periastron will clarify this. With a semi-major axis of about 1\,AU, the companion of Cha\,H$\alpha$\,8 orbits at a much closer orbital distance than most companions detected around BDs so far. In particular, its orbit is much closer than that of recently detected very low-mass companions of BDs, like that of 2M1207 (55\,AU; Chauvin et al.~2005) and that of CHXR\,73 (210\,AU; Luhman 2006; see Luhman, this volume). The favored mechanisms for stellar binary formation, fragmentation of collapsing cloud cores or of massive circumstellar disks, seem to produce preferentially equal mass components, in particular for close separations (e.g. Bate et al.~2003). Thus, they have difficulties to explain the formation of the small mass ratio system Cha\,H$\alpha$\,8. However, we know that close \emph{stellar} binaries with small mass ratios do exist as well (e.g. q=0.2, Prato et al.~2002), and without knowing the exact mechanism by which they form, it might be also an option for Cha\,H$\alpha$\,8. Considering the small mass of the companion of Cha\,H$\alpha$\,8, a planet-like formation could also be possible. Giant planet formation through core accretion might be hampered for low-mass primaries, like M dwarfs, by long formation time scales (Laughlin, Bodenheimer \& Adams 2004; Ida \& Lin 2005), though, recent simulations hint that it can be a faster process than previously anticipated (Alibert et al.~2005). On the other hand, giant planets around M dwarfs might form by disk instability (Boss 2006a, 2006b), at least in low-mass star-forming regions, where there is no photoevaporation of the disk through nearby hot stars (e.g. Cha\,I). The companion of Cha\,H$\alpha$\,8 could have been formed through disk instability, either in situ at 1\,AU or, alternatively, at a larger separation and subsequent inwards migration. Cha\,H$\alpha$\,8 is extremely young (3\,Myr) and its study allows insight into the formation and early evolution at and below the substellar limit. Cha\,H$\alpha$\,8 is only the 2nd known very young BD/VLM spectroscopic binary (after 2M0535-05, Stassun et al.~2006). When combined with angular distance measurements or eclipse detections, spectroscopic binaries allow valuable dynamical mass determinations. The mass is the most important input parameter for evolutionary models, which rely for $<$0.3\,M$_\mathrm{\odot}$, only on the two masses determined for 2M0535-05. In order to measure absolute masses of both components of Cha\,H$\alpha$\,8, it is required to resolve the spectral lines of both components. We will try this with CRIRES/VLT at IR wavelength, were the contrast ratio between primary and secondary is smaller. Having an orbital separation of the order of 13\,milli\,arcsec, the spatial resolution of current imaging instruments is not sufficient to directly resolve Cha\,H$\alpha$\,8. However, it might be possible to detect the astrometric signal caused by the companion, e.g. with NACO/VLT. This would allow measurement of the inclination of the orbital plane and, therefore, breaking the $\sin i$ ambiguity in the companion mass. | 7 | 10 | 0710.2437 |
0710 | 0710.0426_arXiv.txt | We present the results of NICMOS imaging of two massive galaxies photometrically selected to have old stellar populations at $z\sim2.5$. Both galaxies are dominated by apparent disks of old stars, although one of them also has a small bulge comprising about 1/3 of the light at rest-frame 4800 \AA. The presence of massive disks of old stars at high redshift means that at least some massive galaxies in the early universe have formed directly from the dissipative collapse of a large mass of gas. The stars formed in disks like these may have made significant contributions to the stellar populations of massive spheroids at the present epoch. | Considerable observational evidence has built up over the past few years that a substantial fraction of the massive galaxies around us today were already massive at very early epochs. This evidence comes primarily from three sources: \begin{itemize} \item Studies of local massive elliptical galaxies indicate that the stars in the most massive galaxies generally formed very early and over very short time intervals \citep{pee02,tho05,nel05,ren06}. Stars in less massive spheroids formed, on average, later and over longer time spans. \item Massive galaxies in clusters show little evidence for significant evolution up to at least redshift $\sim1$ \citep[\eg][]{deP07,sca07}. \item Direct observations of massive galaxies at redshifts $\gtrsim1.5$ that are dominated by already old stellar populations show that significant numbers of massive galaxies were in place at even earlier epochs \citep*[\eg][]{sto04,mcC04,vanD04,lab05,dad05, red06,pap06,kri06,abr07}. \end{itemize} Although the existence of massive galaxies at high redshifts is now well documented, there have been only a few high-resolution studies of their morphologies (e.g., \citealt{yan03,sto04,zir07,tof07}). Morphologies are important, because they may well retain signs of the formation history of the galaxies. This is particularly true for galaxies that show little or no recent star formation, so that we are able to observe relatively clean examples of the stellar population that formed earliest and that comprises the bulk of the mass of the galaxy. In this paper, we present deep {\it Hubble Space Telescope} ({\it HST}) NICMOS imaging of two galaxies with virtually pure old stellar populations at $z\sim2.5$. In \S~2, we briefly recount how these galaxies were selected. In \S~3, we describe the observations and reduction procedures. In \S~4 and \S~5, we analyze model fits to the images to determine morphologies, and in \S~6 we discuss the implications of our conclusions. We assume a flat cosmology with $H_0 = 73$ \kms\ Mpc$^{-1}$ and $\Omega_M = 0.28$. | } Table~\ref{tab1} summarizes the parameters for the two galaxies. The morphologies of both \ea\ and \eb\ appear to be dominated by disks of old stars. However, the disks are quite different in scale. \eb, at least, also appears to have a small bulge comprising about 1/3 of the total light in the F160W filter ($\sim4800$ \AA, rest frame). We cannot exclude the possibility that \ea\ also has a weak bulge, with up to $\sim15$\% of the total light in the F160W filter; indeed, if the slight apparent difference in morphology between the F160W and F110W images is real, such a difference would seem to favor this possibility. But it is the presence of massive, old disks that continues to give the strongest constraint on formation mechanisms. Such disks also have been seen at redshifts $\sim1.5$, where normal ellipticals with $r^{1/4}$-law profiles are also found \citep*{iye03, cim04, yan04, fu05, sto06, mcG07}. It is difficult to imagine that these massive disks could have formed via any process other than the dissipative collapse of a large cloud of gas. Such disks are also unlikely to have survived major merging events, although the bulge component in \eb\ may testify to either some level of minor merging activity or bulge building via disk instabilities. For galaxies at $z\sim2.5$, the evidence for a dominant old stellar population depends on the inflection in the SED shortward of the $H$ band, and establishing this inflection with optical/near-IR photometry depends on the relatively short baseline from the $H$ to the $K$ band. Furthermore, at the present epoch, essentially all strongly disk-dominated galaxies show evidence for continued star formation. It is therefore not too surprising that claims of passive disks at high redshift should be doubted \citep[\eg][]{pie05}. However, as Fig.~\ref{4c05sed} shows, {\it Spitzer} IRAC data is entirely consistent with the SED of a moderately old stellar population, and no plausible SED incorporating very recent star formation combined with dust would fit the observed photometry. We have recently also obtained IRAC imaging of the field of 4C\,23.56, and our analysis of these data shows that the IRAC photometry falls squarely on our best-fit solar-metallicity BC03 model determined from the optical/near-IR photometry alone: an instantaneous burst with an age of 2.6 Gyr and an extinction $A_V=0.16$ mag \citep{sto07}. Using the more recent preliminary CB07 models, with their improved treatment of AGB stars, we obtain a stellar population age of 2.8 Gyr with $A_V=0$. Again, no plausible model with significant star formation and reddening would fit these data. \begin{deluxetable}{l c c c c} \tablewidth{0pt} \tablecaption{Model Parameters for \ea\ and \eb} \tablehead{ \colhead{Galaxy} & \colhead{Filter} & \colhead{S\'{e}rsic $n$} & \colhead{$r_e$} & \colhead{$r_e$}\\ & & & (\arcsec) & (kpc) } \startdata 4C\,23.56\,ER1 & F110W & $1.03\pm0.10$ & $0\farcs28\pm0\farcs02$ & $2.2\pm0.2$ \\ 4C\,23.56\,ER1 & F160W & $1.52\pm0.06$ & $0\farcs24\pm0\farcs01$ & $1.9\pm0.1$ \\ & & 1.00\tablenotemark{a} & $0\farcs89\pm0\farcs09$ & $7.1\pm0.8$ \\ \raisebox{1.5ex}[0pt]{4C\,05.84\,ER1} & \raisebox{1.5ex}[0pt]{F160W} & 4.00\tablenotemark{a} & $0\farcs37\pm0.20$ & $3.0\pm1.6$ \\ \enddata \tablenotetext{a}{The S\'{e}rsic indices for the two model components for 4C\,05.84\,ER1 have been fixed at these values, which correspond to exponential and $r^{1/4}$-law profiles, respectively.} \label{tab1} \end{deluxetable} Masses for these galaxies can be estimated from the model fits. Assuming solar metallicities and a \citet{cha03} initial mass function, we obtain a mass of $3.9\times10^{11} M_{\odot}$ for \ea\ and $3.3\times10^{11} M_{\odot}$ for \eb (assuming the model at $z=2.4$ with $A_V=0.58$). While the stellar-population age of \eb\ indicates that the last major star-formation episode occurred at $z\sim3.7$, when the universe was $\sim1.8$ Gyr old, \ea\ has a stellar-population age that is formally slightly greater than the age of the universe at $z=2.483$. Clearly, the likely errors in the age determination and the usual caveats regarding the age-metallicity degeneracy mitigate any implied paradox. Nevertheless, this massive galaxy must have formed at a very high redshift. Models with [Fe/H] $= +0.4$ give an age of 1.9 Gyr, but with a significantly worse fit. It therefore seems likely that galaxy formation models will have to allow for the presence of early-forming massive disks. This means that, at least in some dense regions, it has been possible to form $\sim3\times10^{11}$ $M_{\odot}$ of stars within a relatively short time via dissipative collapse and without the aid of major mergers. While our selection criteria have ensured that the galaxies we have discussed here comprise essentially pure old stellar populations, they may well be representative of many massive galaxies at high redshift, most of which would not be in our sample if they retained even tiny amounts of residual star formation or if they had had any significant star formation within a few hundred Myr prior to the epoch at which we observe them. \eb\ has a luminosity and an effective radius that are similar to those of many local galaxies. Our best-fitting S\'{e}rsic model has $r_e=6.3$ kpc. For comparison, for galaxies of similar mass from the Sloan survey with S\'{e}rsic $n<2.5$, \citet{she03} find $r_e=7.2^{+2.9}_{-2.1}$ kpc. This galaxy could become, with passive evolution and perhaps a few minor mergers to increase the bulge-to-disk ratio somewhat, a typical S0 galaxy at the present epoch. On the other hand, we do not see galaxies like \ea\ at the present epoch. By the prescription of \citet{she03}, a low-S\'{e}rsic-index galaxy with the mass of \ea\ would have $r_e=7.6^{+3.1}_{-2.2}$ kpc, but \ea\ actually has $r_e=1.9\pm0.1$ kpc. This means that the stellar mass surface density is much higher than for local galaxies, a result that has also been found for other samples of distant red galaxies \citep[\eg][]{tru06,tof07}. It would seem that the only likely path for such galaxies to evolve to objects consistent with the local population of galaxies is through dissipationless mergers. There is recent evidence that the most massive galaxies in the local universe are likely the result of dry mergers of galaxies with stars that are already old and with very little gas \citep[e.g.,][]{ber07}. With the constraint that these merging components must themselves mostly be fairly massive (to avoid a large dispersion and flattening in the observed color---magnitude relation for present-day massive galaxies, \eg\ \citealt*{bow98}), it seems possible that these early massive disks may well be among the sources for the old stars that today are found in the most massive elliptical galaxies. | 7 | 10 | 0710.0426 |
0710 | 0710.5035_arXiv.txt | Starting from a sample of SDSS quasars appearing also in the 2MASS survey, we study the continuum properties of $\sim 1000$ objects observed in 8 bands, from NIR to UV. We construct the mean spectral energy distribution (SED) and compare and contrast the continua of radio loud (RLQ) and radio quiet (RQQ) objects. The SEDs of the two populations are significantly different, in the sense that RLQs are redder with power law spectral indices $\langle\alpha_{\rm RLQ}\rangle=-0.55\pm0.04$ and $\langle\alpha_{\rm RQQ}\rangle=-0.31\pm0.01$ in the spectral range between $10^{14.5}$ and $10^{15.35}$ Hz. This difference is discussed in terms of different extinctions, different disc temperatures, or slopes of the non-thermal component.% | A substantial effort has been dedicated to construct the spectral energy distributions (SEDs) for sizeable samples of quasars over the whole accessible range of the electromagnetic spectrum % (see e.g.~Sanders et al.~1989; Francis et al.~1991; Elvis et al.~1994; Richards et al.~2006). % However, only a few papers focus on the comparison between the SEDs of radio loud (RLQ) and radio quiet (RQQ) quasars. Elvis et al.~1994 propose an overall spectrum from a sample of 47 quasars, divided in RLQs and RQQs, which shows that no distinction between the SEDs of the two subsamples is apparent in the range $100\mu - 1000$\AA. As noted by these authors, the considered sample is biased towards X-ray and optically bright quasars. Some indication of a possible difference between the NIR to optical colors of RLQs and RQQs in the 2MASS catalogue was reported by Barkhouse \& Hall (2001). Francis, Whiting \& Webster (2000) found that the optical--NIR continuum is significantly redder in radio selected RLQs from the PKS Half-Jansky Flat-Spectrum survey than in optically selected RQQs from the Large Bright Quasar Survey (LBQS). These and other works (e.g.~Kotilainen et al.~2007) indicate that radio loud objects are possibly redder than their radio quiet counterparts, but the samples seriously suffer of selection effects against red radio quiet quasar, because RQQs are mostly collected from optical selected samples (e.g.~the LBQS, the first selection criterion of which is ``blue color'' of candidates; see Hooper, Impey \& Foltz 1997). White et al. suggest that redder quasars from the Sloan Digital Sky Survey (SDSS) are likely to be more radio-powerful than bluer objects, and this is an indication that the red color of radio loud objects is not completely due to the bias against red RQQs in optically selected samples, since the two classes of objects derive from the same survey. Because of the importance of the distinction between RLQs and RQQs in the Unified Models of AGN, these results suggest and motivate a study aimed at investigating the properties of the continuum emission of quasars in the UV to NIR region, in order to compare and contrast the SEDs of radio loud and radio quiet objects. % In Section \ref{secsample} we focus on the quasar sample selection criteria and on the host galaxy contribution. In section \ref{SED} we consider the SED construction and discuss the spectral shape. % The RLQ and RQQ SEDs are then compared and contrasted. In the last section we provide a summary and a discussion of the results. Throughout this paper, we adopt a concordant cosmology with $H_0=70$ ~km~s$^{-1}$~Mpc$^{-1}$, $\Omega_m=0.3$ and $\Omega_{\Lambda}=0.7$. | \label{sum_disc} The aim of the present paper is the study of the continuum emission in the NIR to UV region of a quasar sample, in order to compare and contrast the SEDs of the RLQ and RQQ subsamples. We selected a sample of quasars with both SLOAN and 2MASS detection, to study a spectral range from the NIR (for low-$z$ objects) to the UV (for high-$z$ QSOs). The sample consists of 887 objects, of which 113 RLQs and 774 RQQs. For each subsample we constructed the mean SED and evaluated the average spectral index in the NIR to UV region. The slope of the underlying power-laws is significantly different: the spectral indices of the two population are statistically different at more than 99\% confidence, both for the whole sample and dividing the sample in redshift bins. The difference is present also considering samples well matched in absolute magnitude and redshift. If the blue bump is due to the superposition of black body emissions from an accretion disc, then the color difference between RLQs and RQQs should be interpreted in terms of different mean temperatures, in the sense that RQQs are hotter. \begin{figure} \includegraphics[width=0.45\textwidth]{Fig6} \caption[]{Comparison of the SED of RQQs (dashed line) to the SED of RLQs (dotted line). When RLQs are corrected for an additional dust extinction ($\Delta A_{V}=0.16$mag, solid line), no difference is apparent.} \label{dust} \end{figure} We first examine the possibility that the difference in the SEDs of RLQs and RQQs is due to an enhanced dust extinction in radio loud objects, as it has been suggested e.g.~by Francis et al.~(2000). In Fig. \ref{dust} we compare the SED of RQQs to the SED of RLQs, when RLQs are additionally corrected for dust extinction with respect to RQQs: the offset between the SEDs of the two population is minimized adopting $\Delta A_{V}=0.16$mag. The differences are now completely negligible, supporting the hypothesis that RL objects are more subject to dust extinction than RQQs. The problem would be reconducted to understand why RLQs are more extincted than RQQs. {\it A priori} this could be related to a difference in the inclination angle distribution. % Alternatively, it could be that the conditions of dust production are related to those justifying large radio emission. % We then focus on the possibility that the difference of disc temperature of RLQs and RQQs is real. Since the temperature of the inner disc scales as $M_{\rm BH}^{\,\,\,-0.25}$ (e.g.~Shakura \& Sunyaev 1973), the difference may be attributed to the fact that the black holes of radio loud quasars are supposedly more massive (e.g.~Dunlop et al.~2003; Falomo et al.~2004; Labita et al.~2007). % The color difference may be linked to the BH spin (Stawarz, Sikora \& Lasota 2007), as radio emission is usually ascribed to a faster spinning. However, spinning BHs are expected to have a shorter last stable orbit radius, and then a hotter disc, inconsistently with our results. % Obviously it is possible that the non--thermal (power law) continua of RLQs and RQQs are intrinsically different. Supposing that the non--thermal component accounts for 80\% at $10^{14.5}$Hz, while at $10^{15.35}$Hz (in the Big Blue Bump) the thermal component is dominant and accounts for 80\%, our data are consistent with a picture where the underlying power--laws of RLQs and RQQs have $\alpha_{\rm RLQ}=-1.2$ and $\alpha_{\rm RQQ}=-1.0$ respectively, whereas the thermal bumps are indistinguishable. The problem would be reconducted to understand why the non--thermal continuum of RLQs is softer than RQQs. % It has been suggested (i.e.~Francis et al.~2000) that in radio loud samples there is a significant chance of synchrotron contamination of the rest-frame R--band nuclear luminosities, due to the presence of the relativistic jets. This effect explains the red color of QSOs in high-frequency selected radio samples (i.e. the PKS survey, e.g.~Francis et al.~2001), that suffer from a bias towards pole-on radio sources which are relativistically boosted above the survey flux limit, but no obvious explanation can be invoked to justify the color difference in our sample. | 7 | 10 | 0710.5035 |
0710 | 0710.0424.txt | A description is given of the algorithms implemented in the AstroBEAR adaptive mesh refinement code for ideal magnetohydrodynamics. The code provides several high resolution, shock capturing schemes which are constructed to maintain conserved quantities of the flow in a finite volume sense. Divergence free magnetic field topologies are maintained to machine precision by collating the components of the magnetic field on a cell-interface staggered grid and utilizing the constrained transport approach for integrating the induction equations. The maintenance of magnetic field topologies on adaptive grids is achieved using prolongation and restriction operators which preserve the divergence and curl of the magnetic field across co-located grids of different resolution. The robustness and correctness of the code is demonstrated by comparing the numerical solution of various tests with analytical solutions or previously published numerical solutions obtained by other codes. | The development of efficient and accurate numerical algorithms for astrophysical flow has become of great interest to the astrophysical community. Variable resolution approaches have provided an avenue for efficient simulation of hydrodynamical flow including multi-physical effects which involve substantial variation in length scale. Adaptive Mesh Refinement (AMR) has been recognized as one of the most versatile and efficient approaches to enable the simulation of multi-scale phenomena for which fixed-resolution simulation is either impractical or impossible. AMR discretizations employ a hierarchy of grids at different levels. High resolution is applied only to those regions of the flow which would otherwise be subject to unacceptably large truncation error. The utility of the AMR approach is underscored by the extensive list of codes that are targeted toward astrophysical research which utilize AMR. The list includes AstroBEAR, Enzo \citep{enzo}, Flash \citep{flash}, Orion \citep{truelove,klein,crockett}, Nirvana \citep{nirvana-amr}, Ramses \citep{ramses}, RIEMANN \citep{balsara-amr} and AMRVAC \citep{amrvac} and the list of codes for which the development AMR capability is in progress including Athena \citep{gardiner} and Pluto \citep{pluto}. The simulation of magnetized flow is of particular interest to astrophysical researchers owing to the utility of numerical magnetohydrodynamics (MHD) in modeling a wide range of astrophysical phenomena. The leading line of recent research in this area has focused on the application of higher order Godunov methods to numerical MHD \citep{rj-rs,balsara-rs}. The conservative formulation and proper upwinding employed by these methods enable accurate simulation of strongly supersonic flow. Because of this unique capability, such methods are often referred to ``high resolution shock capturing'' (HRSC) methods. While HRSC methods have long been recognized as the defacto standard for the simulation of supersonic hydrodynamical phenomena, their popularity among researchers interested in magnetized flow has been slowed because standard HRSC approaches to MHD fail to maintain the solenoidality constraint on the magnetic field ($\vec{\nabla} \cdot \vec{B} = 0$). If not corrected, local divergences in the magnetic field arising from this short coming usually grow rapidly, causing anomalous magnetic forces and unphysical plasma transport which eventually destroys the correct dynamics of the flow \citep{brackbill}. Early practitioners of numerical MHD therefore relied heavily on finite difference methods as employed by codes such as Zeus \cite{zeus-mhd} which maintain the solenoidality constraint exactly despite their inferior shock capturing ability \citep{falle}. Later works focused on improved HRSC which either eliminate the development of divergences in the magnetic field, or mitigate the effect of local divergence errors on the dynamics of the flow. In one such approach, a projection operator is devised, usually by solving a Poisson equation, which removes numerical divergences from the grid after each time-step \citep{balsara98,jiang,kim,zachary,rjf}. The primary limitation of this approach is that non-trivial boundary value problems become indeterminate \citep{rj-ct}. In the second so called ``8-wave'' approach first explored by \cite{powell}, alternative formulations of the MHD equations are constructed to prevent the local build up of magnetic divergence by advecting monopoles to other regions of the grid where they are of less consequence to the dynamics of the flow. The work of \cite{dender} augments this approach by adding source terms to the system which act to counter the effect of local divergence in the magnetic field on the dynamics of the flow. A third approach known as constrained transport utilizes a multidimensional, divergence-preserving update procedure for the magnetic field components which are collated on a staggered grid centered on the computational volume interfaces. \citep{evans,rj-ct,balsara-spicer,dai,londrillo,nirvana}. This approach has been shown to provide the most accurate results in the tests of \cite{toth} and \cite{balsara-comp}. The combination of AMR spatial discretizations with HRSC would seem a natural choice in order to satisfy the desire for a high accurate, computationally efficient and versatile strategy for the simulation of magnetized plasma flow. The implementation of the aforementioned adaptations to HRSC methods for MHD in an AMR framework however, poses several challenges. Divergence cleaning schemes utilizing a Poisson projection operator are ill-suited for AMR applications owing to difficulty in handling the projection step along internal boundaries on a patchwork of grids at different resolutions. \cite{powell} found that the application of their 8-wave method on an AMR grid hierarchy, local divergence errors on one level of the AMR hierarchy caused local divergence of comparable magnitude on all levels. The main drawback of this method in an AMR context is that the unphysical effects of local divergences in the magnetic field are not diminished by the application of additional refinement. \cite{crockett}, however, have constructed an approach suitable for AMR applications which combines an approximate projection operator with the divergence advection and dampening the effects of local divergence of \cite{powell} and \cite{dender}. Retaining the divergence-free property of the solution obtained through the application of the constrained transport approach on hierarchical grids requires application of a divergence-preserving prolongation operator which interpolates the magnetic field from a coarse mesh a fine mesh, a divergence-preserving restriction operator which maps the fine mesh magnetic field onto a coarser mesh and furthermore requires that the evolution of the magnetic field be consistent between collocated meshes of different resolution. Two approaches to these challenges have emerged. \cite{balsara-amr} has generalized the divergence free reconstruction procedure of \cite{balsara-spicer} to devise a prolongation operator based on a piece-wise quadratic interpolation procedure that is divergence preserving in the RIEMANN MHD code. \cite{li} present an adaptation of Balsara's procedure that simplifies its implementation for problems involving arbitrary refinement ratios. \cite{toth-roe} have devised a prolongation operator by solving an algebraic system which enforces the maintenance of the volume average curl and divergence between grids of different resolution in the AMR hierarchy. In this paper we provide a concise description of the algorithms and tests of the AstroBEAR HRSC AMR MHD code. AstroBEAR is comprised of several numerical solvers, integration schemes, and radiative cooling modules for astrophysical fluids. The code's AMR capability is derived from the AMR engine of the BEARCLAW boundary embedded adaptive mesh refinement package for conservation laws. This code utilizes the constrained transport approach to adapting HRSC methods to the MHD system of equations. To our knowledge, AstroBEAR is the first AMR code to utilize the prolongation operator of \cite{toth-roe} to maintain the $\vec{\nabla} \cdot \vec{B} = 0$ constraint. By combining multi-physics capabilities relevant to simulation astrophysical plasma flow, AMR, and a wide selection of HRSC integration procedures, AstroBEAR will serve as a valuable research tool. The authors intend that this paper will serve as a reference for future works that apply the code and provide a concise recipe for robust, reliable and accurate HRSC solution strategies for MHD on AMR grid hierarchies. In \S \ref{method} we describe the several HRSC schemes and divergence preservation strategies implemented in the code. In \S \ref{AMR} we provide an overview of the AMR algorithm, highlighting the stages which require special treatment of the magnetic field. In \S \ref{prores} we provide a concise description of the prolongation, restriction and coarse to fine refluxing procedures required to preserve the divergence and consistency of the magnetic field across an AMR hierarchy of grids. In \S \ref{tests} we comment on the results of several test and example problems with particular emphasis on the relative strengths and weaknesses of the various HRSC schemes implemented in the code. In \S \ref{conclusion} we provide a synopsis and discussion of the main results of the paper. | \label{conclusion} The staggered grid constrained transport schemes described in this paper enable the application of high resolution shock capturing methods to magnetized flow. In this paper we have demonstrated that a wide cross section of high resolution shock capturing schemes for general conservation laws may be adapted for magnetized flow while preserving the divergence free constraint on the magnetic field topology exactly by conserving the surface integral of magnetic flux over each computational cell in an upwind fashion. The use of such schemes on multi-resolution AMR grids is encumbered by the requirement that the prolongation and restriction steps preserve the divergence free topology of the magnetic fields. In this paper we have described the application of prolongation and restriction operators which maintain such topologies to machine precision. The numerical schemes discussed here have been implemented and tested in the AstroBEAR adaptive mesh refinement code. The code utilizes a modular design, enabling the user to choose from various methodologies to tailor the numerical integration strategy to the requirements of the application at hand. The robustness of this approach to high resolution, shock capturing MHD on AMR grid structures, and relative advantages of the various numerical schemes implemented in the code are demonstrated in the context of several numerical example problems. The description of the numerical schemes presented in this paper provides a concise recipe for their implementation which will enable the reproduction of these outcomes by other researchers and the interpretation of future works derived from the AstroBEAR code. | 7 | 10 | 0710.0424 |
0710 | 0710.5173_arXiv.txt | The results of a world-wide coordinated observational campaign on the broad-lined Type Ic SN~2003jd are presented. In total, 74 photometric data points and 26 spectra were collected using 11 different telescopes. SN 2003jd is one of the most luminous SN~Ic ever observed. A comparison with other Type Ic supernovae (SNe~Ic) confirms that SN 2003jd represents an intermediate case between broad-line events (2002ap, 2006aj), and highly energetic SNe (1997ef, 1998bw, 2003dh, 2003lw), with an ejected mass of $M_{ej} = 3.0 \pm 1$ M$_{\odot}$ and a kinetic energy of $E_{k}({\rm tot}) = 7_{-2}^{+3} \times ~10^{51}$~erg. SN~2003jd is similar to SN~1998bw in terms of overall luminosity, but it is closer to SNe 2006aj and 2002ap in terms of light-curve shape and spectral evolution. The comparison with other SNe~Ic, suggests that the $V$-band light curves of SNe~Ic can be partially homogenized by introducing a time stretch factor. Finally, due to the similarity of SN~2003jd to the SN~2006aj/XRF~060218 event, we discuss the possible connection of SN 2003jd with a GRB. | \label{parintroduction} In the past decade, the discovery of the connection of some gamma-ray bursts (GRBs) with Type Ic supernovae \citep[SNe~Ic; see][for a review of supernova classification]{filippenko97} boosted interest in the study of this type of SN \citep{galama98,stanek03,hjorth03, malesani04,pian06,soderberg06a,campana06}. In the current paradigm there is a distinction between \emph{normal} SNe~Ic of relatively low kinetic energy \citep[$E_{51} = E_K/(10^{51}~{\rm erg}) \approx 1$;][]{nomoto01}, of which SN~1994I represents a prototype, and high-energy, broad-lined events like SN~1998bw \citep[$E_{51}~>10$;][]{iwamoto98,maeda02} which are associated with GRBs. To understand whether these are separate subclasses or extreme cases of a continuous distribution, it is necessary to study in detail intermediate cases such as the broad-lined SN 2002ap \citep{galyam02,foley03}, which had an intermediate explosion energy \citep[$E_{51} = 4-10$; ][]{mazzali02} and was not connected with a GRB, or the broad-lined, low energy ($E_{51} \approx 2$) SN 2006aj associated with XRF 060218 \citep{mazzali06a,mazzali07,maeda07}. Here we discuss the case of SN~2003jd, a broad-lined object showing clear evidence of an asymmetric explosion \citep{mazzali05} but with no confirmed GRB connection \citep{gcn2434,gcn2439}. Asymmetry is probably a key factor in understanding the diversity of SNe~Ic. SN 2003jd was discovered \citep{burket03} on 2003 Oct. 25 (UT dates are used throughout this paper) with the Katzman Automatic Imaging Telescope (KAIT) during the Lick Observatory Supernova Search \citep{filippenko01,filippenko05}. It is located at $\alpha~$ = $23^h 21^m 03^s.38$ and $\delta$ =-04$\degr 53' 45''.5$ (equinox J2000), which is $8.3''$ E and $7.7''$ S of the centre of the Sb spiral galaxy MCG-01-59-21 \citep{vandenbergh05}. \begin{figure} \psfig{figure=imagesn03jd_3.ps,width=8cm,height=8cm} \caption{The field of SN 2003jd. $B$-band TNG image (14 Nov. 2003) 16.4~d after the $B$ maximum.} \label{figseqstar} \end{figure} The SN was not visible on 16 Oct. in a KAIT unfiltered image (mag $<$ 19), which sets a tight limit to the explosion epoch ($\leq$ 13~d before $B$ maximum). On 28 Oct., SN 2003jd was classified as a peculiar Type Ic event with very broad features analogous to those of SN 1998bw and SN 2002ap \citep{filippenko03}. The new interest in this kind of event prompted an intensive follow-up campaign at different observing sites, lasting for about three months. Few more observations in the late nebular phase were taken about one year later. This paper, which presents and discusses these observations, is organized as follows. In \S \ref{parobs} we describe the observations and data-reduction techniques. We describe the photometric and spectroscopic data of SN 2003jd in \S \ref{parphot} and \S \ref{parspe}. In \S \ref{parintrinsic} we characterize the host galaxy, and in \S \ref{parcomparison} we compare SN 2003jd with a sample of well-studied SNe~Ic: 1994I, 1998bw, 2002ap, 2004aw and 2006aj. \S \ref{parametribolo} presents the bolometric light curve of SN 2003jd, computed using all available photometric data. The similarity with SN 2006aj/GRB 060218 is discussed in \S \ref{parmissinggrb}, together with a discussion of the (lack of) evidence for an associated GRB. Finally, in \S \ref{parphyparam}, using simple bolometric light-curve modelling, we derive some basic explosion parameters. | \begin{tabular}{ll} \hline Parent galaxy & MCG-01-59-021 \\ Galaxy type & Sb pec$^a$ \\ RA (2000) & $23^h 21^m 03^s.38$ \\ Dec (2000) & $-04^\circ53'45.5''$ \\ Recession velocity [\kms] & 5625$^b$ \\ Distance modulus ($H_0 = 72$) & $34.46 \pm 0.20$ mag \\ $E(B-V)_{host}$ & $0.10^{+0.10}_{-0.05}$ mag$^c$\\ $E(B-V)_{Gal}$ & 0.044 mag$^d$ \\ Offset from nucleus & $8.3''$ E, $7.7''$ N \\ \hline \end{tabular}\\ $^a$van den Bergh et al. 2005. \\ $^b$LEDA, corrected for Local Group infall (208 \kms).\\ $^c$Calculated from the equivalent width of the \NaID~ lines \citep{turatto03}.\\ $^d$Schlegel et al. (1998). \end{table} The spectra of the host galaxy taken at +788~d, long after the fading of the SN, can be used to estimate the metallicity and star-formation rate (SFR) of the region were SN~2003jd exploded. Following \cite{pettini04}, from the [O\,{\sc iii}] and [N\,{\sc ii}] line strengths we derive an oxygen abundance of $12 + {\rm log}(O/H) = 8.4 \pm 0.1$ dex. A slightly higher value ($12 + {\rm log}(O/H) = 8.7 \pm 0.1$ dex) is obtained using the $R_{23}$ index (([O\,{\sc iii}] + [O\,{\sc ii}])/H$\beta$) of \cite{pagel79}. The two estimates bracket the measurement by \cite{modjaz07}; they confirm that the metallicity of the SN~2003jd host is below average for local galaxies and similar to that of the SN~1998bw host \citep{modjaz07}. On the other hand, MCG-01-59-021 is a luminous galaxy, $M_{B} = -20.3$ mag (LEDA), which is somewhat at odds with the medium to low metallicity \citep[see Fig.~5 of ][ and Fig.~1 of Prieto et al 2007]{modjaz07}. The SFR at the SN location can be estimated from the \Ha\ flux, $7.1 \pm 0.5~erg~s^{-1}~cm^{-2}$. Following \cite{kennicutt98}, we estimate a SFR $\approx 0.04$~M$_{\odot}$ yr$^{-1}$, which is in fair agreement with the value derived by \cite{modjaz07} (SFR = 0.07~M$_{\odot}$ yr$^{-1}$). This SFR is typical for \HII~ regions, and similar for instance to that derived for the environment of SN 2006aj: SFR $\approx 0.06$~M$_{\odot}$ yr$^{-1}$ \citep{pian06}. We have presented optical photometry and spectroscopy of SN~2003jd spanning from 3~d before $B$-band maximum to $\sim$400~d after maximum. SN 2003jd shows the typical spectral features of a SN~Ic, but broadened as in the case of SNe 2002ap and 2006aj. The light curves are similar in shape to those of SNe 2002ap and 2006aj, but the peak luminosity is rather similar to that of SN 1998bw. The comparison with a sample of well-studied SNe~Ic confirms the heterogeneity of this SN class. However, the application of time stretching (as for SNe~Ia) helps in homogenizing this class of SNe, reducing the differences among the light curves and facilitating spectral comparisons. Different explosion energies and progenitor masses are likely to be the main reasons for the discrepancies. The different viewing angles from which the SNe are observed (given the evidence for asymmetric explosions) may also explain some of the differences, but this effect is probably not large, at least for the light-curve shape \citep{maeda06b}. It may, however, affect the spectra \citep{tanaka07}. The $E_{k}/M_{ej}$ ratio seems to be the main factor in producing broad features. The similarity to SN 2006aj, the presence of broad features, and the asymmetry shown by the oxygen double peak in the nebular spectra are in favour of an asymmetric explosion oriented away from the line of site. The radio observations argue that the asymmetry did not extended to relativistic velocities, although uncertainties in the CSM mean that 2003jd could still be a mis-directed GRB. Finally, we used a simple model for the bolometric light curve to obtain the main physical parameters of SN 2003jd. The derived values confirm that SN 2003jd is a broad-line SN similar to SN 2002ap and SN 2006aj, but with a larger ejected mass ($M_{ej} = 3.0 \pm 1.0$~M$_\odot$) and kinetic energy ($E_{k}(tot) = 7_{-2}^{+3} \times E_{51}$ erg) and producing a large quantity of $^{56}Ni$ ($M_{Ni} = 0.36 \pm 0.04$~M$_\odot$). Comparing the physical parameters of SN 2003jd with those of other SNe~Ic \citep{nomoto07}, SN 2003jd appears to represent a link between broad-lined SNe (2002ap and 2006aj) and GRB-associated supernovae (SNe 1997ef, 1998bw, 2003dh, and 2003lw). | 7 | 10 | 0710.5173 |
0710 | 0710.5345_arXiv.txt | We extend and apply a model-independent analysis method developed earlier by Daly \& Djorgovski to new samples of supernova standard candles, radio galaxy and cluster standard rulers, and use it to constrain physical properties of the dark energy as functions of redshift. Similar results are obtained for the radio galaxy and supernova data sets, which rely upon completely independent methods, suggesting that systematic errors are relatively small for both types of distances; distances to SZ clusters show a scatter which cannot be explained by the quoted measurement errors. The first and second derivatives of the distance are compared directly with predictions in a standard model based on General Relativity. The good agreement indicates that General Relativity provides an accurate description of the data on look-back time scales of about ten billion years. The first and second derivatives are combined to obtain the acceleration parameter $q(z)$, assuming only the validity of the Robertson-Walker metric, independent of a theory of gravity and of the physical nature of the dark energy. The data are analyzed using a sliding window fit and using fits in independent redshift bins. The acceleration of the universe at the current epoch is indicated by the sliding window fit analysis. The effect of non-zero space curvature on $q(z)$ is explored; for a plausible range of values of $\Omega_k$ the effect is small and causes a to shift to the redshift at which the universe transitions from deceleration to acceleration. We solve for the pressure, energy density, equation of state, and potential and kinetic energy of the dark energy as functions of redshift assuming that General Relativity is the correct theory of gravity. Results obtained using a sliding window fit indicate that a cosmological constant in a spatially flat universe provides a good description of each of these quantities over the redshift range from zero to about one. We define a new function, the dark energy indicator, in terms of the first and second derivatives of the coordinate distance and show how this can be used to measure deviations of $w$ from $-1$ and to obtain a new and independent measure of $\Omega_m$. | Understanding of the physical nature of the dark energy which appears to be driving the accelerated expansion of the universe is among the most pressing and important topics in cosmology today. Studies of the expansion history of the universe allow us to constrain the physical nature of its matter and energy constituents. One way that the expansion and acceleration history of the universe can be studied is through the use of a set of coordinate distances and redshifts for some standard set of objects. Type Ia supernovae provide a modified standard candle (e.g. Phillips 1993, Hamuy et al. 1995) that allow the distance modulus, luminosity distance, and coordinate distance to each source to be determined. The recent data sets presented by Astier et al. (2006), Riess et al. (2007), Wood-Vasey et al. (2007), and Davis et al. (2007) have been analyzed by these groups and compared with numerous models by other researchers. In a novel, largely model-independent approach to this problem, it was shown by Daly \& Djorgovski (2003) that the first and second derivatives of the coordinate distance with respect to redshift could be obtained from the coordinate distances and combined to solve for the expansion rate $H(z)/H_0$ and acceleration rate $q(z)$ of the universe. The functions $H(y^{\prime})$ and $q(y^{\prime},y^{\prime \prime})$ are exact, that is, they are not obtained by expansions in terms of derivatives about some point. The only assumptions are that the universe is described by a Robertson-Walker metric and has zero space curvature. The results are independent of the contents of the universe and their physical properties, and even independent of whether General Relativity provides an accurate description of the universe. Here, we drop the assumption of zero space curvature; it turns out that the deceleration parameter at a redshift of zero, $q_0$, remains the same, independent of whether space curvature is zero or not. In this paper we expand on the previous analysis done by Daly \& Djorgovski (2003, 2004). First, we use updated and expanded data sets, as described in section 2.1. Second, we introduce a more direct way to compare the model-independent results obtained from the data with predictions; this is done by directly comparing the first and second derivatives of the coordinate distance with respect to redshift to predicted values in various models, as described in section 2.2. Third, we analyze the data using both a sliding window fit and fits in independent redshift bins. To solve for the physical properties of the dark energy as functions of redshift, a theory of gravity must be specified. To determine the properties of the dark energy, General Relativity is taken to be the correct theory of gravity, allowing us to solve for the pressure, energy density, and equation of state of the dark energy as a function of redshift in section 2.4. Fourth, in section 2.4, we introduce a way to solve for the potential and kinetic energy densities of the dark energy as functions of redshift. In addition, we define a new function, the dark energy indicator, which provides a measure of deviations of $w$ from $-1$ and a new and independent measure of $\Omega_m$. Fifth, in section 2.3, these derivatives are combined to solve for the expansion and acceleration rates of the universe as functions of redshift for both zero and non-zero space curvature; in our previous work we have not considered the effects of non-zero space curvature. The only assumption that must be made to obtain the functions $H(z)/H_0$ and $q(z)$ from the data are that the Robertson-Walker metric is valid in our universe. A discussion and conclusions follow in section 3. | The work presented here improves and extends our previous results. First, expanded and improved data sets are considered: three supernova samples and one radio galaxy sample. The radio galaxy data set has 11 new sources, increasing its size to 30 sources, and the supernovae data sets have increased substantially in size and quality. In addition, SZ cluster distances and gamma-ray burst distances are considered. The dimensionless coordinate distances (obtained directly from the data), and first and second derivatives of the distance are obtained as functions of redshift using a sliding window fit. The good agreement obtained using supernovae and radio galaxies, two completely independent methods, with sources that cover similar redshift ranges, suggests that neither method is strongly affected by systematic effects, and that each method provides a reliable cosmological tool. The first and second derivatives of the distance are combined to obtain the acceleration parameter $q(z)$, allowing for non-zero space curvature. It is shown that the zero redshift value of $q(z)$, $q_o$, is independent of space curvature, and can be obtained from the first and second derivatives of the coordinate distance. Thus, $q_0$, which indicates whether the universe is accelerating at the current epoch, can be obtained directly from the supernova and radio galaxy data; our determinations of $q(z)$ only relies upon the validity of the Robertson-Walker line element, and is independent of a theory of gravity, and the contents of the universe. Each of the supernova samples, analyzed using a sliding window fit, indicate that the universe is accelerating today independent of space curvature, independent of whether General Relativity is the correct theory of gravity, and of the contents of the universe. The effect of non-zero space curvature on $q(z)$ is to shift the redshift at which the universe transitions from acceleration to deceleration, moving this to lower redshift for negative space curvature and to higher redshift for positive space. The zero redshift values of $q$ obtained using a sliding window fit. for the Davis et al. (2007) supernova sample is $q_0(192SN)= -0.48 \pm 0.11$ and that obtained for the radio galaxy sample of Daly et al. (2007) is $q_0(30RG)= -0.65 \pm 0.53$ indicating that the universe is accelerating at the current epoch. The data were also binned so that only certain subsets of the data were used to solve for $y^{\prime}$, $y^{\prime \prime}$, $H(z)/H_0$, and $q(z)$. The results for $y^{\prime}$ and $H(z)/H_0$ indicate that the standard LCDM model provides a good description of the data. The results for $y^{\prime \prime}$ and $q(z)$ are consistent with the standard LCDM model, but do not independently confirm the model or the acceleration of the universe. In addition to the evaluation of the standard cosmological parameters, in an even more direct approach, we compared $y^{\prime}$ and $y^{\prime \prime}$ obtained from the fits to the data to model predictions. Comparisons of $y^{\prime}$ and $y^{\prime \prime}$ with predictions based on General Relativity indicate that General Relativity provides an accurate description of the data on look-back time scales of about ten billion years, thus providing a very large scale test of General Relativity. Another new approach is that the data were analyzed using both a sliding window fit and fits in independent redshift bins. The fits in statistically independent redshift bins are broadly consistent with the sliding window fits, but are generally noisier (as expected). We also explored the effects of non-zero space curvature on determinations of $H(z)$ and $q(z)$. It is shown that the zero redshift value of $q$, obtained by applying equation (4) to $y^{\prime}$ and $y^{\prime \prime}$, is independent of space curvature. This means that our method can be used to determine $q_0$, and thus the degree to which the universe is accelerating at the current epoch, with only one assumption, that the Robertson-Walker line element is valid. In addition, it is found that the effect of space curvature on the shape of $H(z)$ and $q(z)$ is small, relative to the uncertainties arising from the measurement errors. After determining the expansion and acceleration rates of the universe as functions of redshift independent of a theory of gravity, we solve for the pressure, energy density, equation of state, and potential and kinetic energy of the dark energy as functions of redshift assuming that General Relativity is the correct theory of gravity. We also define a new function, the dark energy indicator $s$, which provides a measure of deviations of the equation of state of the dark energy $w$ from $-1$, and provides a new and independent measure of $\Omega_m$ if $w=-1$. The results obtained using a sliding window fit indicate that a cosmological constant in a spatially flat universe provides a good description of each of these quantities over the redshift range from zero to one. The zero redshift values of these quantities obtained with the Davis et al. (2007) supernovae sample are $P_{DE,0}/\rho_{0c} = -0.64 \pm 0.10$, $\rho_{DE,0}/\rho_{0c} = 0.67 \pm 0.05$, $w_{DE,0} = -0.95 \pm 0.08$, $V_{DE,0}/\rho_{0c} = 0.65 \pm 0.05$, $K_{DE,0}/\rho_{0c} = 0.01 \pm 0.03$, and $s_0 = -0.50 \pm 0.08$. In the standard Lambda-Cold Dark Matter Model, $\Omega_{\Lambda} = - P_0/\rho_{0c} = 0.64 \pm 0.1$, obtained using the first and second derivatives of the coordinate distance, provides an independent measure of $\Omega_{\Lambda}$. In addition, in this model, $w=-1$, so $s$ provides a measure of $\Omega_m$, and the value obtained here using the first and second derivatives of the coordinate distance, is $\Omega_m = 0.33 \pm 0.05$. Overall, the shapes of the pressure, energy density, equation of state, and other parameters as functions are redshift are consistent with those predicted in a standard LCDM model. There is a tantalizing hint that there may be divations from the standard model at high redshift; more observations at high redshift will be needed to investigate this further. The results obtained using fits in independent redshift bins are consistent with the standard LCDM model, but do not independently confirm the model. | 7 | 10 | 0710.5345 |
0710 | 0710.3755_arXiv.txt | We argue that the Higgs boson of the Standard Model can lead to inflation and produce cosmological perturbations in accordance with observations. An essential requirement is the non-minimal coupling of the Higgs scalar field to gravity; no new particle besides already present in the electroweak theory is required. | The fact that our universe is almost flat, homogeneous and isotropic is often considered as a strong indication that the Standard Model (SM) of elementary particles is not complete. Indeed, these puzzles, together with the problem of generation of (almost) scale invariant spectrum of perturbations, necessary for structure formation, are most elegantly solved by inflation \cite{Starobinsky:1979ty,Starobinsky:1980te,Mukhanov:1981xt,Guth:1980zm,% Linde:1981mu,Albrecht:1982wi}. The majority of present models of inflation require an introduction of an additional scalar---the ``inflaton''. This hypothetical particle may appear in a natural or not so natural way in different extensions of the SM, involving Grand Unified Theories (GUTs), supersymmetry, string theory, extra dimensions, etc. Inflaton properties are constrained by the observations of fluctuations of the Cosmic Microwave Background (CMB) and the matter distribution in the universe. Though the mass and the interaction of the inflaton with matter fields are not fixed, the well known considerations prefer a heavy scalar field with a mass $\sim \unit[10^{13}]{GeV}$ and extremely small self-interacting quartic coupling constant $\lambda \sim 10^{-13}$ \cite{Linde:1983gd}. This value of the mass is close to the GUT scale, which is often considered as an argument in favour of existence of new physics between the electroweak and Planck scales. The aim of the present Letter is to demonstrate that the SM itself can give rise to inflation. The spectral index and the amplitude of tensor perturbations can be predicted and be used to distinguish this possibility from other models for inflation; these parameters for the SM fall within the $1\sigma$ confidence contours of the WMAP-3 observations \cite{Spergel:2006hy}. To explain our main idea, consider Lagrangian of the SM non-minimally coupled to gravity, \begin{equation} \label{main} L_{\mathrm{tot}}= L_{\mathrm{SM}} - \frac{M^2}{2} R -\xi H^\dagger HR \;, \end{equation} where $L_{\mathrm{SM}}$ is the SM part, $M$ is some mass parameter, $R$ is the scalar curvature, $H$ is the Higgs field, and $\xi$ is an unknown constant to be fixed later.\footnote{In our notations the conformal coupling is $\xi=-1/6$.} The third term in (\ref{main}) is in fact required by the renormalization properties of the scalar field in a curved space-time background \cite{Birrell:1982ix}. If $\xi=0$, the coupling of the Higgs field to gravity is said to be ``minimal''. Then $M$ can be identified with Planck scale $M_P$ related to the Newton's constant as $M_P=(8\pi G_N)^{-1/2}=\unit[2.4\times 10^{18}]{GeV}$. This model has ``good'' particle physics phenomenology but gives ``bad'' inflation since the self-coupling of the Higgs field is too large and matter fluctuations are many orders of magnitude larger than those observed. Another extreme is to put $M$ to zero and consider the ``induced'' gravity \cite{Zee:1978wi,Smolin:1979uz,Spokoiny:1984bd,Fakir:1990iu,Salopek:1988qh}, in which the electroweak symmetry breaking generates the Planck mass \cite{vanderBij:1993hx,CervantesCota:1995tz,Bij1995}. This happens if $\sqrt{\xi}\sim1/(\sqrt{G_N} M_W)\sim10^{17}$, where $M_W\sim\unit[100]{GeV}$ is the electroweak scale. This model may give ``good'' inflation \cite{Spokoiny:1984bd,Fakir:1990iu,Salopek:1988qh,Kaiser:1994wj,% Kaiser:1994vs,Komatsu:1999mt} even if the scalar self-coupling is of the order of one, but most probably fails to describe particle physics experiments. Indeed, the Higgs field in this case almost completely decouples from other fields of the SM\footnote{This can be seen most easily by rewriting the Lagrangian (\ref{main}), given in the Jordan frame, to the Einstein frame, see also below.} \cite{vanderBij:1993hx,CervantesCota:1995tz,Bij1995}, which corresponds formally to the infinite Higgs mass $m_H$. This is in conflict with the precision tests of the electroweak theory which tell that $m_H$ must be below $\unit[285]{GeV}$ \cite{:2005ema} or even \unit[200]{GeV} \cite{PDG2007} if less conservative point of view is taken. These arguments indicate that there may exist some intermediate choice of $M$ and $\xi$ which is ``good'' for particle physics and for inflation at the same time. Indeed, if the parameter $\xi$ is sufficiently small, $\sqrt{\xi} \lll 10^{17}$, we are very far from the regime of induced gravity and the low energy limit of the theory (\ref{main}) is just the SM with the usual Higgs boson. At the same time, if $\xi$ is sufficiently large, $\xi \gg 1$, the scalar field behaviour, relevant for chaotic inflation scenario \cite{Linde:1983gd}, drastically changes, and successful inflation becomes possible. We should note, that models of chaotic inflation with both nonzero $M$ and $\xi$ were considered in literature \cite{Spokoiny:1984bd,Futamase:1987ua,Salopek:1988qh,Fakir1990,Kaiser:1994vs,% Libanov1998,Komatsu:1999mt}, but in the context of either GUT or with an additional inflaton having nothing to do with the Higgs field of the Standard Model. The Letter is organised as follows. We start from discussion of inflation in the model, and use the slow-roll approximation to find the perturbation spectra parameters. Then we will argue in Section \ref{sec:radcorr} that quantum corrections are unlikely to spoil the classical analysis we used in Section \ref{sec:cmb}. We conclude in Section~\ref{sec:concl}. | \label{sec:concl} In this Letter we argued that inflation can be a natural consequence of the Standard Model, rather than an indication of its weakness. The price to pay is very modest---a non-minimal coupling of the Higgs field to gravity. An interesting consequence of this hypothesis is that the amplitude of scalar perturbations is proportional to the square of the Higgs mass (at fixed $\xi$), revealing a non-trivial connection between electroweak symmetry breaking and the structure of the universe. The specific prediction of the inflationary parameters (spectral index and tensor-to-scalar ratio) can distinguish it from other models (based, e.g.\ on inflaton with quadratic potential), provided these parameters are determined with better accuracy. The inflation mechanism we discussed has in fact a general character and can be used in many extensions of the SM. Thus, the $\nu$MSM of \cite{Asaka:2005an,Asaka:2005pn} (SM plus three light fermionic singlets) can explain simultaneously neutrino masses, dark matter, baryon asymmetry of the universe and inflation without introducing any additional particles (the $\nu$MSM with the inflaton was considered in \cite{Shaposhnikov:2006xi}). This provides an extra argument in favour of absence of a new energy scale between the electroweak and Planck scales, advocated in \cite{Shaposhnikov:2007nj}. | 7 | 10 | 0710.3755 |
0710 | 0710.4519_arXiv.txt | In the present paper we report on the difference in angular sizes between radio-loud and radio-quiet CMEs. For this purpose we compiled these two samples of events using Wind/WAVES and SOHO/LASCO observations obtained during 1996-2005. It is shown that the radio-loud CMEs are almost two times wider than the radio-quiet CMEs (considering expanding parts of CMEs). Furthermore we show that the radio-quiet CMEs have a narrow expanding bright part with a large extended diffusive structure. These results were obtained by measuring the CME widths in three different ways. | The relation between coronal mass ejections (CMEs) and type II radio bursts has been studied for a long time, but is not fully understood (see Gopalswamy (2006) for a recent review). Properties of the driving CMEs and the ambient medium through which the CMEs drive shocks show a large variability, which seems to contribute to the difficulties faced in understanding them (Gopalswamy {\it et al.}, 2001). One of the major issues has been the lack of the type II radio emission in the metric (Sheeley {\it et al.}, 1984) and decameter-hectometric (DH) wavelengths (Gopalswamy {\it et al.}, 2001) even for CMEs with speeds exceeding 1000 km s$^{-1}$. Recently, Gopalswamy {\it et al.} (2007) performed a systematic investigation of fast and wide (FW) CMEs that clearly lacked the metric and DH type II radio emission (``radio-quiet''CMEs) and compared them with the ones (``radio-loud'' CMEs) producing detectable radio type II. It was found that the radio-quiet CMEs can be distinguished from the radio-loud CMEs in three aspects: (1) speeds and widths, (2) a fraction of halo CMEs, and (3) solar source location of the CMEs. The radio-quiet CMEs are generally slower and narrower than the radio-loud ones. The fraction of halo CMEs is much larger for the radio-loud CMEs, which is related to the fact that the radio-quiet CMEs are narrower on the average. It is also known that halo CMEs are also faster and wider on the average (Yashiro \emph{et al.}, 2004; Michalek, Gopalswamy, and Yashiro, 2003; Gopalswamy, 2004). When the source locations were examined, Gopalswamy (2006) and Gopalswamy \emph{et al.} (2007) found that more than half of the radio-quiet CMEs were back-sided, while only a small fraction (25) of the radio-loud CMEs were back-sided. They attributed this result to the possibility that only a small fraction of the shock surface is visible to the observer, thereby reducing the possibility of detecting significant radio emission. A fast but narrow CME may have a similar limitation because the CME cross-section and hence the shock surface area are expected to be smaller. One of the suggestions made in Gopalswamy \emph{et al.} (2007) is that most of the radio-quiet CMEs may have a narrow bright part with extended diffuse structure. The purpose of this paper is to examine the evolution of the width of radio-quiet and radio-loud CMEs and compare them to confirm the smaller width of CMEs as a contributor to radio quietness. In Section~2 the procedure for obtaining the samples of the radio-loud and radio-quiet CMEs is presented. In this section three different methods for the determination of CME widths are also explained. In Section~3, we use the measured CME widths to show the spatial difference between the radio-loud and radio-quiet CME populations. | In this study we examined the spatial difference between the radio-loud and radio-quiet CMEs. To get a reliable result, we determined the angular widths of the radio-loud and radio-quiet CMEs using three different methods. In all cases the radio-loud CMEs are wider than the radio-quiet CMEs. When we compare the expanding parts of CMEs (which are responsible for II type radio emission) the width difference between the events is the largest. The expanding structures of the radio-quiet CMEs are narrower ($\approx$40$\%$) in comparison with those of the radio-loud CMEs. The expanding structures of CMEs are also much narrower in comparison with their total widths, especially for the radio-quite events. The ratios of the average catalog width to the average main body width are equal about 2.0 and 2.7 for the radio-loud and radio-quiet CMEs, respectively. The catalog widths for the radio-quiet CMEs are almost three times bigger in comparison with the widths of expanding structures. This means that the radio-quiet CMEs have a narrow expanding bright part with a large extended diffusive structure. It is clear that the spatial size of CME could be one of the most important factors defining the presence of type II radio emission. Our results proved the previous considerations ({\it e.g.} Gopalswamy {\it et al.}, 2001, 2005; Pick and Maia, 2005; Subramanian and Ebenezer, 2006). It is commonly accepted that type II radio bursts are radio signatures of coronal MHD-shock waves (Uchida, 1960; Wild, 1962). Flare-related blast waves and shock driven by CMEs have been considered as two possible pistons of metric type II bursts (see {\it e.g.} Cliver, Webb, and Howard, 1999), while the DH and longer wavelength bursts due to CME-driven shock. CMEless type II bursts (Sheeley {\it et al.}, 1984) and the discrepancy between the metric and IP type ( Reiner {\it et al.}, 2001) were used to argue against the same shock causing the metric and IP type bursts. Gopalswamy (2006) demonstrated that both these discrepancies could be explained. It seems that CME-driven shock works for the entire interplanetary space but additional mechanism (blast waves) may operate for a narrow region ($\approx 1R_{\bigodot}$) close to the solar surface. In both cases, the width of CMEs plays an important role in generation of fast particles and radio bursts. Wider a given CME (more energetic event) wider a shock front and larger area where particles can gain energy. Additionally, larger CMEs could in bigger degree destruct magnetic structures in corona and amplify radio emission ({\it e.g.} Raymond {\it et al.}, 2000; Pick {\it et al.}, 2006). This scenario is confirmed by strong correlation between complex type III and type II radio bursts associated with CMEs (Cane, Erickson, and Prestage, 2002; Gopalswamy, 2004). | 7 | 10 | 0710.4519 |
0710 | 0710.0801_arXiv.txt | Clusters of galaxies are sites of acceleration of charged particles and sources of non-thermal radiation. We report on new constraints on the population of cosmic rays in the Intra Cluster Medium (ICM) obtained via radio observations of a fairly large sample of massive, X--ray luminous, galaxy clusters in the redshift interval 0.2--0.4. The bulk of the observed galaxy clusters does not show any hint of Mpc scale synchrotron radio emission at the cluster center (Radio Halo). We obtained solid upper limits to the diffuse radio emission and discuss their implications for the models for the origin of Radio Halos. Our measurements allow us to derive also a limit to the content of cosmic ray protons in the ICM. Assuming spectral indices of these protons $\delta =2.1-2.4$ and $\mu$G level magnetic fields, as from Rotation Measures, these limits are one order of magnitude deeper than present EGRET upper limits, while they are less stringent for steeper spectra and lower magnetic fields. | Clusters of galaxies are ideal astrophysical environments for particle acceleration. Large scale shocks which form during the process of cluster formation are believed to be efficient particle accelerators (e.g. Sarazin 1999; Gabici \& Blasi 2003; Ryu et al. 2003; Pfrommer et al. 2006). Cosmic rays (CRs) can also be injected into the ICM from ordinary galaxies and AGNs (e.g. V\"olk \& Atoyan 1999) and turbulent eddies may contribute to the particle acceleration process (e.g. Brunetti \& Lazarian 2007). CRs accelerated within the cluster volume would then be confined for cosmological times and the bulk of their energy is expected in protons since they have radiative and collisional life--times much longer than those of the electrons (e.g. Blasi et al. 2007, for a review). While present gamma ray observations can only provide upper limits to the average energy density of CR protons in the ICM (e.g. Reimer et al. 2004), the presence of relativistic electrons in a number of clusters has been ascertained via the detection of a tenuous synchrotron radio emission: giant Radio Halos (RHs) and mini-Radio Halos, fairly symmetric sources at the cluster center, and Radio Relics, elongated sources at the cluster periphery (e.g. Feretti \& Giovannini 2007). It is customary to classify the models for the origin of RHs in {\it secondary electron} (e.g. Blasi \& Colafrancesco 1999) and {\it reacceleration models} (e.g. Brunetti et al. 2001; Petrosian 2001), depending on whether the radiating electrons are produced as secondary products of hadronic interactions or reaccelerated by turbulence from a pre-existing population of non-thermal seeds in the ICM, respectively. These models predict a different connection between radio and X-ray properties of clusters which are discussed in this Letter and compared with new observations: in Sect.2 we review the expectations of the different models, in Sect.3 we briefly present the radio observations of our cluster sample, and in Sects.4 \& 5 we report and discuss our results. Concordance ($H_o$=70, $\Omega_m$=0.3, $\Omega_{\Lambda}$=0.7) cosmology is used. | We have reported on constraints on the origin of RHs and on the CR content in the ICM obtained via radio observations of a fairly large sample of X--ray luminous clusters at $z = 0.2-0.4$. \noindent In the bulk of these clusters we do not find evidence of Mpc--scale radio emission at the level of RHs. Our conclusions become even more stringent considering radio emission on cluster--core scale, typical of smaller RHs and mini-Halos. \noindent We firmly confirm that RH--clusters follow a {\it physical} correlation between synchrotron and X--ray luminosities. We find that clusters have a bimodal distribution in the $P_{1.4}$--$L_X$ plane (Fig.~4); this is in line with the expectation of the {\it re-acceleration scenario}. On the other hand, in order to reconcile these observations with expectations from {\it secondary} models strong dissipation of the magnetic field in the clusters with no radio emission is necessary. \noindent Our measurements allow us to also derive simple limits on the presence of CR protons in the ICM (Fig.~5). In the case of relatively flat spectral energy distribution of these CRs stringent upper limits can be obtained: the energy density of CRs should be $\leq 1$\% of the thermal energy in case of $\geq \mu$G field strength. This would make problematic the detection of gamma rays from $\pi^o$--decay in clusters with GLAST. On the other hand, by assuming steeper spectral energy distributions of these CRs (or lower magnetic fields) our limits become less stringent. | 7 | 10 | 0710.0801 |
0710 | 0710.5938_arXiv.txt | We present IRAC and MIPS images and photometry of a sample of previously known planetary nebulae (PNe) from the SAGE survey of the Large Magellanic Cloud (LMC) performed with the Spitzer Space Telescope. Of the 233 known PNe in the survey field, 185 objects were detected in at least two of the IRAC bands, and 161 detected in the MIPS 24 $\mu$m images. Color-color and color-magnitude diagrams are presented using several combinations of IRAC, MIPS, and 2MASS magnitudes. The location of an individual PN in the color-color diagrams is seen to depend on the relative contributions of the spectral components which include molecular hydrogen, polycyclic aromatic hydrocarbons (PAHs), infrared forbidden line emission from the ionized gas, warm dust continuum, and emission directly from the central star. The sample of LMC PNe is compared to a number of Galactic PNe and found to not significantly differ in their position in color-color space. We also explore the potential value of IR PNe luminosity functions (LFs) in the LMC. IRAC LFs appear to follow the same functional form as the well-established [\ion{O}{3}] LFs although there are several PNe with observed IR magnitudes brighter than the cut-offs in these LFs. | The Large Magellanic Cloud (LMC) has been important for the study of many astrophysical processes and objects because it is one of the nearest galaxies to our own, and due to its location above the Galactic plane and its favorable viewing angle \citep[35$\degr$;][]{vanderm01}, the system can be relatively easily surveyed and many of its global properties determined. These properties are important in particular for the study of planetary nebulae (PNe). The known distance to the LMC removes the relatively large uncertainty in this parameter that affects many Galactic PNe \citep{hajian06}. The distance of $\sim$ 50 kpc allows individual objects to be isolated and in some cases resolved. The effects on PNe of the lower metallicity and dust/gas mass ratio in the LMC can be explored. One can also hope to detect a large fraction of the total number of PNe in the LMC, as opposed to in the Galaxy, where confusion and extinction in the plane allow us to detect only about 10\% of the PNe expected to exist \citep{kwok00,frew05}. An infrared survey of the LMC called Surveying the Agents of a Galaxy's Evolution \citep[][SAGE]{meixner06} has recently been completed using the IRAC \citep{fazio04} and MIPS \citep{rieke04} instruments on the Spitzer Space Telescope \citep{werner04}. SAGE is an unbiased, magnitude-limited survey of a $\sim 7\degr \times 7\degr$ region centered on the LMC. This Spitzer ``Legacy'' survey has provided a tremendous resource for the study of the stellar populations and interstellar medium (ISM) in the LMC. Some early results on the evolved stellar populations were given by \citet{blum06}, who identified $\sim$32,000 evolved stars brighter than the red giant tip, including oxygen-rich, carbon-rich, and ``extreme'' asymptotic giant branch (AGB) stars. In this paper we explore the properties of a sample of known PNe as revealed by the SAGE data. The catalog of 277 LMC PNe assembled by \citet{leisy97} from surveys that cover an area of over 100 square degrees was used for the source of positions of the PNe. Leisy et al. used CCD images and scanned optical plates to obtain accurate positions of the objects to better than 0\farcs5. They point out that the objects are in general PN candidates, with only 139 confirmed at that time with slit spectroscopy. For simplicity we will refer to the objects in the catalog as PNe, even though this caveat still applies for many of the sources. When we began to work with the SAGE data, the Leisy et al. catalog was the largest summary list of the known PNe at the time. During the course of this work, \citet{reid06} published a list of PNe in the central 25 deg$^2$ of the LMC, including 169 of the previously known objects and 460 new possible, likely, or true PNe. We will present our results here for the \citet{leisy97} catalog, and a future paper will include the new objects in the \citet{reid06} survey. | We have presented images and photometry of the \citet{leisy97} sample of PNe in the LMC. Of the 233 known PNe in the survey field, 185 objects were detected in at least two of the IRAC bands, and 161 detected in the MIPS 24 $\mu$m images. Color-color and color-magnitude diagrams were presented using several combinations of IRAC, MIPS, and 2MASS magnitudes. The location of an individual PN in the color-color diagrams was seen to depend on the relative contributions of the spectral components, resulting in a wide range of colors for the objects in the sample. A comparison to a sample of Galactic PNe shows that they do not substantially differ in their position in color-color space. The location of PNe in the various infrared color-color and color-magnitude diagrams are in general well separated from normal stars, but overlap significantly with extragalactic sources and potential YSOs. Any ambiguity between PNe and YSOs or galaxies can be readily resolved by the unique optical characteristics of PNe and their environs. Therefore, an IR color-based search for new PNe in the LMC would be viable in combination with deep optical imaging and spectroscopy. The latter remains the prerequisite to confirm a candidate as a PN. We have offered an exploration of the potential value of IR PNLFs in the LMC. IRAC LFs appear to follow the same functional form as the well-established [\ion{O}{3}] LFs although there are several PNe with observed IR magnitudes brighter than the cut-offs in these LFs. If these objects can be demonstrated to be true PNe and not very-low excitation variants nor symbiotic stars then their existence may confirm the long-standing suggestion that PNe with massive central stars suffer heavy internal extinction. This extinction would eliminate optical outliers beyond the cut-off magnitude but would affect IR LF counts minimally so that all such outliers could be observed. | 7 | 10 | 0710.5938 |
0710 | 0710.0206_arXiv.txt | The commonality of collisionally replenished debris around main sequence stars suggests that minor bodies are frequent around Sun--like stars. Whether or not debris disks in general are accompanied by planets is yet unknown, but debris disks with large inner cavities -- perhaps dynamically cleared -- are considered to be prime candidates for hosting large--separation massive giant planets. We present here a high--contrast VLT/NACO angular differential imaging survey for eight such cold debris disks. We investigated the presence of massive giant planets in the range of orbital radii where the inner edge of the dust debris is expected. Our observations are sensitive to planets and brown dwarfs with masses $>$3 to 7 Jupiter mass, depending on the age and distance of the target star. Our observations did not identify any planet candidates. {\changed We compare the derived planet mass upper limits to the minimum planet mass required to dynamically clear the inner disks. While we cannot exclude that single giant planets are responsible for clearing out the inner debris disks, our observations constrain the parameter space available for such planets.} The non--detection of massive planets in these evacuated debris disks further reinforces the notion that the giant planet population is confined to the inner disk ($<$15~AU). | Collisionally replenished debris dust surrounds about 10--20\% of the main sequence Sun--like stars (e.g. \citealt{Meyer2007}). Such widespread evidence for minor body collisions demonstrates that planetesimals orbit most stars. It is natural to ask whether or not rocky and giant planets are also present in these systems. No convincing correlation could yet be found between close--in exoplanets and the presence of debris (e.g. \citealt{Amaya2007}, but see \citealt{2005ApJ...622.1160B}). However, the presence of massive giant planets has been often invoked to account for the observed azimuthal or radial asymmetries at large radii in many debris disks (e.g. \citealt{Greaves2005,Wilner2002}). While theory offers several alternative mechanisms (e.g. \citealt{Takeuchi2001,Wyatt2005}), dynamical clearing of dust parent bodies by giant planets remains a feasible and exciting theoretical possibility (e.g. \citealt{Amaya2005,Quillen2006,Levison2007, Morbidelli2007}). Examples for such possibly dynamically--cleared disks include two recently identified disks around the young Sun--like stars \object{HD 105} \citep{Meyer2004} and \object{HD 107146} \citep{Williams2004}. Both disks were found to exhibit strong excess emission at wavelengths longer than 30~$\mu$m, while displaying no measurable excesses shortward of 20~$\mu$m. The detailed analysis of the spectral energy distribution of \object{HD 105} suggests that it is consistent with a narrow dust ring ($<$4~AU) with an inner radius of $\sim$42~AU, if the dust grains emit like black bodies \citep{Meyer2004}. Using a similar model \citet{Williams2004} showed that the excess emission from \object{HD 107146} is consistent with arising from cold dust (T=51~K) emitting as a single--temperature black body. The lack of measurable infrared excess shortward of 25~\micron ~illustrates that the inner disk regions are well cleared of dust: for HD~107146 there is at most 140$\times$ less warm dust (T=100~K) than cold dust (T=51~K, \citealt{Williams2004}). The findings of the spectral energy distribution model for \object{HD 107146} have been confirmed by direct imaging with the Hubble Space Telescope, that strengthen the case for a large featureless dust ring outside of an evacuated inner cavity \citep{Ardila2004}. Recent high--contrast imaging surveys have hinted on the general scarcity of giant planets at such large separations (e.g. \citealt{Masciadri2005,Kasper2007,Biller2007,Lafreniere2007}). Quantitative statistical analysis of the non--detections demonstrates that -- at a 90\% confidence level -- the giant planet population cannot extend beyond 30~AU if it follows a $r^{0.2}$ radial distribution, consistent with the radial velocity surveys. The statistical analysis suggests an outer cut--off for the giant planet population at $<$15~AU \citep{Kasper2007}. If so, dynamically cleared cold debris disks may be the ssignposts for rare large--separation giant planets, ideally suited for direct imaging studies. In this paper we report on a VLT/NACO high--contrast imaging survey for large--separation giant planets around HD~105, HD~107146, and six other similar disks. In the following we will review the target stars and disks, the observations, followed by a comparison of our non--detections to lower planet mass limits set by dynamical clearing simulations. \subsection{Targets} Our targets were selected from the sample of 328 Sun--like stars (0.7--2.2 $M_\odot$) targeted in the {\em Formation and Evolution of Planetary Systems} Spitzer Space Telescope Legacy program (FEPS, \citealt{Meyer2006}). From this sample we identified \ntargets southern stars, which: a) display strong infrared excess emission at long wavelengths ($\lambda > 20 \mu$m); b) no measurable excess emission at shorter wavelengths; and, C) are young and close enough to permit the detection of planetary--mass objects within the inner radius of the cold debris. Table~\ref{T:Targets} gives an overview of the key parameters of the target stars. The typical lower mass limit for the debris in the systems is $10^{-4}$ to $10^{-5}$~M$_{\earth}$, making these disks massive analogs of our Kuiper--belt (\citealt{Meyer2004,Kim2005}, Hillenbrand et al., in prep). {\changed The disks of \object{HD 105} and \object{HD 107146} --- included in our sample --- have inner evacuated regions with an estimated radii of $\sim40$~AU and $\sim31$~AU \citep{Meyer2004,Williams2004}. The other six disks exhibit spectral energy distributions similar to \object{HD 105} and \object{HD 107146}. Based on the similarity of the excess emissions and the almost identical spectral types all eight disks are expected to have cleared--out inner disks of similar size.} The only possible exception in this sample is \object{HD 202917}, for which the re--calibration of the IRAC fluxes after our VLT observations revealed a faint, but likely real infrared excess even at wavelengths shortward of 10\micron, suggesting that this inner disk may harbor small, but non--negligible amounts of warm dust. In the following we discuss briefly the results of the age determination for these sources as this has direct impact on the sensitivity of our observations to giant planets. A more detailed discussion of the ages of the whole FEPS sample will be presented in Hillenbrand et al. (in prep). We briefly summarize the upper and lower age estimates ($t_{min}$ and $t_{max}$) for each star along with the most likely age $t_{prob}$, where available. HD~105 has already reached the main sequence ($t_{min}$=27~Myr) and its chromospheric activity suggests a $t_{max}$ of 225~Myr (Hillenbrand et al., in prep.). Very likely a member of the Tuc--Hor moving group \citep{Mamajek2004} its $t_{prob}$ is 30~Myr \citep{Hollenbach2005}. HD~377 is also a main sequence star ($t_{min}>25$~Myr) and the chromospheric activity suggests that $t_{max}$=220~Myr. The median of four other age indicators sets $t_{prob}$=90~Myr. For HD~107146 we adopt the age range of 80--200~Myr. HD~202917 is a likely member of the Tuc--Hor moving group ($t_{min}=t_{prob}=30$~Myr) and its upper age limit is set by its Li--abundance, higher than that of the Pleiades ($t_{max}<100$~Myr). HD~35850 is suggested to be a $\beta$~Pic Moving Group member ($t_{min}$=12~Myr, \citealt{Song2003}) and its observed rotation rate sets a reliable upper age limit of $t_{max}$=100~Myr (Hillenbrand et al. in prep.; cf. \citealt{Barnes2007}). HD~70573 is among the few stars that are known to harbor both a debris disk and a giant planet. \citet{Setiawan2007} found an $m_2 sin i = 6.1$~\mjup possible planet on a 1.76--AU orbit. A combination of different age indicators suggest a $t_{min}=30$~Myr for HD~70573 and a $t_{prob}=60~$Myr; \citet{Setiawan2007} quotes $t_{max}$=125~Myr. {\changed Based on Li--abundance and chromospheric activity, position on the color--magnitude diagram and the analysis of its space motions Mamajek et al. (in prep.) estimates that HD~209253 has $t_{min}=200$~Myr and $t_{max}<1.6$~Gyr with $t_{prob}$=500~Myr.} HD 25457 is a member of the AB~Dor moving group giving a very strong lower age limit ($t_{min}=50$~Myr, \citealt{Zuckermanetal2004}). \citet{Luhman2005} derives an age of 75--125~Myr (we adopt $t_{prob}=75$~Myr), while the upper age limit is set by the chromospheric activity ($t_{max}=170$~Myr). \begin{deluxetable}{lccccccc} \tabletypesize{\scriptsize} \tablecaption{Target parameters. \label{T:Targets}} \tablewidth{0pt} \tablehead{\colhead{Target} & \colhead{R.\,A. (J2000)} & \colhead{Dec. (J2000)} & \colhead{V--mag.$^a$}& \colhead{Dist. [pc]$^a$} &\colhead{Sp. Type} & \colhead{Ages:$^b$ $t_{min}$/$t_{prob}$/$t_{max}$} } \startdata HD 105 & 00 05 52.6 & $-$41 45 11 & 7.51 & 40 & G0V & 27 Myr / 30 Myr / 225 Myr\\ HD 377 & 00 08 25.7 & $+$06 37 01 & 7.59 & 40 & G2V & 25~Myr / 90~Myr / 220~Myr \\ HD 25457 & 04 02 36.8 & $-$00 16 08 & 5.38 & 19 & F5V & 50 Myr / 75 Myr / 170 Myr\\ HD 35850 & 05 27 04.8 & $-$11 54 03 & 6.30 & 27 & F7/8V & 12 Myr / 12 Myr / 100 Myr\\ HD 70573 & 08 22 50.0 & $+$01 51 34 & 8.69 & 70 & G1/2V & 30 Myr / 60 Myr / 125 Myr \\ HD 107146 & 12 19 06.5 & $+$16 32 54 & 7.04 & 29 & G2V & 80 Myr / -- / 200 Myr \\ HD 202917 & 21 20 50.0 & $-$53 02 03 & 8.65 & 46 & G5V & 30 Myr / 30 Myr / 100 Myr \\ HD 209253 & 22 02 33.0 & $-$32 08 02 & 6.63 & 30 & F6/7V & 200~Myr / 500 Myr / 1.6 Gyr \\ \enddata \tablenotetext{a}{All magnitudes and distances from the Hipparcos catalog, except for the distance of HD~70573, which is a main sequence--distance. } \tablenotetext{b}{The age estimates are discussed in the text.} \end{deluxetable} | We present results from a high--contrast angular differential imaging survey of \ntargets cold debris disks, selected to have significantly or totally evacuated inner disks. Our observations searched for massive giant planets that may be responsible for carving out the inner holes in the observed cold debris disks. For most of our targets we reach typical sensitivities of 3 to 7 \mjup between 20 to 50~AU separations, but did not identify any likely planet or brown dwarf candidates. {\changed By comparing the derived planet mass upper limits to lower limits derived from dynamical scattering models (typically 2-5 \mjup between 10 and 30 AU), we limit the parameter space available for any single planet capable of efficiently clearing out the inner planetesimal disks. } Our survey complements recent direct imaging surveys of nearby young stars indicating that massive giant planets at large separations are very rare. Cool debris disks with large inner evacuated cavities remained promising possible exceptions to this rule until now. However, the combination of our observational upper limits and theoretical lower limits strongly suggest that massive giant planets at large separations are not present in most of these systems, reinforcing the finding that the outer cut--off for the giant planet distribution is probably at 15~AU or at even smaller semi--major axes \citep{Kasper2007}. | 7 | 10 | 0710.0206 |
0710 | 0710.1941_arXiv.txt | {Measurements at 100 TeV and above are an important goal for the next generation of high energy $\gamma$-ray astronomy experiments to solve the still open problem of the origin of galactic cosmic rays. The most natural experimental solution to detect very low radiation fluxes is provided by the Extensive Air Shower (EAS) arrays. They benefit from a close to 90$\%$ duty cycle and a very large field of view ($\sim$2 sr), but the sensitivity is limited by their angular resolution and their poor cosmic ray background discrimination. Above 10 TeV the standard technique for rejecting the hadronic background consists in looking for {\it "muon-poor"} showers. In this paper we discuss the capability of a large muon detector (A$_{\mu}$=2500 m$^2$) operated with an EAS array at very high altitude ($>$4000 m a.s.l.) to detect $\gamma$-ray fluxes around 100 TeV. Simulation-based estimates of energy ranges and sensitivities are presented. } \begin{document} | The recent TeV results of the HESS experiment suggest the existence of a population of galactic $\gamma$-ray sources whose emission extends beyond 10 TeV in the 5 to 15$\%$ of Crab flux range (for E $>$ 1 TeV). These sources, associated with nearby shell-type or plerionic SNRs, the most probable factories of galactic cosmic rays, can be studied detecting gamma-rays (and neutrinos) emission in the VHE/UHE energy domain. Therefore, a detector capable to perform a continuous all-sky survey at a level of about a percent of the Crab flux up to 100 TeV is needed. The search and study of {\em "Cosmic PeVatrons"} and their surrounding regions is one of the main scientific issues to be addressed by the next generation of ground-based $\gamma$-ray astronomy detectors \cite{aharonian}. Current experiments are not able to reach 100 TeV because their limited collection area makes the required exposures too long. Extrapolating the galactic source spectra measured by HESS, it appears that future experiments need to achieve at least 100 km$^2\cdot$h in order to obtain meaningful measurements at 100 TeV. The most natural experimental solution that provides such a large exposure is given by the EAS arrays observing each source for $\sim$1500 h/year. As an example, an experiment like ARGO-YBJ already results in an exposure of $\sim$15 km$^2\cdot$h/year, with an angular resolution ($\sim$0.2$^{\circ}$ at 10 TeV) near the best value attainable by a sampling array. In addition, the EAS arrays are the only ground-based detectors allowing simultaneous and continuous coverage of a significant fraction of the sky (about all that overhead). Their large field of view and high duty cycle ($>$90 $\%$) suit to perform a $\gamma$-ray sources population survey at VHE/UHE energies. But more important is their unique potential that allows to have an effective monitoring of the $\gamma$-ray activity of a large number of highly variable sources like blazars and microquasars, as well as the possibility of independent detection and study of GRBs \cite{aharonian}. In addition, the recent observations of unidentified extended sources from the Galactic plane \cite{milagro_galpl} and in the Cygnus region \cite{milagro_cygnus} reported by the Milagro Collaboration demonstrate the strength of EAS arrays in finding diffuse and extended sources. Therefore, the discovery science could be a feature of EAS arrays. The limited sensitivity in detecting $\gamma$-ray point sources, characteristic of EAS experiments, is mainly due to their poor gamma/hadron separation power, limited angular resolution and high energy threshold. Exploiting a full coverage approach at very high altitude leads to the improvement of the angular resolution to $\sim$0.2$^{\circ}$ and to the reduction of the energy threshold well below the TeV region. The standard technique to perform a gamma/hadron discrimination above 10 TeV with EAS arrays consists in looking for {\it "muon-poor"} showers. In this paper we discuss the capability of a large muon detector (A$_{\mu}$=2500 m$^2$) operated with an EAS array at very high altitude ($>$4000 m a.s.l.) to detect $\gamma$-ray fluxes up to 100 TeV. An estimation based on the simulation of energy ranges and sensitivities is reported. | 7 | 10 | 0710.1941 |
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0710 | 0710.1310_arXiv.txt | This Letter reports on initial Expanded Very Large Array (EVLA) observations of the 6035~MHz masers in ON~1. The EVLA data are of good quality, lending confidence in the new receiver system. Nineteen maser features, including six Zeeman pairs, are detected. The overall distribution of 6035~MHz OH masers is similar to that of the 1665~MHz OH masers. The spatial resolution is sufficient to unambiguously determine that the magnetic field is strong ($\sim -10$~mG) at the location of the blueshifted masers in the north, consistent with Zeeman splitting detected in 13441~MHz OH masers in the same velocity range. Left and right circularly polarized ground-state features dominate in different regions in the north of the source, which may be due to a combination of magnetic field and velocity gradients. The combined distribution of all OH masers toward the south is suggestive of a shock structure of the sort previously seen in W3(OH). | \object[Onsala 1]{Onsala 1} (ON 1) is a massive star-forming region with an unusual OH maser spectrum. The ground-state masers, which have been observed interferometrically several times, \citep[e.g.,][]{ho83,argon00,fish05,nammahachak06,fish07}, appear in two disjoint velocity ranges: $< 6$~km\,s$^{-1}$ and 11--17~km\,s$^{-1}$, with no maser emission in between. This pattern is reproduced in 6668~MHz methanol emission, which also appears near 0 and 15~km\,s$^{-1}$ \citep{szymczak00}. The masers presently elude clear interpretation. Comparison of OH maser velocities seen in projection against the ultracompact \ion{H}{2} region \citep[13--14~km\,s$^{-1}$ after correcting for Zeeman splitting; e.g., from][]{fish05} with the LSR velocity of the latter derived from a hydrogen recombination line ($5.1 \pm 2.5$~km\,s$^{-1}$) led \citet{zheng85} to interpret the masers as tracing infall. However, a proper motion study of the OH masers suggests that expansion may dominate the kinematics \citep{fish07}. A more recent model suggests that the OH masers are associated with a molecular outflow \citep{kumar04,nammahachak06}. This model proposes a shocked molecular torus origin for the southern masers, but questions remain regarding the overall morphology and velocity structure of the masers in ON~1. Few northern star-forming regions have been mapped in the 6035~MHz line of OH, in part because few radio arrays in the northern hemisphere have been capable of tuning to the frequency. A previous three-station European VLBI Network (EVN) experiment detected seven 6035~MHz maser features in ON~1 but only observed the redshifted masers \citep{desmurs98}. Single-dish observations confirm that the 6030 and 6035~MHz masers in ON~1 also appear in two disjoint velocity ranges, but Zeeman pairing ambiguities have prevented a definitive measurement of the magnetic field of the blueshifted masers \citep{baudry97,fish06}. The upgrade of the National Radio Astronomy Observatory's (NRAO) Very Large Array (VLA) to the Expanded VLA (EVLA) presents new observational opportunities in North America \citep{mckinnon01,ulvestad07}. Of interest to spectral line observers is the full frequency coverage between 1 and 50~GHz that will become available. This spring, the EVLA for the first time offered observational capabilities in the extended C-band range of 4.2 to 7.7~GHz, which includes key maser frequencies of OH and methanol. This Letter reports on initial observations of 6035~MHz OH masers with the EVLA. | \subsection{Environment of the Masers in ON~1} The distributions of the 1665 and 6035~MHz masers are the most alike of any pair of OH transitions in ON~1. All the 1665~MHz maser regions have associated 6035~MHz masers, apart from in the center and in the extreme south, far from the exciting source where densities are likely to be low. The brightest masers at both transitions occur in the line of maser spots just south of center in Figure~\ref{map}. This contrasts with the 1612, 1667, and 1720~MHz masers, which are only found in the prominent line of masers south of and on the southwestern limb of the continuum source. A similar result was seen at VLBI resolution in W3(OH): 1665 and 6035~MHz masers appear throughout the source, 1612 and 1720~MHz masers appear only in and near the inner edge of an apparent shocked torus that is especially well traced by 6.0~GHz masers (including several of the brightest 6.0~GHz masers detected), and 1667~MHz emission is largely concentrated in areas with an apparent shock morphology \citep{fishevn07}. Higher-resolution observations of the 6.0~GHz masers in ON~1 would be useful to obtain a better alignment of the 6035~MHz frame with respect to the other masers and determine whether the southern masers are located preferentially along the northern edge of the southern shock front. Some of the 6035~MHz masers may be associated with an outflow, such as the H$^{13}$CO$^{+}$ oriented northeast-southwest through ON~1 detected by \citet{kumar04}. \citet{nammahachak06} note that the southern line of OH masers is oriented nearly perpendicular to this structure and propose, based additionally on the linear polarization characteristics of the ground-state masers, that the masers trace a shock in a confining molecular torus. Similar conclusions based on maser distribution and polarization are reached for other OH maser sources hosting outflows \citep{hutawarakorn99,hutawarakorn03,hutawarakorn05,hutawarakorn02}. This interpretation also bears a striking similarity to that proposed for W3(OH), which hosts very strong 6035~MHz masers \citep{fishevn07}, and may therefore suggest a common mechanism for producing strong 6035~MHz maser emission. It is probable that there is a coherent, organized structure in the northern part of ON~1. Many LCP 1665~MHz masers, including one bright one, occur in the region, with a few weak RCP masers offset to the north and west of the LCP emission \citep{fish05,fish07}. Assuming a magnetic field of approximately $-10$~mG in the region, LCP and RCP 1665~MHz maser components would be separated by 6~km\,s$^{-1}$ ($20\, \Delta v$ for an average line width of 0.3~km\,s$^{-1}$) if they both appear in the region. (It is likely that weak RCP 1665~MHz masers, seen near $-2$~km\,s$^{-1}$ in the spectra of \citet{clegg91}, pair with the dominant LCP masers near +4~km\,s$^{-1}$, but no published interferometric observations of the ground-state masers have included the $-2$~km\,s$^{-1}$ features in the observed bandpass.) If the magnetic field and velocity gradients are aligned such that the magnetic field strength decreases (becomes less negative) in the same direction that the velocity increases, amplification of the LCP component would be favored \citep[see][]{cook66}. This effect is much more pronounced for ground-state OH masers than at 6035~MHz, where the Zeeman splitting is a factor of 10 smaller in velocity units, such that the LCP and RCP spectra are separated by $0.6$~km\,s$^{-1} \approx 2\, \Delta v$ per 10~mG, which is not in excess of the turbulent velocity component expected in a maser condensation \citep{reid80}. Hence, 6035~MHz masers appear in both polarizations with similar fluxes, while strong LCP 1665~MHz emission appears at higher velocities (and possibly weaker RCP emission at lower velocities). The key question that remains is how the blueshifted masers in the north connect to the redshifted masers in the south. It is tempting to associate both with the HCO$^+$ outflow, but the northern masers are significantly blueshifted compared to the HCO$^+$ velocity range of 8--16~km\,s$^{-1}$ \citep{kumar04}. These authors also note a CO outflow in the region that spans a larger velocity range, but it is oriented roughly east-west. \citet{kumar04} interpret the outflows as coming from two different sources embedded in the ultracompact \ion{H}{2} region. It is possible that the blue and red OH masers are also associated with two different sources, which would complicate interpretation of their motions and morphology. The large ($|B| > 10$~mG) magnetic fields suggest that the northern masers are near a region of higher density and therefore likely an excitation source. In W3(OH) the methanol masers in the region of highest magnetic field strength ($|B| > 10$~mG) have been modelled as undergoing conical expansion, which \citet{moscadelli02} interpret as possibly being due to an outflow guided by the helical field from a magnetized disk. \subsection{Future Directions} Excluding minor transitional issues, the performance of the EVLA antennas was as expected. Early data from the antennas with upgraded C-band capability are of good quality and are already producing useful science \citep[e.g.,][]{sjouwerman07}. As the upgrade continues, more antennas with 6.0~GHz tuning capability will be added to the array, providing propotionally better sensitivity and imaging characteristics. High spectral resolution VLBI observations of the 1665 and 6035~MHz OH masers, of the sort obtained for W3(OH) \citep{fbs06,fishevn07}, may help answer the questions of whether there is an organized velocity structure in the north and what structure the masers are tracing. If the segregated regions of LCP and RCP emission in the north are due to correlated magnetic and velocity field gradients, small-scale velocity gradients (i.e., on the size scale of a maser spot) of the northern masers may show similar magnitudes and position angles. This would contrast with W3(OH), in which small-scale velocity gradients show no correlation except when masers overlap \citep{fbs06,fishevn07}. Any such observations of the ground-state masers should have a sufficiently wide bandwidth coverage to include the $-2$~km\,s$^{-1}$ 1665~MHz masers in order to be able to identify magnetic field strengths (and variations thereof) throughout the northern masers. While observations of linear polarization exist for the ground-state masers \citep{nammahachak06}, full-polarization observations at 6035~MHz would help to understand the magnetic field geometry, since linear polarization vectors are much less subject to being corrupted by external and internal Faraday rotation at the higher frequency \citep[e.g.,][]{fishreid06}. | 7 | 10 | 0710.1310 |
0710 | 0710.1917_arXiv.txt | {The standard cooling flow model has predicted a large amount of a cool gas in the clusters of galaxies. The failure of the Chandra and XXM-Newton telescopes to detect a cooling gas (below 1-2 keV) in the clusters of galaxies has suggested that some heating process must work to suppress the cooling. The most likely heating source is the heating by AGNs. There are many heating mechanisms, but we will adopt the effervescent heating model which is a result of the interaction of the bubbles inflated by AGN with the intra-cluster medium(ICM). Using the FLASH code, we have carried out 1D- time dependent simulations to investigate the effect of the heating on the suppression of the cooling in cooling flow clusters. We have found that the effervescent heating model can not balance the radiative cooling and it is an artificial model. Furthermore, the effervescent heating is a function of the ICM pressure gradient but the cooling is proportional to the gas density square and square root of the gas temperature.} {}{}{}{} | % According to the steady flow assumption of the standard cooling flow model, we must find a cool gas and a multi-phase medium within the cooling core in clusters of galaxies which are not observed in any wavebands (X-ray and non X-ray). This is known as the cooling flow problem in clusters of galaxies. In others words, there is a discrepancy between standard cooling flow model and observations (X-ray and non X-ray). This strong discrepancy is interpreted as either the gas is being prevented from cooling by some heating process, or it cools without any spectroscopic signature~\citep{fabian01} which is difficult. But, we see that this discrepancy between the standard cooling flow model and X-ray observations indicates that either the cooling in the center of cooling flow clusters must be suppressed by any heating mechanism, or the steady flow assumption of the standard cooling flow model is not appropriate. Somewhere else, we have concentrated in the second point which it is found that the steady flow with cooling is impossible, i.e the cooling flow problem is due to the wrong steady flow assumption. In this work we will concentrate in the first point, heating the cooling flow.\\ The failure of the multi-phase model has revived the idea of a heating mechanism which can suppress the cooling. There are five main conditions for the heating \citep[see][ for review]{gardini04}~: \begin{enumerate} \item The heating must be fine tuned and distributed to get the smooth observed temperature profile. Too much heating would result a outflow from the center region. Too little heating is not sufficient to suppress the cooling. Moreover, the heating process must be self regulated: the mass flow rate triggers the heating and the heating reduces the mass flow rate~\citep{hans02}. That mechanism is called a heating with a feedback. \item The heating mechanism by AGN must be sporadic because the radio activities are not observed in every cooling flow clusters. \item The kinetic energy injection must be subsonic and the ICM must not be shocked, as observed in general case. The shock could compress the gas producing the catastrophic cooling much faster. \item The heating mechanism must preserve the observed entropy profiles of the ICM, which decrease toward the center. \item The heating must not destroy the metallicity profiles which are peaked toward the centers \citep{tamura02,grandi01,irwin01}. \end{enumerate} A giant elliptical or cD galaxy sits at the potential well of every cooling flow cluster~\citep{mathews03,eilek04}. The popular heating mechanisms are a heating by AGN~\citep{binney95,binney93,br02}, thermal conduction, cosmic rays, galaxies motions, magnetic field reconnection \citep{soker90} and turbulent mixing \citep{kim03}. About 71 percent of the cDs in the cooling flow clusters are radio loud compared to only 23 percent of non-cooling flow clusters ~\citep{burns90}. This result suggests that there a relationship between the AGN activities and the presence of cooling flow. In some of cooling flow clusters, the recent X-ray observations reveal holes in the X-rays surface brightness coincident with the radio lobes, known as X-ray cavities or bubbles~\citep[ see][ as example]{fabian06}. The bubbles or cavities are not a universal phenomenon, suggesting a duty cycle. {}~Radio sources are not even detected in some cooling-flow clusters. The observed metallicity gradients make a constraint on the ability of baubles to mix the ICM~\citep{bohringer04}. ~\citet{bm102} have run many 1D simulations of clusters with various levels of heating. They concluded that the best fit of temperature profiles with the real clusters is without heating. The AGNs are assumed to inject buoyant bubbles into the ICM, which heat the ambient medium by doing work $PdV$ as they rise and expand. This mechanism is called the effervescent heating model or mechanism~\citep{begelman01}. In this work, we will concentrate on this mechanism, using time dependent hydrodynamics simulations, FLASH code~\citep{fry02}. \Rem{ | \label{sec:sm_heat_model} % We have carried out time dependent simulations with cooling and heating. The general result is that the best fit to the observations is the model without heating (model A). The effervescent heating is a function of the ICM pressure gradient but the cooling is proportional to the gas density square and square root of the gas temperature.~Furthermore, the conclusions are~: \\ \\ 1- From our simulations, the cooling can not be a steady flow as assumed by standard cooling flow model. The gas density must increase with the time; i.e the ICM must be compressed under the force of the inflowing gas.\\ \\ 2- The effervescent heating model profile (without the smoothness function) is more steeper than the radiative cooling profile, which makes it is very difficult to balance the cooling. \\ \\ 3- The function $1-e^{-r_o/r}$ is not a cutoff function but it is a smoothing function in order to control the steepness of the heating function letting it ~balance the radiative cooling; i.e. it is an artificial model. \\ \\ 4- At inner radius close to the realistic value of the accretion radius (model D), the accretion mass rate due to flowing gas is not enough to fuel the black hole in the center of cluster. It's not reasonable to set the accretion radius of the black hole at large radius, (for example 1-1.5 kpc, as used for the three dimension simulations). \\ \\ 5- In some cooling flow clusters, the observed cooling time scale is very short, in the range $10^{8}$ yr to $10^{9}$ yr, but the cooling gas below 1-2 keV is not observed. This result is interpreted as indication that there must be a heating mechanism to stop the cooling. As in figure~\ref{fig:a1991_purecooling_tmscl}, the cooling time scale inside a radius of 10 kpc is shorter than $1$ Gy but there is no a cool gas. Moreover, from our simulations (model B), we found that the heating process increases the cooling time scale. If cooling time scale is a good approximation for the actual cooling time and the cluster is heated, then we must find that the cooling time scale is larger than observed.\\ \\ | 7 | 10 | 0710.1917 |
0710 | 0710.3586_arXiv.txt | We investigate large-amplitude baryon acoustic oscillations (BAO's) in off-diagonal entries of cosmological power-spectrum covariance matrices. These covariance-matrix BAO's describe the increased attenuation of power-spectrum BAO's caused by upward fluctuations in large-scale power. We derive an analytic approximation to covariance-matrix entries in the BAO regime, and check the analytical predictions using $N$-body simulations. These BAO's look much stronger than the BAO's in the power spectrum, but seem detectable only at about a one-sigma level in gigaparsec-scale galaxy surveys. In estimating cosmological parameters using matter or galaxy power spectra, including the covariance-matrix BAO's can have a several-percent effect on error-bar widths for some parameters directly related to the BAO's, such as the baryon fraction. Also, we find that including the numerous galaxies in small haloes in a survey can reduce error bars in these cosmological parameters more than the simple reduction in shot noise might suggest. | Humans are fortunate that there are baryons in the Universe. Not only are we made of them, but baryons are responsible for features in the shape of cosmological power spectra and two-point correlation functions that are quite valuable to cosmologists. These features are called baryon acoustic oscillations (BAO's), and are imprints of acoustic oscillations in the gas of the early universe \citep{peeyu, sz, holtz, ehu, mwp}. Their presence in galaxy clustering statistics provides a standard ruler to measure the relation between distance and redshift, and the expansion of the universe at late times \citep{bg, se}. They provide one of the main tools currently proposed for studying the effects of dark energy. BAO's have been detected in modern low-redshift galaxy surveys, both in the 2dFGRS \citep{c05} and SDSS \citep{e05,hutsi,pshape07}. These detections were made using the (two-point) correlation function, or its Fourier dual, the power spectrum. Conveniently, the BAO regime is on large-enough scales that non-linear effects are mild. Still, for precision cosmology, these mild effects must be understood, and are the topic of much recent research \citep[e.g.][]{mill,jk,huff,sw,aea,se07,sssbao}. Non-linear evolution of the matter power spectrum tends to dampen or smear BAO's. For example, \citet{esw} found that large-scale bulk flows and cluster formation produce motions that smear out the BAO peak in the correlation function, but that these motions are confined to relatively small scales of $\sim 10 \hMpc$ in Lagrangian space. They argued that these motions roughly preserve wiggles on the largest scales of the power spectrum, but wipe out wiggles on smaller scales. The effects we describe in this paper, using Eulerian perturbation theory, likely arise physically from the same large-scale bulk motions. The attenuation of BAO's can also be understood by considering an additive mode-coupling power spectrum, which rises on small scales as structure develops, along with a function which attenuates the linear power spectrum on non-linear scales. In the halo model \citep[HM, reviewed in][]{cs}, this small-scale contribution is the one-halo (1h) term, and comes from pairs of objects within single haloes. A qualitatively similar effect occurs in renormalized perturbation theory \citep[RPT;][]{csrpt,mrpt}, which is less empirical than the HM, and seems more accurate through translinear scales. For example, \citet{csrpt} show that, more physically than in the HM, the mode-coupling power spectrum in RPT goes to zero on small scales. In this paper, we show that wiggles exist in off-diagonal terms of matter and galaxy power spectrum covariance matrices, almost entirely out-of-phase with BAO's in the power spectrum. We interpret these BAO's in the covariance matrix as manifestations of the suppression exacted on power-spectrum BAO's by power on large scales. Regions of the Universe with upward fluctuations in large-scale power have more-suppressed BAO's. The body of the paper is organized as follows. In Section \ref{anal}, we discuss analytic predictions for the BAO's in the covariance matrix of matter. In Section \ref{nbody}, we test the analytic predictions against $N$-body simulations, and investigate the detectability of the wiggles in the covariance matrix. Finally, in Section \ref{cosmomatter}, we investigate the effect of these covariance-matrix BAO's on cosmological parameter estimation, focusing on parameters directly related to BAO's in the power spectrum. We perform this analysis for both matter and galaxy power spectra, using the HM framework. | Our main points are the following: \begin{itemize} \item In off-diagonal entries in power spectrum covariance matrices, BAO's exist which appear much stronger than the BAO's in the power spectrum. These wiggles are a manifestation of the suppression which large-scale power does to BAO's in the power spectrum, and originate in the perturbation-theory trispectrum. We give a simple analytic approximation to these wiggles in terms of the linear power spectrum and its first two derivatives, and check the analytical predictions using $N$-body simulations. \item These wiggles are potentially detectable in current and upcoming surveys, but because of the large noise in a covariance matrix measurement, they are only detectable at a small significance level. We estimate that a one-sigma detection could be done with a survey of size a couple of $h^{-3}{\,}{\rm Gpc}^3$, and that a three-sigma detection could require a survey of volume $\sim30\hGpcV$. \item The wiggles make a modest difference in cosmological parameter error bars from analysing galaxy and matter power spectra. For example, using the true, wiggly covariance matrix in estimating the baryon fraction results in error bars several percent tighter than a no-wiggle covariance matrix. Doing so in estimating the baryon acoustic scale has a smaller effect. \item At fixed volume, including galaxies in smaller-mass haloes provides tighter error bars on parameters such as the baryon fraction and the baryon acoustic scale than analysing galaxies in only large haloes. In the context of the HM, this effect goes beyond the simple gains from analysing a sample with smaller shot noise. We attribute this to the one-halo term of the galaxy power spectrum becoming dominant at smaller scales for small haloes than large ones. \end{itemize} \begin{sloppypar} Most of the calculations in this paper made use of our package of Python code for cosmology, called {\scshape CosmoPy}. It can be downloaded from \url{http://www.ifa.hawaii.edu/cosmopy/}. \end{sloppypar} | 7 | 10 | 0710.3586 |
0710 | 0710.3854_arXiv.txt | We investigate the evolution of the properties of model populations of ultraluminous X-ray sources (ULXs), consisting of a black-hole accretor in a binary with a donor star. We have computed models corresponding to three different populations of black-hole binaries, motivated by our previous studies. Two of the models invoke stellar-mass ($\sim10\,M_\odot$) black-hole binaries, generated with a binary population synthesis code, and the third model utilizes intermediate-mass ($\sim1000\,M_\odot$) black-hole accretors (IMBHs). For each of the three populations, we computed 30,000 binary evolution sequences using a full Henyey stellar evolution code. A scheme for calculating the optical flux from ULXs by including the reprocessed X-ray irradiation by, and the intrinsic viscous energy generation in, the accretion disk, as well as the optical flux from the donor star, is discussed. We present color-magnitude diagrams (CMDs) as ``probability images'' for the binaries as well as for the donor stars alone. ``Probability images'' in the plane of orbital period and X-ray luminosity are also computed. We show how a population of luminous X-ray sources in a cluster of stars evolves with time. The most probable ULX system parameters correspond to high-mass donors (of initial mass $\gtrsim 25~M_\odot$) with effective O through late B spectral types and equivalent luminosity classes of IV or V. We also find the most probable orbital periods of these systems to lie between 1-10 days. Estimates of the numbers of ULXs in a typical galaxy as a function of X-ray luminosity are also presented. From these studies we conclude that if the stellar-mass black-hole binaries are allowed to have super-Eddington limited X-ray luminosities: (i) the value of the binding energy parameter for the stellar envelope of the progenitor to the black-hole accretor must be in the range of $0.01 \lesssim \lambda \lesssim 0.03$ in order not to overproduce the ULXs, and (ii) the stellar-mass black-hole models still have a moderately difficult time explaining the observed ULX positions in the CMD. Other possible explanations for the apparent overproduction of very luminous X-ray sources in the case of stellar-mass black-hole accretors are discussed. Our model CMDs are compared with six ULX counterparts that have been discussed in the literature. The observed systems seem more closely related to model systems with very high-mass donors in binaries with IMBH accretors. We find that a significant contribution to the optical flux from the IMBH systems comes from {\em intrinsic} accretion disk radiation whose source is viscous dissipation of gravitational potential energy. In effect, the IMBH systems, when operating at their maximum luminosities ($10^{41}-10^{42}$ ergs s$^{-1}$), are {\em milli-AGN}. With regard to the IMBH scenario, while attractive from the aspect of binary evolution models, it leaves open the larger question of how the IMBHs form, and how they capture massive stellar companions into just the correct orbits. | 7 | 10 | 0710.3854 |
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0710 | 0710.0585_arXiv.txt | We describe efforts over the last six years to implement regularization methods suitable for studying one or more interacting black holes by direct N-body simulations. Three different methods have been adapted to large-N systems: (i) Time-Transformed Leapfrog, (ii) Wheel-Spoke, and (iii) Algorithmic Regularization. These methods have been tried out with some success on GRAPE-type computers. Special emphasis has also been devoted to including post-Newtonian terms, with application to moderately massive black holes in stellar clusters. Some examples of simulations leading to coalescence by gravitational radiation will be presented to illustrate the practical usefulness of such methods. | In the study of strong gravitational interactions, utilization of the chain data structure can be very beneficial. Over many years, the original chain regularization method (Mikkola \& Aarseth 1993) has proved to be effective in star cluster simulations containing binaries. As we shall see in the following, it is also a useful tool in connection with time transformations which do not employ the usual coordinates. By introducing one or more dominant masses these advantages become more apparent. Such problems fall naturally into three classes according to the number of massive objects and each class requires special attention. At the simplest level we have the case of one central massive body which dominates the motion of other members within a certain distance. The role of the reference body is readily seen in the case of three interacting particles which can be studied by three-body regularization (Aarseth \& Zare 1974). This idea was extended to an arbitrary membership (Zare 1974). However, a natural application was lacking until the problem of black holes (BHs) became a challenge for simulators in recent years. The aptly named wheel-spoke regularization (Aarseth 2003a) has now been adapted to study compact subsystems containing a single massive object. Historically speaking, a special method for a BH binary was implemented in an $N$-body code first. Here the main idea is based on a time-transformed leapfrog scheme (TTL) suitable for dealing with large mass ratios (Mikkola \& Aarseth 2002). Remarkably, this method yields machine precision for unperturbed two-body motion. Although regularized chain coordinates are not employed directly, the accuracy is improved by using relative quantities with respect to the nearest massive body. Alternative methods may be needed for problems involving more than two massive objects. The recent algorithmic regularization (Mikkola \& Tanikawa 1999, Preto \& Tremaine 1999, Mikkola \& Merritt 2006) appears to be a promising way of studying such systems. Indeed, the masses only play a kinematical role in one formulation, suggesting that it may be applicable to systems with large mass ratios. However, it should be emphasized that for practical reasons any of the above methods are of necessity limited to relatively small particle numbers, or in other words, compact subsystems. | 7 | 10 | 0710.0585 |
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0710 | 0710.2313_arXiv.txt | In differentially rotating discs with no self-gravity, density waves cannot propagate around the corotation, where the wave pattern rotation speed equals the fluid rotation rate. Waves incident upon the corotation barrier may be super-reflected (commonly referred to as corotation amplifier), but the reflection can be strongly affected by wave absorptions at the corotation resonance/singularity. The sign of the absorption is related to the Rossby wave zone very near the corotation radius. We derive the explicit expressions for the complex reflection and transmission coefficients, taking into account wave absorption at the corotation resonance. We show that for generic discs, this absorption plays a much more important role than wave transmission across the corotation barrier. Depending on the sign of the gradient of the vortensity of the disc, $\zeta=\kappa^2/(2\Omega\Sigma)$ (where $\Omega$ is the rotation rate, $\kappa$ is the epicyclic frequency, and $\Sigma$ is the surface density), the corotation resonance can either enhance or diminish the super-reflectivity, and this can be understood in terms of the location of the Rossby wave zone relative to the corotation radius. Our results provide the explicit conditions (in terms of disc thickness, rotation profile and vortensity gradient) for which super-reflection can be achieved. Global overstable disc modes may be possible for discs with super-reflection at the corotation barrier. | Differentially rotating fluid discs, ubiquitous in astrophysics, are known to exhibit rich dynamics and possible instabilities (e.g. Papaloizou \& Lin 1995; Balbus \& Hawley 1998). While local instabilities, such as Rayleigh's centrifugal instability (for discs with specific angular momentum decreasing outwards), gravitational instability (for self-gravitational discs with too large a surface density, or more precisely, Toomre $Q\lo 1$), and magnetorotational instability (for discs with a sub-thermal magnetic field), are well understood (at least in the linear regime), global effects and instabilities are more subtle, since they involve couplings and feedbacks of fluid at different locations (see Goldreich 1988 for an introduction/review). A well-known example is the corotation amplifier (e.g. Mark 1976; Narayan, Goldreich \& Goodman 1987), which arises from the interaction across the corotation between waves carrying opposite signs of angular momentum. Much stronger corotation amplifications (WASER -- wave amplification by the stimulated emission of radiation, and SWING amplifiers) can be achieved for self-gravitating discs (e.g., Goldreich \& Lynden-Bell 1965; Julian \& Toomre 1966; Lin \& Lau 1975; see Shu 1992 for a review). Another well-known example is the Papaloizou-Pringle instability in finite accretion tori (confined between two free surfaces), in which coupling between waves inside the corotation with those outside, combined with reflecting inner and outer boundaries, leads to violent overstable modes (Papaloizou \& Pringle 1984; Goldreich et al.~1986). Recent works on global disc instabilities include the Rossby wave instability (for discs with a strong enough density or vortensity bump; Lovelace et al.~1999; Li et al.~2000) and the accretion-ejection instability (for magnetized discs; Tagger \& Pellat 1999, Tagger \& Varniere 2006). In this paper we are interested in 2D fluid discs without self-gravity and magnetic field. For disturbances of the form $e^{im\phi-i\omega t}$, where $m>0$ and $\omega$ is the wave (angular) frequency (and thus $\omega_p=\omega/m$ is the pattern frequency), the well-known WKB dispersion relation for density waves takes the form (e.g., Shu 1992) \be (\omega-m\Omega)^2=\tomega^2=\kappa^2+k_r^2c^2, \label{eq:disp} \ee where $\Omega$ is the disc rotation frequency, $\tomega=\omega-m\Omega$ is the Doppler-shifted wave frequency, $\kappa$ is the radial epicyclic frequency, $k_r$ is the radial wavenumber, and $c$ is the sound speed. Thus waves can propagate either inside the inner Lindblad resonance radius $r_{\rm IL}$ (defined by $\tomega=-\kappa$) or outside the outer Lindblad resonance radius $r_{\rm OL}$ (defined by $\tomega=\kappa$), while the region around the corotation radius $r_c$ (set by $\tomega=0$) between $r_{\rm IL}$ and $r_{\rm OL}$ is evanescent. Since the wave inside $r_{\rm IL}$ has pattern speed $\omega_p$ smaller than the fluid rotation rate $\Omega$, it carries negative wave action (or angular momentum), while the wave outside $r_{\rm OL}$ carries positive wave action. As a result, a wave incident from small radii toward the corotation barrier will be super-reflected, (with the reflected wave having a larger amplitude than the incident wave amplitude) if it can excite a wave on the other side of the corotation barrier. If there exists a reflecting boundary at the inner disc radius $r_{\rm in}$, then a global overstable mode partially trapped between $r_{\rm in}$ and $r_{\rm IL}$ will result (see, e.g. Narayan et al.~1987 for specific examples in the shearing sheet model, and Goodman \& Evans,1999 and Shu et al.~ 2000 for global mode analysis of singular isothermal discs). The simple dispersion relation (\ref{eq:disp}), however, does not capture an important effect in the disc, i.e., corotation resonance or corotation singularity. Near corotation $|\tomega|\ll\kappa$, the WKB dispersion relation for the wave is [see equation (\ref{eq:disper}) below] \be \tomega={2\Omega k_\phi\over k_r^2+k_\phi^2 +\kappa^2/c^2}\left({d\over dr}\ln{\kappa^2 \over 2\Omega\Sigma}\right)_c, \label{eq:disp2}\ee where $k_\phi=m/r$ and $\Sigma$ is the surface density, and the subscript ``c'' implies that the quantity is evaluated at $r=r_c$. The quantity \be \zeta\equiv {\kappa^2\over 2\Omega\Sigma}= {(\nabla\times {\bf u_0})\cdot {\hat z}\over\Sigma} \label{eq:zeta}\ee is the vortensity of the (unperturbed) flow (where ${\bf u_0}$ is the flow velocity). The dispersion relation (\ref{eq:disp2}) describes Rossby waves, analogous to those studied in geophysics (e.g. Pedlosky 1987) \footnote{A Rossby wave propagating in the Earth's atmosphere satisfies the dispersion relation $\tomega=(2k_\phi/k^2R)(\partial\Omega_3/ \partial\theta)$, where $k^2=k_\phi^2+k_\theta^2$, $\Omega_3=\Omega\cos\theta$ is the projection of the rotation rate on the local surface normal vector and $\theta$ is the polar angle (co-latitude).}. For $k_r^2\gg \kappa^2/c^2$ and $k_r^2\gg k_\phi^2$, we see that Rossby waves can propagate either outside the rotation radius $r_c$ (when $d\zeta/dr>0$) or inside $r_c$ (when $d\zeta/dr<0$). In either case, we have $k_r\rightarrow\infty$ as $r\rightarrow r_c$. This infinite wavenumber signifies wave absorption (cf. Lynden-Bell \& Kalnajs 1972 in stellar dynamical context; Goldreich \& Tremaine 1979 in the context of wave excitation in discs by a external periodic force; see also Kato 2003, Li et. al. 2003, and Zhang \& Lai 2006 for wave absorption at the corotation in 3D discs). At corotation, the wave pattern angular speed $\omega/m$ matches $\Omega$, and there can be efficient energy transfer between the wave and the background flow, analogous to Landau damping in plasma physics. Narayan et al.~(1987) treated this effect as a perturbation of the shearing sheet model, and showed that the corotational absorption can convert neutral modes in a finite shearing sheet into growing or decaying modes. Papaloizou \& Pringle (1987) used a WKB method to examine the effect of wave absorption at corotation on the nonaxisymmetric modes in an unbound (with the outer boundary extending to infinity) cylindrical torus. In this paper, we derive explicit expressions for the complex reflection coefficient and transmission coefficient for waves incident upon the corotation barrier. We pay particular attention to the behavior of perturbations near the corotation resonance/singularity. Our general expressions include both the effects of corotation amplifier and wave absorption at corotation (which depends on $d\zeta/dr$). We show explicitly that depending on the sign of $d\zeta/dr$, the corotation resonance/singularity can either enhance or diminish the super-reflectivity, and this can be understood in terms of the location of the Rossby wave zone relative to the corotation radius. Our paper is organized as follows. After presenting the general perturbation equations (section 2), we discuss the the wave dispersion relation and propagation diagram, and derive the local solutions for the wave equation around the Lindblad resonances and corotation resonance (section 3). We then construct global WKB solution for the wave equation, and derive the wave reflection, transmission and corotational damping coefficients in section 4. An alternative derivation of the wave damping coefficient is presented in section 5. Readers not interested in technical details can skip Sections 2-5 and concentrate on Section 6, where we illustrate our results and discuss their physical interpretations. Section 6.1 contains a numerical calculation of the wave reflectivity across corotation and discusses the limitation of the WKB analysis. We discuss how global overstable modes may arise when super-reflection at the corotation is present in section 7 and conclude in section 8. | In this paper we have derived explicit expressions for the reflection coefficient, transmission coefficient and wave absorption coefficient when a wave is scattered by the corotation barrier in a disc. These expressions include both the effects of corotation amplifier (which exists regardless of the gradient of the vortensity $\zeta=\kappa^2/\Omega\Sigma$ of the background flow) and wave absorption at the corotation (which depends on $d\zeta/dr$). They demonstrate clearly that the corotation wave absorption plays a dominant role in determining the reflectivity and that the sign of $d\zeta/dr$ determines whether the corotation singularity enhances or diminishes the super-reflectivity. Our result can be understood in terms of the location of the Rossby wave zone relative to the corotation radius. We also carried out numerical calculations of the reflectivity. Our result provides the conditions (in terms of disc thickness, rotation profile and surface density profile) for which super-reflection is achieved and global overstable modes in discs are possible. In future works we will explore global oscillation modes and their stabilities in a variety of astrophysical contexts, ranging from accreting white dwarfs to accreting black hole systems. The possible overstabilities of these modes are directly linked to the effects studied in this paper and may provide explanations for some of the quasi-periodic variabilities observed in these systems. | 7 | 10 | 0710.2313 |
0710 | 0710.4123_arXiv.txt | We present new estimates of ages and metallicities, based on FORS/VLT optical (4400--5500\,\AA) spectroscopy, of 16 dwarf elliptical galaxies (dE's) in the Fornax Cluster and in Southern Groups. These dE's are more metal-rich and younger than previous estimates based on narrow-band photometry and low-resolution spectroscopy. For our sample we find a mean metallicity ${\rm [Z/H]} = -0.33$\,dex and mean age 3.5\,Gyr, consistent with similar samples of dE's in other environments (Local Group, Virgo). Three dE's in our sample show emission lines and very young ages. This suggests that some dE's formed stars until a very recent epoch and were self-enriched by a long star formation history. Previous observations of large near-infrared ($\sim 8500$\,\AA) \ionCaT\ absorption strengths in these dE's are in good agreement with the new metallicity estimates, solving part of the so-called Calcium puzzle. | The observed strength of the near-infrared \ionCaT\ triplet absorption lines in early-type galaxies has presented astronomers with an interesting puzzle over the past couple of years. \citet{cenarro01} defined a new CaT* index, carefully correcting for the underlying H Paschen absorption. Whereas other metallicity tracers, such as Mg$_2$, correlate with velocity dispersion $\sigma$, it was found that CaT* \emph{anti-correlates} with $\sigma$ in elliptical galaxies (E's) and in bulges of spiral galaxies \citep{saglia02, cenarro03, falcon03}. Population synthesis model predictions also show that, for sub-solar metallicities, CaT* should be sensitive to metallicity but virtually independent of age, while at super-solar metallicity, the CaT* saturates \citep{vazdekis03}. However, taking metallicities estimated from optical spectra, \citet{saglia02} reported that the measured CaT* values in E's are smaller by 0.5\,{\AA} than those predicted by population synthesis models. \citet{michielsen03} (hereafter Paper~1) showed that the anti-correlation of CaT* with $\sigma$ extends into the dwarf elliptical (dE) regime. These dE's were expected to have metallicities of the order of [Z/H]$\,\sim -1$, and ages of the order of 10\,Gyr \citep{heldmould94, rakos01}. The measured CaT* values were significantly larger than those expected for such old, metal-poor stellar systems. All of the proposed solutions to this conundrum, such as variations of the initial-mass function or the calcium yield as a function of metallicity or velocity dispersion, require considerable fine-tuning. None of them satisfactorily explains both the small CaT* in E's and the large CaT* in dE's without creating other difficulties, e.g.\ with the FeH~$\lambda$9916 index values observed in bright ellipticals (FeH is strong in dwarf stars but nearly absent in giants) and with stellar mass-to-light ratios \citep{saglia02, cenarro03}. However, the stellar populations of E's and dE's are most likely not single-age, single-metallicity populations, or SSPs, as was implicity assumed in essentially all age and metallicity estimates \citep[see e.g.][]{pasquali05}. Moreover, the stars that dominate the blue spectral range (mostly hot dwarf stars) do not necessarily have the same mean ages/metallicities as the stars producing the red light (mostly cool giants). These issues, together with systematic uncertainties inherent to population synthesis tools, potentially contribute to the CaT puzzle. Because of their low surface brightness, accurate estimates of the ages and metallicities of dE's are still scarce. Recent studies of dE's in the Virgo cluster \citep{g03, v04} report younger ages and higher metallicities than found in Fornax dE's. In Paper~1, the ages and metallicities of the dE's were taken from the literature, where low-resolution, modest S/N spectroscopic \citep{heldmould94} or narrow-band photometric \citep{rakos01} techniques were used. As a sanity check, we have now acquired high-resolution, high S/N optical spectra with FORS/VLT of all the dE's for which we presented CaT* measurements in Paper~1. This puts us in a position where we can for the first time compare the CaT* measurements with model predictions based on robust age and metallicity estimates. | \label{disc} At least in the dE regime, the \ionCaT\ Triplet puzzle seems to be solved. With the new age and metallicity estimates presented in this Letter, the predicted and observed CaT* indices are in good agreement for all sample galaxies but two. The fact that CaT* values predicted using age/metallicity estimates that were derived from spectral features in the blue part of the spectrum agree with the observed CaT* values in the NIR indicates that the SSP assumption is not the cause of the CaT puzzle for dE's. Rather, the CaT puzzle in the dE regime was caused by the spuriously low metallicities and high ages, derived from lower resolution spectra using less sophisticated theoretical models, assigned to Fornax dE's. This shows that the CaT* index, and, in old stellar systems in which the PaT index is small, the CaT index as well, is indeed a good tracer of metallicity. This also solves the apparent dichotomy between dE's on the one hand and globular clusters, Local Group dwarf spheroidals (dSph), and ultra-compact dwarfs (UCD) on the other hand. CaT line-strengths measured in individual stars of Local Group dSphs have been shown to be a very accurate tracer of metallicity \citep{b06,to03} and [Fe/H]-values derived from CaT measurements have been used extensively to construct metallicity distributions of the stars in dSphs. Also, for UCDs \citep{ev07} and globular clusters \citep{saglia02} the CaT index has proved to be an excellent tracer of metallicity. Here, we have shown that in dE's as well, the CaT* index measured from integrated-light spectra can be used as a tracer of metallicity. To summarise, we derive new age and metallicity estimates for 16 dE's in the Fornax Cluster and in Southern Groups using high S/N optical VLT/FORS1+2 spectra. We have measured the H$\beta$, Mg$b$, Fe5270, and Fe5335 indices in the Lick/IDS system and applied the TMB03 models to them. We find that these dE's have solar [$\alpha$/Fe] abundance ratios. A full-spectrum fit using Pegase-HR with the ELODIE.3.1 stellar library provides us with a second, independent age and metallicity estimate for these galaxies. We find both approaches to be in excellent agreement. With mean metallicity ${\rm [Z/H]} = -0.35$\,dex and ages younger than $\approx 7$\,Gyr, these dE's are more metal-rich and younger than previously thought. Some even show strong emission lines, an indication of on-going star formation, in agreement with previous H$\alpha$ imaging of dE's \citep{derijcke03,michielsen04}. The ages and metallicities we derive for the dE's in the Fornax cluster and in Southern groups fall in roughly the same range as those derived by \citet{g03} and \citet{v04} for dE's in the Virgo cluster. This is at variance with previous estimates for Fornax dEs which yielded lower metallicities and higher ages \citep{heldmould94, rakos01}, based on lower resolution spectra and less sophisticated theoretical models. The new age and metallicity estimates are in good agreement with the observed \ionCaT\ triplet absorption strengths, solving the Calcium puzzle for low-mass systems. | 7 | 10 | 0710.4123 |
0710 | 0710.4854_arXiv.txt | Big Bang nucleosynthesis (BBN) is the earliest sensitive probe of the values of many fundamental particle physics parameters. We have found the leading linear dependences of primordial abundances on all relevant parameters of the standard BBN code, including binding energies and nuclear reaction rates. This enables us to set limits on possible variations of fundamental parameters. We find that ${^7}$Li is expected to be significantly more sensitive than other species to many fundamental parameters, a result which also holds for variations of coupling strengths in grand unified (GUT) models. Our work also indicates which areas of nuclear theory need further development if the values of ``constants'' are to be more accurately probed. | 7 | 10 | 0710.4854 |
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0710 | 0710.3190_arXiv.txt | We set sensitive upper limits to the X-ray emission of four Type Ia supernovae (SNe~Ia) using the \cxo. SN~2002bo, a normal, although reddened, nearby SN~Ia, was observed 9.3 days after explosion. For an absorbed, high temperature bremsstrahlung model the flux limits are $3.2\times 10^{-16}$ ergs cm$^{-2}$ s$^{-1}$ (0.5--2 keV band) and $4.1\times 10^{-15}$ ergs cm$^{-2}$ s$^{-1}$ (2--10 keV band). Using conservative model assumptions and a 10 km s$^{-1}$ wind speed, we derive a mass loss rate of $\dot{M} \sim 2\times 10^{-5}\, M_\odot$ yr$^{-1}$, which is comparable to limits set by the non-detection of H$\alpha$ lines from other SNe~Ia. Two other objects, SN~2002ic and SN~2005gj, observed 260 and 80 days after explosion, respectively, are the only SNe~Ia showing evidence for circumstellar interaction. The SN~2002ic X-ray flux upper limits are $\sim$4 times below predictions of the interaction model currently favored to explain the bright optical emission. To resolve this discrepancy we invoke the mixing of cool dense ejecta fragments into the forward shock region, which produces increased X-ray absorption. A modest amount of mixing allows us to accommodate the \chandra\ upper limit. SN~2005gj is less well studied at this time. Assuming the same circumstellar environment as for SN~2002i, the X-ray flux upper limits for SN~2005gj are $\sim$4 times below the predictions, suggesting that mixing of cool ejecta into the forward shock has also occurred here. Our reanalysis of Swift and \chandra\ data on SN~2005ke does not confirm a previously reported X-ray detection. The host galaxies NGC~3190 (SN~2002bo) and NGC~1371 (SN~2005ke) each harbor a low luminosity ($L_X \sim 3-4 \times 10^{40}$ ergs s$^{-1}$) active nucleus in addition to wide-spread diffuse soft X-ray emission. | Type Ia supernovae (SN~Ia) are an important subclass of supernova (SN) that are thought to arise from explosions of white dwarfs in binary systems, although the exact nature of their progenitor systems is largely a mystery. This ignorance is not due to a lack of effort since SN~Ia are the subject of intense scrutiny, not least of all because of their importance as cosmological probes. SN~Ia have been used to measure the Hubble constant \citep{ham95, riess96} and have even given strong evidence for a non-zero cosmological constant \citep{riess98, perl99}. Our failure to identify SN~Ia progenitors highlights a major gap in our understanding of stellar evolution in binary systems, and presents a stumbling-block to understanding the chemical evolution of galaxies, since Ia's are known to be efficient producers of iron. In addition, until SN~Ia progenitor systems are clearly identified, it will remain difficult to convincingly eliminate evolutionary sources of systematic error in the observed trend of SN brightness with redshift. The current model framework for Type Ia SNe involves carbon deflagration/detonation in a white dwarf driven close to the Chandrasekhar limit by accretion. The most likely type of progenitor system \citep{branch95} is a C-O white dwarf accreting H/He-rich gas from a companion, either from its wind or through Roche lobe overflow. Double degenerate scenarios with coalescing pairs of C-O white dwarfs are also possible, while sub-Chandrasekhar-mass white dwarfs are less likely \citep{branch95}. In the non-coalescing scenarios, there will be circumstellar gas whose composition and geometry depend on the nature of the progenitor system. If the circumstellar medium (CSM) emits radiation, or absorbs radiation from the SN, this can be used to distinguish between types of progenitor systems. Thermal X-ray emission is expected to arise from the hot gas in the interaction region between the rapidly moving supernova ejecta and the wind from presupernova mass loss. Detecting this X-ray emission offers a direct and potentially quite sensitive probe of the amount of CSM. Because of its high sensitivity, broad bandwidth, unprecedented imaging capability, and ability to respond rapidly to a target of opportunity (ToO) request, the \cxo\ is the premier instrument to search for faint X-ray emission from transient point sources. In the following we present the analysis and interpretation of \chandra\ observations of four recent SNe~Ia. Two are fairly normal SNe~Ia, while the others are peculiar cases showing strong H$\alpha$ emission from circumstellar interaction. The following section describes the observed targets, \S3 presents the observations and techniques, \S4 uses the upper limits on X-ray flux to constrain the nature of the CSM, and the final section concludes. In an Appendix we describe the X-ray properties of the host galaxies for the two nearby systems, NGC~3190 and NGC~1371. | With our \chandra\ observation of SN~2002bo we have set the most sensitive X-ray flux upper limits at an earlier epoch than for any previous SN~Ia. The previous best case was that of SN~1992A for which a \rosat\ upper limit of $\sim$$10^{-14}$ ergs cm$^{-2}$ s$^{-1}$ was set on the 0.5-2 keV band X-ray flux $\sim$35 days after explosion \citep{schpet93}. Our \chandra\ limit in the same energy band is an order of magnitude lower and was set only 9 days after explosion. We also set a sensitive upper limit in the important 2--10 keV band. Converting our flux upper limits to constraints on the density of the CSM in the system depends on uncertain assumptions about the thermodynamic state of the hot plasma in the expanding ejecta and shocked wind. As we have shown in this paper there is roughly an order of magnitude difference in the inferred wind density under the assumption of fully equilibrated electron and ion temperatures vs.\ the case with non-equilibrated temperatures. The latter case yields the more conservative (i.e., higher) constraint: a wind density parameter $w < 1.2 \times 10^{15}$ g cm$^{-1}$. In terms of a slow wind with velocity of $v_{w10}=10$ km s$^{-1}$ this corresponds to a upper limit on the mass-loss rate of $\dot{M} < 2 \times 10^{-5}~M_{\odot}$ yr$^{-1}$. The much higher sensitivity of the \chandra\ data notwithstanding, this value is {\em larger} by a factor of 10 or so than the mass loss upper limit derived from \rosat\ data by \citet{schpet93} for SN~1992A. This discrepancy is due to the simplicity of the model used by \citet{schpet93} to evaluate their upper limit: they do not include absorption by residual wind material above the forward shock, assume electron-ion temperature equilibration, compare to model luminosities without applying a bolometric correction, and calculate the SN age from the time of maximum rather than from the time of explosion. In other comparisons our \chandra\ upper limit is comparable to those found previously from limits on the H$\alpha$ flux for the normal type Ia supernovae SN~1994D \citep{cumm96} and SN~2001el \citep{matetal05}. In general limits on $\dot{M}$ of $\sim$$10^{-5}~M_{\odot}$ yr$^{-1}$ are not stringent enough to rule out the class of symbiotic-type binaries as SN~Ia progenitors at least not for these particular cases. The value of these results lies in our ability to calculate, using well understood physics, the expected X-ray emission from hot gas. Viewed in this light, we briefly discuss limits on the CSM obtained from radio nondetections of a number of nearby SNe~Ia \citep{panagetal06}. The radio results rely on semi-empirical parameterized functional forms for the time- and frequency-dependence of synchrotron emission and free-free absorption, whose relevant parameters are assumed to have values given by radio results for SNe~Ib/c. Likewise the essential parameter, i.e., the one linking the wind density parameter to the radio luminosity of the SN, cannot be calculated from theory with any accuracy, and an empirical calibration, again drawn from measurements of SNe~Ib/c, must be used. Although the sensitive radio flux limits clearly argue for low density environments around SN~Ia, the extremely low numerical limits on the mass-loss rates claimed by \citet{panagetal06} cannot yet be considered definitive. The other normal SN~Ia we study here is SN~2002ke. We re-examine the claim by \citet{immler06} of a tentative X-ray detection by Swift and find that we cannot substantiate it. We pay particular attention to the astrometric and photometric calibration of the Swift X-ray data by comparing to a \chandra\ observation done several months later. We find no evidence for a significant X-ray detection of SN~2005ke by either Swift or \chandra, to a flux limit that is several factors higher than what we obtained for SN~2002bo. Since this limit is at a much later epoch, when the intensity of the CSM/ejecta interaction should be much reduced, we did not attempt to determine numerical limits to the wind density parameter for this SN. We have also presented \chandra\ upper limits for the two examples of SNe Ia with clear evidence for circumstellar interaction: SN~2002ic and SN~2005gj. The upper limit on the X-ray luminosity of SN~2002ic in the $0.5-6$ keV band unexpectedly reveals a serious drawback to the interaction model proposed previously (CCL): the predicted X-ray flux turns out to be larger by at least a factor of four. We identified the major missing element of the model responsible for the controversy: macroscopic mixing of cool metal-rich ejecta fragments into the hot gas of the forward shock, which results in strong absorption of the X-rays emitted by the forward shock. Interestingly, the absorption of X-rays by mixed shocked ejecta should have the effect of decreasing the required CS density in the model to explain the late time optical luminosity. This effect together with the higher radiation efficiency of the interaction with a clumpy CS matter should result in a slightly higher expansion velocity of the shocked SN ejecta in the interaction model. These outcomes appear to be preferred by the optical spectra of SN~2002ic (see CCL). SN~2005gj appears to belong to the SN~2002ic-like family of SN~Ia with dense CS envelopes \citep{aldetal06,prietoetal07}. The \chandra\ upper limit from this object at about day 80 is a factor of four larger than the flux predicted by the interaction model of SN~2002ic recomputed for the corresponding epoch. Although the strength of the interaction argues against it, one possible explanation is that the CS density around SN~2005gj is just somewhat lower than around SN~2002ic. A better possibility invokes mixing of the shocked ejecta, which can reduce the emergent X-ray flux. Which of these is more likely to be the correct explanation requires a better understanding of the environment of SN~2005gj than we have at present. Future studies of the optical light curve and spectra of this interesting object are strongy encouraged. Resolving the issue of the X-ray non-detection of SN~2002ic-like objects has its dark side due to our introduction of the parameter $\tau_{\rm oc}$, which is essentially incalculable. Even three-dimensional hydrodynamical simulations are unlikely to be able to recover this value in a fully self-consistent way. We, therefore, can predict in detail neither the spectrum nor the evolution of the X-ray flux from SN~2002ic-like supernovae. Two relevant remarks can be made, however. First, in the case of SN~2002ic the model without mixing predicts an increasing X-ray flux in the band below 10 keV up until day $\sim$400 (CCL). Therefore, X-ray detection at earlier epochs is unlikely to be more favorable, as the lack of detection of SN~2005gj at roughly 80 days after explosion (versus 260 days for SN~2002ic) tends to support. The second remark relates to the spectrum of the emergent X-ray emission: hard X-rays (i.e., with photon energies greater than $\sim$20 keV) are not affected by the absorption that mutes the lower energy flux (see Fig.~\ref{ic_mod}). Future sensitive X-ray observations covering a wider energy band ($1-30$ keV), therefore, should reveal a SN~2002ic-like event as a strongly absorbed hard X-ray source and thus verify the proposed mixing scenario. | 7 | 10 | 0710.3190 |
0710 | 0710.3159_arXiv.txt | We present a spatially resolved comparison of the stellar-mass and total-mass surface distributions of nine early-type galaxies. The galaxies are a subset of the Sloan Lens ACS survey \citep[or SLACS;][]{2006ApJ...638..703B}. The total-mass distributions are obtained by exploring pixelated mass models that reproduce the lensed images. The stellar-mass distributions are derived from population-synthesis models fit to the photometry of the lensing galaxies. Uncertainties -- mainly model degeneracies -- are also computed. Stars can account for all the mass in the inner regions. A Salpeter IMF actually gives too much stellar mass in the inner regions and hence appears ruled out. Dark matter becomes significant by the half-light radius and becomes increasingly dominant at larger radii. The stellar and dark components are closely aligned, but the actual ellipticities are not correlated. Finally, we attempt to intuitively summarize the results by rendering the density, stellar-vs-dark ratio, and uncertainties as false-colour maps. | In the current paradigm of galaxy formation, the building block of structure is a dark matter halo, consisting mainly of non-baryonic dark matter together with $\sim 15\%$ baryons in the form of gas. Dark halos are thought to originate from the collapse of primordial density fluctuations, growing unimpeded until virial equilibrium is reached. Within these halos the baryonic component dissipates energy, collapsing further towards the center and eventually forming the visible galaxy. Subsequent mergers redistribute the matter within halos. To work out the details of the basic picture, which goes back to \cite{1978MNRAS.183..341W}, it is essential to determine the connection between the visible galaxies and dark halos. Strong gravitational lensing by galaxies is potentially a very useful way of doing this, since the total mass of a lensing galaxy is relatively easy to constrain. Also, since lensing tends to be more effective for distant galaxies ($z_{\rm lens}\sim0.1$ to 1), it nicely complements the stellar-dynamical techniques applicable in nearby galaxies. The difficulty with lensing galaxies (that is to say, galaxies producing multiple images of background sources) is that they are relatively rare. Till recently, only about 80 were known. But recently, 28 new galaxy lenses have been discovered by the Sloan Lens ACS Survey \citep[SLACS;][]{2006ApJ...638..703B} thanks to a new survey strategy, which eventually may double the number of strong lenses. The method is to select galaxies from the Sloan survey \citep{2000AJ....120.1579Y} that have emission lines indicating high-redshift background objects, and then to image these candidates using the Advanced Camera for Surveys on the Hubble Space Telescope. A basic analysis of a sample of galaxy lenses is to fit simple lens models and then compute $M/L$ and its evolution with redshift. This was first done by \cite{1998ApJ...509..561K}. An extension is to place lensing galaxies on the Fundamental Plane of ellipticals, using a measured or model-derived velocity dispersion \citep{2000ApJ...543..131K,2003ApJ...587..143R,2006ApJ...640..662T}. Other work \citep{2004ApJ...611..739T,2006ApJ...649..599K} compares lens models with the measured dispersions to constrain the mass profiles of the galaxies. These studies have found no unexpected features or trends with redshift, and argue in favor of passive evolution. A more detailed analysis involves modeling both the lens mass distribution and the stellar population. The star-formation history is not well-constrained by the observed fluxes and colours and must be marginalized over, but the stellar mass is fairly insensitive to model assumptions apart from the initial mass function (IMF). The lensing mass distribution, when aggressively modeled, turns out to have much larger uncertainties than simple models assume; nevertheless, the uncertainties can be estimated and useful conclusions drawn. \cite{2005ApJ...623L...5F} found massive ellipticals to show a transition from no significant dark matter within $\Re$ to dark-matter dominance by $5\Re$, whereas lower-mass galaxies showed no significant dark halos even at $5\Re$. The radial gradient in the dark-matter fraction agreed with the results on nearby galaxies derived from stellar dynamics \citep{2005MNRAS.357..691N}. In this paper we extend the detailed comparison of stellar and total mass to two dimensions, using a subsample of the SLACS lenses. All the SLACS objects have a background galaxy that is lensed into two or four extended images. In nine of the objects, we can identify small features within the extended images. For lenses showing point-like multiply-imaged features there is a well-developed technique for reconstructing the projected mass distribution, along with uncertainty estimates \citep{2004AJ....127.2604S}. Accordingly we take these nine lenses as our sample. The stellar mass content is estimated by combining the available photometry of the lensing galaxies with population-synthesis models \citep{2003MNRAS.344.1000B}. \begin{figure} \begin{center} \includegraphics[width=3.4in]{2DLens_f1.eps} \caption{Stellar masses derived from the available photometry. The predictions of the Bruzual \& Charlot (2003) models are shown for a Chabrier (2003) IMF. A galaxy with apparent magnitude $\rm F814W=20$ is considered at two different redshifts as labelled. The shaded regions span the model predictions when assuming a wide range of $\tau$ models (i.e., an exponentially decaying star formation history at fixed metallicity). The dark/light shaded regions correspond to a formation redshift of $z_F=5$ and $2$, respectively, and all span a range of star formation timescales ($-1<\log \tau ({\rm Gyr}) < +1$) and metallicities ($-1<\log Z/Z_\odot <+0.3$).} \label{fig:stellar} \end{center} \end{figure} | The potential usefulness of early-type lensing galaxies for understanding the interdependence of baryons and dark matter in galaxy formation and evolution is widely appreciated. Lensing can be used to map the total mass and starlight can be used to map the baryons. In this paper we do this for nine galaxies from the SLACS survey by \cite{2006ApJ...638..703B}. From the lensing data we derived free-form pixelated models for the total mass, and from the galaxy photometry we computed stellar population-synthesis models. In both casses we generated Monte-Carlo ensembles of models, in order to marginalize over unknowns such as lensing degeneracies and star-formation histories, thus obtaining realistic uncertainties. The technique is basically the same as in \cite{2005ApJ...623L...5F}, but whereas the earlier work only compared radial profiles now we compare stellar and total mass in 2D. Related work has been done on larger samples of galaxies, but is limited to fitting simple parameterized models for the lens, and does not model the stellar populations at all. The 2D mass maps are shown in Figures \ref{fig:mass} and \ref{fig:cmaps}. The former overlays contours of $\Sigtot$ on a grayscale of $\Sigstel$. The latter shows the same information with uncertainties as well, all encoded in false colour: red for stellar mass, blue for dark, and pale versus coloured for uncertainty. It is evident that (a)~these galaxies are dominated by stars in the inner regions, but mainly dark matter in the outer regions, and (b)~stellar and dark components are well-aligned, but neither has a simple elliptical shape. One can of course still compute an ellipticity defined as a moment, and this is shown in Figure~\ref{fig:ellip}. We see that ellipticities of the stellar and total mass are uncorrelated in magnitude but are almost perfectly aligned. It would be interesting to see if this is true of galaxy-formation simulations. The profiles (Figure~\ref{fig:profile}) show that dark matter halos begin to dominate around the half-light radius, although some galaxies seem to present a stronger contribution from dark matter even inside $\Re$ (e.g., J2300). The present sample is too small to extract any strong correlation of the dark matter distribution with global properties such as total mass or luminosity. However, the trend found previously \citep{2005ApJ...623L...5F}, namely that there should be more dark matter in more massive galaxies, with a scaling of roughly $\Mtot\propto\Mstel^{1.2}$ (which is equivalent to the tilt of the Fundamental Plane) is compatible with the combined data of lensing galaxies (labelled CASTLES and SLACS) along with the dynamical analysis of local galaxies with the SAURON integral field unit \citep{2006MNRAS.366.1126C}. The top-left panel of Figure~\ref{fig:profile} also shows that a Salpeter (1955) IMF cannot be used to estimate stellar masses as the population synthesis models predict too much stellar mass compared to the total mass obtained from our lensing studies. This result also agrees with the analysis of the SAURON sample. We emphasize that stellar and total masses are obtained from different data through models of different physical processes. There is no tuning to give similar results. That the stellar and total densities come out compatible at the centers ---where the baryon content is expected to dominate the mass budget--- indicates that systematic effects are not significant. | 7 | 10 | 0710.3159 |
0710 | 0710.4065_arXiv.txt | {The majority of X-ray burst sources do not display a burst rate that increases with luminosity as expected, but this is seen in the two clocked bursters XB\th 1323-619 and GS\th 1826-24. We present a detailed investigation of these two sources which in the case of the first source, spans 18 years. Based on measurements of the burst rate, X-ray luminosity, the $\alpha$-parameter and the two time constants generally present in the burst decays, we demonstrate the importance of the {\it rp} nuclear burning process. A detailed comparison with theory shows that although the burst rate in each source agrees well with the theoretical value, there is a difference of more than a factor of 5 in the burst rate at a given luminosity between the sources. We show that the main reason for this is that the two sources have substantially different emitting areas on the neutron star in non-burst emission, a factor often neglected. Variation of this area may explain the inverse relation of burst rate with luminosity in the majority of burst sources. | % \label{sect:intro} X-ray bursting has been known for many years (Grindlay et al. 1976) as a phenomenon taking place in many low mass X-ray binaries (LMXB) of luminosities less than 10$^{38}$ erg s$^{-1}$ consisting of a rapid rise in intensity by a factor of $\sim$20, followed by an exponential decay over about 50 seconds and repeating on a timescale of hours. Bursting is usually not seen in higher luminosity sources forming the Z-track class which exhibit intensity increases as flaring over much longer timescales. Thus if we divide LMXB by luminosity into Atoll or Z-track sources, X-ray bursting is generally observed in the Atoll class. Bursting appears unrelated to inclination angle since it is seen not only in Atoll sources not displaying orbital-related behaviour but also in dipping sources which are seen at high inclination. X-ray bursting has been extensively studied (Lewin et al. 1995; Strohmayer \& Bildsten 2006) and it was realized at an early stage to consist of unstable nuclear burning. Measured values of the $\alpha$-parameter, defined as the ratio of the energy released in steady mass accretion to the integrated energy of the burst, agreed well with values expected for unstable burning. Thus X-ray bursting was recognized as explosive burning of recently accreted material on the surface of the neutron star (Woosley \& Taam 1976), the material building up on the neutron star over a period of hours. In spite of this, understanding of X-ray bursting is relatively poor. On the above basis the rate of X-ray bursting would be expected to be generally stable when the luminosity of a sources is stable since then the mass accretion rate $\dot M$ is stable, but bursting rate is rather erratic in most sources. Even worse, the rate of bursting should increase when source luminosity increases, as less time is required to accumulate the necessary mass and achieve the required plasma density and temperature on the surface of the neutron star for unstable burning, but many sources display the opposite of this (van Paradijs et al. 1988). In two sources, which may be called the ``clocked bursters'', bursting is close to being exactly periodic and the burst rate increases with luminosity: GS\th 1826-24 (the original clocked burster) and XB\th 1323-619 which we have studied over an extended period as one of the $\sim$10 dipping LMXB. In the present work, we make detailed comparisons of these two well-behaved sources with each other, and with the theory of unstable burning, to identify the reasons why the sources differ so much from each other, and to try to understand their behaviour in terms of theory. This would facilitate the understanding of the majority of burst sources that are not well-behaved. \begin{figure*}[!ht] % \begin{center} \includegraphics[width=40mm,height=140mm,angle=270]{balucinska_2007_revised_fig1} \caption{{\it Exosat} lightcurve of XB\th 1323-619. Regular bursting takes place every 5.33 hr whereas X-ray dipping occurs at the orbital period of 2.93 hr.} \label{} \end{center} \end{figure*} | We have shown that marked differences in the burst rate occur between XB\th 1323-619 and GS\th 1826-24. and have shown that the difference is mostly due to the very different emitting areas of non-burst emission on the neutron star, a factor that has not generally been allowed for in discussion of X-ray bursting. If we view the burst rate as a function of $\dot m$, the accretion rate per unit emitting area, we find that the sources follow relations that agree within 40\%, so that both can be explained to this accuracy by the same model. It is expected that the non-burst emitting area will be an important factor in the other burst sources and in future work we will test the hypothesis that in sources where the burst rate decreases with $L$, this behaviour can be explained in terms of the non-burst area increasing with $L$ so that $\dot m$ decreases. | 7 | 10 | 0710.4065 |
0710 | 0710.5580_arXiv.txt | Ram pressure stripping can remove significant amounts of gas from galaxies that orbit in clusters and massive groups, and thus has a large impact on the evolution of cluster galaxies. In this paper, we reconstruct the present-day distribution of ram-pressure, and the ram pressure histories of cluster galaxies. To this aim, we combine the Millennium Simulation and an associated semi-analytic model of galaxy evolution with analytic models for the gas distribution in clusters. We find that about one quarter of galaxies in massive clusters are subject to strong ram-pressures that are likely to cause an expedient loss of all gas. Strong ram-pressures occur predominantly in the inner core of the cluster, where both the gas density and the galaxy velocity are higher. Since their accretion onto a massive system, more than 64 per cent of galaxies that reside in a cluster today have experienced strong ram-pressures of $>10^{-11}$ dyn cm$^{-2}$ which most likely led to a substantial loss of the gas. | In clusters, galaxies can lose some or all of their gas by ram pressure stripping (RPS) due to their motion through the intracluster medium (ICM). Both analytical estimates and hydrodynamical simulations show that RPS can remove a significant amount of gas from galaxies, and can thus explain observations such as the HI deficiency of cluster disc galaxies (see e.g.~\citealt{Haynes_Giovanelli_1986,Solanes_etal_2001,cayatte94}), and the truncated star forming discs in the Virgo cluster (\citealt{Warmels_1988,koopmann04b}). RPS only affects the gaseous component of the galaxy so that a distinct feature of ram pressure stripped galaxies is that their gas discs are distorted or truncated while their stellar discs appear undisturbed. An increasing number of observations of ram-pressure stripped galaxies have become available in the last years (see for example \citealt{Irwin_etal_1987}, \citealt{Veilleux_etal_1999}, \citealt{kenney04}, \citealt{vollmer04a}, \citealt{chung:07}, \citealt{sakelliou:05,sun:07}). RPS is commonly cited in early work as a possible explanation for the increased fraction of S0 galaxies in rich clusters relative to the field (\citealt{Biermann_Tinsley_1975}, \citealt{Butcher_Oemler_1978}). This explanation was dismissed in the original paper by \cite{dressler_1980} on the basis of the observation that the relation between different galaxy populations and local density appears to hold independently of the cluster richness. Later studies pointed out that additional mechanisms that lead to a significant redistribution of mass and/or new star formation are required to explain the entire S0 population of galaxy clusters (see for example \citealt{moran:07} and references therein). The role of RPS in the chemical enrichment of the ICM has also been discussed. Observational data suggest that the ICM is enriched with metals to approximately one third of the solar value, suggesting that some fraction of the metals must have originated from cluster galaxies and since been removed from them. Processes responsible for the supply of this enriched gas include AGN feedback (see e.g. \citealt{roediger:07c} and references therein), galactic winds driven by supernovae explosions, and ram-pressure stripping (\citealt{white:91,mori:00,schindler05,domainko:06}). It should be noted, however, that although RPS certainly affects the metallicity of the ICM, it may not be the dominant mechanism. Numerical simulations by \cite{domainko:06} indicate that RPS can account for only about 10 per cent of the ICM metal content within a radius of 1.3 Mpc. The first analytical estimate of RPS dates back to the paper by \citet{gunn72} who proposed that for galaxies moving face-on through the ICM the success or failure of RPS can be predicted by comparing the ram pressure with the galactic gravitational restoring force per unit area. Later hydrodynamical simulations of RPS (\citealt{abadi99}, \citealt{quilis00}, \citealt{schulz01}, \citealt{marcolini03,acreman03}, \citealt{roediger:06a}, \citealt{roediger:06b}, \citealt{roediger:07}) suggest that this analytical estimate fares fairly well as long as the galaxies are not moving close to edge-on. The ICM-ISM interaction is, however, a complex process influenced by many parameters. Different aspects have been studied by several groups. \cite{roediger05} and others have shown that the ICM-ISM interaction is a multistage process: The most important phases are the instantaneous stripping, on a time-scale of a few 10 Myr, an intermediate phase, on a time-scale of up to a few 100 Myr, and a continuous stripping phase that, in principle, could continue until all gas is lost from the galaxy. While numerical simulations and observations indicate that RPS has important consequences on the amount of gas in cluster galaxies, this physical process is usually not included in semi-analytic models of galaxy formation. The effect of ram-pressure stripping has been discussed only in a couple of studies using semi-analytic models (\citealt{okamoto:03, lanzoni:05}), and is shown to produce only mild variations in galaxy colours and star formation rates. This happens because the stripping of the hot gas from galactic haloes (strangulation) suppresses the star formation so efficiently that the effect of ram-pressure is only marginal. We note that the studies mentioned above include ram-pressure stripping based on the analytical criterion formulated originally in Gunn \& Gott (1972). In recent numerical work, \cite{roediger:07} have shown that this formulation often yields incorrect mass loss rates. Other numerical studies (e.g. \cite{vollmer01a}) have argued that ram-pressure stripping can also temporarily enhance star formation. An updated modelling of ram-pressure stripping that takes into account these results has not been included yet in semi-analytic models. We plan to address this in future studies. For a study of the ram pressure distribution, it is necessary to have information on the dynamics of galaxies and on the properties of the ICM. The dynamics of dark matter halos has been studied in a number of papers using numerical simulations (e.g. \citealt{benson:05,khochfar:05, diemand:04}). However, we know of no study on the distribution and history of ram pressures experienced by galaxies in clusters. If ram-pressure plays some role in establishing the observed morphological mix in galaxy clusters and/or the observed radial trends, it is important to quantify the distribution and history of ram pressures experienced by galaxies that reside in clusters today. This is the subject of this paper. | In this study we have taken advantage of the Millennium simulation by \cite{springel:05} and of the publicly available semi-analytic model by De Lucia \& Blaizot (2007) to reconstruct the ram-pressure distribution and ram-pressure history of galaxies that reside in clusters at the present epoch. The gas profile in dark matter is described through two analytic models which give fairly similar results. We have not included ram-pressure stripping self-consistently in the semi-analytic model employed for our study. Instead we have simply used the available information about the orbital distribution and galaxy merging trees to estimate the importance of ram-pressure stripping on galaxies that reside in massive haloes at the present epochs. We find that more than half of the galaxies in a $M=10^{15}M_{\odot}$ cluster, have experienced ram pressures of $> 10^{-11}$ dyn cm$^{-2}$ after their accretion onto a massive system. This fraction is only slightly lower for $M=10^{14}M_{\odot}$ clusters, implying that a significant fraction of galaxies in clusters at the present epoch suffered substantial gas loss due to ram pressure stripping. The fraction of galaxies that suffered significant ram-pressure after accretion increases with decreasing distance from the cluster centre, in qualitative agreement with the observed increase of early-type galaxies. As expected, strong episodes of ram-pressure occur predominantly in the inner core of galaxy clusters, and are restricted to within the virial radius. Our results show, however, that virtually all the galaxies in clusters suffered weaker episodes of ram pressure, suggesting that this physical process might have a significant role in shaping the observed properties of the entire cluster galaxy population. The limited number of simulation outputs does not allow us to reconstruct accurately the orbit of the cluster galaxies, and therefore their ram-pressure histories. Our result show that ram pressure fluctuates strongly so that episodes of strong ram-pressure alternate to episode of weaker ram-pressure. During these time intervals, the gaseous reservoir could be replenished and new episodes of star formation could occur. Our results indicate that ram-pressure stripping must play a significant role in the evolution of galaxies residing in massive clusters. A more self-consistent modelling is however required in order to draw more quantitative conclusions about the importance and effects of this physical process. | 7 | 10 | 0710.5580 |
0710 | 0710.5307_arXiv.txt | Several kinds of astronomical observations, interpreted in the framework of the standard Friedmann--Robertson--Walker cosmology, have indicated that our universe is dominated by a Cosmological Constant. The dimming of distant Type Ia supernovae suggests that the expansion rate is accelerating, as if driven by vacuum energy, and this has been indirectly substantiated through studies of angular anisotropies in the cosmic microwave background (CMB) and of spatial correlations in the large-scale structure (LSS) of galaxies. However there is no compelling {\em direct} evidence yet for (the dynamical effects of) dark energy. The precision CMB data can be equally well fitted without dark energy if the spectrum of primordial density fluctuations is not quite scale-free and if the Hubble constant is lower globally than its locally measured value. The LSS data can also be satisfactorily fitted if there is a small component of hot dark matter, as would be provided by neutrinos of mass $\sim0.5$ eV. Although such an Einstein--de Sitter model cannot explain the SNe~Ia Hubble diagram or the position of the `baryon acoustic oscillation' peak in the autocorrelation function of galaxies, it may be possible to do so e.g. in an inhomogeneous Lemaitre--Tolman--Bondi cosmology where we are located in a void which is expanding faster than the average. Such alternatives may seem contrived but this must be weighed against our lack of any fundamental understanding of the inferred tiny energy scale of the dark energy. It may well be an artifact of an oversimplified cosmological model, rather than having physical reality. | Following his formulation of general relativity Einstein \cite{Einstein:1917} boldly applied the theory to the universe as a whole. The first cosmological model was {\em static} to match the known universe, which at that time was restricted to the Milky way, and to achieve this Einstein introduced the `cosmological constant' term (for a historical perspective, see \cite{Straumann:2002he}). Within a decade however Slipher and Hubble demonstrated that the nebulae on the sky are in fact other `island universes' like the Milky Way and that they are mainly receeding from us --- the universe is expanding. Einstein wrote to Weyl in 1933: {\em ``If there is no quasi-static world, then away with the cosmological term''}. This however is not a matter of choice since general coordinate invariance, which Einstein's equation is based on, permits an arbitrary constant (multiplied by the metric tensor) to be added to the lhs: \begin{equation} R_{\mu\nu} - \frac{1}{2}g_{\mu\nu}R + \lambda_\mathrm{metric} g_{\mu\nu} = \frac{-T_{\mu\nu}}{M_\mathrm{P}^2}. \label{gr} \end{equation} Here we have written Newton's constant, $G_\mathrm{N} \equiv 1/8\pi\,M_\mathrm{P}^2$ where $M_\mathrm{P} \simeq 2.4\times10^{18}$~GeV is the (reduced) Planck mass in natural units ($\hbar = k_\mathrm{B} = c = 1$). With the subsequent development of quantum field theory it became clear that the energy-momentum tensor on the rhs can also be freely scaled by another additive constant multiplying the metric tensor, which reflects the (Lorentz invariant) energy density of the vacuum: \begin{equation} \langle T_{\mu\nu}\rangle_{\rm fields} = -\langle\rho\rangle_{\rm fields}g_{\mu\nu}. \end{equation} This contribution from the matter sector adds to the ``bare'' term from the background geometry, yielding an effective cosmological constant: \begin{equation} \Lambda = \lambda_\mathrm{metric} + \frac{\langle\rho\rangle_{\rm fields}}{M_\mathrm{P}^2}, \label{tuning} \end{equation} or, correspondingly, an effective vacuum energy: \begin{equation} \rho_\mathrm{v} \equiv \Lambda M_\mathrm{P}^2. \label{rhov} \end{equation} Einstein {\em assumed} without any observational evidence that the universe is perfectly homogeneous. We know that the universe is quite isotropic about us so this is in fact likely if we are not in a special location --- an assumption later dignified by Milne as the `Cosmological Principle'. Then using the {\em maximally symmetric} Robertson-Walker metric to describe space-time \begin{equation} \mathrm{d}s^2 \equiv g_{\mu\nu} \mathrm{d}x^\mu \mathrm{d}x^\nu = \mathrm{d}t^2 - a^2(t)[ \mathrm{d}r^2/(1 - kr^2) + r^2 \mathrm{d}\Omega^2], \label{rw} \end{equation} we obtain the Friedmann equations describing the evolution of the cosmological scale-factor $a (t)$: \begin{equation} H^2 \equiv \left(\frac{\dot{a}}{a}\right)^2 = \frac{\rho}{3M_\mathrm{P}^2} - \frac{\kappa}{a^2} + \frac{\Lambda}{3}, \qquad \frac{\ddot{a}}{a} = -\frac{1}{6M_\mathrm{P}^2}(\rho + 3p) + \frac{\Lambda}{3}, \label{fried} \end{equation} where $\kappa = 0, \pm1$ is the 3-space curvature signature and for ordinary matter (`dust') and radiation we have used the `ideal gas' form: $T_{\mu\nu} = p g_{\mu\nu} + (p + \rho) u_\mu u_\nu$, with $u_\mu \equiv \mathrm{d}x_\mu/\mathrm{d}s$. The conservation equation $T^{\mu\nu}_{;\nu} = 0$ implies $\mathrm{d}(\rho a^3)/\mathrm{d}a = -3 p a^2$, so given the `equation of state parameter' $w \equiv p/\rho$, the evolution history can now be constructed. Since the redshift is $z \equiv a/a_0 - 1$, for non-relativistic particles with $w \simeq 0$, $\rho_\mathrm{NR} \propto (1 + z)^{-3}$, while for relativistic particles with $w = 1/3$, $\rho_\mathrm{R} \propto (1 + z)^{-4}$, but for the cosmological constant, $w = -1$ and $\rho_\mathrm{v} = $constant. Thus radiation was dynamically important only in the early universe (for $z \geqsim 10^4$) and for most of the expansion history only non-relativistic matter is relevant. The Hubble equation can be rewritten with reference to the present epoch (subscript 0) as \begin{eqnarray} H^2 &=& H_0^2 \left[\Omega_\mathrm{m} (1 + z)^3 + \Omega_\kappa (1 + z)^2 + \Omega_\Lambda \right], \\ \Omega_\mathrm{m} &\equiv& \frac{\rho_{\mathrm{m}0}}{\rho_\mathrm{c}}, \quad \Omega_\kappa \equiv -\frac{\kappa}{a_0^2 H_0^2}, \quad \Omega_\Lambda \equiv \frac{\Lambda}{3 H_0^2}, \label{hub} \end{eqnarray} where $\rho_\mathrm{c} \equiv 3H_0^2M_\mathrm{P}^2/8\pi \simeq (3 \times 10^{-12} \,\mathrm{GeV})^4 h^2$ is the `critical density' and the present Hubble parameter is $H_0 \equiv 100h~\mathrm{km}~\mathrm{s}^{-1}~\mathrm{Mpc}^{-1}$ with $h \simeq 0.7$, i.e. about $10^{-42}$~GeV. This yields the sum rule \begin{equation} \Omega_\mathrm{m} + \Omega_\kappa + \Omega_\Lambda = 1, \label{sumrule} \end{equation} so cosmological models can be usefully displayed on a `cosmic triangle' \cite{Bahcall:1999xn}. As emphasised in an influential review \cite{Weinberg:1988cp}, given that the density parameters $\Omega_\mathrm{m}$ and $\Omega_\kappa$ were observationally constrained already to be not much larger than unity, the two terms in eq.(\ref{tuning}) are required to somehow {\em conspire} to cancel each other in order to satisfy the approximate constraint \begin{equation} |\Lambda| \leqsim H_0^2, \end{equation} thus bounding the present vacuum energy density by $\rho_\mathrm{v} \leqsim 10^{-47}\,\mathrm{GeV}^4$ which is a factor of over $10^{120}$ below its ``natural'' value of $\sim M_\mathrm{P}^4$ --- the `cosmological constant problem'. Subsequently, several types of evidence have been advanced to argue that this inequality is in fact saturated with $\Omega_\Lambda \simeq 0.7$ ($\Rightarrow \Lambda \simeq 2H_0^2$), $\Omega_\mathrm{m} \simeq 0.3$, $\Omega_\kappa \simeq 0$ (see \cite{Sahni:1999gb,Peebles:2002gy}), i.e. there is non-zero vacuum energy of {\em just the right order of magnitude so as to be detectable today}. In the Lagrangian of the Standard $SU(3)_\mathrm{c} \otimes SU(2)_\mathrm{L} \otimes U(1)_Y$ Model (SM) of electroweak and strong interactions, the term corresponding to the cosmological constant is one of the two `super-renormalisable' terms allowed by the gauge symmetries, the second one being the quadratic divergence in the mass of fundamental scalar fields due to radiative corrections (see \cite{Zwirner:1995iw}). To tame the latter sufficiently in order to explain the experimental success of the SM has required the introduction of a `supersymmetry' between bosonic and fermionic fields which is `softly broken' at about the Fermi scale, $M_\mathrm{EW}\sim\,G_\mathrm{F}^{-1/2}\simeq246$~GeV Thus the cutoff scale of the SM, viewed as an effective field theory, can be lowered from $M_\mathrm{P}$ down to $M_\mathrm{EW}$, at the expense of introducing over 100 new parameters in the Lagrangian, as well as requiring delicate control of the non-renormalisable operators which can generate flavour-changing neutral currents, nucleon decay etc, so as not to violate experimental bounds. This implies a {\em minimum} contribution to the vacuum energy density from quantum fluctuations of ${\cal O}(M_\mathrm{EW}^4)$, i.e. halfway on a logarithmic scale down from $M_\mathrm{P}$ to the energy scale of ${\cal O}(M_\mathrm{EW}^2/M_\mathrm{P})$ corresponding to the observationally indicated vacuum energy. Thus even the introduction of supersymmetry cannot eradicate a discrepancy by a factor of at least $10^{60}$ between the natural expectation and observation. It is likely that a satisfactory resolution of the cosmological constant problem can be achieved only in a quantum theory of gravity. Recent developments in string theory and the possibility that there exist new dimensions in Nature have generated many interesting ideas concerning possible values of $\Lambda$ (see e.g. \cite{Witten:2000zk,Padmanabhan:2002ji,Nobbenhuis:2004wn}). Nevertheless it is the case that there is no accepted solution to the enormous discrepancy discussed above. The problem is not new but there has always been the hope that some day we would understand why $\Lambda$ is exactly zero, perhaps due to a new symmetry principle. However if it is in fact non-zero and dynamically important {\em today}, the crisis is even more severe since this also raises a cosmic `coincidence' problem, viz. why is the present epoch special?~\footnote{While $\Lambda \sim H_0^2$ today, we cannot have $\Lambda \sim H^2$ {\em always} since this would amount to a substantial renormalisation of $M_\mathrm{P}$ in eq.(\ref{fried}) (taking $\kappa=0$ as suggested by inflation); this is inconsistent with primordial nucleosynthesis which requires the ``cosmic'' Newton's Constant to be within a few \% of its laboratory value \cite{Cyburt:2004yc}. Thus arguably the most natural solution to the `coincidence problem' is ruled out empirically.} It has been suggested that the `dark energy' may not be a cosmological {\em constant} but rather the slowly evolving potential energy $V (\phi)$ of a hypothetical scalar field $\phi$ named `quintessence' which can {\em track} the matter energy density (see \cite{Copeland:2006wr}). This however is equally fine-tuned since we need $V^{1/4} \sim 10^{-12}$~GeV but $\sqrt{\mathrm{d}^2 V/\mathrm{d}\phi^2} \sim H_0 \sim 10^{-42}$~GeV (in order that the evolution of $\phi$ be sufficiently slowed by the Hubble expansion), and moreover does not address the fundamental issue,\footnote{Admittedly this criticism applies also to attempts to do away with dark energy by interpreting the data in terms of modified cosmological models, but less so since these do not invoke gravitating vacuum energy at all.} namely why are all the other possible contributions to the vacuum energy absent? Given the no-go theorem against any dynamical cancellation mechanism in eq.(\ref{tuning}) in the framework of general relativity \cite{Weinberg:1988cp}, it might appear that solving the problem will necessarily require modification of our understanding of gravity. Interesting suggestions have been made in this context e.g. the DGP model in which gravity alone propagates in a new dimension that opens up on distance scales of ${\cal O}(H_0^{-1})$ (see \cite{Lue:2005ya}). However since this is an unnaturally large scale for a fundamental theory, this model is clearly just as fine tuned as the cosmological constant and moreover suffers from intrinsic theoretical difficulties such as violation of unitarity due to `ghosts' (see \cite{Koyama:2007za}). The situation is so desperate that `anthropic' arguments have been advanced to explain why the cosmological constant is of just the right order of magnitude to allow of our existence today --- if it was much higher then galaxies would never have have formed (see \cite{Weinberg:2000yb}). However a recent analysis \cite{Tegmark:2005dy} shows that the most likely value from such Bayesian arguments is in fact 20--30 times bigger than the observationally suggested value. Moreover, one cannot claim to have a rigorous understanding of the prior probability distribution, lacking a theory of the cosmological constant itself. It is commonly assumed that the prior is flat in $\Lambda$, but if it is instead flat in say $\log(\Lambda)$, the most likely value of $\Lambda$ would be zero. A flat prior in $\Lambda$ is indeed expected in quintessence-like models with a very flat potential \cite{Weinberg:2000qm}, but such models are highly fine-tuned and have little physical basis. Whereas an uniform distribution of $\Lambda$ does arise in the 'landscape' of the large number ($\sim 10^{500}$) of possible vacua in string theory (see \cite{Douglas:2006es}), there is no accepted measure for how these vacuua may come to be populated through cosmological evolution. Given this situation, we can ask whether the observations really require dark energy or whether this is just an artifact of interpreting the data using an over-idealised description of the universe. For the FRW model we see from eq.(\ref{fried}) that deducing $\Lambda$ to be of ${\cal O}(H_0^2)$ from observations should not be unexpected, this being its {\em natural} scale in such a model universe (what seems quite unnatural is the {\em implied} fundamental energy scale of $\sqrt{H_0 M_\mathrm{P}}$ since $H_0$ is so much smaller than $M_\mathrm{P}$). From this perspective, it is easy to see why there have been recurring claims for $\Lambda \sim H_0^2$ --- this is effectively forced upon us by the theoretical framework in which the data is interpreted. | We now have a `cosmic concordance' model with $\Omega_\mathrm{m} \sim 0.3, \Omega_\Lambda \sim 0.7$ which is supposedly consistent with all astronomical data but has no satisfactory explanation in terms of fundamental physics. One might hope to eventually find explanations for the dark matter (and baryonic) content of the universe in the context of physics beyond the Standard Model but there appears to be little prospect of doing so for the apparently dominant component of the universe --- the dark energy which behaves as a cosmological constant. Cosmology has in the past been a data-starved science so it has been appropriate to consider the simplest possible cosmological models in the framework of general relativity. However now that we are faced with this serious confrontation between fundamental physics and cosmology, it would seem prudent to reconsider the underlying assumptions, especially that of exact homogeneity. The observed isotropy of the CMB along with a belief in the Cosmological Principle has generally been taken to imply homogeneity but this has come to be questioned of late. There is growing interest in the inhomogeneous but isotropic Lemaitre-Tolman-Bondi model of a local void (see \cite{Krasinski:1997}). It has been shown \cite{Tomita:1999qn,Tomita:2000rf,Tomita:2000jj,Tomita:2001gh,Tomita:2002df} that a LTB model can in principle give an explanation of the oddities in the local Hubble flow as well as account for the SNe~Ia Hubble diagram without invoking acceleration (see also \cite{Celerier:1999hp,Alnes:2005rw,Enqvist:2006cg,Alnes:2006uk}). Recently it has been claimed \cite{Biswas:2006ub} that with a large local void of size $\sim 450h^{-1}$~Mpc the angular diameter distance of the BAO peak can also be matched in this model. Such local inhomogeneity may be responsible for generating the mysteriously aligned low multipoles in the {\sl WMAP} sky \cite{Inoue:2006rd}. A void of this size does seem extremely unlikely, but one has recently been actually detected at $z \sim 1$ and appears to be responsible for the anomalously `cold spot' seen by {\sl WMAP} \cite{Rudnick:2007kw}. Landau famously said {\em ``Cosmologists are often wrong, but never in doubt''}. The situation today is perhaps better captured by Pauli's enigmatic remark --- the present interpretation of the data may be {\it ``\ldots not even wrong''}. However we are certainly not without doubt! The crisis posed by the recent astronomical observations is not one that confronts cosmology alone; it is the spectre that haunts any attempt to unite two of the most successful creations of 20th century physics --- quantum field theory and general relativity. It would be fitting if the cosmological constant which Einstein allegedly called his {\it ``biggest blunder''} proves to be the catalyst for triggering a new revolution in physics in this century. | 7 | 10 | 0710.5307 |
0710 | 0710.0378_arXiv.txt | Various mechanisms have been proposed to explain the inflated size of HD~209458b after it became clear that it has no companions capable of producing a stellar reflex velocity greater than around $5\,\ms$. Had there been such a companion, the hypothesis that it forces the eccentricity of the inflated planet thereby tidally heating it may have been readily accepted. Here we summarize a paper by the author which shows that companion planets with masses as low as a fraction of an Earth mass are capable of sustaining a non-zero eccentricity in the observed planet for at least the age of the system. While such companions produce stellar reflex velocities which are fractions of a meter per second and hence are below the stellar jitter limit, they are consistent with recent theoretical work which suggests that the planet migration process often produces low-mass companions to short-period giants. | Of the fourteen transiting extrasolar planetary systems for which radii have been measured, at least three appear to be considerably larger than theoretical estimates suggest. It has been proposed by \citet{b4} that undetected companions acting to excite the orbital eccentricity are responsible for these oversized planets, as they find new equilibrium radii in response to being tidally heated. In the case of HD 209458, this hypothesis has been rejected by some authors because there is no sign of such a companion at the 5 ms$^{-1}$ level, and because it is difficult to say conclusively that the eccentricity is non-zero. Transit timing analysis as well as a direct transit search have further constrained the existence of very short-period companions, especially in resonant orbits. Whether or not a companion is responsible for the large radius of HD 209458b, almost certainly some short-period systems have companions which force their eccentricities to nonzero values. This paper is a summary of \citet{m1} which is dedicated to quantifying this effect. The eccentricity of a short-period planet will only be excited as long as its (non-resonant) companion's eccentricity is non-zero. In fact, \citet{m1} shows that the latter decays on a timescale which depends on the structure of the {\it interior} planet, a timescale which is often shorter than the lifetime of the system. {\it This includes Earth-mass planets in the habitable zones of some stars}. On the other hand, there exist parameter combinations ($Q$-value of the innermost planet, companion mass and semimajor axis) for which significant eccentricity in the short-period planet can be sustained for at least the age of the system, and these include systems with companion masses as low as a fraction of an Earth mass. | 7 | 10 | 0710.0378 |
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0710 | 0710.2461_arXiv.txt | We present a measure of the hard (2--8 keV) X-ray luminosity function (XLF) of Active Galactic Nuclei up to $z\sim5$. At high redshifts, the wide area coverage of the {\em Chandra} Multiwavength Project is crucial to detect rare and luminous ($L_X>10^{44}$ erg s$^{-1}$) AGN. The inclusion of samples from deeper published surveys, such as the {\em Chandra} Deep Fields, allows us to span the lower $L_X$ range of the XLF. Our sample is selected from both the hard ($z<3$; $f_{2-8~{\rm kev}}>6.3\times10^{-16}$ erg cm$^{-2}$ s$^{-1}$) and soft ($z>3$; $f_{0.5-2.0~{\rm kev}}>1.0\times10^{-16}$ erg cm$^{-2}$ s$^{-1}$) energy band detections. Within our optical magnitude limits ($r^{\prime},i^{\prime}<24$), we achieve an adequate level of completeness ($>50\%$) regarding X-ray source identification (i.e., redshift). We find that the luminosity function is similar to that found in previous X-ray surveys up to $z\sim3$ with an evolution dependent upon both luminosity and redshift. At $z>3$, there is a significant decline in the numbers of AGN with an evolution rate similar to that found by studies of optically-selected QSOs. Based on our XLF, we assess the resolved fraction of the Cosmic X-ray Background, the cumulative mass density of Supermassive Black Holes (SMBHs), and the comparison of the mean accretion rate onto SMBHs and the star formation history of galaxies as a function of redshift. A coevolution scenario up to $z\sim2$ is plausible though at higher redshifts the accretion rate onto SMBHs drops more rapidly. Finally, we highlight the need for better statistics of high redshift AGN at $z\gtrsim3$, which is achievable with the upcoming {Chandra} surveys. | Our present understanding of the evolution of accreting SMBHs over cosmic time comes from the measure of the luminosity function (i.e., the number undergoing a detectably luminous phase within a specific co-moving volume as a function of luminosity and redshift) of Active Galactic Nuclei (AGN). Energy production through mass accretion onto SMBHs allows us to observationally identify these sites as the familiar AGN with Quasi-stellar objects (QSOs) the most luminous example. Hence, the AGN luminosity function provides a key constraint to discern the underlying physical properties of the population (i.e., black hole mass and accretion rate distributions as a function of redshift) and thereby elucidate the mechanisms (i.e., galaxy mergers and/or self-regulated growth) that are instrumental in their formation and evolution. To date, an enormous effort has been undertaken to measure the luminosity function over the wide range in luminosity spanned by AGN out to high redshift. The bright end has been well established to $z\sim5$ by optical surveys \citep[e.g.,][]{ri06,cr04,wo03} which primarily select QSOs using a multi-color photometric criteria. The most dramatic feature found is the rise and fall of the co-moving space density with peak activity at $z\sim2.5$. With an unprecedented sample of over 20,000 QSOs in the 2dF QSO Redshift Survey (2QZ), \citet{cr04} convincingly show a systematic decrease in luminosity (pure luminosity evolution; PLE) from $z=2$ to the present, in agreement with past surveys \citep[e.g.,][]{sc83,bsp88,hfc93}, which find very few bright QSOs in the local universe. This fading of the luminous QSO population is thought to be due to a decrease in the mass accretion rate \citep[e.g.,][]{ca00} that appears to be intimately related to the order-of-magnitude decline of the cosmic star formation rate from $z\sim1$ to the present \citep{bo98,fr99,me04b}. The dropoff in the space density at $z>3$ \citep{os82,wa94,sc95,fa01,wo03} may be indicative of either the detection of the onset of accretion onto young SMBHs or a high-redshift population that has been missed, possibly under a veil of obscuration \citep[e.g.,][]{al05,ma05}. Excessive amounts of dust and gas may be ubiquitous in galaxies at early epochs due an increase in the merger rate \citep{ka07} that induces high star formation rates \citep[e.g.,][]{ch05}. It has been evident for quite some time that optical surveys miss a significant fraction of the AGN population. They fail to find the majority of AGN due to a steeply declining luminosity function with the low luminosity end severely affected by host-galaxy dilution. Though current optical selection techniques do show considerable improvement \citep{ri05,ji06}, they still fail to account for many low luminosity AGN. Of equal significance, many AGN (e.g., Seyfert 2s, narrow line radio galaxies) are missed due to dust obscuration (causing their optical properties to differ from the type 1 QSOs) and can only be adequately selected in the low redshift universe \citep{hao05} based on their highly ionized, narrow emission lines. Fortunately, AGN can now be efficiently accounted for by selection techniques in other wave bands such as the X-ray as demonstrated in this work, radio \citep[e.g.,][]{wa05} and infrared \citep[e.g.,][]{po06}. As further elaborated below, current models based on recent observations continue to attribute the bulk of the Cosmic X-ray Background (CXRB), the previously unresolved X-ray emission, to these various types of obscured AGN. Over more than two decades, X-ray surveys have been improving and extending the known AGN luminosity function by including sources at low luminosity, with or without optical emission lines, and hidden by a dense obscuring medium. The Extended Medium Sensitivity Survey \citep{gi90} was one of the first surveys to measure the X-ray luminosity function (XLF thereafter) out to the QSO peak using a sample of just over 400 AGN detected by the $EINSTEIN$ Observatory. Since the survey only probed the more luminous AGN ($L_X>10^{44}$ erg s$^ {-1}$) above $z=0.3$ due to the bright flux limit ($f_{0.3-3.5~{\rm keV}}>5\times10^{-14}$ erg cm$^{-2}$ s$^{-1}$), it was quite understandable that \citet{ma91} and \citet{de92} found the XLF to behave similarly to the optical luminosity function (i.e., PLE) with a decreasing space density from $z\sim2$ to the present. The $ROSAT$ satellite with its increase in flux sensitivity ($f_{0.5-2.0~\rm{keV}}> 10^{-15}$ erg cm$^{-2}$ s$^{-1}$) enabled AGN to be detected at lower luminosities, and out to higher redshifts ($z\sim4.5$). Surveys ranged from the wide area and shallow $ROSAT$ Bright Survey \citep[f$_{lim}\sim10^{-12}$ erg cm$^{-2}$ s$^{-1}$; 20000 deg$^{2}$;][]{sc00} to the deep and narrow Lockman Hole \citep[f$_{lim}\sim10^{-15}$ erg cm$^{-2}$ s$^{-1}$; 0.3 deg$^{2}$;][]{le01}. With a compilation of 690 AGN from these fields and other available $ROSAT$ surveys \citep{bow96,ap98,mc98,za98,ma00} that effectively fill in the parameter space of flux and area, \citet{mi00} were able to resolve $\approx$ 60-90\% of the soft CXRB into point sources and generate a soft XLF that extended beyond the QSO peak. They found that the XLF departed from a simple PLE model, now well known to describe X-ray selected samples as further elaborated below. A $luminosity$-dependent $density$ evolution (LDDE thereafter) model was required due to the slower evolution rate of lower luminosity AGN compared to that of the bright QSOs. The limited sky coverage at the faintest X-ray fluxes achievable with $ROSAT$ prevented an accurate measure of both the faint-end slope at $z\gtrsim1$ and the overall XLF at high redshift since only a handful of AGN were identified above a redshift of 3. In the current era of {\em Chandra} and {\em XMM-Newton}, X-ray surveys are now able to detect AGN and QSOs not only enshrouded by heavy obscuration but those at high redshift ($z>3$) with greatly improved statistics due to the superb spatial resolution and sensitivity between 0.5 to 10 keV of these observatories. Previous X-ray missions such as $EINSTEIN$ and $ROSAT$ as described above were limited to the soft band which biases samples against absorption. The $ASCA$ observatory \citep[e.g.,][]{ak03} successfully found many nearby absorbed AGN but lacked the sensitivity to detect the fainter sources contributing most of the 2--8 keV CXRB. With deep observations of the {\em Chandra} Deep Field North \citep[CDF-N;][]{br01}, Deep Field South \citep[CDF-S;][]{ro02} and Lockman Hole \citep{ha01}, a large fraction, between $\sim70\%$ \citep{wo05} and $\sim89\%$ \citep{mo03} of the hard (2--8 keV) CXRB has been resolved into point sources. Many of the hard X-ray sources found so far arise in optically unremarkable bright galaxies \citep[e.g.,][]{ba03b,to01}, which can contain heavily obscured AGN. A more robust luminosity-dependent evolutionary scheme has emerged from recent measures of the XLF. With the inclusion of absorbed AGN from {\em Chandra} and {\em XMM-Newton} surveys, lower luminosity AGN are clearly more prevalent at lower redshifts ($z<1$) than those of high luminosity that peak at $z\sim2.5$. This behavior is due to the flattening of the low luminosity slope at higher redshifts that has been well substantiated with hard (2--8 keV) X-ray selected surveys \citep{cow03,ue03,fi03,ba05,si05b,laf05}. Using a highly complete soft (0.5--2.0 keV) band selected sample of over 1000 type 1 AGN, \citet{ha05} has shown that this LDDE model accurately fits the data and shows a gradual shift of the peak in the co-moving space density to lower redshifts with declining luminosity. This behavior may be evidence for the growth of lower mass black holes emerging in an ``anti-hierarchical'' or ``cosmic downsizing'' fashion while accreting near their Eddington limit \citep{ma04,me04a}, or the embers of a fully matured SMBH population with sub-Eddington accretion rates. The former scenario has been substantiated by \citet{ba05} based on the optical luminosities of the galaxies hosting X-ray selected AGN, and \citet{he04} using type 2 AGN from the SDSS with the [OIII] emission line luminosity as a proxy for the mass accretion rate and an estimate of the black hole mass from the $M-\sigma$ relation. Even though there has been much progress, there are remaining uncertainties in the current measure of the XLF. (1) A significant number of X-ray sources in the recent surveys with {\em Chandra} and {\em XMM-Newton} are not identified. \citet{main05} find that the peak of the redshift distribution shifts to higher values ($z\sim1.2-1.5$) when incorporating photometric redshifts for objects too faint for optical spectroscopy. (2) \citet{ba05} demonstrate that the XLF can be fit equally well by a PLE model at $z<1.2$. These models only begin to substantially differ for low luminosity AGN at $z>1.5$ where AGN statistics are quite low with most being provided by the deep CDF-N and CDF-S observations. New moderate depth surveys such as the Extended {\em Chandra} Deep Field-South \citep[E-CDF-S;][]{le05} and the Extended Groth Strip \citep[EGS;][]{na05} will provide additional AGN at these luminosities and redshifts but await optical followup. (3) How does the AGN population behave at redshifts above the peak ($z\sim2.5$) of the optically-selected QSO population? We have presented evidence \citep{si04,si05b} for a similar evolution of luminous X-ray selected QSOs to those found in the optical surveys \citep[e.g.,][]{ri06} with a decline in the co-moving space density at $z>3$ but these AGN are mainly type 1. We don't expect the inclusion of luminous absorbed QSOs to drastically alter our measure of the XLF since they may be at most as numerous as the type 1 QSOs \citep{gi07}. Recent radio \citep{wa05} and near infrared \citep{br06} surveys are further supporting a strong decline in the co-moving density at high redshift. In the present study, we measure the XLF of AGN in the hard X-ray (intrinsic 2--8 keV) band with an emphasis on reducing uncertainties at high redshift ($3<z<5$). This paper is an extension to the preliminary results on the co-moving space density of AGN as reported by the ChaMP survey \citep{si04,si05b}. As previously described, these early epochs represent the emergence of the luminous QSOs and the rapid growth phase of young SMBHs. To date, the limited numbers of X-ray selected AGN at $z>3$ have constrained current measures \citep{laf05,ba05,ue03} to lower redshifts. Motivated by \citet{ba05}, we use the observed soft X-ray band for AGN selection above $z=3$ where we measure the rest-frame energies above 2 keV. Due to the rarity of luminous high redshift AGN, such an endeavor requires a compilation of surveys that covers a wide enough area at sufficient depths. As previously mentioned, a large dynamic range from the deep, narrow pencil beam to the wide, shallow surveys is required to measure a luminosity function that spans low and high luminosities at a range of redshifts. Currently, there are a handful of deep surveys with {\em Chandra} (i.e., CDF-N, CDF-S) and {\em XMM-Newton} (Lockman Hole) that have published catalogs with a fair sample of low luminosity ($42<log~L_X<44$) AGN out to $z\sim5$. To provide a significant sample of more luminous AGN ($log~L_X>44$), the {\em Chandra} Multiwavelength Project (ChaMP) is carrying out a wider area survey of archived {\em Chandra} fields and the CLASXS \citep{ya04,st04} survey is imaging a contiguous area with nine {\em Chandra} pointings. The statistics of high redshift AGN are sure to improve with the anticipated results from the SWIRE/{\em Chandra} (Wilkes et al., in preparation), XBootes \citep{mu05}, E-CDF-S \citep{le05}, EGS \citep{na05}, XMM/COSMOS \citep{ha07}, and the newly approved {\em Chandra}/COSMOS surveys. We organize the paper as follows: Section~\ref{compile}, we describe the various surveys used in this analysis including X-ray sensitivity, sky area coverage, incompleteness as a function of not only X-ray flux but optical magnitude, and AGN selection. Our method for measuring the luminosity function is presented in Section~\ref{methods} and the results, including best-fit analytic models, in Section~\ref{results} for all AGN types. In Section~\ref{text:cxrb}, we address the resolved fraction of the CXRB and any underrepresented source populations. We directly compare our luminosity function to that of optically-selected samples in Section~\ref{text:compare_opt}. Based on our luminosity function, we derive in Section~\ref{history} the accretion rate distribution as a function of redshift and the cumulative mass density of SMBHs. Section~\ref{coevolution}, we compare the global mass accretion rate of SMBHs to the star formation history of galaxies out to $z\sim5$. We end in Section~\ref{predict} with some predictions of the numbers of high redshift ($z>3$) AGN expected in new surveys that effectively enable us to extend the luminosity function to these high redshifts with accuracy. Throughout this work, we assume H$_{\circ}=70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\Lambda}=0.7$, and $\Omega_{\rm{M}}=0.3$. | We have presented an extension of the hard (2--8 keV) X-ray luminosity function of AGN up to $z\sim5$. The ChaMP effectively covers a wide area (1.8 deg$^{2}$) at sufficient depths \hbox{($f_{0.5-2.0~{\rm keV}}\sim10^{-15}$ erg cm$^{-2}$ s$^{-1}$)} to significantly improve the statistics of luminous ($L_{\rm X}>10^{44}$ erg s$^{-1}$) AGN at high redshift. In total, we have amassed a sample of 682 AGN with 31 at $z>3$. The addition of lower luminosity AGN from the {\em Chandra} Deep Fields is instrumental to characterize the faint end slope at $z\gtrsim1$. We have corrected for incompleteness (i.e., fraction of sources without a redshift) as both a function of X-ray flux and optical magnitude, as dictated by the limitations of optical spectroscopic followup. Significant optical followup is still required to accurately account for the obscured population, especially at high redshifts. We have measured the hard XLF with both a binned ($1/V_a$) and an unbinned method that fits an analytic model to our data using a maximum likelihood technique. We found that the luminosity function is similar to that found in previous studies \citep{ue03,laf05,ba05} up to $z=3$, with an evolution dependent upon luminosity. At higher redshifts, there is a significant decline in the numbers of AGN with an evolution rate similar to that found with optically-selected QSO samples. We further show that the strong evolution above the cutoff redshift may cause the LDDE model to underpredict the number of low luminosity AGN at $z>2$. A PLE model with a faint-end slope dependent on redshift agrees better with the binned ($1/V_a$) data at $z\sim2.5$ though it may overpredict the number of faint AGN at higher redshifts. We highlight the need to improve the statistics of both high redshift AGN at $z>4$ and lower luminosity (i.e., below the break in the luminosity function) AGN at $z>3$. Such improvements are feasible from our predictions of the numbers AGN that will be found in the latest {\em Chandra} surveys (E-CDF-S, EGS and COSMOS) Our new luminosity function accounts for $\sim52\%$ of the 2--8 keV Cosmic X-ray Background. The integrated emission from these AGN give a $z=0$ mass density of SMBH of $1.64\times10^5$ \Msun~Mpc$^{-3}$, lower than other published values using X-ray selected AGN samples and the local value measured using a galaxy luminosity function and a bulge-velocity relation, possibly due to unaccounted optically-faint AGN ($r^{\prime},i^{\prime}>24$), and Compton-thick AGN. Further, a comparison of the mean mass accretion rate of SMBHs to the star formation history of galaxies out to $z\sim5$ shows the familiar co-evolution scheme up to $z\sim2$ and a divergence at higher redshifts with perhaps star formation preceeding the formation of SMBHs. | 7 | 10 | 0710.2461 |
0710 | 0710.1181_arXiv.txt | {We study the evolution and fate of solar composition super-massive stars in the mass range 60 --- 1000 \ms. Our study is relevant for very massive objects observed in young stellar complexes as well as for super-massive stars that could potentially form through runaway stellar collisions.} {We predict the outcomes of stellar evolution by employing a mass-loss prescription that is consistent with the observed Hertzsprung-Russell Diagram location of the most massive stars.} {We compute a series of stellar models with an appropriately modified version of the Eggleton evolutionary code.} {We find that super-massive stars with initial masses up to 1000 \ms\ end their lives as objects less massive than $\simeq 150$\,\ms. These objects are expected to collapse into black holes (with $M \la 70 \ms$) or explode as pair-instability supernovae.} {We argue that if ultraluminous X-ray sources (ULXs) contain intermediate-mass black holes, these are unlikely to be the result of runaway stellar collisions in the cores of young clusters.} | \label{sec:intro} In this paper, we study the structure and evolution of very massive stars (VMS), defined as objects with masses of 60 up to 150\ms, as well as super-massive stars (SMS) with masses in the range 150 - 1000 \msun. The interest in the upper limit of stellar masses and the evolution and fate of the most massive stars in the Universe was greatly boosted by the discovery of ultraluminous X-ray sources \citep[ULX, see][ and references therein]{fabbiano89,colbert04,fabbiano03,fabbiano_x_sources06,soria_recipes06}. These objects are most commonly interpreted as binaries involving either sub-Eddington accretion onto an \textit{intermediate mass black hole} (IMBH) with mass $\sim (10^2 - 10^5)$\,\ms\ or super-Eddington accretion onto a \textit{stellar mass black hole} with mass $\sim10$\,\ms. In the latter case, beaming or support of super-Eddington luminosity by an accretion disk is required. Currently, the issue of the black hole masses in ULXs remains unresolved \citep[see, e.g.,][]{fabbiano_x_sources06}. It has been argued that IMBHs may also reside in the cores of some globular clusters \citep[see, e.g.,][]{gerssen02,gerssen03,gebhardt05,Patruno:2006bw,feng_ulx_imbh06}, but see \citet{baumgardt_m15_2003,baumgardt_g1_2003} for counter-arguments. The masses of the putative black holes in the cores of well-known globular clusters such as 47~Tuc and NGC~6397, are expected to be $\sim (10^2 - 10^3)$\,\ms\ \citep{DeRijcke:2006tu}. ULXs are observed in both spiral and elliptical galaxies, i.e. in environments with diverse metal abundances and star formation rates. In this paper, we focus on stars with solar initial composition ($X$=0.7,$Z$=0.02). This choice is partly motivated by the high metallicity of the starburst galaxy M82 -- with $Z\simeq Z_{\odot}$ \citep{mcleod_m82_solar93,origlia_m82_chem04,mayya_m82_chem06} -- which contains one of the most promising candidate ULXs, M82~X-1, and that may host a black hole of intermediate mass, as argued by e.g. \citet{ptak_m82_imbh99,kaaret_m82_imbh01,matsumoto_m82_ulx01,strohmayer_m82_imbh03}, but see \citet{okajima_noimbh_m82} for arguments in favour of a stellar mass black hole in M82~X-1. The current observational estimate of the upper cut-off of stellar masses is $\sim 150$\,\ms\ \citep{massey_hunter98,weidner_kroupa_upper_m04,figer_upper05}. However, for solar chemical composition, it is expected that even such massive stars rarely produce black holes more massive than $\simeq 20$\,\ms\ \citep{maeder92,fryer_kal_bh01}, due to copious mass loss during both the hydrogen and helium burning stages. If ULXs and young globular clusters really harbour black holes with masses exceeding several 10\,\ms, the problem of their formation becomes a challenge for the theory of stellar evolution with mass loss. Related intriguing problems involve the formation and evolution of the luminous stars found in the central parsec of the Galaxy, e.g., the S-stars and the IRS13 and IRS16 conglomerates \citep{schodel_irs13,lu_irc16,maillard_13e}\footnote{Note that a high stellar mass inferred from luminosity may be an artifact of unresolved binarity. For instance, IRS16 SW turned out to be a binary with two almost identical components of $\sim 50$\,\ms, \citep{depoy_irs16,martins_irs16sw_bin06,peeples_irc16sw_bin06}.}, or in the R136 stellar complex of the 30 Doradus nebula in the Large Magellanic Cloud \citep{walborn97,massey_hunter98}. Returning to the problem of IMBH formation, there are currently two models in favour. One involves the gradual accumulation of mass by accretion onto a seed black hole, which, while swallowing gas and stars in a stellar cluster, may grow to the intermediate mass range \citep[see, e.g., ][ and references therein]{miller_imbh_form04}. The other scenario considers the collapse of an object which descends from an SMS formed by hierarchical runaway merger of ordinary stars in a young dense stellar cluster, during or after core collapse \citep[see, e.g., ][]{pmmh99,pm_imbh_clust02,mbphm04,gfr04,pbhmm_clust04,fgr06}. The motivation for the current study stems from the latter scenario. It assumes that stars are born with a wide distribution of masses ranging from the hydrogen-burning limit up to a maximum of about 150\,\msun. More massive stars observed in the Galaxy may originate from stellar mergers. Most of these may be the result of coalescence of components of binaries in common envelopes. The mass of binary merger products may be up to 300\msun. \citet{pmmh99,miller_hamilton02,pbhmm_clust04,gfr04} have demonstrated by means of detailed N-body simulations, in which simple stellar evolution was taken into account, that there is a range of initial conditions where stellar coagulation drives the mass of a star to $\sim1000$\,\msun. The resulting star burns-up quite quickly, and the process of hierarchical merging in the cluster core terminates as soon as the first massive stars experience supernovae and collapse into black holes. The collision runaway process terminates as the mass loss from the explosions of massive stars in the cluster center drives the expansion of the cluster core. This scenario was used to explain the large black hole mass of M82 X-1 \citep{pbhmm_clust04}. The studies of hierarchical mergers take into account the possibilities of collisional mass loss, mass loss from stellar winds, and rejuvenation of merger products due to fresh hydrogen supply in collisions. However, stellar evolution is treated rather crudely in these simulations, using extrapolations by several orders of magnitude for stars that typically have zero-age main-sequence (ZAMS) mass $\leq 100$\,\ms. The possible formation of core-halo configurations and the existence of an upper stellar mass limit are generally ignored. The treatment of mass loss in stellar winds is particularly uncertain. In principle, the winds of very massive and super-massive stars may be so strong that the cluster core expands in such a dramatic way that it stops being dominated by collisions, in which case the hierarchical merger is terminated \citep{pmmh99,vanbev_dyn05}. We aim to refine evolutionary calculations for merger products, and as a first step, we study the evolution of solar composition VMS and SMS over the mass range 60-1000\,$\ms$. We construct chemically homogeneous models for these stars and confirm the existence of their upper mass limit ($\simeq 1000$\,\ms), and study their evolution with mass loss through the core hydrogen and core helium-burning stages. This allows us to determine the mass and nature of the pre-supernova objects and to predict the fate of these objects. The results of our evolutionary computations are presented and discussed in Sect.~\ref{sec:evol}, a comparison with observations is given in Sect.~\ref{sec:obs}, the fate of SMS is discussed in Sect.~\ref{sec:fate}. Conclusions of our study follow in Sect.~\ref{sec:concl}. In the Appendix, we discuss a number of individual stars in close proximity to the Humphreys-Davidson (HD) limit. | \label{sec:concl} In this paper we discuss the evolution of stars with initial masses in the range 60\,\msun\, to 1001\,\msun. Our study was motivated by the results of the direct $N$-body simulations of dense star clusters by Portegies Zwart et al. (1999), in which a star grows by repeated collisions to well beyond 100\,\msun. In later studies, Portegies Zwart et al. (2004) proposed that the collision runaway in star clusters could explain the presence of a black hole of $\apgt 600$\,\msun\, in the star ULX M82~X-1 in cluster MGG11, which supposedly could have formed from a $\apgt 1000$\,\msun\, star. According to these models, the mass of the VMS increases over the stellar lifetime, starting as a homogeneous massive $\sim 100$\,\msun\, star that grows to $\apgt 1000$\,\msun\, within the core-hydrogen burning stage of evolution. We have assumed that our stellar evolution calculations are representative for stars that grow in mass via the collision runaway process. This is not necessarily correct, as the hydrogen reservoir of the collision runaway product will continuously be replenished by repeated collisions, whereas in our simulations we start with a high-mass homogeneous model. Furthermore, rotation may play a relevant role in the evolution of these objects, which was also ignored in our study. Having noted this, we may nonetheless expect these massive objects to be subject to heavy mass loss, and we may provide meaningful results in terms of the fate of the objects under consideration. We have confirmed previous results on the existence of an upper mass limit for chemically homogeneous stars, which is reached when the luminosity in the outermost layers of the stars approaches the Eddington luminosity, at which point gravity is unable to balance radiation pressure. For non-rotating solar composition stars, this limit is reached at about 1000\,\ms. We evolved the stars of 60\,\ms\ to 1001\,\ms\ to the end of the core helium-burning stage and calibrated the stellar mass-loss prescription adopted by enforcing the condition that no star should spend more than a few percent of its life above the HD-limit. We found that mass-loss rates and HR-diagram positions of the earliest O-type stars, and stars classified as transitional between O- and WR-stars, may be consistent with our computed tracks. Based on our mass-loss prescription and stellar evolution calculations, we argue that the observed massive stars in the Galaxy and the Magellanic Clouds had birth masses of up to $\simeq200$\,\ms. For stars in the lower part of this mass-range, mass loss during the MS-stage leads to the exposure of layers moderately enriched in He, and they may be observed as late-type WN stars with hydrogen features in their spectra. In the upper part of this range, stars become hydrogen-deficient WR stars already on the main-sequence. Recently, \citet{belkus07} published a study of the evolution of stars with masses up to 1000\,\ms. Our study differs from that of Belkus et al. in two important aspects: (i) they provided a simple evolutionary recipe based on similarity theory, assuming that stars are homogeneous throughout the entire course of their evolution due to vigorous convection and stay in thermal equilibrium, whereas we present detailed stellar structure models; (ii) Belkus et al. used extrapolated mass-loss rates as predicted from radiation-driven wind theory, whilst we employed a mass-loss prescription that is consistent with the location of the most massive stars in the Hertzsprung-Russell diagram. In both studies, it is found that SMS are subject to dramatic mass loss that probably inhibits the formation of IMBH by the runaway stellar collision scenario. Star clusters that experience core collapse before the most massive stars have left the main-sequence can develop a supermassive star via collision runaway. The mass of such an object is accumulated in subsequent collisions on a time scale of less than 3\,Myr. The mass which can be grown in this time interval can be estimated using Eq.~(2) of \citet{spz_etal06_imbh}\footnote{These results were obtained using standard models for runaway mass accumulation without accounting for the possibility of vigorous mass loss as considered in the present paper.}: \begin{equation} m_{\rm r} \simeq 0.01 m \left( 1 + {t_{\rm rl} \over 100\,{\rm Myr}} \right)^{-1/2}, \end{equation} where $m_{\rm m}$ is the mass of the object formed by runaway collision, $m$ is the system mass, and $t_{\rm rl}$ is the system relaxation time. The average mass increase per collision is about $\sim 20$\,\msun\ \citep{pmmh99} (for a more complete discussion, however, see \citet{pm_imbh_clust02}). The supermassive star that accumulates $\sim 1000$\msun\, has experienced some 45 collisions between the moment of gravothermal collapse of the cluster and the moment that the supermassive star dies. The mean time between collisions for this model is $\aplt 6.0 \times 10^4$\,years, resulting in a mass accretion rate of $\apgt 3\times 10^{-4}$\,\msun/yr. This rate of mass accretion is lower than the rate of mass loss that we get for the most massive stars ($\sim$ $3.8\times 10^{-3}$\,\msun/yr), but since these numbers of mass-accretion and mass-loss rate are quite uncertain, it is not inconceivable that, in spite of the strong stellar winds, the objects may nonetheless experience a net gain in mass during their lifetimes. This conclusion was drawn by \citet{suzuki07} who recently studied stellar evolution with mass loss of collisionally merged massive stars. They concluded that stellar winds would not inhibit the formation of SMS. We note, however, that the Suzuki et al. (2007) calculations were limited to masses up to $\sim$ 100 \ms\ and their calculations did not consider the more important effect of mass loss for the final mass growth process, where the mass is supposed to increase beyond this value of 100 \ms\ by an order of magnitude. In this study, we have found that a super-massive star is likely to shed most of its mass well before it experiences a supernova explosion. In several cases the supernova progenitor still had $\apgt 100$\,\msun, and such objects could possibly collapse to black holes of intermediate mass but with $M_{bh} << 1000$\,\ms. Our calculations suggest the possibility of the formation of objects that experience pair-instability supernovae, which would be interesting phenomena to observe. The recently discovered, very luminous supernova 2006gy may have been a PISN \citep{smith_07,ofek_07,langer_07}, although this is still under debate. Obviously there is the possible caveat that the quantitative difference between starting as a homogeneous high-mass star (as we considered in our paper), or growing to one over the main-sequence lifetime by repeated collisions may be significant, but at this stage we cannot provide further insights about the consequences of these potential differences. We therefore draw our conclusions on the calculations at hand, and we argue that the majority of supermassive stars probably end up as black holes of $M \lesssim 70$\,\ms, with the possible exception of some of them exploding as pair-instability supernovae. With the approximate nature of our applied mass-loss algorithm in mind, and also taking into account the (also approximate) results of \citet{belkus07}, we infer that the accuracy on the limit for black hole masses is $\sim$ $\pm 30$\,\ms. We therefore conclude that most supermassive stars end their lives as $70\pm 30 $\,\msun\, black holes. | 7 | 10 | 0710.1181 |
0710 | 0710.1148_arXiv.txt | We present detailed geometry and kinematics of the inner outflow toward HL Tau observed using Near Infrared Integral Field Spectograph (NIFS) at the Gemini-North 8-m Observatory. We analyzed H$_2$ 2.122 $\micron$ emission and [\ion{Fe}{2}] 1.644 $\micron$ line emission as well as the adjacent continuum observed at a $<$0\farcs2 resolution. The H$_2$ emission shows (1) a bubble-like geometry to the northeast of the star, as briefly reported in the previous paper, and (2) faint emission in the southwest counterflow, which has been revealed through careful analysis. The emission on both sides of the star show an arc 1\farcs0 away from the star, exhibiting a bipolar symmetry. Different brightness and morphologies in the northeast and southwest flows are attributed to absorption and obscuration of the latter by a flattened envelope and a circumstellar disk. The H$_2$ emission shows a remarkably different morphology from the collimated jet seen in [\ion{Fe}{2}] emission. The positions of some features coincide with scattering continuum, indicating that these are associated with cavities in the dusty envelope. Such properties are similar to millimeter CO outflows, although the spatial scale of the H$_2$ outflow in our image ($\sim$150 AU) is strikingly smaller than the mm outflows, which often extend over 1000--10000 AU scales. The position-velocity diagram of the H$_2$ and [\ion{Fe}{2}] emission do not show any evidence for kinematic interaction between these flows. All results described above support the scenario that the jet is surrounded by an unseen wide-angled wind, which interacts with the ambient gas and produce the bipolar cavity and shocked H$_2$ emission. | Jets or outflows are associated with young stellar objects (YSOs) at a variety of masses and evolutionary stages. There is growing evidence that these are powered by mass accretion, although the details of the flow launching mechanism are not yet clear. Understanding this issue is hampered by the angular resolution of the telescopes and interferometers, which is not yet high enough to resolve the central engine of the driving source. Instead, investigating flow geometry and kinematics close to the driving source provides useful information for understanding the relation between stellar mass accretion and outflows. In particular, recent observations of near-infrared H$_2$ emission have provided new clues to understand the activities of mass ejection close to ($<$300 AU) low-mass protostars. Pioneering work includes echelle spectroscopy and Fabry-Perot imaging of several Class I protostars by Davis et. al. (2001, 2002). Davis et al. (2001) revealed the presence of two blueshifted components at high (50--150 km s$^{-1}$) and low velocities (5--20 km s$^{-1}$). Such a kinematic structure in H$_2$ emission is similar to the forbidden emission line regions close to T Tauri stars (see e.g., Hartigan et al. 1995). While the high-velocity component is clearly associated with a collimated jet, the nature of the low velocity component is less clear. This low velocity component could be either a magnetohydrodynamically driven gas (e.g., Pyo et al. 2003, 2004; Takami et al. 2004) or a component entrained by an unseen wide-angled wind (e.g., Pyo et al. 2006; Takami et al. 2006). Davis et al. (2002) revealed H$_2$ emission associated with a collimated jet and cavity walls toward some YSOs. Such emission is excited by shocks at a temperature of $\sim$2000 $K$ (Takami et al. 2004, 2006; Beck et al. 2007). Beck et al. (2007, hereafter Paper I) present the morphology and excitation of the near-infrared H$_2$ emission toward six T Tauri stars using NIFS (Near Infrared Integral Field Spectograph) at the Gemini North Telescope on Mauna Kea. NIFS is an image slicing IFU fed by Gemini's near-infrared adaptive optics system, Altair, to obtain integral field spectroscopy with a resolving power of $R \sim 5000$. These observations revealed a variety of morphological distribution within 200 AU of T Tauri stars, and showed that H$_2$ molecules are presumably excited by shocks. The study we present shows a detailed analysis of the emission associated with one of the best studied YSOs, HL Tau. The YSO is in transient phase from a Class I protostar to T Tauri star (see Pyo et al. 2006, and references therein), and is known to host an extended collimated jet (Mundt et al. 1990), millimeter CO outflow (Monin et al. 1996; Cabrit et al. 1996), circumstellar disk (e.g., Wilner et al. 1996; Kitamura et al. 2002), and also a flattened gas envelope (Sargent \& Beckwith 1990) which may be infalling toward the star (Hayashi et al. 1993). Paper I revealed the presence of a bubble-like morphology in H$_2$ emission extending toward the northeast. In this paper we analyze the results for H$_2$ 1-0 S(1) emission (2.122 $\micron$) together with the [\ion{Fe}{2}] 1.644 $\micron$ line and also adjacent continuum. Our results suggest the presence of an unseen wide-angled wind, interacting with the ambient gas and produce a bipolar cavity and shocked H$_2$ emission. Throughout the paper, we adopt the distance to the target of 140 pc (Elias 1978). | Mass ejection from low-mass YSOs is often observed in two different manners: (1) collimated jets often seen in optical-IR wavelengths, or millimeter wavelengths in some cases (see e.g., Bally et al. 2007 for review) (2) molecular bipolar outflows, which show a variety of morphologies in millimeter CO lines (see e.g., Arce et al. 2007 for review). The relation between these types of flows remain unclear. In particular, two major scenarios have been discussed over decades for driving of the molecular outflows. These are: (1) jet-driven scenario, for which the molecular outflow results from interaction between a collimated jet and ambient material (e.g., Raga \& Cabrit 1993); and (2) wind-driven scenario, for which the molecular outflow results from an unseen wide-angled wind (e.g., Shu et al. 1991). Studies to date suggest that neither the jet-driven nor wind-driven models can explain a wide range of morphologies and kinematic properties observed in all outflows (see Cabrit et al. 1997; Arce et al. 2007). The observed morphology and kinematics in H$_2$ emission from HL Tau are similar to millimeter CO outflows. It is striking that the H$_2$ flow observed in HL Tau is smaller than a typical spatial scale of the millimeter CO outflows. While the CO outflows often extend over 1000--10000 AU scales (see e.g., Lee et al. 2000, 2002), including HL Tau itself (Monin et al. 1996; Cabrit et al. 1996), the H$_2$ flow in HL Tau shows a spatial scale of only $\sim$150 AU. The H$_2$ emission shows a remarkably different geometry and kinematics from the [\ion{Fe}{2}] emission associated with the collimated jet. Furthermore, there is no evidence for kinematic interaction between the H$_2$ flow and the collimated jet: i.e., there is no clear evidence of acceleration of deceleration of the H$_2$ and [Fe II] flows where these are overlapped in Figure 2. These results support the scenario that the jet is surrounded by an unseen wide-angled wind, which interacts with the ambient gas and produces the bipolar cavity and also shocked H$_2$ emission. The presence of such a wide-angled wind indeed agrees with proposed magneto-hydrodynamically driven wind models (see e.g., Brandford \& Payne 1982; Uchida \& Shibata 1985; Shang et al. 2006). Alternatively, Close et al. (1997) suggests that the cavities associated with HL Tau could be produced by a precessing jet. However, we emphasize that our results do not show any evidence for kinematic interaction between the [\ion{Fe}{2}] emission associated with the collimated jet and the H$_2$ emission at the cavity walls. Furthermore, a typical cooling time scale of the near-infrared H$_2$ emission line regions is less than a year (see e.g.,Takami et al. 2004), and the gas associated with the jet would move only $<$0\farcs2 in such a period (Mundt et al. 1990). Thus, the jet would have to be spatially located much closer to the H$_2$ emission line regions if it was responsible for the H$_2$ line excitation. The arcs in H$_2$ emission 1''.0 away from the star suggests that the ejection activity is time-variable. Assuming a flow velocity of 10 km s$^{-1}$ inclination angle of 34$^\circ$ from the plane of the sky (Mundt et al. 1990), we estimate the dynamical age of these arcs of $\sim$70 yrs. Periodic monitoring observations spread over year-long timescales are necessary to measure their proposer motions and determine the accurate dynamical age. Such a time-variable wind may be a rather common activity associated with low-mass YSOs. Indeed, a wind bubble similar to the northwest H$_2$ emission is also observed in the neighboring YSO XZ Tau in optical emission lines (Krist et al. 1997, 1999; Coffey et al. 2004). Furthermore, some millimeter CO outflows show multiple shell structures extending over 1000--10000 AU scales (see e.g., Lee et al. 2002). Arce \& Goodman (2001) show that episodic outflows well explain the presence of ``Hubble wedges'' (i.e., a jagged profile) in the position-velocity diagram, and also steep power-low slopes of mass-velocity relation ($dM(v)/dv \propto v^{-\gamma}$, $\gamma > 2$) observed toward some outflows. Time-variable mass mass ejection may be related to episodic mass accretion, observed in extreme cases, as FU Orionis (FUor) or EX Orionis (EXor) outbursts (see e.g., Hartmann \& Kenyon 1996; Herbig 1989). Indeed, optical-IR absorption lines of FUors suggest the presence of an energetic disk wind (e.g., Calvet et al. 1993). Furthermore, spectroscopic monitoring of CO overtone absorption suggests that one of the FUors ejected an expanding shell of dense, low-temperature material within a few decades ago (Hartmann et al. 2004). Such a link between episodic mass ejection and accretion has also been proposed to explain the morphology of the collimated jet associated with YSOs (e.g., Dopita 1978; Reipurth 1989; Zinnecker et al. 1998). | 7 | 10 | 0710.1148 |
0710 | 0710.3651_arXiv.txt | In this paper we continue our study of CMB TE cross correlation as a source of information about primordial gravitational waves. In an accompanying paper, we considered the zero multipole method. In this paper we use Wiener filtering of the CMB TE data to remove the density perturbation contribution to the TE power spectrum. In principle this leaves only the contribution of PGWs. We examine two toy experiments (one ideal and one more realistic), to see how well they constrain PGWs using the TE power spectrum. We consider three tests applied to a combination of observational data and data sets generated by Monte Carlo simulations: (1) Signal-to-Noise test, (2) sign test, and (3) Wilcoxon rank sum test. We compare these tests with each other and with the zero multipole method. Finally, we compare the signal-to-noise ratio of TE correlation measurements first with corresponding signal-to-noise ratios for BB ground based measurements and later with current and future TE correlation space measurements. We found that an ideal TE correlation experiment limited only by cosmic variance can detect PGWs with a tensor-to-scalar ratio $r=0.3$ at $98\%$ confidence level with the $S/N$ test, $93\%$ confidence level with the sign test, and $80\%$ confidence level for the Wilcoxon rank sum test. We also compare all results with corresponding results obtained using the zero multipole method. We demonstrate that to measure PGWs by their contribution to the TE cross correlation power spectrum in a realistic ground based experiment when real instrumental noise is taken into account, the tensor-to-scalar ratio, $r$, must be approximately four times larger. In the sense to detect PGWs, the zero multipole method is the best, next best is the S/N test, then the sign test, and the worst is the Wilcoxon rank sum test. | Primordial gravitational waves (PGWs) (tensor) generate negative temperature-polarization (TE) correlation for low multipoles, while primordial density (scalar) perturbations generate positive TE correlation for low multipoles (see \cite{crittenden94,baskaran06,grishchuk07,pmkpaper1}; and references in \cite{pmkpaper1}). This signature can be to detect PGWs. The test based on this signature (the \textit{zero multipole method}, see \cite{pmkpaper1}) is useful as an insurance against false detection or as a monitor of imperfectly subtracted systematic effects. In this paper, we analyze an alternative method. This method uses Wiener filtering to remove the contribution of density perturbations to the TE cross correlation power spectrum for small $\ell$, leaving only the negative residual component of the TE power spectrum due to PGWs. Actually, this method can be treated a test of the (negative) contribution to the TE correlation power spectrum due to PGWs on large scales using uncertainties in the measurements consistent with the total TE power spectrum. We use Monte Carlo simulations to analyze the probability of detecting PGWs using this method. By detection of PGWs we mean in this paper, the measurement of the parameter $r$, the ratio of the primordial tensor power spectrum, $P_t(k)$, to the primordial scalar power spectrum, $P_s(k)$, taken at wavenumber, $k_0$: \begin{equation} r = \frac{P_t(k_0)}{P_s(k_0)} = \frac{A_t}{A_s} \end{equation} where $k_0 = 0.05 \text{ Mpc$^{-1}$}$ (see \cite{Smith2006}). For this paper we only consider the tensor-to-scalar ratio, $r$. All other parameters are assumed to be at their WMAP3 values (\cite{Spergel2007WMAP}). The only other parameter which might affect the results is the tensor spectral index, $n_t$, however we assume, in this paper, that $n_t$ is very close to zero (\cite{peiris03}). The problem of $n_t < 0$ will be considered in another paper. The plan of this paper is the following. In Section \ref{wienfilt}, we describe the method for detection of PGWs based on measurements of the TE power spectrum. In Section \ref{simul}, we describe the numerical Monte Carlo simulations we tested. In Sections \ref{montecarlo}, \ref{signintro}, and \ref{wilcoxonintro}, we introduce the statistical tests used to contstrain PGWs. In Section \ref{comptests}, we compare the three different statistical tests used in this analysis. Section \ref{resul} gives the results for the two toy experiments described in \cite{pmkpaper1}. The only uncertainty in the first toy experiment is due to cosmic variance (\ref{toy1res}). In the second toy experiment, along with cosmic variance, we take into account instrumental noise (\ref{toy2res}). We present results of Monte Carlo simulations for WMAP (\ref{wmapres}) and Planck (\ref{planckres}). In Section \ref{bbcomp}, we compare signal-to-noise ratio for BB and TE measurements. | The method described here is one in which we filter out the signal due to density perturbations, leaving only the contribution to the TE power spectrum due to PGWs. We then test the resulting TE power spectrum to see if it is negative. Three different statistical tests were used to see if there was a significant detection of PGWs. The $S/N$ test can give a value for $r$ using a comparison with Monte Carlo simulations, while the Wilcoxon rank sum test can only give an allowable range for $r$. The sign test will only tell us if $r \not= 0$. From \cite{pmkpaper1}, we saw that we could detect $r=0.3$ to $3\sigma$ with a measure of $\ell_0$, the position where the TE power spectrum first changes sign, for the ideal experiment. Using the method discussed in this paper, we see that we are unable to make this significant of a detection. The best result was for the $S/N$ test which would give a $2.3\sigma$ detection of $r=0.3$. To detect PGWs on the level of $3\sigma$, the tensor-to-scalae ratio $r$ should be $r \ge 0.4$. The sign test would give $2\sigma$ detection for $r=0.3$ and a $3\sigma$ detection for $r=0.45$. The Wilcoxon ranked sum test gives only a $1.2\sigma$ detection for $r=0.3$ and a $3\sigma$ detection for $r=0.7$. Similar results were gotten for the other three experiments tested. Thus in the sense of potential to detect PGWs, the zero multipole method is the best, next best is the $S/N$ test, then the sign test, and the worst is the Wilcoxon ranked sum test. \cite{baskaran06} present illustrative examples in which high $r$ is consistent with measured TT, EE, and TE correlations. The value of $r$ is so high in these examples that if PGWs with such $r$ really existed, current BB experiments would already detect PGWs. All models predict that the TE cross correlation power spectrum change sign only once for $\ell < 100$. The fact WMAP cannot exclude several multipoles with $C_{\ell}^{TE} > 0$ in between multipoles of $C_{\ell}^{TE} < 0$ means that the TE cross correlation power spectrum either changes sign several times for $\ell < 100$ or there is some instrumental noise which causes some anticorrelation measurements. Using instrumental noise consistent with WMAP, our Monte Carlo simulations give $\Delta \ell_0 \approx 16$ and $\ell_0 > 40$, which means that there is no evidence of PGWs in the TE correlation power spectrum. | 7 | 10 | 0710.3651 |
0710 | 0710.3698_arXiv.txt | The physical processes that define the spine of the galaxy cluster X-ray luminosity -- temperature (L-T) relation are investigated using a large hydrodynamical simulation of the Universe. This simulation models the same volume and phases as the Millennium Simulation and has a linear extent of $500h^{-1}$Mpc. We demonstrate that mergers typically boost a cluster along but also slightly below the L-T relation. Due to this boost we expect that all of the very brightest clusters will be near the peak of a merger. Objects from near the top of the L-T relation tend to have assembled much of their mass earlier than an average halo of similar final mass. Conversely, objects from the bottom of the relation are often experiencing an ongoing or recent merger. | Since the launch of the XMM-Newton and Chandra satellites \citep{XMM,Chandra}, measurements of the X-ray emission from hot gas in clusters of galaxies have achieved unprecedented levels of accuracy and depth. However, the physical origin of the scaling relations between observable quantities, such as the luminosity of the X-ray emitting gas and its temperature, remain only partly understood. There are currently a number of surveys \citep{Romer, Schwope, Pierre} in progress with the potential to greatly expand our understanding of the processes that define correlations such as the luminosity-temperature (L-T) relation of clusters. For this potential to be realised we require a sound theoretical basis upon which to work. To this end, numerical hydrodynamical simulations have become indispensable tools and continue to grow in size and complexity \citep{Pearce, Kay07, Faltenbacher} but they have to date lacked a sufficiently large dynamic range in mass. In this work we use a hydrodynamical model of a large volume that contains over a hundred galaxy clusters. For the first time we are able to study the evolutionary processes within a cosmological context as we have hundreds of well resolved objects spanning a large dynamic range rather than the more typical handful \citep{Rowley} (hereafter R04), or idealised models \citep{Ritchie,Poole}. This paper is organised as follows: in the remainder of this section we summarise the work done to date on defining the physical processes that define the shape of the L-T relation. Then, in section 2, we give an account of the simulations we have undertaken, explain how our cluster sub-sample was selected and how the properties of these clusters were derived. Section 3 details our results before we discuss their implications and conclude in section 4. X-rays are chiefly emitted from the hot gas in clusters via thermal bremsstrahlung (for dark matter halos more massive than $10^{14} h^{-1} {\rm M_\odot}$ their temperature is typically above $2{\rm keV}$). For such a homologous population \citet{Kaiser} showed that simple scaling relations were expected between bulk properties such as the mass, temperature and luminosity. Observational work subsequently found that the properties of X-ray clusters where indeed related but the slopes of the relations were not those derived by Kaiser. \citet{Kaiser} assumed that galaxy clusters were self-similar entities and that therefore only a single property, such as the mass, was required in order to describe the other bulk properties. Such a homology results in an L-T relation with a power-law slope of 2. However, as figure~\ref{obsdata} demonstrates, X-ray observations of clusters with a median redshift of $\sim 0.07$ found that the slope was closer to 3 \citep{Markevitch,Arnaud,Wu99, Xue,Horner,Mulchaey, Osmond} and perhaps became even steeper on group scales \citep{Helsdon}. \begin{figure*} \begin{minipage}{150mm} \begin{center} \opt{bw}{ \includegraphics[angle=0, width=\textwidth]{figs/bw_obsdatafinalwkey} } \opt{col}{ \includegraphics[angle=0, width=\textwidth]{figs/obsdatafinalwkey} } \caption{ Compilation of low-redshift observed group and cluster X-ray luminosities within $r_{500}$ compared to their emission-weighted temperature. $r_{500}$ is a radius enclosing an overdensity of 500. The small points are the simulated groups and clusters used in this work. The data was taken from variously: \citet{Markevitch,Arnaud,Wu99, Helsdon,Xue,Horner,Mulchaey,Osmond}.} \label{obsdata} \end{center} \end{minipage} \end{figure*} Hydrodynamical simulations performed in the absence of cooling or any additional heat sources other than compression and shock heating have long been known to reproduce the self-similar hierarchy well \citep{ENF,NFW}. Unfortunately they do not reproduce either the slope or the normalisation of the observations, producing clusters that are too bright for any given temperature, even at the bright end. Following this work, simulations with limited physics within a cosmological volume have been used in an attempt to reconcile the apparent discrepancy between theory and observation regarding the slope of the L-T relation \citep{Pearce, Muanwong, Bialek,Borgani02}. These models showed that a simple cooling or preheating scheme was sufficient to match the simulated L-T relation to that observed at redshift zero. More recently \citet{Kay07} investigated the effects of feedback on the X-ray properties of clusters in hydrodynamical simulations, and demonstrated that their results were in good agreement with both the observed scaling relations and structural properties (e.g. entropy and temperature profiles), particularly for cool-core clusters. \citet{Balogh06} investigated the role that preheating, cooling and concentration of the halo profile can have on the scaling relations. They found that, for a realistic range of halo concentrations, the scatter generated was minimal in comparison with observed values. Variations in the cooling time of the gas in the centre of clusters could account for much of the scatter but is limited by the age of the universe and so could not explain the whole range. Finally, varying feedback from supernovae and AGN could explain the entire range, but required an order of magnitude difference in energy injection to cover the whole envelope. Their result implies that it is processes in the cores of clusters that are primarily responsible for driving the scatter in the scaling relations. This confirms earlier work by \citet{Fabian, Markevitch} and \citet{McCarthy}. \citet{Kay07} identify the scatter with strong cool core clusters, and expect the scatter to be smaller at high redshift due to the diminished prevalence of such systems. Nowadays, the general consensus is that the scatter is largely due to the strength of the X-ray core. In this work, which includes strong preheating, X-ray cores are absent. This allows us to study the shape of the relation without the additional complication of a large intrinsic scatter. In this work we will use a sample of halos identified from the full simulation volume. With these we will show that because mergers tend to move clusters up the L-T relation they extend it beyond the point where the most massive, relaxed clusters are expected to lie. Thus many of the brightest, most luminous objects are ongoing or recent merger events which (as R04 point out) may be difficult to resolve observationally if they are close to the peak of the merger. In addition, because we have many closely spaced outputs we can track the motion of each of our clusters on the L-T plane, allowing us to define a ``mean merger'' vector. As this vector is not perfectly parallel to the L-T relation but rather falls slightly below it, a gentle roll in the relation naturally arises. | This work examines the physics that underlies the spine of the X-ray L-T relation. Due to our strong preheating prescription our halos do not have strong cores and as such do not reproduce the large scatter in the observed L-T relation, allowing us a clear window into the basic physics. We intend to examine the physical origin of the observed scatter in future work \citep{Gazzola} where a more physically motivated energy feedback prescription will be used and bright cooling cores are present. Preheating schemes such as the one used here are well known to accurately reproduce the slope and normalisation of the L-T relation as a whole \citep{Pearce, Muanwong, Bialek,Borgani02}. The model we have implemented also accurately reproduces the mean location of halos on the L-T plane at the present day but in a much larger volume than has typically been used previously. In the real world bright cooling cores will further complicate matters but the processes discussed here which relate to the outer halo properties will underlie these, with the variation in core properties leading to a scatter about the relation discussed here. By identifying mergers using the mass accretion histories of our objects and matching these episodes to the motion of each object on the L-T plane we have derived a ``mean merger'' vector in this plane. This vector lies largely parallel to the cluster L-T relation, as previously noted by RO4. At any particular time the mass function of the dark matter haloes present within a volume will be exponentially truncated at the high mass end above some characteristic mass scale. The large boost generated during a merger will produce points on the L-T plane appearing to lie above this characteristic mass, where there should be few objects. We therefore expect the majority of the brightest objects to be experiencing ongoing mergers, although they may be difficult to identify if they are close to their peak. The mean merger vector we have derived is not exactly parallel to the L-T relation but rather lies slightly below it. This behaviour leads to all bar one of our low-scattered objects being obvious recent or ongoing merger events (figure~\ref{lm}). We also note that at the high mass end the vast majority of our haloes lie below the mean relation shown on figure~\ref{sample}. The fact that the mean merger vector lies slightly below the mean relation provides a natural explanation for the slight curvature evidenced in the simulated relation. In summary, while it is straightforward to reproduce the observed slope and normalisation of the X-ray luminosity--temperature relation using a simple preheating scheme, such a scheme does not reproduce the observed scatter. As a preheating model includes the full underlying framework of the hierarchical build up of structure bulk mergers are not significant drivers of this scatter. Mergers can, however, produce objects that are brighter and hotter than would be expected from the cluster mass as merger events drive objects along the L-T relation towards the bright end. We find that a typical merger track does not exactly parallel the L-T relation but rather lies slightly below it, leading to a prevalence of recent or ongoing merger events on the low-scattered side of the relation. This process also leads to a slight curvature of the mean relation at the high mass end. \begin{figure} \begin{center} \opt{bw}{ \includegraphics[angle=0, width=250pt]{figs/bw_meanmerg} } \opt{col}{ \includegraphics[angle=0, width=250pt]{figs/meanmerg} } \caption{Relative motion on the L--T plane during each of the mergers defined in section 3.3. Each line represents a single merging event. The long dashed line indicates the mean L--T relation whereas the dotted line indicates the mean merger direction. } \label{mergefit} \end{center} \end{figure} | 7 | 10 | 0710.3698 |
0710 | 0710.1654_arXiv.txt | We report an SMA interferometric identification of a bright submillimeter source, GOODS 850-5. This source is one of the brightest 850 $\mu$m sources in the GOODS-N but is extremely faint at all other wavelengths. It is not detected in the GOODS \emph{HST} ACS images and only shows a weak 2 $\sigma$ signal at 1.4 GHz. It is detected in the \emph{Spitzer} IRAC bands and the MIPS 24 $\mu$m band, however, with very low fluxes. We present evidence in the radio, submillimeter, mid-IR, near-IR, and optical that suggest GOODS 850-5 may be a $z>4$ galaxy. | The Rayleigh-Jeans portion of the dust spectral energy distribution (SED) produces a strong negative $K$-correction and makes the observed submillimeter flux of a dusty galaxy almost invariant at $z>1$ to $z\sim10$ \citep{blain93}. This makes the submillimeter wavelength a potentially powerful probe to the high-redshift universe. However, to date, all the identified submillimeter galaxies are at redshifts lower than 4, likely because of the limited resolution of current submillimeter instruments and the limited sensitivity of current radio instruments. The Submillimeter Common-User Bolometer Array (SCUBA) on the single-dish James Clerk Maxwell Telescope resolved 20\%--30\% of the submillimeter extragalactic background light into point sources brighter than $\sim2$ mJy at 850 $\mu$m (\citealp{barger99,eales99}). Because of the low resolution ($\sim15\arcsec$) of SCUBA, identifications of the submillimeter sources have to assume the radio--FIR correlation in local galaxies \citep[see, e.g.,][]{condon92} and rely on radio interferometry to pinpoint the location of the submillimeter emission. Optical spectroscopy of radio identified submillimeter sources shows that they are ultraluminous ($>10^{12}$ $L_{\sun}$, corresponding to star formation rates of $10^2$--$10^3$ $M_{\sun}$ yr$^{-1}$) sources at $z\sim2$--3 and that they dominate the total star formation at this redshift \citep{chapman05}. However, the positive $K$-correction of the radio synchrotron emission makes the radio wavelength insensitive to high-redshift galaxies, and radio observations can only identify 60\%--70\% of the blank-field submillimeter sources \citep[e.g.,][]{barger00,ivison02}. The radio unidentified submillimeter sources are commonly thought to be at redshifts higher than the radio detection limit (typically $z\sim3$--4) but there is no direct evidence for such a high-redshift radio-faint submillimeter population. With recent development in submillimeter interferometry, it is now possible to locate submillimeter sources directly without relying on radio observations. We have begun a program to target radio-faint submillimeter sources with the Submillimeter Array (SMA\footnotemark[7]) to determine whether there is a high-redshift ($z>4$) tail in the redshift distribution of the submillimeter sources. Here we report our first identification in this program, GOODS 850-5. GOODS 850-5 was detected by our SCUBA jiggle-map survey in the Great Observatories Origins Deep Surveys-North (GOODS-N, \citealp{giavalisco04a}) with an 850 $\mu$m flux of $12.9\pm2.1$ mJy \citep{wang04}. It was also detected in the combined jiggle and scan map of GOODS-N (GN 10, see \citealp{pope06} and references therein). It is the second brightest submillimeter source in our jiggle-map catalog of the GOODS-N and has an IR luminosity of $\sim2\times10^{13}$ $L_{\sun}$. It does not have a $5 \sigma$ radio counterpart in the deep Very Large Array (VLA) 1.4 GHz catalogs of \citet{richards00} and \citet{biggs06}. We obtained an unambiguous identification of GOODS 850-5 with the SMA and found that the counterpart to this source is remarkably faint at \emph{all} other wavelengths. All data in the optical, near-IR, mid-IR, submillimeter, and radio point to a source at $z>4$. This is important evidence for the existence of high-redshift submillimeter sources. In this letter, we report the SMA observation and the multiwavelength photometry (\S~\ref{observation}) and the redshift constraints derived from the photometric data (\S~\ref{z_constraints}). We discuss the implication for the star formation history in \S~\ref{discussion} and summarize in \S~\ref{summary}. We adopt $H_0=71$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\Lambda}=0.73$, and $\Omega_M=0.27$. \begin{figure*} \epsscale{1.17} \plotone{f1.eps} \caption{Multi-wavelength images of GOODS 850-5. Each panel has a size of $24\arcsec$. North is up. All four of the ACS (F435W, $b$; F606W, $v$; F775W, $i$; and F850LP, $z$) bands and all four of the IRAC channels (1, 2, 3, and 4) are included in the color pictures. The color codes are labeled in the images. Grayscale images have inverse scales. The SMA position is labeled with $2\arcsec$ diameter circles. The contours in the SMA image have levels of -2, 2, 4, and 6 $\sigma$ where 1 $\sigma$ is 1.4 mJy beam$^{-1}$. \label{thumbnail}} \end{figure*} | \label{summary} We obtained accurate astrometry for the bright submillimeter source GOODS 850-5 with the SMA interferometer at 850 $\mu$m. The counterpart is extremely faint in the optical, near-IR, mid-IR, and radio. Both the radio and 24 $\mu$m faintness of GOODS 850-5 suggests a high redshift of $z>4$. The very red SED between 3.6 and 8.0 $\mu$m and the nondetection in the optical and $K_s$ bands can only be fitted by stellar continuum at $z>4$. This discovery provides important evidence that some of the radio undetected submillimeter sources are at high redshift and extends the redshift distribution of bright submillimeter sources to $z>4$. It suggests that a great fraction of star formation at high redshift is hidden from optical surveys. | 7 | 10 | 0710.1654 |
0710 | 0710.4137_arXiv.txt | The median observed velocity width $\v90$ of low-ionization species in damped Ly$\alpha$ systems is close to $90\kms$, with $\sim 10\%$ of all systems showing $\v90\gt210\kms$ at $z=3$. We show that a relative shortage of such high-velocity neutral gas absorbers in state-of-the-art galaxy formation models is a fundamental problem, present both in grid-based and particle-based numerical simulations. Using a series of numerical simulations of varying resolution and box size to cover a wide range of halo masses, we demonstrate that energy from gravitational infall alone is insufficient to produce the velocity dispersion observed in damped Ly$\alpha$ systems, nor does this dispersion arise from an implementation of star formation and feedback in our highest resolution ($\sim45\pc$) models, if we do not put any galactic winds into our models by hand. We argue that these numerical experiments highlight the need to separate dynamics of different components of the multiphase interstellar medium at $z=3$. | Damped Ly-alpha absorbers (DLAs) provide us with a high-resolution probe of galaxy formation processes up to redshift $z\sim5$, allowing us to study the neutral hydrogen distribution and kinematics of young galaxies and their surroundings on scales from several pc to several kpc, although detailed information about the spatial extent of the absorbing gas cannot be readily extracted from the observed velocity line profiles. DLAs contain a large fraction of high-redshift baryons potentially available for star formation (SF; Prochaska et al. 2005), however, their in-situ SF efficiency appears to be a factor of 20-30 lower than in the present-day galaxies of comparable gas surface density \citep{wolfe.06}. Even though DLAs are presumably associated with fairly compact structures, they cover about a third of the sky, so there is a high probability that a line of sight to a remote quasar will contain such an absorber. Currently, just over a 1000 such systems have been studied, and their number will undoubtedly increase in coming years \footnote{http://www.ucolick.org/$\sim$xavier/SDSSDLA/}. One of the unsolved long-standing puzzles in our understanding of DLAs is their large neutral gas velocity dispersion \citep{prochaska.97}. The median value of $\v90$, the velocity interval encompassing 90\% of the optical depth, in absorption lines of low ions associated with cold neutral gas (such as SiII, NI, FeII, etc.) is close to $90\kms$, and 10\% of all systems have $\v90\gt210\km/s$. Observations indicate that there is virtually no correlation between the column density $\nhi$ of the absorber and its absorption line velocity width, with even lower column density systems at the DLA threshold $\nhi=10^{20.3}\cmsq$ demonstrating velocity widths up to $400\kms$. On the other hand, a correlation is observed between the velocity widths and metallicities of the DLAs. In fact, the correlation is remarkably tight between the equivalent width of the \ion{Si}{2}~1526 transition (a kinematic diagnostic because the line is saturated) and gas metallicity \citep{prochaska....07} that shows a slope matching the local velocity/metallicity relation in dwarf galaxies \citep{dekel.03}. These correlations may be a consequence of a relation between the mass of the galaxies and their metallicities \citep{wolfe.98,ledoux....06}, or an indication that feedback from SF has a direct effect on the observed velocity dispersion \citep{nulsen..98}. To date all DLAs exhibit metal-line absorption \citep{prochaska....03}, and one expects that they will all show significant MgII equivalent widths. The opposite is not true; the majority of MgII-selected absorbers are not DLAs. Unlike traditional low ions, MgII traces both cold and warm neutral material, as well as warm partially photoionized gas. Recently, \citet{murphy....07} found a correlation between metallicity and the rest-frame MgII equivalent width, suggesting a link between kinematics and the metal-enrichment history of the absorber. They stressed that the absorption-line kinematics should be viewed separately from the host galaxy kinematics, in other words, the observed velocity widths need not scale directly with the mass of the DLA-hosting halos. \citet{bouche....06} used the cross-correlation between 1806 MgII absorbers and $\sim250,000$ luminous red galaxies from the Sloan Digital Sky Survey to find that the absorber halo mass is anti-correlated with the Mg II equivalent width, suggesting that at least part of the observed velocity dispersion in MgII regions is produced by supernova-driven winds and/or other feedback mechanisms which could drive cold and warm gas efficiently out of lower mass ($\mhalo\lt10^{11.5}\msun$) galaxies. On the other hand, this anti-correlation could also be interpreted as a transition from cold to shock-heated gas in more massive $\sim10^{12.5}\msun$ halos accompanied by a drop in the cumulative cross-section of neutral clouds found in these halos \citep{tinker.07}. At least for now, observational data are insufficient to find a relation between the absorbing halo mass and the velocity width $\v90$ of neutral absorption in DLAs. Therefore, we should consider at least three distinct physical mechanisms which could contribute especially to the high-end tail of the DLA velocity distribution: 1) If cold clouds accreting onto massive galaxies are dense enough to survive collisional heating while falling into the hot ($\sim10^6{\rm\,K}$) virialized region, they could retain a sufficiently high neutral fraction with column densities above $10^{20.3}\cmsq$ \citep{mcdonald.99}. Note that for this mechanism to be viable, the accreting gas should be already sufficiently metal-rich, as even the most metal-poor DLAs have been found to have metallicities $Z\gsim-2.8$, i.e. it would imply a model in which SF was ubiquitous at some time before $z=3$ in field galaxies of mass $10^8-10^{10}\msun$. In this scenario a large fraction of metals observed in DLAs was produced far from massive ($\gsim10^{12}\msun$) halos where we see these metals in absorption. On the other hand, it is also possible that the low-metallicity accreting material could mix efficiently with the more metal-rich environment of the massive virialized halo. 2) Conversely, DLA kinematics could be dominated by outflow velocities of supernova-driven winds, i.e. we could see feedback directly \citep{nulsen..98}. This hypothesis is indirectly supported by the high detected outflow velocities in Lyman-break galaxies (LBGs), the higher effective widths $W_{1526}$ of the SiII $1526\AA$ transition in GRB-DLAs, and the anti-correlation between the absorbing halo mass and the effective line width in MgII absorbers. It would also naturally explain the metallicity-velocity correlation, since metals should be produced in the same regions which drive the winds. 3) The apparent inefficiency of in-situ SF in DLAs, coupled with their relatively high inferred cooling rates suggests a picture in which very compact star-forming regions, perhaps associated with LBGs, illuminate much more extended neutral gas structures \citep{wolfe.06}. \citet{maller...01} used a toy model of DLA absorption to show that a large covering factor of the cold gas in protogalactic clumps is consistent with observations. Moreover, numerical simulations confirm that massive halos at $z=3$ often consist of multiple compact galaxies embedded into larger neutral clouds. Combined outflows from stellar winds and SNe in these galaxies could drive large chunks of surrounding neutral material to larger galactic radii where they could pick up a higher velocity dispersion from the local galaxy group. Although we now have absorption data on $\sim$1000 DLAs, very few groups have attempted to model DLA kinematics in recent years. Both analytical and numerical models in the literature tend to rely on a velocity dispersion put in by hand, in part because of our lack of understanding of the physical processes of SF and feedback on sub-galactic scales, and in part due to numerical resolution limitations. \citet{mcdonald.99} used an analytical model combining the Press-Schechter formalism with a picture of individual spherically symmetric halos consisting of multiple absorbing clouds to show that the rate of energy dissipation corresponding to the velocity dispersion of these clouds necessary to produce the observed line profiles far exceeds the rate at which energy can be supplied to these clouds by gravitational collapse and mergers of halos. They noted that a large fraction of DLAs show multiple components in absorption profiles of the low-ionization lines, and that this observation can be reproduced by a model in which the missing dissipation energy comes from supernova explosions. In fact, the energy injection rate of $1.8\times10^{50}{\rm\,ergs\,yr^{-1}}/(10^{12}\msun)$ reproduces well the fraction of multi-component DLAs and the overall absorption line velocity width distribution with the median value of $\sim90\kms$. However, their model does not provide the physical mechanism for transferring the feedback energy to the turbulent motions in gas clouds, apart from specifying the source of this energy. In addition, it is a highly approximate model which assumes spherical exponential halos consisting of discrete neutral clouds moving with a given velocity dispersion. Moreover, the same feedback energy per unit mass is injected into all halos, independently of their mass, an assumption which probably puts too much velocity dispersion into small halos. On the numerical front, \citet{nagamine...07} used cosmological smoothed particle hydrodynamics (SPH) simulations coupled to a phenomenological galactic wind model to compute the rate of incidence of DLAs as a function of halo mass, galaxy apparent magnitude, and impact parameter, for a variable strength of hydrodynamical feedback. In their model, gas particles were driven out of dense SF regions by hand, by assigning a momentum in random directions. The wind mass loss rate was assumed to be twice the SF rate, and the wind carried a fixed fraction of the SN feedback energy, corresponding to the fixed wind velocities of $242\kms$ (weak wind) and $484\kms$ (strong wind). Depending on the strength of feedback, they found that it evacuated gas from predominantly low-mass galaxies and increased the cross-section and hence the rate of incidence of more massive galaxies. In their weak feedback model $\sim10^{9.6}\msun$ halos contributed the most to the cross-section, whereas with strong feedback the peak shifted to $\sim10^{11}-10^{12}\msun$ halos, depending on numerical resolution, as feedback became more and more efficient at removing gas from low-mass systems. Although, \citet{nagamine...07} did not analyze the velocity width statistics, one would expect to see numerous line profiles with $\v90\gt100\kms$ in their models, due the increased incidence rate of massive galaxies. The purpose of this paper is twofold. First, we test a hypothesis that the observed kinematics is driven primarily by energy coming from gravitational infall in the process of hierarchical buildup of galaxies. This is essentially an extension of the idea put forward by \citet{haehnelt..98} that the observed DLA line profiles are caused by a combination of random halo motions, rotation, and infall, coupled to the paradigm of several distinct modes of accretion onto growing proto-galaxies \citep{dekel.06}. We use the approach developed in our earlier paper \citep[][hereafter Paper I]{razoumov...06}, postprocessing high resolution adaptive mesh refinement (AMR) simulations of galaxy formation with high-angular resolution radiative transfer of UV ionizing photons. We improved our algorithm in several ways including a better treatment of hydrodynamical heating during the radiative transfer stage, as well as taking into account SF and feedback. We experimented with a number of SF models, including the standard four-criterion model of \citet{cen.92} and the two-mode SF model of \citet{sommer-larsen..03}, both of which can be tuned to give the observed SF rates in the absence of or with mild feedback. Including strong feedback in our experience tends to always suppress an ongoing SF, either unless relevant scales in the clumpy interstellar medium (ISM) are resolved, or interaction between the wind and the ambient medium has been turned off. It is well known that without any special treatment the feedback energy is quickly lost away in cooling. Several prescriptions have been suggested to alleviate this problem; popular algorithms used in particle-based galaxy formation models are (1) turning off cooling of gas particles in the feedback regions \citep{thacker.00,sommer-larsen..03,stinson.....06} and (2) using a subresolution model for the multiphase ISM \citep{springel.03a,nagamine...07} usually coupled with kinematic feedback in which individual multiphase gas particles are driven out of the star-forming regions in random or specified directions. Grid-based AMR simulations in which feedback energy is simply converted into the gas thermal energy seem to be quite effective in limiting SF, although this efficiency certainly depends on numerical and temporal resolution. In the grid-based simulations presented in Section~\ref{amrModels} we do not suppress cooling or use any kinematic outflow model, in other words, we enforce a fairly conservative approximation to the role of SF in DLA kinematics, largely limited to conversion of some fraction of gas into stars. In the second half of this paper (Sec.~\ref{sphModels}) we examine DLA kinematics with TreeSPH models in which radiative cooling is suppressed in regions surrounding active SF sites. | We have performed a numerical study of DLA kinematics using high-resolution grid-based AMR simulations with moderate feedback, as well as particle-based SPH models with much more efficient feedback. With grid-based models our goal was to test the hypothesis that the observed velocity dispersion in DLAs could be explained by energy coming from gravitational collapse of cosmic structures. The broadest line profiles would then come from neutral clouds surviving infall into the massive virialized ($\gsim10^{11}\msun$) systems, as a $10^{11}\msun$ halo at $z=3$ has a circular velocity already in excess of $100\kms$. We used a series of models with a comoving volume size ranging from $4h^{-1}$ to $32h^{-1}\mpc$ and the maximum physical resolution of $100\pc$ to build a DLA population covering halos in the mass range $10^{8.5}-10^{13}\msun$. Although our models can reproduce well the total incidence rate of DLAs at $z=3$, our median $\v90$ velocities are of order $40-50\kms$, well below the observed value of $90\kms$. We can point out three possibilities to resolve this discrepancy: 1. {\it More massive environments --} The velocity dispersion in individual clouds is a strong function of mass. Therefore, one could expect that including more massive halos through sampling of the rare density peaks could alleviate the velocity problem. On the other hand, all our 4-16 Mpc high resolution models feature similar velocity tails (Fig.~\ref{colVel-boxSize}) even though the mass function in the $16\mpc$ model L1 extends to nearly 50 times more massive halos than in the $4\mpc$ model N1. Therefore it seems unlikely that the more massive ($\mhalo\gt8\times10^{12}\msun$) and consequently much rarer halos could make any substantial contribution to the overall statistics, even though each individual halo could have a velocity dispersion of several hundred $\kms$. Note that including only massive halos in larger simulation volumes might produce $f(N,X)$ at low $\nhi$ which is too small \citep{jedamzik.98} or too large, depending on the technique. In our simulations neglecting low-mass halos can actually raise $f(N,X)$ at low column densities, as all the gas which would otherwise be pulled in by low-mass halos stays in extended filaments which sometimes can produce column densities $\nhi\gt10^{20.3}\cmsq$. On the other hand, our larger $32\mpc$ models H0 and H1 do not have enough grid resolution to resolve these filaments, and hence produce small line densities $\ell_{\rm DLA}(X)$. 2. {\it Resolution --} Our current simulations are already very expensive computationally as we refine adaptively everywhere in the volume, and the $128^3$ base grid models have in practice $300^3$ to $400^3$ resolution elements. If we start from $256^3$ or a larger base grid, many more halos will form drawing in a larger fraction of baryons at earlier redshifts, whereas inside individual galaxies we will start resolving clumpy interstellar gas. However, without running these simulations and sampling a large number of galaxies at higher resolution, it is impossible to predict reliably the effect on the velocity and column density distributions. In addition, higher resolution would allow us to better model propagation of supernova-driven winds in the clumpy interstellar gas. The key question is the relative contribution of lower mass ($\lsim10^{10.5}\msun$) and more massive ($\gsim10^{10.5}\msun$) halos. In our current models these two populations contribute about equally to the total ``galactic'' DLA column density (Fig.~\ref{incidenceMass}). It is possible that at higher resolution the relative contribution of the more massive halos might grow, as they will feature more numerous spatially separated components falling on the same line of sight, perhaps with neutral gas tidal tails and bridges, whereas isolated lower mass galaxies could become even more compact and pull the remaining cold neutral gas from extended filamentary structures decreasing its covering factor on the sky. Unfortunately, the spatial distribution of absorbing clouds cannot be reconstructed from the observed line profiles in the velocity space, which would otherwise point to the minimum scale needed to resolve muli-component DLAs. \citet{lopez...02} studied an almost dust free DLA at $z=2.33$ featuring 14 resolved metal line components. They found very similar Z/Fe ratios of low-ionization species in all 14 components, suggesting that these components could originate in clouds physically located in a small region of space, perhaps much smaller than the resolution limit of our current simulations. 3. {\it Local microphysics --} The most intriguing possibility is that the observed velocity dispersion could arise as a result of feedback from SF. More efficient feedback could disrupt the highest column density absorbers ($\nhi\gsim10^{21.5}\cmsq$) in the lower panel of Fig.~\ref{scatter-colVel-boxSize} creating a population of low $\nhi$ systems with high velocity widths. With such a disruption mechanism put in by hand with the galactic wind model, \citet{nagamine...07} saw efficient gas removal from lower-mass ($\lsim10^{10}\msun$) halos and increased cross-sections in $\gsim10^{11}\msun$ halos. In view of our results, a model in which more efficient feedback is obtained without resorting to the unphysical assumptions of kinematic winds and/or suppressing cooling in feedback regions would seem particularly compelling. The key challenge here is to come up with a model which describes dynamics of individual components of the multiphase ISM separately from each other, either resolving them explicitly, or using a subresolution model in which different components are advected differently on a grid. In these models supernova winds would channel most of their energy into the lower density regions between the star-forming clouds, whereas the clouds would be more likely to survive the disruptive effect of the winds. In Section~\ref{sphModels} we used TreeSPH models to test the effect that the suppression of cooling has on neutral gas kinematics. We found the median velocity widths $\v90$ of only 31 to $34\kms$, even though that these models have been shown to expel gas efficiently from the star-forming regions at high redshifts and produce realistic galactic disks at $z=0$. In our view, such low velocity dispersion could be explained by a number of factors. A. {\it Resolution --} When numerical models do not resolve density inhomogeneities in the ISM, putting a certain amount of supernova feedback into the thermal energy of the gas and turning off cooling leads to an efficient conversion of this energy into the superwind expansion into a nearly uniform medium. However, the mass loading of the wind in a uniform medium is much larger than in a clumpy medium of the same average density, since in the latter case the wind will predominantly expand into the lower density voids between the clumps creating a high velocity outflow which will occasionally sweep chunks of neutral material off the edges of denser clouds. Clumps of gas entrained in winds are often observed in low-redshift outflows \citet{rupke..05}. On the other hand, at low resolution winds will displace most gas in the galaxy moving it to somewhat higher galactic radii, producing ``puffy'' galaxies with a relatively low velocity dispersion. This is exactly what we see in SPH simulations at $45\pc$ physical resolution. The expectation is that at higher resolution we should see multiphase outflows, with higher overall velocities in hot ionized gas, and numerous dense clouds giving rise to multi-component low-ionization absorption lines. Getting such galactic winds from first principles is notoriously difficult in numerical simulations. In this paper, we argue that the necessary condition for obtaining the correct DLA velocity dispersion is a separate dynamical treatment of different components of the multiphase interstellar medium at $z=3$. In addition to the resolution effects, two other factors could have affected our SPH results. Here we just mention these effects briefly, as clearly their study goes beyond the reach of this paper. B. {\it Mean redshift of feedback --} A realistic population of galaxies at $z=0$ can be obtained in cosmological simulations invoking early ($z\gsim 4-6$) episodes of star formation with energetic feedback \citep{sommer-larsen..03}. In this scenario, the radial distribution of gas and its angular momentum arise as a cumulative result of successive episodes of star formation driving gas out of the galaxies. Other model parameters might result in similar star formation histories; not all parameters can be constrained uniquely from observations. One could speculate that a different star formation history, e.g. one featuring energetic feedback down to redshift $z\sim3$, would produce a higher neutral hydrogen velocity dispersion at such a redshift. We note that from the spectra of $z\sim3-4$ LBGs, outflow velocities of $300-400 \kms$ are routinely inferred \citep[e.g.,][]{pettini.......01, shapley...03}. Note, however, also that \citet{laursen.07} were able to match both the magnitude and radial fall-off of Ly$\alpha$ surface brightness of typical $z\sim3$ galaxies, indicating that the spatial distribution of neutral hydrogen is correctly predicted by current SPH models. C. {\it AGN feedback --} It is possible that the core of each massive galaxy hosts a supermassive black hole which might have shown an AGN-type activity at some time in the past. Neither grid-based nor particle-based simulations presented in this paper include feedback from AGNs which could have a profound impact on gas kinematics in host galaxies. Feedback from star formation is generally thought to play an important role in galaxy evolution. If galactic winds give rise to the observed neutral gas kinematics in DLAs, the same winds should affect the column density distribution. In Paper I and in this paper, we have attempted to look at both distributions through large-scale, ``blind-survey'' models, in which we simulate cosmological volumes using aggressive grid refinement for all galaxies. However, these models are very expensive to compute, especially as we get closer to resolving the typical scales of the ISM. Perhaps, a more efficient approach is to combine these large-scale numerical surveys with simulations of the ISM in isolated galaxies, whether or not accounting for mergers and cosmic infall. Fortunately, there is sufficient amount of observational data which can be used to constrain these models. Combining traditional QSO-DLA data with a closer look at the high-redshift star-forming regions with GRB-DLAs and with emission line studies should allow us to learn a lot more about young proto-galaxies in coming years. | 7 | 10 | 0710.4137 |
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