{ "0607/astro-ph0607512_arXiv.txt": { "abstract": "Using a sample of 19,464 galaxies drawn from the DEEP2 Galaxy Redshift Survey, we study the relationship between galaxy color and environment at $0.4 < z < 1.35$. We find that the fraction of galaxies on the red sequence depends strongly on local environment out to $z > 1$, being larger in regions of greater galaxy density. At all epochs probed, we also find a small population of red, morphologically early--type galaxies residing in regions of low measured overdensity. The observed correlations between the red fraction and local overdensity are highly significant, with the trend at $z > 1$ detected at a greater than $5\\sigma$ level. Over the entire redshift regime studied, we find that the color--density relation evolves continuously, with red galaxies more strongly favoring overdense regions at low $z$ relative to their red--sequence counterparts at high redshift. At $z \\gtrsim 1.3$, the red fraction only weakly correlates with overdensity, implying that any color dependence to the clustering of $\\sim L^{*}$ galaxies at that epoch must be small. Our findings add weight to existing evidence that the build--up of galaxies on the red sequence has occurred preferentially in overdense environments (i.e., galaxy groups) at $z \\lesssim 1.5$. Furthermore, we identify the epoch $(z \\sim 2)$ at which typical $\\sim L^{*}$ galaxies began quenching and moved onto the red sequence in significant number. The strength of the observed evolutionary trends at $0 < z < 1.35$ suggests that the correlations observed locally, such as the morphology--density and color--density relations, are the result of environment--driven mechanisms (i.e., ``nurture'') and do not appear to have been imprinted (by ``nature'') upon the galaxy population during their epoch of formation. ", "introduction": "The galaxy population both locally and out to $z \\sim 1$ is found to be effectively described as a combination of two distinct galaxy types: red, early--type galaxies lacking much star formation and blue, late--type galaxies with active star formation \\citep[e.g.,][]{strateva01, baldry04, bell04, menanteau06}. The spatial distribution of this bimodal galaxy population is frequently phrased today in terms of the so--called morphology--density or color--density relation. As first quantified by \\citet{oemler74}, \\citet{davis76}, and \\citet{dressler80}, the morphology--density relation holds that star--forming, disk--dominated galaxies tend to reside in regions of lower galaxy density relative to those of red, elliptical galaxies. Many physical mechanisms that could be responsible for this correlation between galaxy morphology, star--formation history, and environment have been proposed \\citep[see][for a review of probable mechanisms]{cooper06a}. Are the morphology-density and color--density relations a result of environment--driven evolution, or were these trends imprinted upon the galaxy population during their epoch of formation? Only through comprehensive studies of galaxy properties and environments, both locally and at high redshift, will we be able to understand the role of local density in determining the star--formation histories and morphologies of galaxies. While the close relationship between galaxy type and density was primarily uncovered via the study of nearby clusters, recent work using the 2--degree Field Galaxy Redshift Survey \\citep[2dFGRS,][]{colless01, colless03} and the Sloan Digital Sky Survey \\citep[SDSS,][]{york00} has established that the connections between local environment and galaxy properties such as morphology, color, and luminosity extend over the full range of densities, from rich clusters to voids \\citep[e.g.,][]{kauffmann04, balogh04a, blanton05, croton05, rojas05}. Furthermore, high--resolution imaging and spectroscopic data in increasingly more distant clusters $({\\rm to}\\ z \\sim 1)$ have shown that the trends observed locally persist to higher $z$, at least in the highest density environments \\citep[e.g.,][]{balogh97, treu03, poggianti06}. Using a large sample of galaxies drawn from the DEEP2 Galaxy Redshift Survey, \\citet{cooper06a} extended the understanding of the color--density relation at $z \\sim 1$ across the full range of environments, from voids to rich groups, showing that the correlation between galaxy color and mean overdensity found locally is in place, at least in a global sense, when the universe was half its present age. While the role of environment appears to have been very critical at $z \\sim 1$ and perhaps at earlier times, quantitative measures of the evolution of environmental influences on the galaxy population or of correlations such as the morphology--density and color--density relation are limited. Comparisons of local results with studies of high--redshift clusters have pointed towards significant evolution in the relationship between galaxy properties and local environment from $z \\sim 1$ to $z \\sim 0$ \\citep[e.g.,][]{dressler97, couch98, smith05}. While these results indicate an environment--driven evolution in the galaxy population (i.e., pointing towards nurture versus nature as the origin of the galaxy bimodality), such work has been limited to the vicinity of rich clusters, and thus we know little about evolution in the relationship between galaxy type and density across the full scope of galaxy environments. Furthermore, clusters include only a relatively small fraction of the total galaxy population at any epoch by number (and an even smaller fraction by volume). Thus, the evolution of the color--density relation among the vast majority of the galaxy population remains unprobed. In large clusters, the physical mechanisms at work (e.g., galaxy harassment, ram--pressure stripping, and global tidal interactions) go beyond those acting in group--sized systems and the field. Results from the first comprehensive study of galaxy environment over a broad range of densities at high $z$ indicate that such cluster--specific physical mechanisms cannot explain the global color--density relation as found at $z \\sim 1$ \\citep{cooper06a}. Accordingly, in looking for evolution in the relationship between the bimodal nature of galaxy properties and the local galaxy environment, we must turn our attention to the entire dynamic range of galaxy overdensities at high redshift. In this vein, recent work employing a sample of low-- and high--redshift field galaxies from the VIMOS VLT Deep Survey \\citep[VVDS,][]{lefevre05} has found a strong evolutionary trend in the color--density relation for galaxies spanning the redshift range $0.25 < z < 1.5$, with the color--magnitude diagrams for galaxies at $0.9 < z < 1.5$ showing no significant dependence on environment \\citep{cucciati06}. Also working at high redshift, a study of the blue fraction (that is, the fraction of galaxies that have blue color) in galaxy groups and in the field population by \\citet{gerke06b} finds instead that the field blue fraction significantly differs from that of the group population out to $z \\sim 1.3$. In this paper, we use the large sample of high--$z$ galaxies obtained by the DEEP2 survey to conduct a detailed study of the color--density relation at $0.4 < z < 1.35$. In \\S 2, we discuss the data sample employed along with our measurements of galaxy environments and colors. Our main results regarding the relationship between color and environment are presented in \\S 3. Finally, in \\S 4 and \\S 5, we discuss our findings alongside other recent results and summarize our conclusions. Throughout this paper, we assume a flat $\\Lambda$CDM cosmology with $\\Omega_m = 0.3$, $\\Omega_{\\Lambda} = 0.7$, $w = -1$, and $h = 1$. ", "conclusions": "In this paper, we present a detailed study of the evolution in the color--density relation at $0.4 < z < 1.35$. Using a sample of galaxies drawn from the DEEP2 Galaxy Redshift Survey, we estimate the local overdensity about each galaxy according to the projected $3^{\\rm rd}$--nearest--neighbor surface density. From this, we measure the evolution of the red fraction with environment across time. Our principal results are as follows: \\begin{itemize} \\item We find that the color--density relation observed locally still exists at $z > 1$; the fraction of galaxies on the red sequence increases with local galaxy overdensity to nearly the redshift limits of the DEEP2 survey. \\item At all epochs probed $(0.4 < z < 1.3)$, we find there exists a population of red, morphologically early--type galaxies residing in the the most underdense environments. \\item The color--density relation evolves with redshift, growing weaker with lookback time such that at $z \\gtrsim 1.3$ there is no detectable dependence of galaxy color on local environment in the DEEP2 sample. \\item Our results support a picture in which the red sequence grew preferentially in dense environments (i.e., galaxy groups) at $z \\lesssim 1.5$. Clearly, the local environment plays an important role in ``nurturing'' galaxies, establishing the existence of correlations such as the morphology--density and color--density relation over cosmic time. The strength of evolutionary trends suggests that the correlations observed locally do not appear to have been imprinted (by ``nature'') upon the galaxy population during their epoch of formation. \\item Our findings imply that there should be little color dependence in the clustering of $\\sim L^{*}$ galaxies at $z \\gtrsim 1.3$. \\end{itemize}" }, "0607/astro-ph0607038_arXiv.txt": { "abstract": "{Radio surveys at frequencies of $\\sim 1$~GHz allow to map the synchrotron emission (SE) in a frequency range where (except for very low Galactic latitudes or towards localized regions) it dominates over the other radio components. New all sky total intensity and polarization data at 1.4 GHz have been recently collected. We focus on the Galactic radio emission correlation properties described in terms of angular power spectrum (APS). We present for the first time the APS, in both total intensity and polarization modes, for some representative Galactic cuts and suitable APS power law (PL) parametrizations. ", "introduction": "\\label{sec:intro} \\vskip -0.15cm In the recent years a complete coverage of the radio sky at 1420~MHz, both in total intensity and in polarization intensity, has been achieved. It derives from the combination of total intensity surveys (Reich 1982, Reich \\& Reich 1986, Reich et al. 2001) and of the DRAO (Wolleben et al. 2004) and the Villa Elisa (Testori et al. 2004) polarization surveys, covering respectively the Northern and Southern celestial hemisphere, sensitive to the Stokes parameters $Q$ and $U$. All these surveys have a FWHM resolution of $36'$. The sky has been sampled at steps of $\\simeq 15'$ in the total intensity and Villa Elisa surveys. The preliminary version of the DRAO survey used in this work has a sky sampling much better than that of the Leiden surveys (Brouw \\& Spoelstra 1976). The typical sensitivities of these surveys are of few tens of mK. These properties allow a reliable study of the Galactic correlation properties up to a multipole $\\ell_{max} \\sim 250$ ($\\ell_{max} \\sim 180/\\vartheta_{min}(^\\circ)$, where $\\vartheta_{min}$ is the smallest angular scale at which accurate information can be extracted). \\vskip -0.3cm ", "conclusions": "" }, "0607/astro-ph0607454_arXiv.txt": { "abstract": "We present the clustering of DEEP2 galaxies at $0.71$ \\mpch, but rises on smaller scales due to merging events. Their model indicates that the relative bias of quasars to galaxies decreases at higher redshift and is $\\sim1$ at $z=1$, in accord with the results presented here. Recent measurements of the quasar auto-correlation function do not show a strong dependence of clustering on luminosity \\citep{Croom05}, which may be problematic for this paradigm, although the errors on the observations are still large enough ($\\sim$30\\% in the 2dF data) to be consistent with its predictions. More recently, \\cite{Hopkins05a} present an alternative model for quasar lifetimes in which bright and faint quasars are in similar physical systems but are in different stages of their life cycles. This work is inspired by numerical simulations of galaxy mergers \\cite{Springel05} which incorporate black hole growth and feedback. Whereas the \\cite{Kauffmann00} model assumes an exponential decline of the quasar luminosity with time, in the \\cite{Hopkins05a} scenario quasars spend more time on the lower-luminosity end of their light curves. The Hopkins et al. approach is also able to explain both the optical and X-ray quasar luminosity functions \\citep{Hopkins05b}. \\cite{Lidz06} build on this model, postulating from simulations that halo mass is strongly correlated with the peak quasar luminosity, but only indirectly connected to the highly variable instantaneous luminosity. This leads to predictions that faint and bright quasars reside in similar--mass dark matter halos and that quasar clustering should depend only weakly on luminosity. Based on this hypothesis, they estimate the mass distribution of dark matter halos that host active quasars at $z=2$ and the characteristic dark matter halo masses of active quasars at $01.3$. On the other hand, a number of tests based on information theory and Bayesian statistics show only marginal evidence for luminosity dependent clustering. Anyway, the quality of the data is not good enough to accurately quantify how quasar biasing depends on luminosity. We critically discuss the limitations of our dataset and show that a much larger sample is needed to rule out current models for luminosity segregation. Studying the evolution of the clustering amplitude with redshift, we detect an increase of the quasar correlation length with lookback time at the 99.3 per cent confidence level. Adopting the concordance cosmological model, we discuss the evolution of quasar biasing with cosmic epoch and show that quasars are typically hosted by dark matter haloes with mass $\\sim 10^{13} M_\\odot$. ", "introduction": "It is widely believed that quasars are powered by accretion onto supermassive black holes. However, a detailed understanding of the physical processes leading to quasar activity (and their connection with galaxy formation) is still lacking. Simple semi-analytic models associate quasars with galaxy major mergers and assume a tight relation between their instantaneous luminosity and the mass of the central black-hole, $M_{\\rm bh}$ (Kauffmann \\& Haehnelt 2000; Wyithe \\& Loeb 2003; Volonteri et al. 2003). The fraction of gas accreted onto the black hole during each merger is chosen to match the observed relation between the velocity dispersion of the bulge and $M_{\\rm bh}$ (Ferrarese \\& Merritt 2000). This ends up producing a correlation between the quasar luminosity and the mass of the host dark-matter halo. Since the clustering properties of dark-matter halos strongly depend on their mass, the quasar clustering amplitude is thus expected to sensibly depend on luminosity. Recent numerical simulations of galaxy mergers including black-hole accretion and feedback have cast some doubts on this picture (Springel et al. 2005). These numerical experiments suggest that a given black-hole produces quasar activity with a wide range of luminosities (Hopkins et al. 2005). During its active phase, the black-hole is most likely observed as a relatively low-luminosity quasar with a small Eddington ratio. For a short period of time, however, its emission reaches its peak value (close to the Eddington luminosity) which is indeed proportional to the mass of the powering black-hole. Based on these models, Lidz et al. (2006) conclude that quasar clustering should depend only weakly on luminosity. From the observational point of view, only recently quasar samples have grown big enough (in terms of number of objects) to attempt the study of the clustering amplitude as a function of luminosity. By analyzing the galaxy-quasar cross-correlation at $1.8\\simlt z\\simlt 3.5$ Adelberger \\& Steidel (2005) found no evidence for luminosity-dependent clustering. They used 79 quasars spanning 4.4 orders of magnitude in absolute luminosity which have been divided into 2 luminosity bins. Larger samples are obviously required to confirm this result. Croom et al. (2002, 2005) studied the redshift-space clustering amplitude of 2dF quasars as a function of their apparent magnitude. Even though these authors initially found weak evidence for brighter QSOs being more strongly clustered, their most recent analysis shows no indication of luminosity-dependent clustering. These two studies, however, do not address the issue of luminosity dependent clustering. In fact, they consider magnitude-limited samples within a broad redshift range ($0.31.3$, a frequentist model selection technique (the F-test) prefers a multi-parameter fit to the data at the 95 per cent confidence level (ii) A number of statistical tests based on information theory and Bayesian techniques show weak evidence for luminosity dependent clustering at high redshift ($z>1.3$) and no evidence at low redshift ($z<1.3$). These results somewhat depend on the number of principal components used in the fitting procedure. Accounting for a larger fraction of the bootstrap variance increases the significance of the detection of clustering segregation with luminosity. (iii) Larger datasets, possibly with a deeper coverage, are needed to discriminate among current models of quasar formation and to pin-point the detailed quasar clustering trends as a function of luminosity at a given redshift. (iv) Splitting the sample into six complementary redshift bins, we find strong evidence for an increase of the clustering amplitude with lookback time. We detect pure quasar-clustering evolution between $z_{\\rm eff}=0.93$ and $z_{\\rm eff}=1.99$ at the $2.7\\sigma$ confidence level. A linear fit for the evolution of $r_0$ with redshift is given in equation (\\ref{eq:r0fit}). (v) Accounting for the evolution of the mass density in a $\\Lambda$CDM model, we find that the high-redshift quasars ($z_{\\rm eff}=1.99$) are $\\sim 2.6$ times more biased than their low-redshift counterparts ($z_{\\rm eff}=0.93$). Evolution in $b$ is detected at the $4.3\\sigma$ confidence level. (vi) The clustering amplitude of optically selected quasars suggests that they are hosted by halos with mass $M\\sim 10^{13} M_\\odot$ (see also PMN04)." }, "0607/astro-ph0607296.txt": { "abstract": "{ We consider the production of positrons in microquasars, i.e. X-ray binary systems that exhibit jets frequently, but not continuously. We estimate the production rate of positrons in microquasars, both by simple energy considerations and in the framework of various proposed models. We then evaluate the collective emissivity of the annihilation radiation produced by Galactic microquasars and we find that it might constitute a substantial contribution to the annihilation flux measured by INTEGRAL/SPI. We also discuss the possible spatial distribution of Galactic microquasars, on the basis of the (scarce) available data and the resulting morphology of the flux received on Earth. Finally, we consider nearby ``misaligned\" microquasars, with jets occasionally hitting the atmosphere of the companion star; these would represent interesting point sources, for which we determine the annihilation flux and the corresponding light curve, as well as the line's spectral profile. We discuss the possibility of detection of such point sources by future instruments. } ", "introduction": " ", "conclusions": "" }, "0607/gr-qc0607143_arXiv.txt": { "abstract": " ", "introduction": "\\label{ch:Introd} Gravitational waves are the most elusive prediction of Einstein's theory of gravity. The indirect evidence of their existence relies on the observations of the binary pulsar PSR 1913+16 (Hulse and Taylor~\\cite{1975ApJ...195L..51H}), that shows a decay of the orbital period consistent with the loss of angular momentum and energy due to the emission of gravitational waves. The prospect of starting a new astronomy based on gravitational radiation and providing a new corroboration of General Relativity has motivated many theoretical and experimental researches. As a result, the detection of gravitational waves seems feasible in the next decade by an international network of Earth-based laser interferometer detectors (LIGO, VIRGO, TAMA300, and GEO600)~\\cite{virgoetal}, bar resonant antennas (EXPLORER, AURIGA, NAUTILUS, ALLEGRO)~\\cite{bars} and by the Laser Interferometer Space Antenna (LISA)~\\cite{lisa}. Three scientific runs have been so far carried out by the LIGO detectors, in collaboration with GEO and TAMA detectors for two of the three runs and with the bar detector ALLEGRO for the last run. The data analysis of the first and second science run sets upper limits on the gravitational signal emitted by a number of possible sources, such as stochastic background, coalescing binary stars, pulsars~\\cite{2004PhRvD..69l2004A, 2004PhRvD..69l2001A, 2004PhRvD..69j2001A, 2004PhRvD..69h2004A,2005PhRvD..72f2001A, 2005PhRvL..94r1103A}. The third science runs have been performed with a higher sensitivity and the data analysis leads to a significant improvement of the gravitational radiation upper limits~\\cite{Abbott:2005ez}. Meanwhile, a second generation of detectors is already in the design stage for the exploration of the high frequency band, up to several kHz (advanced GEO600~\\cite{schnabel-2004-21}, wide-band dual sphere detectors~\\cite{Cerdonio:2000bh}), with an improvement of sensitivity up to two orders of magnitude with respect to the first-generation instruments. Gravitational radiation could provide new information about the nature of astrophysical sources and help in the interpretation of the dynamical evolution of many such systems. Among the many sources of gravitational waves, the oscillations of compact stars are considered of great interest by astrophysicists and nuclear physicists. The extreme conditions present in the core of compact stars make them a unique laboratory, where nearly all the modern theoretical areas of research in physics can be tested. In many astrophysical scenarios, compact stars may undergo oscillating phases. After violent events such as core collapse of a massive star, an accretion-induced-collapse of a white dwarf, or a binary white dwarf merger, the newly born protoneutron star is expected to pulsate non-linearly before various dissipative mechanisms damp the oscillations. Another system where pulsating phases may occur is a massive meta-stable compact object, which is born after the merger of a binary neutron star system. The gravitational signal emitted by the stellar oscillations lies in the high frequency band ($\\nu \\gtrsim 1 kHz$) and strongly depends on the structure and physics of the star, for instance on the equation of state, rotation, crust, magnetic fields as well as on the presence of dissipative effects such as viscosity, shock formation, magnetic breaking, convective outer layers, etc. With a detailed analysis of the gravitational wave spectrum emitted by stellar pulsations we could infer through asteroseismology the fundamental parameters of neutron stars, such as mass, radius and rotation rate~\\cite{1998MNRAS.299.1059A, Andersson:1998ak}. This information is necessary for the nuclear physicists as they can test the equations of state proposed for the description of matter at supra-nuclear densities. However, the weakness of the gravitational signal and the noise associated with the location and technology of the detectors compels theorists to provide more and more accurate models to predict the spectral and wave form properties of the gravitational signal. These templates are indispensable for enhancing the chances of detection, by extracting the signal from the noise with statistical methods. The spectral and dynamical properties of the oscillations of compact stars have been extensively investigated during the last forty years in Newtonian and Einstein theories of gravity. Linear perturbative techniques are appropriate for the analysis of small amplitude pulsations both in the frequency~\\cite{Thorne:1967th, kokkotas-1999-2} and the time domain approach~\\cite{allen-1998-58, Ruoff:2001ux, Seidel:1987in, Seidel:1990xb, Nagar:2004ns}. In General Relativity, the oscillation spectrum of a compact object, such as black holes and neutron stars, is characterized by a discrete set of quasi-normal modes (QNM). These modes have complex eigenfrequencies whose real part describes the oscillation frequency, and the imaginary part the damping time due to the emission of gravitational waves. The classification of QNM is well known for a large set of stellar models and can be divided schematically in fluid and spacetime modes. The fluid modes have a Newtonian counterpart and can be sub-classified by the nature of the restoring force that acts on the perturbed fluid element. The spacetime modes are purely relativistic and are due to the dynamical role assumed by the spacetime in General Relativity (more details are given in chapter~\\ref{ch:4Lin_Pert} and reference therein). Rotating stars in General Relativity can be described with various approximations, such as the slow-rotation approximation~\\cite{Hartle:1967ha} or recently with codes developed in numerical relativity~\\cite{2003LRR.....6....4F}. The former approach is based on a perturbative expansion of the equations in powers of the dimensionless rotation parameter $\\epsilon = \\Omega / \\Omega_{K}$, where $\\Omega$ is the uniform angular rotation and $\\Omega_{K}$ is the Keplerian angular velocity, which is defined as the frequency of a particle in stable circular orbit at the circumference of a star. The measured period of the fastest rotating pulsar corresponds to a relatively small rotation parameter $\\epsilon \\sim 0.3$, which may suggest that the slow rotation approximation provides an accurate description of rotating stars even for high rotation rates. However, in these cases the accuracy of this perturbative approach is different for the various physical stellar quantities. For instance, the quadrupole moment shows an accuracy to better than twenty percent, while the radius of the corotating and counterrotating innermost stable circular orbits is accurate to better than one percent~\\cite{2005MNRAS.358..923B}. Nevertheless, a protoneutron star may be expected to have a higher rotation rate that is not possible to describe with the slow rotation approximation. These regimes can be better addressed in numerical relativity by evolving the full set of non-linear Einstein equations~\\cite{Dimmelmeier:2004prep, Stergioulas:2000vs, 2003LRR.....6....3S}. Furthermore, recent works on the core collapse~\\cite{Dimmelmeier:2002bk, Dimmelmeier:2002bm}, r-mode instability~\\cite{2000ApJ...531L.139R, 2001MNRAS.322..515L, 2001PhRvL..86.1152L}, accretion from a companion~\\cite{1993ApJ...419..768F}, or supernova fall back material~\\cite{2002MNRAS.333..943W}, show that a neutron star manifests a degree of differential rotation. \\\\ These studies have clarified the effects of rotation on the dynamics of the oscillations, as well as showed the importance of using a relativistic treatment that takes into account the effects of the dragging of inertial frames. In particular, the presence in rotating stars of instabilities due to emission of gravitational waves has gained great attention. Almost all classes of oscillations of rotating stars, such as the f- and r-modes, are potentially unstable to the so-called Chandrasekhar- Friedman-Schutz (CFS) instability~\\cite{1970PhRvL..24..762C, 1978ApJ...222..281F}. This appears because beyond a critical value of the stellar angular velocity, a mode that in a corotating frame is retrograde and then has negative angular momentum, may appear moving forward in the inertial frame of a distant observer. As a result, this inertial observer will detect gravitational waves with positive angular momentum emitted by this mode. Thus, the gravitational radiation removes angular momentum by the retrograde mode by making it increasingly negative and then leading to instabilities. The losses of the angular momentum through gravitational waves slow down the star on secular timescales; eventually the star rotates slower than a critical value and the mode becomes stable. These instabilities could be strong sources of gravitational radiation and also limit the rotation rate of neutron stars, providing a possible explanation for the measured rotation period of pulsars. Many studies are currently dedicated to understand whether viscosity, magnetic fields, shock waves on the stellar surface or non-linear dynamics of oscillations may saturate this instability. The non-linear analysis of stellar oscillations is more complex and only recently are some investigations being carried out, due to improvements achieved by the non-linear codes in numerical relativity~\\cite{2003LRR.....6....4F}. Different methods have been used to investigate the properties of non-linear oscillations, such as for instance 3-dimensional general relativistic hydrodynamics code in Cowling or conformal flatness approximations~\\cite{Stergioulas:2003ep, Dimmelmeier:2004prep}, a combination of linear perturbative techniques with general relativistic hydrodynamics simulations~\\cite{2005MNRAS.356.1371Z}, or a new method where the non-linear dynamics is studied as a deviation from a background, which is described by a stellar equilibrium configuration~\\cite{Sperhake:2001xi}. These works, which have been dedicated to investigate the non-linear dynamics of different astrophysical systems: non-linear oscillations of a torus orbiting a black hole~\\cite{2005MNRAS.356.1371Z}, non-linear axisymmetric pulsations of uniform and differential rotating compact stars~\\cite{Dimmelmeier:2004prep} and non-linear radial oscillations of non-rotating relativistic stars~\\cite{Sperhake:2001xi}, have revealed a new phenomenology associated with the non-linear regimes, the presence in the spectrum of non-linear harmonics. These harmonics arise from the coupling between different classes of linear modes or from non-linear self-couplings~\\cite{Sperhake:2001xi}, and have a characteristic that could be appealing for the detection of gravitational waves: their frequencies appear as linear combinations of the linear oscillation modes. Therefore, some of these non-linear harmonics (sub-harmonics) can emerge at lower frequencies than the related linear modes, and then be within the frequency range where the detectors have higher sensitivity. However, since the amplitude of non-linear perturbations is usually the product of the amplitudes of first order perturbations, in order to have a detectable gravitational wave strain one needs non-linear effects that can enhance the gravitational signal, such as resonances, parameter amplifications or instabilities. \\\\ Strong non-linear regimes are adequately studied with a fully non-perturbative approach. However, many interesting physical effects of mild non-linear dynamics can be well addressed by second order perturbative techniques. An example is given by the analyses of black hole collisions in the so-called ``close limit approximation'', where the second order perturbations of Schwarzschild black holes~\\cite{Gleiser:1995gx, Garat:2000gp} have provided accurate results even in non-linear regimes where the perturbative methods are expected to fail. Non-linear perturbative methods have been successfully used also for studying linear perturbations of rotating stars, where the rotation is treated perturbatively with the slow rotation approximation ~\\cite{Hartle:1967ha, 2003LRR.....6....3S}. \\\\ \\indent An important aspect of second order perturbative analyses is that of providing an estimate of the error associated with the first order treatment, as there is not an \\emph{a priori} method to determine the accuracy of the linear perturbative results. Thus, the convergence and the corrections associated with any term of the perturbative series can be determined only by investigating higher perturbative orders. Furthermore, non-linear perturbative equations are usually a system of partial differential equations, thus their numerical integration is computationally less expensive than the full Einstein equations which are treated in numerical relativity. This relative simplicity of the perturbative approach may then provide accurate results and can also be used to test the full non-linear simulations. However, an extension to second order perturbative investigations is not always straightforward~\\cite{Gleiser:1995gx, Garat:2000gp}. Some issues may arise from the identification of the physical quantities among the second order perturbative fields or from the movement of the stellar surface in non-linear stellar oscillations~\\cite{Sperhake:2001si, Sperhake:2001xi}. \\\\ \\indent In physical systems where the perturbative analysis can be described by more than a single parameter, as for stellar oscillations of a slowly rotating star or mode coupling between linear perturbations, the multi-parameter relativistic perturbation theory~\\cite{Bruni:2002sm, sopuerta-2004-70} can help the interpretation of the gauge issues of non-linear perturbations. The identification of gauge invariant quantities allows us to have direct information about the physical properties of the system under consideration. The construction of such quantities is not in general simple, but recent works~\\cite{Nakamura:2003wk, Nakamura:2004gi} show how to build second order gauge-invariant perturbative fields from the knowledge of the associated first order gauge invariant perturbations. For a specific class of astrophysical systems a gauge invariant and coordinate independent formalism has been introduced nearly thirty years ago ~\\cite{Gerlach:1979rw, Gerlach:1980tx} for the analysis of one-parameter non-radial perturbations on a time dependent and spherically symmetric background. Recently, this formalism has been further developed~\\cite{Martin-Garcia:1998sk, Gundlach:1999bt}, and has been used to study non-radial perturbations on a collapsing star~\\cite{2003PhRvD..68b4002H} and for linear perturbation on a static star~\\cite{Nagar:2004ns}. The research project we have been working on aims to extend the perturbative analysis of compact stars at non-linear orders, in order to have a more comprehensive understanding of stellar oscillations and the related gravitational radiation. In particular, this thesis presents a gauge invariant formalism and a numerical code for studying the coupling between the radial and non-radial perturbations of a perfect fluid spherical star. The formalism for the polar perturbations has been worked out in a first paper~\\cite{Passamonti:2004je}. The formalism and applications to axial perturbations are presented in~\\cite{Passamonti:2005axial}. Work in progress on the applications of the polar perturbative formalism will be presented in a future work. Radial and non-radial oscillations can be excited in the aftermath of a core collapse or by accreting matter on a neutron star. Radial perturbations of a non-rotating star are not damped by emission of gravitational radiation, but they can emit gravitational waves at non-linear order through the coupling with the non-radial perturbations. This picture changes in the presence of rotation, where the radial pulsations become sources of gravitational radiation and form a new class of modes, called quasi-radial modes. This coupling may be interesting for instance during a core bounce, where it is expected that an excitation prevalently of the quasi-radial and quadrupole modes. Even though the quadrupole component provides the dominant contribution to the gravitational radiation, the radial pulsations may store a considerable amount of kinetic energy and transfer a part of it to the non-radial perturbations. As a result, this non-linear interaction could produce a damping of the radial pulsations and an interesting gravitational signal. The strength of this signal depends naturally on the efficiency of the coupling, which is an effect worth exploring. In this thesis, we start to investigate this non-linear effect for small oscillations of a non-rotating star with the aim of including rotation in future works. \\\\ \\indent The polar coupling, i.e. between the radial and the polar sector of non-radial perturbations, is expected to be \\emph{a priori} more effective than the axial case. Indeed, the linear polar modes have a richer spectrum than the axial sector and from the values of the frequencies and the damping times of the fluid QNMs, the resonances and composition harmonics should be more probable between the polar fluid and the radial modes. However, for the purposes of this thesis we have implemented a numerical code for the axial coupling for mainly two reasons: \\emph{i)} the axial coupling can have interesting physical effects, \\emph{ii)} the perturbative equations are simpler and enable us to understand better the issues related to the numerical stability and accuracy of the code as well as the effects of the low density regions near the stellar surface on the non-linear simulations, etc. \\\\ \\indent When we consider the first order axial non-radial perturbations, we see that the only fluid perturbation is the axial velocity, which can be interpreted to describe a stationary differential rotation. Therefore, the linear axial gravitational signal does not have any dependence on the dynamics of the stellar matter. This picture changes at coupling order, where the differential rotation and first order metric perturbations can couple with the radial pulsations and source the axial gravitational waves. We will see in chapter~\\ref{sec:NumInt} that this axial coupling produces a new class of quasi-radial modes, which can exist only for differentially rotating stars. This thesis is organized in seven chapters. In chapter~\\ref{ch3:Non-Lin_Per}, we introduce the perturbative formalism used in this work, i.e. the multi-parameter relativistic perturbation theory and the gauge invariant formalism introduced by Gerlach and Sengupta and further developed by Gundlach and Martin Garcia~(GSGM). In chapter ~\\ref{ch:4Lin_Pert}, we describe the linear perturbations of a spherical star, i.e. the radial pulsations, the polar and axial non-radial perturbations. The equations for describing the coupling between the radial and axial and polar non-radial oscillations are presented in chapter~\\ref{ch:5_NL}, where in addition we discuss also the boundary conditions. In chapter~\\ref{ch:GINLP}, we present the proof of the gauge invariance of the perturbative tensor fields that describe the non-linear perturbations for this coupling. Chapter~\\ref{sec:NumInt} is dedicated to the numerical code that simulates in the time domain the evolution of the coupling between the radial and axial non-radial perturbations. In this chapter we give all the technical details relating to the code and the results of the simulations. Finally in chapter~\\ref{ch:conclusions}, the conclusions and possible future developments are discussed. The appendix has seven sections. We have reported the source terms of the equations derived by Gundlach and Mart\\'{\\i}n Garc\\'{\\i}a in section~\\ref{full-equations}. In section~\\ref{AppSW11MG}, we write the full expressions of the sound wave equation for the fluid variable $H$, while in sections~\\ref{AppSources} and~\\ref{AppSources_axial} we respectively present the source terms of the perturbative equations that describe the coupling between the radial pulsations and polar and axial non-radial perturbations. In addition, in section~\\ref{sec:Tens_Harm} we give the tensor harmonics, while some of the numerical methods used in the numerical code are given in sections~\\ref{sec:finit_appr} and~\\ref{sec:Num_meth}. \\chapter{Non-Linear Relativistic Perturbation Theory} \\label{ch3:Non-Lin_Per} Exact solutions of the equations of physics may be obtained for only a limited class of problems. This aspect is particularly present in General Relativity, where the complexity of the astrophysical systems and the non-linear Einstein field equations allows us to describe exactly only simplified and highly symmetric cases. Among various approximation techniques, perturbation methods are appropriate whenever the problem under consideration closely resembles one which is exactly solvable. It assumes that the difference from the exactly solvable configuration is small and that one may deviate from it in a gradual fashion. Deviations of the physical quantities from their exact solutions are referred to as perturbations. Analytically, this is expressed by requiring that the perturbation be a continuous function of a parameter, measuring the strength of the perturbation. Although perturbative techniques are more appropriate for small values of the perturbative parameter, sometimes they can give reliable results also for mildly non-linear regimes as shown for example in the analysis of black-hole collision \\cite{1999bhgr.conf..351P}. Hence, in many cases the validity limit of perturbative methods cannot be determined \\emph{a priori}. A more accurate estimation can be reached by studying the convergence of the perturbative series, which then involves the analysis of the second or higher perturbative orders. The gauge issue arises in General Relativity as in any other theory based on a principle of general covariance. The perturbative description of a physical system is not unique due to the presence of unphysical degrees of freedom related to the gauge, i.e. to the system of coordinates chosen for the analysis. This ambiguity can be eliminated either by fixing a particular gauge or by constructing perturbative variables which are invariant for any gauge transformation. In the former case, the properties and symmetries of the physical systems can help us to decide an appropriate gauge. In the latter approach, the identification of the gauge invariant fields is the difficult task. In this section we review the perturbative framework we have used for investigating the coupling between the radial and non-radial perturbations. In section~\\ref{sec:3.1_MPPT} we report the main results of the multi-parameter perturbation theory introduced by Bruni et al. \\cite{Bruni:2002sm} and Sopuerta et al.~\\cite{sopuerta-2004-70}. Section~\\ref{sec:GSGM} is dedicated to the formalism introduced by Gerlach and Sengupta \\cite{Gerlach:1979rw, Gerlach:1980tx}, which has been further developed by Gundlach and M. Garcia \\cite{Gundlach:1999bt}, while in section~\\ref{sec:3.3-NLFW} we outline the perturbative structure of our work which is based on the 2-parameter expansion of a static background. ", "conclusions": "" }, "0607/astro-ph0607332_arXiv.txt": { "abstract": "The formation of jets such as dynamic fibrils, mottles, and spicules in the solar chromosphere is one of the most important, but also most poorly understood, phenomena of the Sun's magnetized outer atmosphere. We use extremely high-resolution observations from the Swedish 1-m Solar Telescope combined with advanced numerical modeling to show that in active regions these jets are a natural consequence of upwardly propagating slow mode magneto\\-acoustic shocks. These shocks form when waves generated by convective flows and global p-mode oscillations in the lower lying photosphere leak upward into the magnetized chromosphere. We find excellent agreement between observed and simulated jet velocities, decelerations, lifetimes and lengths. Our findings suggest that previous observations of quiet sun spicules and mottles may also be interpreted in light of a shock driven mechanism. ", "introduction": "The solar chromosphere is sandwiched between the surface, or photosphere, and the hot and tenuous outer corona. This highly structured region, on average 2000~km thick, is constantly perturbed by short lived (3 -- 10 minutes), jet-like extrusions that reach heights of 2000 -- 10000~km above the photosphere. These thin jets are formed in the vicinity of photospheric magnetic field concentrations. Until recently, their small size and short lifetimes have made detailed analysis difficult \\citep[]{Beckers1968,Suematsu+etal1995}, which has led to a multitude of poorly constrained theories of their formation \\citep[]{Sterling2000}. In addition, there has been considerable confusion about the relationship between spicules at the quiet Sun limb, mottles observed on the quiet Sun disk, and dynamic fibrils (DFs) found in the vicinity of active region plage \\citep[]{Grossmann-Doerth+Schmidt1992}, although the similarity in many of their properties strongly suggests some of these phenomena are related \\citep{Tsiropoula+etal1994}. We focus on observations of DFs (\\S 2), compare them to advanced numerical simulations (\\S 3), report on regional differences of DF properties (\\S 4), and finish with a comparison to quiet Sun jets (\\S 5). ", "conclusions": "There are many striking similarities between quiet Sun mottles and the active region DFs studied here. Both phenomena appear as highly dynamic, dark features in the wings and core of H$\\alpha$, and are associated with magnetic flux concentrations. More importantly, \\citet{Suematsu+etal1995} found evidence that quiet Sun mottles also follow parabolic paths with decelerations that are too small to be consistent with a purely ballistic flight at solar gravity. While the interpretation of their observations proved difficult without detailed numerical models, \\citet{Suematsu+etal1995} note that the apparent velocity profiles in mottles are fully compatible with impulsive acceleration followed by a constant deceleration, with maximum upward velocities usually about equal in amplitude to the maximum downward velocities. The velocities they report for mottles, of order 10-30~km/s, are similar to those we find in our SST observations of DFs. \\citet{Suematsu+etal1995} also find that the largest Doppler velocities in mottles appear at the beginning of the ascending phase (blue-shifts) and at the end of receding phase (red-shifts), with downward red-shifted motion sometimes occurring close to their base during the ascending phase. These observations agree well with the properties of our simulated jets (Fig. 2, lower panel). In addition, \\citet{Christopoulou+etal2001} observe limb spicules with clear parabolic paths with decelerations and maximum velocities similar to those for mottles and DFs. All of these strong similarities between previous mottle and spicule observations and our modeling and SST observations of DFs seem to imply that highly dynamic chromospheric shock waves cause significant up- and downward excursions of the upper chromosphere in both active region and quiet Sun, as proposed by \\citet{DePontieu+etal2004}. Some unresolved issues remain, such as the longer lifetimes of quiet sun mottles and spicules (2-10 minutes), and the greater heights of 2-10 Mm that spicules reach at the limb. % Preliminary analysis of our simulations suggests that these differences could be related to large scale differences in magnetic topology. Further numerical simulations of various magnetic topologies will help resolve these issues. For example, it is possible that spicules reach slightly greater heights because they consist of two populations: jets that are driven by shocks (as described here), and jets caused by reconnection. The latter jets could form a subset that on average is taller than the shock driven jets, and perhaps be part of a continuous spectrum of reconnection jets that includes surges, macrospicules and H$\\alpha$ upflow events \\citep[]{Chae+etal1998}. Whatever the role of reconnection in quiet Sun, our findings indicate that, at least in active regions, most jets are caused by chromospheric shocks driven by convective flows and oscillations in the photosphere." }, "0607/astro-ph0607568_arXiv.txt": { "abstract": "{ Oscillations of stellar $p$~modes, excited by turbulent convection, are investigated. In the uppermost part of the solar convection zone, radiative cooling is responsible for the formation of turbulent plumes, hence the medium is modelled with downdrafts and updrafts.} {We take into account the asymmetry of the up- and downflows created by turbulent plumes through an adapted closure model. In a companion paper, we apply it to the formalism of excitation of solar $p$~modes developed by Samadi \\& Goupil (2001).} {Using results from 3D numerical simulations of the uppermost part of the solar convection zone, we show that the two-scale mass-flux model (TFM) is valid only for quasi-laminar or highly skewed flows (Gryanik \\& Hartmann 2002) and does not reproduce turbulent properties of the medium such as velocity-correlation products. We build a generalized two-scale mass-flux Model (GTFM) model that takes both the skew introduced by the presence of two flows \\emph{and} the effects of turbulence in each flow into account. In order to apply the GTFM to the solar case, we introduce the plume dynamics as modelled by Rieutord \\& Zahn (1995) and construct a closure model with plumes (CMP).} { The CMP enables expressing the third- and fourth-order correlation products in terms of second-order ones. When compared with 3D simulation results, the CMP improves the agreement for the fourth-order moments by a factor of two approximately compared with the use of the quasi-normal approximation or a skewness computed with the classical TFM.} { The asymmetry of turbulent convection in the solar case has an important impact on the vertical-velocity fourth-order moment, which has to be accounted for by models. The CMP is a significant improvement and is expected to improve the modelling of solar $p$-mode excitation. ", "introduction": "In the uppermost part of the solar convective zone, turbulent entropy fluctuations and motions of eddies drive acoustic oscillations. 3D numerical simulations of the stellar turbulent outer layers have been used to compute the excitation rates of solar-like oscillation modes \\cite{Stein01A}. As an alternative approach, semi-analytical modelling can provide an understanding of the physical processes involved in the excitation of $p$~modes: in this case, it is indeed rather easy to isolate the different physical mechanisms at work in the excitation process and to assess their effects. Various semi-analytical approaches have been developed by several authors \\citep{GK77,GK94,B92,Samadi00I}; they differ from each other by the nature of the assumed excitation sources, by the assumed simplifications and approximations, and also by the way the turbulent convection is described (see the review by \\citealt{Stein04}). Among the different theoretical approaches, that of \\cite{Samadi00I} includes a detailed treatment of turbulent convection, which enables us to investigate different assumptions about turbulent convection in the outer layers of stars \\citep{Samadi05c}. In this approach, the analytical expression for the acoustic power supplied to the $p$~modes involves fourth-order correlation functions of the turbulent Reynolds stress and the entropy source term, which for the sake of simplicity are expressed in terms of second-order moments by means of a closure model. The most commonly used closure model at the level of fourth-order moments (FOM) is the {\\it Quasi-Normal Approximation} (QNA), which is valid for a Gaussian probability distribution function \\citep[see][]{Lesieur97} and was first introduced by \\cite{Million41}. The QNA is rather simple and convenient to implement. However, \\cite{Ogura63} has shown that such a closure could lead to part of the kinetic energy spectrum becoming negative. In this paper, we confirm the results of \\cite{KR2006} (hereafter KR2006), namely that this approximation indeed provides a poor description of the physical processes involved in solar turbulent convection. Mass flux models (e.g., \\citealt{Randall92}, \\citealt{Ab}) explicitly take the effects of {\\it updrafts} and {\\it downdrafts} on the correlation products into account. The presence of two well-defined flow directions then introduces an additional contribution when averaging the fluctuating quantities, since averages of fluctuating quantities over each individual flow differ from averages over the total flow. For applications in atmospheric sciences, the mass-flux model for convection has recently been improved by \\citet[][hereafter GH2002]{GH2002}. Their motivation has been to account for the fact that horizontal scales of temperature and velocity fluctuations are different (hence their improvements lead to a `two-scale mass-flux model' (TFM)) as well as to understand and measure the effects of the skewness of their distribution. According to GH2002, mass-flux models, which also include the TFM, underestimate the FOM by as much as 70\\%. Therefore, such models clearly miss some important physical effects present in convective flows. \\cite{GH2002} and \\cite{GH2005} studied the asymptotic limits of TFM which led the authors to propose an interpolation between the QNA and the limit of large skewness provided by the TFM. This new parametrization permits a much better description of the FOM for convection in the atmosphere of the Earth (GH2002). We show that for their parametrization to be applicable to the case of solar convection, a more realistic estimate for the skewnesses of velocity and temperature fluctuations is required than that provided by the TFM itself (Sect.~\\ref{Sect_TFM}). The parametrization of GH2002 requires the knowledge of the skewnesses and second-order moments to compute FOM. These have to be provided either by measurements, by another model, or by numerical simulations. In the present paper we do not aim to construct a complete model to compute these quantities, which is the goal of the Reynolds stress approach (e.g., \\citealt{Canuto92}; \\citealt{Canuto98}). Rather, we aim to analyze the shortcomings of the TFM and suggest improvements using numerical simulations of solar convection as a guideline. The conclusions drawn from this analysis are used to derive a model for fourth-order moments in terms of second-order moments that can be used in computations of solar $p$-mode excitation rates. To proceed with the latter, we developed a formulation of the TFM that takes the effects of turbulence in each flow into account. This generalized TFM model (hereafter GTFM) is useful for both the superadiabatic and adiabatic outer solar layers. This formulation can actually be applied in other contexts than just the excitation of solar $p$~modes as long as the convective system is composed of two flows. The GTFM is more general and realistic than the TFM, but it requires the knowledge of additional properties of both the turbulent upwards and downwards flows. We choose to determine these properties by means of a plume model. Turbulent plumes are created at the upper boundary of the convection zone, where radiative cooling becomes dominant and where the flow reaches the stable atmosphere. In this region the updrafts become cooler and stop their ascent. This cooler flow is more dense than its environment and it triggers the formation of turbulent plumes \\citep{Stein98}. As shown by \\cite{RZ95}, these structures drive the dynamics of the flow; hence, to construct a closure model, we study the plume dynamics developed by \\cite{RZ95} (hereafter RZ95). This makes it possible to build a {\\em closure model with plumes (CMP)}, which is valid in the solar quasi-adiabatic convective region. In a companion paper \\citep[][hereafter Paper~II]{Belkacem06b}, we generalize this one-point correlation model to a two-points correlation model and calculate the power injected into solar $p$~modes. The paper is organised as follows: Sect.~2 introduces the TFM. Its validity is then tested with a 3D numerical simulation of the uppermost part of the solar convection region. In Sect.~3, we extend the TFM formulation (GTFM) in order to take into account turbulent properties of both upward and downward flows. We next investigate the asymptotic limits of the GTFM. In Sect.~4, we construct the CMP with the help of the RZ95 plume model. We test the validity of this model with results from the 3D simulation and show that the use of the plume model limits the validity of the CMP to the quasi-adiabatic zone. The CMP is then used to obtain analytical expressions for the third and fourth moments. Section~5 is dedicated to discussions and conclusions. ", "conclusions": "With the help of 3D numerical simulations of the upper part of the solar convective region, we have shown that the QNA and the TFM fail to describe the fourth-order velocity and temperature correlation moments, if merely used on their own. These results confirm KR2006 and geophysical studies \\citep{GH2002} and led us to generalize the TFM in order to take the effects of the turbulent properties of the up- and downflows explicitly into account (GTFM). We point out that the GTFM can be used in other contexts than the solar one as long as the convective system can be described with two turbulent flows. One might wonder whether it is likely that the CMP and the model for $p$~mode excitation developed in Paper~II are generally applicable to solar-like stars. To answer this question requires further work, but results on important ingredients of these models are encouraging. The case of convection in the planetary boundary layer of the atmosphere of the earth was already discussed in GH2002. Their interpolation model for FOMs has meanwhile been investigated for the case of convection in the ocean \\citep{Losch04} and solar granulation (\\citealt{KR2006}, who also study the case of a K~dwarf; preliminary results were published in \\citealt{Kupka05}). We corroborate the latter here with simulations for solar granulation based on more realistic boundary conditions. The overall conclusion that can be drawn from these studies is that, at least away from the boundary layers of convection zones, the FOMs in purely convective flows can be estimated according to the interpolation model by GH2002 with an accuracy typically in the range of 20\\% to 30\\%, whereas the QNA is off by a factor of two to three. For the superadiabatic layer, the discrepancies of the QNA remain the same in any case of the same size. We focused here on the solar case, more precisely a region that is nearly adiabatic, just below the superadiabatic zone where the acoustic modes are excited. As indicated by the 3D simulations, the coherent downdrafts, called plumes, are more turbulent than the upflow. In addition, we use the plume model developed by RZ95 to estimate the upward and downward mean velocities. With these additional approximations, the GTFM yields a closure model, the CMP, which can be applied in the quasi-adiabatic zone (located just below the superadiabatic one). Comparisons of calculations based on the CMP with direct calculations from the 3D numerical simulations show a good agreement. Hence, the CMP provides an analytical closure for third- and fourth-order moments. These moments are expressed in a simple way and require only the knowledge of the second-order moments and the parameters of the plume model. We stress that the CMP involves four parameters: the number of plumes in the considered shell (i.e., near the photosphere), the exponent of the power law for the mean vertical velocity of plumes, the law to describe the temperature difference between the two flows, and the mean fractional area of the updrafts and hot drafts. A study of the dependence of the results on these parameters is in progress. For instance, an increase of $a$ will imply an increase of $S_w$ in Eq.~(\\ref{Sw_eff}), and hence of the fourth-order moment $$. Nevertheless, it is extremely difficult to deduce the behaviour of the system, since from Eq.~(\\ref{conserv_mass}) a variation of $a$ changes the velocities of the flows. Instead, one could use a set of numerical simulations to study the effect of a change of the parameter $a$. In a companion paper, we use the CMP in a semi-analytical approach to calculate the power supplied to the solar $p$~modes. It is found that the power is quite significantly affected by the adopted closure model. Our final aim is to apply the CMP to the study of stochastic excitation of solar-like $p$~modes in stars other than the Sun. It will be necessary to assess the validity of the CMP approximations to extend their application to stellar conditions different from the solar case. This will also require investigating the dependence of the parameters entering the CMP, for instance, on the effective temperature of the star (work which is in progress). As pointed out in Sect.~\\ref{CMP}, the CMP is valid only in the quasi-adiabatic zone due to the power laws used to model the plume dynamics. This will be discussed further in the companion paper in which the present model will be used in the superadiabatic zone in order to propose a new closure for the calculation of stellar $p$~modes. Finally, we note that in the present work we do not take the effect of differential rotation and meridional circulation into account. However, recent helioseismic investigations \\citep{Schou02,Zhao04} have shown that variability of those large-scale flows gradually affects wavelength and frequencies, leading to a redistribution of the observed power spectrum \\citep{Sher05,Hindman05}. Hence, it could have an indirect effect on the amplitudes of $p$~modes. Furthermore, large-scale laminar non-uniform flows can have a significant effect on the formation of the coherent structures and intrinsic turbulence \\citep{Miesch00,Brun02,Rempel05}. To what extent they can affect solar $p$~mode amplitudes, through the closure model and the Reynolds stresses, remains to be investigated.\\\\" }, "0607/astro-ph0607042_arXiv.txt": { "abstract": "We investigate contributions to the extragalactic gamma-ray background (EGB) due to neutralino dark matter (DM) pair-annihilation into photons, from DM density enhancements (minispikes) surrounding intermediate-mass black holes (IMBHs). We focus on two IMBH formation scenarios; our conservative scenario where IMBHs are remnants of Population-III stars, and our optimistic scenario here IMBHs are formed in protogalactic disks. In both scenarios, their formation in pregalactic halos at high redshift lead to the formation of minispikes that are bright sources of gamma-ray photons. Taking into account minispike depletion processes, we only sum contributions from a cosmological distribution of IMBHs with maintained minispikes. Our conservative scenario (BH mass $10^2 M_\\odot$ with a $r^{-3/2}$ minispike) predicts gamma-ray fluxes that are an order larger than the equivalent flux, using the same DM parameters (mass $100\\,\\mathrm{GeV}$ and annihilation cross-section $3\\times10^{-26}\\,\\mathrm{cm^3\\,s^{-1}}$), from the host halo without IMBH minispikes. Our optimistic scenario (BH mass $10^5 M_\\odot$ with a $r^{-7/3}$ minispike) predicts fluxes that are three orders larger, that can reach current EGB observations taken by EGRET (DM parameters as above). This fact may serve interesting consequences for constraining DM parameters and elucidating the true nature of IMBHs. Additionally, we determine the spectra of DM annihilation into monochromatic gamma-rays, and show that its flux can be within observational range of GLAST, providing a potential `smoking-gun' signature of DM. ", "introduction": "Introduction} Despite compelling indirect evidence, from galactic to cosmological scales, the fundamental nature of the dominant non-baryonic component in the matter density of the universe (dark matter, hereafter DM) remains unknown. Intriguingly, extended models of particle physics independently provide us with a host of particle candidates for this as yet unknown matter, of which the most popular is the supersymmetric neutralino (see reviews \\cite{JungmanKamionGriest,Bergstrom, BertoneHooperSilk} for details). Upgrades of underground direct detectors looking for scattering of DM particles from nuclei, together with future neutrino, antimatter, and gamma-ray detectors looking for products of DM annihilation, will dramatically enhance our chances of understanding the true nature of DM. In particular, the forthcoming launch of the Gamma Ray Large Area Space Telescope (GLAST) \\cite{GLAST} and numerous ground based Atmospheric Cerenkov Telescopes make indirect gamma-ray search especially promising. Since the DM annihilation rate scales as the DM density squared, there is great advantage in observing areas where the DM density is believed to be high. The galactic centre (GC) is the immediate choice, and indeed strong gamma-ray emission has been observed and its nature and origin have been investigated by many researchers \\cite{Bengtsson,Berezinsky, BergstromUllio,BergstromUllioBuckley,BergstromEdsjoGunnarsson,Cesarini_etal,Hooper_etal,Fornengo_etal,Horns}. However, the DM density in the GC is highly uncertain, making accurate predictions difficult. For example, DM enhancements called `spikes' can form during the formation of a central supermassive-BH (SMBH) \\cite{GondoloSilk}, but it can also be depleted by various processes by varying degrees \\cite{UllioZhaoKamion,Merritt,Merritt_etal,BertoneMerritt}. In addition, nearby astrophysical (non-DM) gamma-ray sources make a potential DM detection impossible for all but a narrow range of DM parameters \\cite{ZaharijasHooper}. Intermediate-mass black holes (IMBHs, see e.g.~\\cite{MillerColbert}) provide an alternative source that may work positively for DM detection. Bertone et al.~\\cite{BertoneZentnerSilk} recently investigated the possibility of detecting a `smoking gun' gamma-ray signature of DM using IMBHs in the Milky-Way as point sources. They showed that IMBH formation increases the DM density in its vicinity to produce a `minispike', and also that DM enhancement depletion processes are generally less significant for IMBHs due to their roughly spherical distribution about the GC. They conclude that under optimistic circumstances, the Energetic Gamma Ray Experimental Telescope (EGRET) may have already seen a few of the IMBH minispikes as unidentified sources. Another avenue of indirect DM search is via the extragalactic gamma-ray background (hereafter EGB) measured over a wide energy range \\cite{Sreekumar_etal,StrongMoskalenkoReimer}. The origin of this background is currently unknown, and it has been speculated that DM annihilation gamma-rays from cosmological distributions of DM contribute to some degree \\cite{ BergstromEdsjoUllio,Ullio,TaylorSilk,ElsasserMannheim,ElsasserMannheim2,Ando,OdaTotaniNagashima,AndoKomatsu}. Since the DM annihilation cross-section is so small, a consideration of DM enhancements is crucial for meaningful gamma-ray flux predictions. A popular DM enhancement is those at the centres of galactic DM halos. However, Ando \\cite{Ando} has recently shown that they are strongly constrained by observations of our galaxy. The author assumes universality of galactic DM halo profiles, and shows that DM annihilation cannot significantly contribute to the EGB without exceeding gamma-ray observations from our GC \\cite{Ando}. The author also points out that this constraint could be loosened when one takes DM substructures into account. In this paper we argue IMBHs minispikes as a \\emph{substructure in the DM halo, that can lead to enhancements that do not conflict with current observations of our galaxy}. We determine contributions to the EGB from IMBH minispikes by summing gamma-ray fluxes from all redshifts. IMBH minispikes are not expected to greatly suffer from depletion processes, but we do take into account BH-BH mergers, which are known to strongly deplete minispikes and do occur in IMBHs. We also consider a conservative case ($10^2 M_\\odot$ BHs of Population-III origin \\cite{scenarioA} with a $r^{-3/2}$ minispike) and an optimistic case ($\\sim 10^5 M_\\odot$ BHs formed in the centres of protogalactic disks \\cite{ scenarioB} with a $r^{-7/3}$ minispike). Our result is that contributions to the EGB are increased by $1-3$ orders in magnitude. In particular, our optimistic case predicts fluxes that can reach current EGB observations. As this has interesting implications for constraining DM parameters and IMBH scenarios, we critically assess uncertainties in our calculation. We then determine the flux of DM annihilation into line gamma-rays, and show that under optimistic conditions, it is observable by GLAST. This provides a potential `smoking-gun' signature of DM. This paper is structured as follows. In Sec.~\\ref{sec:imbh} we introduce IMBHs, starting with their existence, followed by their formation scenarios, and finishing off with a summary of recent numerical studies. Then in Sec.~\\ref{sec:formulations} we develop our calculation frameworks, first for the EGB, followed by DM annihilation, then minispike formation, moving finally on to our IMBH number density fitting. Calculations and results are in Sec.~\\ref{sec:calc}, and discussions and conclusions in Sec.~\\ref{sec:conc}. In all our calculations we adopt the standard flat cosmological constant plus cold DM ($\\Lambda$CDM) cosmology, with $\\Omega_M = 0.3$, $\\Omega_\\Lambda = 0.7$, $h = 0.7$, and $\\sigma_8=0.9$. ", "conclusions": "" }, "0607/astro-ph0607274_arXiv.txt": { "abstract": "We explore the environment of $z \\approx 1$ AGN using a sample of 53 spectroscopically identified X-ray sources in the All-wavelength Extended Groth strip International Survey. We quantify the local density in the vicinity of an X-ray source by measuring the projected surface density of spectroscopically identified optical galaxies within a radius defined by the 3rd nearest neighbour. Our main result is that X-ray selected AGN at $z \\approx 1$ avoid underdense regions at the 99.89\\% confidence level. Moreover, although we find that the overall population shares the same (rich) environment with optical galaxies of the similar $U-B$ and $M_B$, there is also tentative evidence (96\\%) that AGN with blue colors ($U-B \\la 1$) reside in denser environments compared to optical galaxies. We argue that the results above are a consequence of the whereabouts of massive galaxies, capable of hosting supermassive black holes at their centers, with available cold gas reservoirs, the fuel for AGN activity. At $z\\approx1$ an increasing fraction of such systems are found in dense regions. ", "introduction": "In recent years there has been increasing evidence that the formation of spheroids and the build-up of supermassive black holes at their centers are strongly interconnected (Ferrarese \\& Merritt 2000; Gebhardt et al. 2000; Alexander et al. 2005). Moreover, it is now well established that galaxy properties, such as morphology, color and star-formation, strongly depend on environment (e.g. Butcher \\& Oemler 1978; Lewis et al. 2002; Gomez et al. 2003; Hogg et al. 2004), suggesting a close link between local density and the evolution of individual systems. Putting the evidence above together, it is natural to assert that AGN activity, being strongly coupled to galaxy formation and evolution, should also depend on environment. Despite significant observational progress however, the link between local density and AGN remains controversial. For example, at low redshift ($z\\approx0.1$) Miller et al. (2003) found no dependence on environment of the fraction of spectroscopically identified AGN in the Sloan Digital Sky Survey (SDSS; Schneider et. al. 2005). More recent studies, also using SDSS data, suggest that it is only when the AGN population is split into subsamples based on optical classification and/or luminosity, that environmental differences become apparent. For example powerful AGN ($L[{\\rm O\\,III}] > 10^{7} \\, L_{\\odot}$) and/or narrow-line Seyferts are found in increasingly {\\it less} dense regions at $z \\approx 0.1$, while less luminous AGN and/or LINERs show no dependence on local density (e.g. Kauffmann et al. 2004; Wake et al. 2005; Constantin \\& Vogeley 2006). The above low-$z$ results however, appear to be in conflict with observations suggesting that at least certain classes of powerful AGN, such as radio galaxies, reside in relatively rich environments (e.g. Zirbel 1997). Additionally, the large scale distribution of optically and X-ray selected AGN at $z\\ga 1$ is consistent with correlation lengths in the range $r_0 = 5 - 10 \\, h^{-1} \\rm \\, Mpc$ (e.g. Croom et al. 2005; Basilakos et al. 2004; Gilli et al. 2005; Adelberger \\& Steidel 2005). This indicates that AGN have local density distribution similar to early-type systems at $z \\approx 1$ ($r_0 \\approx \\rm 6.6 \\,h^{-1}\\,Mpc$; Coil et al. 2004b) and that they avoid poor environments at these redshifts (e.g. emission-line galaxies, $r_0 \\approx \\rm 3.2 \\,h^{-1}\\,Mpc$; Coil et al. 2004b). Contrary to these results, Coil et al. (2006) show that the clustering amplitude of broad-line QSOs at $0.7 5 \\times 10^{43} \\rm \\, erg \\, s^{-1}$ and therefore may be dominated by AGN emission at optical wavebands. Alternatively the above tentative trend with color may suggest that for these systems environment plays a role in the observed activity. A larger X-ray sample is required to further explore this. The SDSS spectroscopic survey currently provides the only direct estimate of the environment of AGN, albeit at much lower redshift, $z\\approx0.1$, compared to this study. Kauffmann et al. (2004) using the 1st release of the SDSS found a strong dependence on environment for powerful AGN ($L {\\rm [OIII]} > 10^7 L_{\\odot}$) at $z<0.1$, in the sense that the most luminous systems are found in the field. Lower luminosity AGN in this study show no dependence on local density. The results above on the environment of low-$z$ AGN are clearly different from those reported here. Kauffmann et al. (2003) also explored the host galaxy properties of powerful AGN in the SDSS. They find that they are associated with massive ($ M \\ga 10^{10}\\, M_{\\odot}$) early type galaxies which are however, distinct from the bulk of the optically selected early-type population, in that they show evidence for on-going or recent star-formation activity. Based on the evidence above Kauffmann et al. (2003) suggest that there are two essential ingredients for strong AGN activity: a massive central black hole and abundant gas supply to fuel it. Only massive early-type galaxies have large enough bulges to host massive black holes. From these early-type galaxies those which show evidence for substantial amounts of young stellar populations have sufficient gas supply to both produce young stars and to feed the central engine. Such galaxies are relatively rare in the present day Universe and are preferentially found in low-density regions. Only in these environments can gas-rich star-forming galaxies survive today (e.g. Kauffman et al. 2004; Gomez et al. 2003; Lewis et al. 2002; Poggianti et al. 2006). At higher redshifts however, there is an increasing fraction of massive galaxies with sufficient cold gas reservoirs that can potentially produce young stars and also fuel luminous AGN. For example, there is accumulating evidence that the number density of luminous blue galaxies, which are rare locally (e.g. Kauffman et al. 2004), substantially increases to $z\\approx1$ (Bell et al. 2004; Cooper et al. 2006). These systems have blue colors, most likely because of star-formation, while their luminosities suggest high stellar masses. These galaxies are clearly prime candidates for powerful AGN hosts. Interestingly, these luminous blue galaxies, contrary to their low-$z$ counterparts, are shown to reside in regions of enhanced density (Cooper et al. 2006). As demonstrated in Fig. \\ref{hist_con_den_color.fig} our X-ray selected AGN sample is likely to include such systems. Moreover, the bluer X-ray sources in the sample appear to reside in higher density regions compared to optical galaxies on average. The evidence above suggests that our finding for an association between X-ray selected AGN and higher density regions at $z \\approx 1$ is related to the whereabouts of massive galaxies with available cold gas reservoirs to sustain accretion of material on the central black hole. Observations show that such systems are also found in denser environments at $z\\approx1$, contrary to the local Universe. A central question in AGN studies is the triggering mechanism of the observed activity. Recent results suggest that mergers are not the main process for activating a supermassive black hole (Grogin et al. 2005; Pierce et al. 2006). Is it possible then, that the environment plays a role in triggering the AGN activity suggesting a causal link between the two? The enhanced density of the bluer X-ray sources in our sample compared to optical galaxies, if confirmed with larger samples, supports such an association. For example it is plausible that the central black hole becomes active during the infall of the host galaxy to the overdense region, as it experiences the gravitational potential of the structure or frequent interactions with other galaxies in that region. In the framework of hierarchical models massive galaxies, which can potentially host luminous AGN, are more often found in denser regions. At $z \\approx 1$, richer environments were substantially more active compared to the local Universe (Bundy et al. 2006), with the higher contrast occurring for intermediate-mass group-like systems with velocity dispersions $\\sigma \\approx 500-600 \\rm \\, km \\, s^{-1}$ (Poggianti et al. 2006), close to the upper mass limit that is well sampled by the DEEP2 (Gerke et al. 2005). According to Poggianti et al., this enhanced activity is associated with infalling systems many of which are likely to have sufficient mass to harbour a central black hole and cold-gas reservoirs to sustain star-formation activity and possibly also feed the black hole. As these massive galaxies further grow (e.g. by mergers) from high redshift to the present day, the supply of cold gas is cut-off (e.g. Croton et al. 2005; Cooper et al. 2006) leading to quenching of the star-formation and possibly the AGN activity. These massive systems will therefore, appear as red-and-dead at low-$z$. A possible link between AGN triggering and environment, if confirmed, would suggest that the enhanced local density in the vicinity of X-ray sources at $z \\approx 1$ is a consequence of the hierarchical evolution of structures in the Universe and its impact on the cold gas reservoirs of individual galaxies. The AEGIS survey provides a unique dataset that can potentially test this scenario. Open questions include: what fraction of the X-ray population is found in optically selected groups (e.g. Gerke et al. 2005)? Are there morphological and/or color gradients with distance from the centre of the group, indicating enhanced activity for infalling members? Are there differences in the color, stellar mass and/or X-ray luminosity between AGN locked in groups and those that are not? Addressing these points requires a larger X-ray sample than that presented here. A full discussion of these issues is therefore referred to a future publication using the full AEGIS X-ray sample. \\\\ \\\\ Financial support has been provided through PPARC and the Marie-Curie Fellowship grant MEIF-CT-2005-025108 (AG) the Leverhulme trust (KN), the Hubble Fellowship grants HF-01165.01-A (JAN) and HF-01182.01-A (ALC), the NSF grants AST00-71198 and AST0071048. The W.M. Keck Observatory, a scientific partnership among Caltech, the University of California and NASA. The Observatory was made possible by the generous financial support of the W.M. Keck Foundation. The authors wish to acknowledge the very significant cultural role that the summit of Mauna Kea has within the indigenous Hawaiian community; we are fortunate to be able to conduct observations from this mountain." }, "0607/astro-ph0607260_arXiv.txt": { "abstract": "{We have analyzed 5.5 years of timing observations of 7 ``slowly'' rotating radio pulsars, made with the Westerbork Synthesis Radio Telescope. We present improved timing solutions and 30, mostly small, new glitches. Particularly interesting are our results on PSR~J1814$-$1744, which is one of the pulsars with similar rotation parameters and magnetic field strength to the Anomalous X-ray Pulsars (AXPs). Although the high-B radio pulsars do not show X-ray emission, and no radio emission is detected for AXPs, the roughly similar glitch parameters provide us with another tool to compare these classes of neutron stars. Furthermore, we were able to detect glitches one to two orders of magnitude smaller than before, for example in our well-sampled observations of PSR~B0355$+$54. We double the total number of known glitches in PSR~B1737$-$30, and improve statistics on glitch sizes for this pulsar individually and pulsars in general. We detect no significant variations in dispersion measure for PSRs B1951$+$32 and B2224$+$65, two pulsars located in high-density surroundings. We discuss the effect of small glitches on timing noise, and show it is possible to resolve timing-noise looking structures in the residuals of PSR~B1951$+$32 by using a set of small glitches. ", "introduction": "Two sorts of irregularities, in the otherwise very stable pulsar rotation rates exist, which limit the accuracy to which pulse arrival times can be measured: timing noise and glitches.\\\\ Timing noise is seen as random fluctuations in the rotation rate of the pulsar on timescales of days to years. It is largest in young pulsars and pulsars with large period derivatives \\citep{ch80,antt94}. Glitches on the other hand are characterized by a sudden increase of the pulsar rotation frequency $(\\nu)$, accompanied by a change in spindown rate $(\\dot\\nu)$ and sometimes followed by relaxation or exponential decay to the previous rotation state. Typical magnitudes of glitches are from $10^{-10}\\nu$ to $10^{-6}\\nu$ and steps in slowdown rate are on the order of $10^{-3}\\dot\\nu$. Glitches give an unique opportunity to study the internal structure of neutron stars, as they are believed to be caused by sudden and irregular transfer of angular momentum from the superfluid inner parts of the star to the more slowly rotating crust \\citep{rzc98}. They are mostly seen in pulsars with characteristic ages ($\\tau_c$) around $10^4-10^5$ yr and can occur up to yearly in some pulsars. Most of the youngest pulsars, with $\\tau_c\\lesssim2000$~yr show very little glitch activity. This could be because they are still too hot which allows the transfer of angular momentum to happen more smoothly \\citep{ml90}. A lot of new glitches have been found in the last decades \\citep{lsg00,uo99,wmp+00}. This makes statistical analysis of glitch sizes and activity possible. However, no clear relations between glitch parameters or dependence of glitch parameters on rotation properties have been found so far. Extending the sample of glitches, especially with previously unknown small-magnitude glitches will lead to a better insight into the glitch mechanism and the structure of the neutron star. ", "conclusions": "\\subsection{Glitch sizes} Little is known about the range of glitch sizes and their distribution within this range. \\cite{lsg00} showed a distribution for 48 glitches in 18 pulsars. They claimed there was a peak around $\\Delta\\nu/\\nu\\approx10^{-6.5}$, but they also noted that the glitch sample could be incomplete at the lowest level. How small do we expect glitches to be? Glitches are believed to be the result of the unpinning of neutron superfluid vortices which causes transfer of angular momentum of the interior superfluid to the solid crust of the star. There is no report of restrictions on a minimum stress needed before vortices unpin, so it is not unlikely that the observed lower limit on glitch sizes is mainly due to the detection limits of the observing systems and the timing precision achievable for a given pulsar. In this paper we measure glitch sizes down to $\\Delta\\nu/\\nu=10^{-11}$, which provides the first evidence that such small glitches occur and can be measured in slowly rotating pulsars. These glitches are then of a similar size to the one reported by \\cite{cb04} in a millisecond pulsar, and thus perhaps provide further evidence for a continuous distribution of glitch sizes. Let us now consider how these small glitches affect the observed glitch size distribution. In Fig. \\ref{fig:glitches2}, a histogram is shown for all now known glitch sizes. New glitches found in this study are shown added on top of the old glitch distribution. A Kolmogorov-Smirnov test shows that over the whole range, the distribution has only a probability of $0.001$\\% to be consistent with a flat distribution in log space of glitch sizes. But if we consider only the part of the diagram between $10^{-9}~<~\\Delta\\nu/\\nu~<~10^{-5.5}$, where the statistics are better, another KS test shows that the distribution has a $29.8$\\% chance to be drawn from a flat distribution. The increased number of glitches with sizes around $\\Delta\\nu/\\nu\\approx10^{-9}$, now comparable to the amount of larger glitches observed, suggests again that the lack of the smallest glitches at the lower end of the distribution is due to observing limits. The lack of glitches at the upper end of the distribution can not be due to observing limits. Apparently there is some physical restriction to the maximum size of a glitch, and we can consider the boundary of $\\Delta\\nu/\\nu\\approx10^{-5}$ as the natural upper limit of glitch sizes. The glitch size distribution may be biased because there is a whole range of different pulsars included, which can all show different glitching behaviour. For example the Vela pulsar suffers from a large number of huge glitches. The fact that those are more easily detectable, will have a large influence on the overall glitch size distribution. We take a frequently glitching pulsar, PSR~B1737$-$30 and make a similar histogram as before to see what happens if we study the glitch size distribution for one pulsar. The result is shown in the bottom panel of Fig. \\ref{fig:glitches2}. As we now have almost doubled the number of glitches detected for this pulsar, and we measure glitches an order of magnitude smaller than before, we can do a statistical analysis of the glitch size distribution for a single pulsar with glitches covering almost the whole range of sizes observed. Using a Kolmogorov-Smirnov test, we calculated that in the observed size range there is a $90.2$\\% probability that the glitches of PSR~B1737$-$30 are drawn from a flat distribution (in log space) of glitch sizes. The greater glitch rate of PSR~B1737$-$30 allows us to obtain sufficient glitches to show that the distribution of glitches in this pulsar are likely drawn from a flat log-space distribution of glitch sizes. Whether this is representative of the glitch size distribution of pulsars in general is more difficult to determine based on just this pulsar and a larger sample of pulsars where such fits can be determined is required. \\addtocounter{footnote}{1} \\begin{figure} \\centering \\includegraphics[width=6cm,angle=270]{janssenfig6a.ps} \\includegraphics[width=6cm,angle=270]{janssenfig6b.ps} \\caption[]{Upper panel: histogram of all known glitch sizes, from the ATNF glitch table\\addtocounter{footnote}{-1}\\footnotemark\\, \\citep{mhth05} and \\cite{uo99}. New glitches found in this study are added on top of the known glitches. The lower panel shows the glitches of PSR~B1737$-$30. \\label{fig:glitches2}} \\end{figure} \\footnotetext{http://www.atnf.csiro.au/research/pulsar/psrcat/glitchTbl.html} \\subsection{Glitch activity} An indication of glitch activity was first introduced by \\cite{ml90} as the mean fractional change in the rotation frequency per year due to glitches: \\begin{equation} A_g=\\frac{1}{t_g}\\sum\\frac{\\Delta\\nu}{\\nu}, \\end{equation} where $\\Sigma \\frac{\\Delta\\nu}{\\nu}$ is the sum of all glitches occurring in the observed time span $t_g$, which is measured in years. They found that the glitch activity is highest for pulsars with low characteric ages and high frequency derivatives. Most young pulsars ($\\tau_c \\lesssim 2000$ yr) show no glitch activity, which is believed to be an effect of higher temperatures reducing the effect of vortex pinning. Apparently for older pulsars, the glitch activity is higher for pulsars with higher spindown rates (e.g. \\cite{wmp+00} and \\cite{uo99}). Our values for glitch activity for pulsars B0740$-$28 and B1737$-$30 are the same as given by \\cite{uo99}, see Table~\\ref{tab:parameters}. However, due to the four mini-glitches in PSR~B0355$+$54 and no large glitches in our data span, we calculate a much lower value for the glitch activity of this pulsar than the previously published value including the large glitch seen in this pulsar. Although we find quite a large number of small glitches in the older pulsars of our sample, the derived glitch activity is still low. \\subsection{Glitch vs timing noise} Apart from glitches, irregularities in the rotation of the pulsar are usually described as timing noise. Like glitches, timing noise is also seen mostly in the younger pulsars with high spin frequency derivatives. Timing noise is seen as random wanderings in the rotation rate and the timescales range from months to years. Relaxations after glitches are also supposed to be on those timescales. As we are now able to measure smaller and smaller glitches, although not yet including relaxations, it is important to investigate whether small glitches can mimic timing noise. \\cite{antt94} use the parameter $\\Delta_8$ to quantify the amount of timing noise in a pulsar. They define the parameter as \\begin{equation} \\Delta(t) = \\log \\left( \\frac{1}{6\\nu}|\\ddot\\nu| t^3 \\right), \\end{equation} which is usually quoted as $\\Delta_8$ using a standard time interval of $10^8$ seconds. They note that values of $\\Delta_8$ are expected to vary for non-overlapping data sets. Therefore it is difficult to make a statement about the influence of small glitches on this parameter. What we can do, is compare the value of this parameter for our own data sets, using solutions with glitches and without glitches. If the parameter stays the same, we can conclude that it is not influenced by glitches and the glitches are a different phenomenon than timing noise. Unfortunately, the variation in $\\Delta_8$ is so large, even for adjacent or partially overlapping intervals of $10^8$ seconds in our data, that it is impossible to draw a conclusion from the values. To make a better distinction, if possible, between timing noise and glitches, more modelling is needed, both on the expected glitch size distributions, as well as on the exact influence on timing parameters of small glitches and recoveries from large glitches. We have seen that for frequently glitching pulsars, it can be difficult to resolve glitches that occur close together in time. This effect is probably more important for small glitches, as they appear to occur more often and thus are more likely to merge together. There are many manifestations of timing noise and some have a form which clearly cannot be explained as being due to glitches. However our discovery of small glitches and the way in which we were able to improve a ''timing-noise-like'' set of residuals for PSR B1951$+$32 by including glitches in the solution indicates that they may play a role, and that improved sensitivity and more frequent observations may be required to find more such instances. \\begin{table}% \\caption{Glitch parameters for the 30 new glitches found in this study. Where no error in the glitch epoch is quoted, it was not possible to constrain the epoch by using the extra glitch parameter GLPH. In these cases the middle of the interval between adjacent datapoints is quoted as the glitch epoch. \\label{tab:glitch}}% \\begin{tabular}{llll}% \\hline \\hline Pulsar name & Epoch (MJD)& $\\Delta\\nu/\\nu (10^{-9})$ & $\\Delta\\dot\\nu/\\dot\\nu (10^{-3})$ \\\\ \\hline B0355$+$54 & 51673(15) & \\phantom{00}0.04(2) & \\\\ & 51965(14) & \\phantom{00}0.030(2)& \\phantom{00}$-$0.102(7) \\\\ & 52941(9) & \\phantom{00}0.04(1) & \\phantom{$-$00}0.13(4) \\\\ & 53216(11) & \\phantom{00}0.10(2) & \\phantom{00}$-$0.03(4) \\\\ B0525$+$21 & 52289(11) & \\phantom{00}1.46(5) & \\phantom{$-$00}0.6(1) \\\\ & 53379 & \\phantom{00}0.17(5) & \\\\ B0740$-$28 & 51770(20) & \\phantom{00}1.0(3) & \\phantom{$-$00}0.9(2) \\\\ & 52027(5) & \\phantom{00}2.1(2) & \\phantom{00}$-$1.1(2) \\\\ & 53090.2(2.6) & \\phantom{00}2.9(1) & \\phantom{$-$00}0.39(3) \\\\ & 53469.7(8.1) & \\phantom{00}1.1(2) & \\\\ B1737$-$30 & 51685(21) & \\phantom{00}0.7(4) & \\phantom{$-$00}0.09(7)\\\\ & 51822(7) & \\phantom{00}0.8(3) & \\phantom{00}$-$0.07(7)\\\\ & 52007(6) & \\phantom{00}0.7(1) & \\phantom{00}$-$0.08(2)\\\\ & 52235 & \\phantom{0}42.1(9) & \\\\%124(2)\\\\ & 52271 & 444(5) & \\\\%-75(1)\\\\ & 52344 & 220.6(9) & \\\\%-49.4(6) \\\\ & 52603(5) & \\phantom{00}1.5(1) & \\phantom{00}$-$0.60(3)\\\\ & 52759(5) & \\phantom{00}1.6(3) & \\phantom{00}$-$0.59(8)\\\\ & 52859 & \\phantom{0}17.6(3) & \\phantom{$-$00}0.9(1)\\\\ & 52943.5 & \\phantom{0}22.1(4) & \\\\ J1814$-$1744 & 51700(16) & \\phantom{00}5(2) & \\\\ & 52117(6) & \\phantom{0}33(2)& \\phantom{00}$-$0.5(4)\\\\ & 53302(22) & \\phantom{00}7(2) & \\phantom{00$-$}2(1) \\\\ B1951$+$32 & 51967(9) & \\phantom{00}2.25(9) & \\phantom{00}$-$0.2(1) \\\\ & 52385(11) & \\phantom{00}0.72(9) & \\phantom{00}$-$0.04(8)\\\\ & 52912(5) & \\phantom{00}1.29(7) & \\phantom{00$-$}0.30(9) \\\\ & 53305(6) & \\phantom{00}0.51(9) & \\phantom{00$-$}0.11(7) \\\\ B2224$+$65 & 51900 & \\phantom{00}0.14(3) & \\phantom{00}$-$2.9(2) \\\\ & 52950& \\phantom{00}0.08(4) & \\phantom{00}$-$1.4(2) \\\\ & 53434(13)& \\phantom{00}0.19(6) & \\\\ \\hline \\end{tabular} \\end{table} \\begin{table} \\caption{Glitches in AXPs and high-B radio pulsars. The differences in $\\Delta\\dot\\nu/\\dot\\nu$ for the second glitch in 1RXS~J1708$-$4009 are due to a different time scale fitted in the papers. \\cite{dis+03} fit short term change, \\cite{kg03} long-term. See Fig.1b of \\cite{kg03}.\\newline References: 1. \\cite{dis+03}, 2. \\cite{kg03} 3.\\cite{kgw+03}, 4. \\cite{ckl+00}, 5. this paper. \\label{tab:axp}} \\begin{tabular}{lllll} \\hline \\hline AXP name & Epoch (MJD) &$\\Delta\\nu/\\nu~(10^{-9})$ & $\\Delta\\dot\\nu/\\dot\\nu~(10^{-3})$& ref.\\\\ \\hline 1RXS & 51459.0 & \\phantom{0}643(5) &\\phantom{$-$00}17.2(5) &1\\\\ J1708$-$4009 & 51444.601 & \\phantom{0}549(20) & \\phantom{$-$00}10.0(4) & 2\\\\ & 52015.65& \\phantom{0}330(20) &\\phantom{$-$0}330(30)$^a$ &1\\\\ & 52014.177 & \\phantom{0}141(27) & \\phantom{$-$000}0.1(4)$^a$ & 2\\\\ \\hline 1E 2259$+$586 &52443.9& 4100(30) & \\phantom{$-$}1110(70) & 3\\\\ \\hline\\hline Pulsar name & Epoch (MJD) &$\\Delta\\nu/\\nu~(10^{-9})$ & $\\Delta\\dot\\nu/\\dot\\nu~(10^{-3})$& ref.\\\\ \\hline J1119$-$6127 & 51398 &\\phantom{000}4.4(4) &\\phantom{$-$000}0.039(5) & 4\\\\ \\hline J1814$-$1744 & 51700 & \\phantom{000}5(2) & &5 \\\\ & 52117 & \\phantom{00}33(2) &\\phantom{000}$-$0.5(4) &5 \\\\ & 53302 &\\phantom{000}7(2) &\\phantom{$-$000}2(1) & 5\\\\ \\hline \\end{tabular} \\end{table}" }, "0607/astro-ph0607056_arXiv.txt": { "abstract": "We present high-speed, three-colour photometry of the eclipsing cataclysmic variable SDSS J170213.26+322954.1 (hereafter SDSS J1702+3229). This system has an orbital period of 2.4 hours, placing it within the ``period gap'' for cataclysmic variables. We determine the system parameters via a parameterized model of the eclipse fitted to the observed light curve by $\\chi^2$ minimization. We obtain a mass ratio of $q = 0.215 \\pm 0.015$ and an orbital inclination $i = 82^{\\circ}.4 \\pm 0^{\\circ}.4$. The primary mass is $M_{\\rmn{w}} = 0.94\\pm0.01 M_{\\sun}$. The secondary mass and radius are found to be $M_{\\rmn{r}} =0.20\\pm0.01 M_{\\sun}$ and $R_{\\rmn{r}} = 0.243 \\pm 0.013 R_{\\sun}$ respectively. We find a distance to the system of $440 \\pm 30$\\, pc, and an effective temperature for the secondary star of $3800\\pm100$K (corresponding to a spectral type of M0$\\pm$0.5V). Both the distance and effective temperature are consistent with previous values derived via spectroscopy of the red star. The secondary star is significantly less massive than expected for the orbital period, and significantly warmer than expected for its mass. This can be explained if the secondary star is significantly evolved: the mass and effective temperature are consistent with a secondary star that began mass transfer with a greatly reduced central hydrogen fraction. The nature of the secondary star in SDSS J1702+3229 supports predictions that CVs with evolved secondary stars might be found accreting within the period gap. ", "introduction": "\\label{sec:introduction} Cataclysmic variable stars (CVs) are a class of interacting binary system undergoing mass transfer via a gas stream and accretion disc from a Roche-lobe filling secondary to a white dwarf primary. A bright spot is formed at the intersection of the disc and gas stream, giving rise to an `orbital hump' in the light curve at phases $0.6-1.0$ due to foreshortening of the bright-spot. \\citet{warner95a} gives a comprehensive review of CVs. The light curves of eclipsing CVs can be quite complex, with the accretion disc, white dwarf and bright-spot all being eclipsed in rapid succession. With sufficient time-resolution, however, this eclipse structure allows the system parameters to be determined to a high degree of accuracy \\citep{wood86a}. SDSS J1702+3229 is a deeply eclipsing CV, first discovered through the Sloan digital sky survey \\citep{szkody04}. The spectrum is highly suggestive of a dwarf-nova type system, a fact confirmed by its recent outburst during which it exhibited 0.3 mag superhumps (VSNET alert 8715), placing it amongst the SU UMa sub-class of Dwarf Novae. The system is particularly worthy of study, as its orbital period of 2.4 hours places it squarely within the period gap. Furthermore, its deeply eclipsing nature allows accurate system parameters to be derived. As such, SDSS J1702+3229 constitutes an excellent test of evolutionary models for CVs, which predict secondary star masses and radii for systems within the gap. In this paper we present {\\sc ultracam} $u'g'r'$ lightcurves of SDSS J1702+3229, and use these lightcurves to derive the system parameters. The observations are described in section~\\ref{sec:obs}, the results are presented in section~\\ref{sec:results}, and discussed in section~\\ref{sec:disc}. ", "conclusions": "\\label{sec:disc} \\subsection{Superhump period excess and mass ratio} \\label{subsec:superhumps} On the 3$^{rd}$ Oct 2005, SDSS J1702+3229 entered its first recorded outburst (VSNET Alert 8709); On Oct 7$^{th}$, 0.3 mag superhumps were observed, with a superhump period $P_{sh}$ of 0.1056 days (VSNET Alert 8715). The detection of superhumps in this object is significant in the context of the well-established relationship between mass ratio and superhump period excess, as very few systems are available to calibrate the relationship at high mass ratios \\citep{patterson05}. From the reported superhump period and the orbital period reported in this paper we calculate a superhump period excess, $\\epsilon = 0.0551 \\pm 0.0005$. This is entirely consistent with the $\\epsilon$--$q$ relationship established by \\cite{patterson05} which predicts $\\epsilon = 0.052 \\pm 0.005$ for the mass ratio of SDSS J1702+3229. \\subsection{Evolutionary status of the secondary star} \\label{subsec:evolved} Several independent theoretical studies predict the existence of a population of CVs with substantially evolved secondary stars \\citep[e.g][]{baraffe00,andronov04,schenker02,podsiadlowski03}. Under the disrupted magnetic braking model of CV evolution, the period gap is caused by a cessation of magnetic braking which occurs when the secondary star becomes fully convective. The small central hydrogen abundance and higher central temperature of evolved stars imply lower radiative opacities; this naturally favours radiative transport. Sufficiently evolved secondary stars do not, therefore, become fully convective until they have evolved to shorter periods than the upper edge of the period gap. For sufficiently evolved secondary stars, the system can continue accreting throughout the gap \\citep{baraffe00}. We might therefore expect that a sizeable proportion of CVs with periods inside the period gap would possess evolved secondary stars. Evidence that the secondary star in SDSS J1702+3229 has undergone significant nuclear evolution comes from two independent lines of argument; the secondary star is both too warm and insufficiently massive to be an ordinary star, given its orbital period. In this paper, we find an effective temperature for the secondary star of $3800\\pm100$\\,K, which agrees well with the reported spectral type of M1.5$\\pm$1, or $3600\\pm200$\\,K \\citep{szkody04}, derived from the TiO band spectral index derived by \\cite{reid95}. An effective temperature of 3800\\,K is too warm for a non-evolved CV secondary star at an orbital period of 2.4 hours \\citep{baraffe00}. At periods of around 3 hours, a typical spectral type is approximately M4 (c.\\,f.\\, U Gem and IP Peg), whereas a spectral type of M1 is more typical of periods near 5 hours \\cite[see][for examples]{sad98}. On the other hand, the early spectral type {\\em is} consistent with a secondary star which has undergone significant nuclear evolution. From figure~3 of \\cite{baraffe00}, it is clear that the effective temperature of the secondary star in SDSS J1702+3229 is consistent with models of a secondary star that began mass transfer with a greatly reduced central hydrogen fraction. The mass of the secondary star at a given orbital period is also indicative of the evolutionary state; the secondary mass is, to first order, a simple function of the period and the mass-radius relationship \\citep[see][for example]{howell01a}. The consequence of this is that an evolved secondary star will be under-massive for a given period, compared with the expected mass of a main-sequence secondary star. In fact, this is precisely what we observe; adopting the main-sequence mass-radius relationship of \\cite{chabrier97}, we find the secondary star mass at a period of 2.4 hours should be $0.25M_{\\sun}$. In this paper we find a secondary star mass of $M_{\\rmn{r}} =0.20\\pm0.01 M_{\\sun}$, which is significantly lower than expected for a main sequence star. It is worth spending some time speculating upon the origin and subsequent evolution of SDSS J1702+3229. Clearly the secondary star has undergone significant nuclear evolution; the requirement for this to happen within a Hubble time implies that the initial mass of the secondary should be greater than 0.8$M_{\\sun}$. What happened to the system after contact depends upon the secondary mass and upon the evolutionary state of the secondary at contact; for a primary mass of 0.94$M_{\\sun}$, an initial secondary mass of 1.2$M_{\\sun}$ or less implies that mass transfer {\\em from a main sequence star} will be both thermally and dynamically stable \\citep{politano96}. However, if the secondary star had already left the main sequence before contact, or was more massive than 1.2$M_{\\sun}$, then the system would have undergone a phase of thermal timescale mass transfer. During this phase the orbit would shrink rapidly. Once mass transfer regained stability, the system would become recognisable as a CV at a much lower period. This channel raises the possibility that SDSS J1702+3229 might have emerged from thermal timescale mass transfer within the period gap. The accurate values of secondary mass, radius and effective temperature presented here should allow detailed modelling to recover the evolutionary past of SDSS J1702+3229, however this is beyond the scope of this paper. The future evolution of SDSS J1702+3229 depends sensitively on the evolutionary state of the secondary star. However, comparison with the models of \\cite{podsiadlowski03} suggests that the system will evolve below the observed ``period-minimum'', and may even be a progenitor of an AM CVn system. \\subsection{Evolved secondary stars in CVs} A small number of CV systems now show very strong evidence for the presence of an evolved secondary star; both QZ Ser and EI Psc have unusually hot secondaries for their orbital period \\citep{thorst02a,thorst02b}, and the secondary star in EI Psc has a very large N/C abundance \\citep{gaensicke03}. Indeed, the period of EI Psc is substantially below the period minimum, which is in principle only possible for an evolved secondary star. Given the difficulty in obtaining information about the secondary stars in CVs, it is at least plausible that systems with evolved secondary stars constitute a significant fraction of the CV population, although it is not yet possible to state whether the fraction is as high as 10\\%, as suggested by \\cite{podsiadlowski03}. The existence of evolved secondaries amongst the cataclysmic variable population, and the fact that such systems may not show a period gap provides a natural explanation for the non-magnetic CVs found within the gap. It is tempting to speculate that most non-magnetic CVs within the period gap are explained by systems with evolved secondary stars. Given that the first such system which allows accurate determination of system parameters has an evolved secondary there may be some support for this hypothesis. Furthermore, such a hypothesis might explain the large fraction of systems within the period gap which exhibit superhumps \\citep{katysheva03}; superhumping systems must have mass ratios below a critical value of approximately 0.3, and the undermassive secondary stars in evolved systems would increase the proportion of systems with mass ratios which are unstable to superhumps. We conclude that the secondary star in SDSS J1702+3229 shows evidence for significant nuclear evolution. The existence of a CV with an evolved secondary within the period gap supports predictions \\citep[e.g.][]{baraffe00} that CVs with evolved secondaries can continue accreting inside the period gap, and in some cases might show no period gap at all." }, "0607/astro-ph0607110_arXiv.txt": { "abstract": "We present a detailed spectral analysis of the prompt and afterglow emission of four nearby long-soft gamma-ray bursts (GRBs~980425, 030329, 031203, and 060218) that were spectroscopically found to be associated with type Ic supernovae, and compare them to the general GRB population. For each event, we investigate the spectral and luminosity evolution, and estimate the total energy budget based upon broadband observations. The observational inventory for these events has become rich enough to allow estimates of their energy content in relativistic and sub-relativistic form. The result is a global portrait of the effects of the physical processes responsible for producing long-soft GRBs. In particular, we find that the values of the energy released in mildly relativistic outflows appears to have a significantly smaller scatter than those found in highly relativistic ejecta. This is consistent with a picture in which the energy released inside the progenitor star is roughly standard, while the fraction of that energy that ends up in highly relativistic ejecta outside the star can vary dramatically between different events. ", "introduction": "The discovery in 1998 of a Gamma-Ray Burst (GRB) in coincidence with a very unusual supernova (SN) of Type Ic (GRB\\,980425 -- SN\\,1998bw; spectroscopically identified) was a turning point in the study of GRBs, offering compelling evidence that long-soft GRBs \\citep{kou93} are indeed associated with the deaths of massive stars \\citep{gal98}. The large energy release inferred for the supernova suggested that GRBs are potentially associated with a novel class of explosions, having unusual properties in terms of their energy, asymmetry, and relativistic ejecta. More importantly, however, GRB\\,980425 provided the first hint that GRBs might have an intrinsic broad range of energies: the total energy output in $\\gamma$-rays (assuming an isotropic energy release) was only $\\egiso \\approx 7 \\times 10^{47}$\\,erg, some four orders of magnitude less energy than that associated with typical GRBs \\citep{blo03}. Finally, the fact that GRB\\,980425/SN\\,1998bw was located in a nearby galaxy with a redshift $z$ = 0.0085 (\\citealt{tin98}; at\\footnote{Throughout the paper we assume a cosmology with $H_0 = 72\\;{\\rm km\\; s^{-1}\\;Mpc^{-1}}$, $\\Omega_{M} = 0.27$, and $\\Omega_{\\Lambda} = 0.73$.} $35.6\\;$Mpc it remains the closest GRB to the Earth] gave rise to the possibility that such lower-energy bursts might be more common than had previously been thought, but harder to detect due to instrumental sensitivity. Unfortunately, during the elapsed eight years, very few SNe have been observed simultaneously with GRBs. To date, three more nearby GRBs have been unambiguously, spectroscopically identified with supernovae, two of which were discovered in 2003 (GRB\\,030329 -- SN\\,2003dh; \\citealt{sta03,hjo03}, GRB\\,031203 -- SN\\,2003lw; \\citealt{tag03}) and one in 2006 (GRB\\,060218 -- SN\\,2006aj; \\citealt{pia06}). Each of these SNe is of the same unusual type as SN\\,1998bw. Note, however, that there is weaker photometric evidence that many other GRBs are accompanied by SNe, mainly by identification of a late time ``SN bump'' in the GRB optical afterglow lightcurve \\citep[e.g.][]{blo99,gal00,zeh04,woo06}. In this paper, we address only the four events with spectroscopically verified SN associations. Among these, only GRB\\,060218 (interestingly at 143.2\\,Mpc for $z = 0.0335$, the closest after GRB\\,980425; \\citealt{mir06,pia06}) had $\\gamma$-ray energetics somewhat comparable to GRB\\,980425, and could possibly be added to the intrinsically-faint GRB sample. In the case of GRB\\,030329, the total energy release was at the low end of the typical range ($\\egiso \\sim 10^{52}\\;$erg), much higher than in the other three events. In fact, SN\\,2003dh was obscured by the extreme optical brightness of the GRB afterglow and was only detected spectroscopically in the GRB optical lightcurve. Finally, the total energy release of GRB\\,031203 was intermediate between that of GRB\\,980425 and regular GRBs. \\citet{ram05} argued that the faint GRB\\,031203 was a typical powerful GRB viewed at an angle slightly greater than about twice the half-opening angle of the central jet. In the present study, we consider the energy released during the GRBs, the afterglow, and the SN explosion for these four events in all wavebands, from $\\gamma$-rays to radio waves, and we also estimate their kinetic energy content. The properties of the prompt and afterglow emission are described in detail in \\S\\S~\\ref{sec:prompt} and \\ref{sec:AG}, respectively, and are compared in \\S~\\ref{sec:comparison}. The bolometric energy calculation and the evolution of the explosion responsible for their associated SNe are presented in \\S~\\ref{sec:SN}, while \\S~\\ref{sec:dis} discusses the combined GRB-SN properties and their potential implications. We conclude in \\S~\\ref{sec:summary} with a brief summary of our primary results and their implications. ", "conclusions": "\\label{sec:summary} One of the liveliest debated issues associated with GRBs is on the total energy released during the burster explosion: are GRBs standard candles? The GRB community has vacillated between initial claims that the GRB intrinsic luminosity distribution was very narrow \\citep{h94}, to discounting all standard candle claims, to accepting a standard total GRB energy of $\\sim 10^{51}$ ergs \\citep{f01}, and to diversifying GRBs into ``normal'' and ``sub-energetic'' classes. The important new development is that we now have significant observational support for the existence of a sub-energetic population based on the different amounts of relativistic energy released during the initial explosion. A network of theoretical tests lends credence to this idea. The existence of a wide range of intrinsic energies that we presented in this work may pose challenges to using GRBs as standard candles -- it is also worth stating explicitly that, when viewed together, these four events fall away from the Amati relation. Our results are consistent with the emerging hypothesis that GRBs and XRFs share a common origin in massive WR stars. The central engine gives rise to a polar outflow with two components \\citep{woo06}. One large angle outflow (the SN), containing most of the energy and mass, is responsible for exploding the star and producing the $^{56}$Ni to make the SN bright. Only a tiny fraction of the material in this component reaches mildly relativistic velocities, which is more narrowly focused. A second outflow component (the GRB jet) occupies a narrower solid angle, probably contains smaller energy (which can range from comparable to much smaller), and most of its energy is in material with relativistic velocities (where the typical Lorentz factor of the material that carries most of the energy in this component can vary significantly between different SN-GRBs). After it exits the star, internal shocks within this jet and external shocks with the residual wind material around the star make the GRB or XRF and its afterglow. Apparently, the properties of the broad component are not nearly so diverse as those of the core jet \\citep{R-RM04,sod06,woo06}. We have argued, using well-known arguments connected with parameters such as opacity and variability timescales, that these less-energetic events do not require a highly-relativistic outflow. Our best estimates of Lorentz factors, $\\Gamma$, for these events are in the range of 2$-$10. Indeed, it is much more difficult to produce a jet with very high Lorentz factor -- i.e., a high energy loading per baryon -- than with low Lorentz factor. A jet with low Lorentz factor could result even if a jet of relatively pure energy is produced, since it may be loaded with excess baryons by instabilities at its walls as it passes through the star, or if it does not precisely maintain its orientation \\citep{R-RCR02,aloy,ZWH04}. The above suggest that GRBs made by jets with lower Lorentz factor should be quite common in the universe \\citep{GR04}. Continued advances in the observations will surely yield unexpected revisions and additions in our understanding of GRBs in connection with SNe: currently, we are attempting to draw large conclusions from limited observations of exceedingly complex phenomena. However, the big surprise at the moment is that these SN-GRB events appear to be intrinsically different from and much more frequent \\citep{GR04,guetta,pia06} than luminous GRBs, which have been observed in large numbers out to higher redshifts. We are very grateful to Scott Barthelmy and Takanori Sakamoto for their help with the {\\it Swift} BAT data analysis, to Rob Preece and Michael Briggs for their help with the BATSE-WFC joint analysis, and to Ersin G\\\"o\\u{g}\\\"u\\c{s} for helpful discussions. We also thank the WSRT staff, in particular Tony Foley. This work is supported by IAS and NASA under contracts G05-6056Z (YK) and through a Chandra Postdoctoral Fellowship award PF3-40028 (ERR), by the Department of Energy under contract DE-AC03-76SF00515 (JG), and by PPARC (ER). SEW acknowledges support from NASA (NNG05GG08G), and the DOE Program for Scientific Discovery through Advanced Computing (SciDAC; DE-FC02-01ER41176). RAMJW is supported by the Netherlands Foundation for Scientific Research (NWO) through grant 639.043.302. This paper benefited from collaboration through an EU-funded RTN, grant number HPRN-CT-2002-00294. The Westerbork Synthesis Radio Telescope is operated by ASTRON (Netherlands Foundation for Research in Astronomy) with support from NWO." }, "0607/hep-ex0607010_arXiv.txt": { "abstract": "A search has been made for neutrinos from the {\\it hep} reaction in the Sun and from the diffuse supernova neutrino background (DSNB) using data collected during the first operational phase of the Sudbury Neutrino Observatory, with an exposure of 0.65 kilotonne-years. For the {\\it hep} neutrino search, two events are observed in the effective electron energy range of 14.3~MeV $<\\mathrm{T_{\\mathrm{eff}}}<$ 20~MeV where 3.1 background events are expected. After accounting for neutrino oscillations, an upper limit of $2.3\\times10^4$~cm$^{-2}$s$^{-1}$ at the 90$\\%$ confidence level is inferred on the integral total flux of {\\it hep} neutrinos. For DSNB neutrinos, no events are observed in the effective electron energy range of 21~MeV $<\\mathrm{T_{\\mathrm{eff}}}<$ 35~MeV and, consequently, an upper limit on the $\\nu_e$ component of the DSNB flux in the neutrino energy range of 22.9~MeV $ 19.3$~MeV, based on measurements with the Super-Kamiokande detector~(\\cite{sk-dsnb}). While an indirect limit on the $\\nu_e$ component of the DSNB flux can be inferred from this (\\cite{lunardini}), the previous best direct upper limit is $6.8\\times10^{3}$~cm$^{-2}$s$^{-1}$ for neutrino energies 25~MeV $8 \\times 10^{38}$ ergs s$^{-1}$) X-ray sources that were present in the first epoch observations were still in outburst in all of the following observations. Many of these probable long-duration outburst BHLMXBs reside within globular clusters of the galaxies. Conversely, no definitive short-duration outburst BHLMXBs were detected in any of the observations. This places an upper limit on the ratio of short--to--long-duration outbursters that is slightly lower, but consistent with what is seen in the Milky Way. The fact that none of the luminous sources turned off between the first and last epochs places a 95 per cent lower limit of 50 yr on the mean burst duration of the long-duration outburst sources. The most likely scenario for the origin of these sources is that they are long-period ($>$30 d) black hole binaries with a red giant donor, much like GRS1915+105. However, unlike GRS1915+105, most of the sources show only modest variability from epoch to epoch. ", "introduction": " ", "conclusions": "" }, "0607/astro-ph0607440_arXiv.txt": { "abstract": "{ Hard X-rays and $\\gamma$-rays are the most direct signatures of energetic electrons and ions in the sun's atmosphere which is optically thin at these energies and their radiation involves no coherent processes. Being collisional they are complementary to gyro-radiation in probing atmospheric density as opposed to magnetic field and the electrons are primarily 10--100~keV in energy, complementing the ($>$100~keV) electrons likely responsible for microwave bursts. The pioneering results of the Ramaty High Energy Solar Spectroscopic Imager (RHESSI) are raising the first new major questions concerning solar energetic particles in many years. Some highlights of these results are discussed -- primarily around RHESSI topics on which the authors have had direct research involvement -- particularly when they are raising the need for re-thinking of entrenched ideas. Results and issues are broadly divided into discoveries in the spatial, temporal and spectral domains, with the main emphasis on flare hard X-rays/fast electrons but touching also on $\\gamma$-rays/ions, non-flare emissions, and the relationship to radio bursts. } ", "introduction": "\\label{intro} Major observational results from RHESSI and instrumental details have been extensively described elsewhere (e.g. \\cite{lin:al-02} and other articles in that volume, and \\cite{den:al-05}) and will not be repeated here. Based on results from numerous earlier spacecraft from OGOs, OSOs and TD1A through SMM, Hinotori and Yohkoh (these three giving the first HXR images), the conventional wisdom prior to RHESSI envisaged electron and ion acceleration high in a loop near a reconnection site. Most of the hard X-rays (HXRs) and $\\gamma$-rays were believed to originate in two bright loop footpoints by collisional thick target deceleration of fast particles with a near power-law spectrum in the dense chromosphere \\cite{bro-71}, plus occasional fainter emission at or above the looptop as seen in Yohkoh \\cite{mas:al-94} and sometimes even higher as seen in limb occulted flares \\cite{kan-83}. Until RHESSI, apart from one balloon flight \\cite{lin:sch-87}, spectral resolution was very limited, particularly in images and in (non-imaged) $\\gamma$-rays. RHESSI has transformed this via Ge detector spectrometry, yielding high resolution spectra and spectral images in HXRs, high resolution $\\gamma$-ray line spectroscopy, and the first $\\gamma$-ray line images. RHESSI also excels in having an unsaturated spectral range from a few keV to ten of MeV, thus yielding data on the hot SXR plasma as well as on fast particles (see articles in special issues of Solar Phys.\\ vol.\\ 210, 2002 and Astrophysical Journal Letters vol. 595, 2003). While many of the RHESSI data show events with some resemblance to the canonical thick target footpoint scenario, with near power-law spectra, there are many examples deviating from this simple picture. Here the main emphasis is on these new features as they are the driving force behind the need for a rethink. ", "conclusions": "\\label{conclude} RHESSI data constitute the greatest breakthrough in flare fast particle studies since the first HXR detectors were launched over 30 years ago. The results will pose `rethink' challenges for an entire new generation of solar physicists, all the more so when considered in the wider context of multi-wavelength data, especially in the complementary radio regime to which CESRA is dedicated. \\acknowledgement{This work was supported by a PPARC Rolling Grant (JCB, EPK) and a Visitor Grant from the Royal Society of Edinburgh (AMV). We are grateful to Alec MacKinnon for help with the manuscript.}" }, "0607/astro-ph0607489_arXiv.txt": { "abstract": "S Mus is the Cepheid with the hottest known companion. The large ultraviolet flux means that it is the only Cepheid companion for which the velocity amplitude could be measured with the echelle mode of the HST GHRS. Unfortunately, the high temperature is difficult to constrain at wavelengths longer than 1200 \\AA\\/ because of the degeneracy between temperature and reddening. We have obtained a FUSE spectrum in order to improve the determination of the temperature of the companion. Two regions which are temperature sensitive near 16,000 K but relatively unaffected by H$_2$ absorption (940 \\AA, and the Ly $\\beta$ wings) have been identified. By comparing FUSE spectra of S Mus B with spectra of standard stars, we have determined a temperature of 17,000 $\\pm$ 500 K. The resultant Cepheid mass is 6.0 $\\pm$ 0.4 M$_\\odot$. This mass is consistent with main sequence evolutionary tracks with a moderate amount of convective overshoot. ", "introduction": "Observational determinations of Cepheid masses are a long-standing goal both in order to have a thorough understanding of these primary distance indicators and also because they provide an excellent benchmark for stellar evolutionary calculations. The most important uncertainty in evolutionary tracks of massive stars near the main sequence is the importance of core convective overshoot, which determines the lifetime on the main sequence and the luminosity in subsequent phases. When the mass of a Cepheid can be measured, it can be combined with an accurate luminosity, and compared with theoretical predictions. Ultraviolet high resolution spectroscopy has provided a group of double-lined spectroscopic binaries containing a Cepheid. Specifically, the orbital velocity amplitudes of the hot companions of Cepheids could be measured originally with IUE, and until recently with the Hubble Space Telescope (HST) Space Telescope Imaging Spectrograph (STIS) or Goddard High Resolution Spectrograph (GHRS). This orbital velocity amplitude can be combined with the orbital velocity amplitude of the Cepheid from a ground-based orbit and the mass of the companion to produce the mass of the Cepheid. Typically, a very accurate temperature or spectral type for the hot companion can be obtained from IUE low resolution spectra from 1200 to 3200 \\AA, from which a mass can be accurately inferred. For the S Mus system, the orbit of the Cepheid S Mus A has been determined several times with increasing accuracy as more data have been obtained (Evans, 1990; B\\\"ohm-Vitense et al. 1997; and Petterson et al. 2004). The hot companion of the Cepheid, S Mus B, is sufficiently bright at 1720~{\\AA} that it could be observed with the echelle mode of HST/GHRS which provided a resolution of 80,000 (B\\\"ohm-Vitense et al. 1997). The orbital velocity amplitude of S Mus B they found from two GHRS observations is 30.6 km s$^{-1}$ with an uncertainty of 5\\%. The uncertainty is dominated by the centering of the star in the large science aperture for the first observation. This is the most accurate velocity amplitude measured for a Cepheid companion. The high temperature of S Mus B, however, means that its temperature (and hence its inferred mass) is less accurately determined than that for cooler companions in other systems. For late B stars, the energy distribution turns over between 1200 and 1400 \\AA, making that region of the spectrum extremely temperature sensitive. For an early B star, the spectrum rises monotonically toward shorter wavelengths to the end of the IUE spectral range ($\\sim$1200 \\AA). This means the effects of reddening and temperature are much more difficult to disentangle, and hence the temperature is less accurately determined. The reddening of the system is E(B-V) = 0.21 mag (Evans, Massa and Teays 1994), which is large enough that it must be taken into account. A number of approaches to determining the temperature in the wavelength range 1200 to 3200 \\AA\\/ have been used, as summarized by B\\\"ohm-Vitense et al (1997), including energy distributions from IUE low resolution spectra, and Si lines near 1300 \\AA\\/ from IUE high resolution spectra (Evans, Massa and Teays, 1994). In addition, two Voyager spectra were obtained to extend the energy distribution to 950 \\AA\\/ (Evans, Holberg and Polidan, 1996). The difficulty in interpreting these spectra comes from the heavy absorption by H$_2$ molecular absorption bands. In the low resolution Voyager spectra, approximate corrections had to be incorporated to compensate for this absorption. As a substantial refinement to this basic approach, we obtained a high-resolution FUSE spectrum of S Mus B in order to determine its temperature more precisely, with the ultimate goal of improving the estimated mass of its Cepheid companion. ", "conclusions": "In order to determine the mass of S Mus B corresponding to this temperature, we used masses from the compilation by Andersen (1991) derived from very accurate eclipsing binary solutions. We have combined these with recent temperatures from Ribas et al. (2000). These temperatures are based on Str\\\"omgren photometry, and should be comparable to the temperatures of the standard stars. In Fig. 5 we show the relation between temperature and mass for the stars in the Anderson list more massive than 2.5 M$_\\odot$ (O and B stars). In order to obtain a mass corresponding to the temperature of S Mus B, we did a linear fit to the data for $\\log T_{\\rm eff}$ 4.27 to 4.17 (T 18,500 to 14,800 K). The resulting mass for 17,000 K is 5.3 M$_\\odot$. A change of 500 K results in a change of 0.26 M$_\\odot$. The orbital velocity amplitude ratio of the Cepheid to the hot companion was found to be 1.14 $\\pm$ 0.06 from the GHRS echelle observations of the companion and the ground-based Cepheid orbit (B\\\"ohm-Vitense et al. 1997). Combining this with the companion mass found here results in a Cepheid mass of 6.0 $\\pm$ 0.4 M$_\\odot$. This value only differs slightly from the previous determination (5.9 $\\pm$ 0.7 M$_\\odot$ B\\\"ohm-Vitense et al., 1997) but the error bars are significantly reduced. This temperature and mass determination supersedes previous estimates, because it lifts the degeneracy in the energy distribution longward of 1200 {\\AA} and avoids the need for the approximate corrections for H$_2$ absorption required to interpret Voyager spectra. Fig. 6 summarizes the information currently available for Cepheid masses. (Luminosities are taken from Evans et al, 1998). For comparison, the luminosity predicted for the tips of the blue loops the evolutionary tracks from several several groups is shown. The two lines in the center are from the Padua and Geneva groups for moderate overshoot. To the left is prediction from the Padua tracks for their maximum overshoot. The line on the right is from Becker (1981) with no overshoot. The mass we have determined for S Mus clearly favors moderate overshoot." }, "0607/hep-ph0607343_arXiv.txt": { "abstract": "The quark matter equation of state (EOS) derived from the standard Nambu - Jona-Lasinio (NJL) model is soft enough to render neutron stars (NS) unstable at the onset of the deconfined phase, and no pure quark matter can be actually present in its interior. Since this is a peculiarity of the NJL model, we have studied a modified NJL model with a momentum cut-off which depends on the density. This procedure, which improves the agreement between QCD and NJL model at large density, modifies the standard NJL equation of state, and then it is potentially relevant for the stability analysis of neutron stars. We show that also within this approach, the NS instability still persists, and that the vacuum pressure, as a signal of quark confinement, has a fundamental role for the NS stability. In this respect, our conclusions point to a relationship between confinement and NS stability. ", "introduction": "In the core of astrophysical compact objects, like neutron stars (NS) or proto-neutron stars, nuclear matter is expected to reach a density which is several times nuclear saturation density. Calculations based on microscopic Equation of States (EoS), which include only nucleonic degrees of freedom, show that the central density of the most massive neutron stars can be from seven to ten times the nuclear saturation density \\cite{bbb}. In such configurations the nucleons are closely packed, and to consider them as separate entities becomes highly questionable. Unfortunately, it is difficult to calculate accurately the transition point from nucleonic to quark matter, and only rough estimates have been given in the literature \\cite{glu}. The microscopic theory of the nucleonic Equation of States has reached a high degree of sophistication, and different many-body methods have been developed. They show a substantial agreement among each others \\cite{tri}, and therefore the main uncertainty of the transition point stays on the quark matter EoS. Assuming a first order phase transition, as suggested by lattice calculations, one can use different models for quark matter to estimate the transition point, and calculate the compact object configuration. This approach has been followed by several authors, and a vast literature exists on this subject. The applications of quark matter models to the study of the NS structure have used different versions of the MIT bag model \\cite{bag}, the color-dielectric model \\cite{col} and different formulations of the NJL model \\cite{bub}. In general, it seems that the maximum mass of NS that contain quark matter in their interior is bounded to be less than 1.6 - 1.7 solar masses. However, more recently it has been shown in \\cite{alf} that if one corrects the MIT bag model by introducing additional terms suggested by perturbative QCD, the maximum mass can reach values close to 1.9 solar masses, similar to the ones obtained with nucleonic degrees of freedom. Despite the similarity of the results on the value of the maximum NS mass, the predictions on the NS configurations can differ substantially from model to model. The most striking difference is in the NS quark matter content, which can be extremely large in the case of EoS related to the MIT bag model or the color-dielectric model, but it is vanishingly small in the case of the original version of the NJL model \\cite{bub,sch}. In the latter case it turns out that, as soon as quark matter appears at increasing NS mass, the star becomes unstable towards collapse to a black hole, with only the possibility of a small central region with a mixed phase of nucleonic and quark matter. This result can be quite relevant for the physics of NS, since the NJL model is the only model which is based on phenomenological low energy data, i.e. on hadron properties. The main drawback of the model is the absence of confinement, since the gap equations which determine the quark masses as a function of density cannot incorporate a confining potential. One then assumes that the chiral phase transition marks also the confinement transition, as indicated by all lattice calculations. More recently a confining potential has been introduced in the NJL model \\cite{thom}, which is simply switched off at the chiral phase transition. In this case indeed no sharp instability of the NS is observed at the onset of quark matter in the central core. This is suggestive of a connection between the presence of confinement and the possibility of NS with a quark matter central core. However it is not clear if the instability of the NS is related or not to confinement for (at least) two reasons. Indeed, the introduction of the confining potential requires several parameters and, moreover, the NJL model at large density suffers of another drawback pointed out in \\cite{casa}. In fact, it has been shown that the standard NJL model is not able to reproduce the correct QCD behavior of the gap for large density, and therefore a different cut-off procedure at large momenta has been proposed. More precisely, a density dependent cut-off has been introduced, and this strongly modifies the standard NJL model thermodynamics. Therefore, the stability analysis of NS by the new EoS, which follows from this modified treatment of NJL model at large density, is a preliminary step to understand whether confinement is an essential ingredient in the stabilization mechanism. \\par It is the purpose of this paper to clarify this point and to identify the origin of the NS instability at the quark onset within the original NJL model, which sharply distinguishes this model from all the others. We show that with the density dependent cut-off procedure the NS instability still persists and that the vacuum pressure, as signal of quark confinement, has a fundamental role for the NS stability, as yet observed in the MIT bag model. In this respect, our conclusions point to an indirect relationship between confinement and NS stability. ", "conclusions": "" }, "0607/astro-ph0607241_arXiv.txt": { "abstract": "We present a new, simple, fast algorithm to numerically evolve disks of inelastically colliding particles surrounding a central star. Our algorithm adds negligible computational cost to the fastest existing collisionless N-body codes, and can be used to simulate, for the first time, the interaction of planets with disks over many viscous times. Though the algorithm is implemented in two dimensions---i.e., the motions of bodies need only be tracked in a plane---it captures the behavior of fully three-dimensional disks in which collisions maintain inclinations that are comparable to random eccentricities. We subject the algorithm to a battery of tests for the case of an isolated, narrow, circular ring. Numerical simulations agree with analytic theory with regards to how particles' random velocities equilibrate; how the ring viscously spreads; and how energy dissipation, angular momentum transport, and material transport are connected. We derive and measure the critical value of the coefficient of restitution above which viscous stirring dominates inelastic damping and the particles' velocity dispersion runs away. ", "introduction": "How does a disk of collisional particles surrounding a star evolve in the presence of planets? The answer to this question has important implications. For example, after the planets of our Solar System accreted most of their mass, many small, rocky and icy bodies remained orbiting the Sun. Somehow, the planets eliminated most of these remnant planetesimals, while leaving some behind to form the asteroid belt, the Kuiper belt, and the Oort cloud. In the vicinity of Uranus and Neptune, the small bodies must have been highly collisional. Otherwise, these planets would have taken $10^{12}$ yr to form {\\it in situ} \\citep{GLS04}.\\footnote{ In the terrestrial zone, the small bodies were also likely collisional, although the case is not as convincing there as it is in the outer Solar System \\citep{GLS04}.} Yet virtually all simulations of the late stages of planet formation in the outer Solar System---such as those that model the migration of the ice giants, the resulting trapping of Kuiper belt objects into resonances, and the ejection of small bodies to the Oort cloud---neglect collisions. When the effects of collisions are accounted for, the current picture of the formation of planetary systems might change drastically. Planetary rings provide another setting in which interparticle collisions play a crucial role. What are the origins of narrow rings shepherded by satellites? How do narrow rings settle into their special apse and node-aligned states \\citep[e.g.,][]{CC04}? And how do rings back-react upon and shape the orbits of shepherd satellites? Our understanding of satellite-ring interactions bears on mysteries such as the origin of eccentricities of extra-solar giant planets \\citep[e.g.,][]{GS03}. Despite its importance, the behavior of particle disks in the presence of perturbing bodies is poorly understood. Numerical simulations can help to further understanding. But until now, simulations of collisional disks have been too inefficient to follow, say, how disks viscously spread in the long term. Collisions are traditionally simulated with a brute-force method \\citep[e.g.,][]{Bra77,WT88}: at each time step of the integration of the gravitational equations of motion, it is determined which pairs of particles might collide before the next timestep. These potential collision pairs are then integrated forwards in time with a much smaller timestep, to see if they really do collide. But this method is inefficient: a brute-force search for collision partners requires around $N_{\\rm tp}^2$ operations at each timestep, where $N_{\\rm tp}$ is the number of test particles. In addition, most potentially colliding pairs do not collide, particularly in optically thin disks. Hence much computing time is wasted on missed collisions. More complex algorithms have been devised to reduce computing time \\citep[e.g.,][]{LS00,CTB01}. But these are still not nearly as fast as the fastest collisionless N-body codes, such as SWIFT \\citep{LD94}. We sought a collision algorithm that (i) could be added to any N-body code, such as the freely-available SWIFT; (ii) contributes negligibly to the computational cost; (iii) is simple conceptually; (iv) is easy to code; and (v) follows correctly the long-term viscous evolution of disks in the presence of planets. We designed our algorithm to simulate a vertically optically thin disk of identical, collisional, massless, inelastic but indestructible test particles that feel the gravity of the Sun and of multiple planets. Complications that we do not include, such as the self-gravity of the particles, order-unity optical depths, and particles with differing sizes, spins, and cohesive strengths, could all affect the viscous evolution in ways that are not currently understood. But at this stage it seems wisest to ignore these complications, even though the algorithm could be modified to handle them. Viewed in its most basic terms, inelastic collisions dampen random velocities and act as a source of friction between neighboring streamlines. As long as our algorithm preserves this behavior, while conserving angular momentum and accounting for the loss of energy in inelastic collisions, it seems likely that it will properly model the long-term evolution of collisional disks. In the present paper, we test this assertion thoroughly when there are no planets, comparing in detail the results of our simulations with those of analytic theory. In a future paper, we shall test our algorithm in the presence of planets. ", "conclusions": "We have introduced an algorithm to simulate collisions between inelastic particles in an optically thin disk orbiting a central mass. The algorithm is simple to implement and adds negligible running time to existing collisionless N-body codes. A major feature of the algorithm is that the disk particles' motions need only be tracked in a plane. Yet the algorithm transcends its two-dimensional appearance to simulate a three-dimensional disk of particles whose random velocity distribution tends to be isotropized by collisions. We have performed a battery of tests of the algorithm for the case of an isolated, narrow, circular ring. Numerical simulations agree with analytic theory with regard to how the particles' velocity dispersion equilibrates, how the ring viscously spreads, how energy and angular momentum are transported, and how energy dissipation relates to the viscous angular momentum flux and to the background shear. Angular momentum transport arises not only from particle advection ($HF_n$), but also from correlations in the random velocity field ($F_{r\\phi}$) and from finite particle sizes ($F_{\\rm NL}$). The relative magnitudes of each of these three terms can be measured from simulations. In making these and other measurements, we sought ways to minimize noise introduced by finite particle numbers (Poisson fluctuations). For example, when measuring viscous fluxes of angular momentum and energy, it proves useful to consider only those particles that actually collide during the measurement interval. The stage is now set for simulating more complicated systems---narrow eccentric rings (like the Maxwell and Titan ringlets of Saturn, or the Epsilon ring of Uranus), and circumstellar disks with embedded planets. Among the phenomena we are interested in exploring numerically are the formation of sharp edges by shepherd satellites, the evolution of narrow rings into states of rigid apsidal precession, and the eccentricity evolution of planets as driven by disks." }, "0607/astro-ph0607288_arXiv.txt": { "abstract": "We study a possibility to use the octopole moment of gravitationally lensed images as a direct measure of the third-order weak gravitational lensing effect, or the gravitational flexion. It turns out that there is a natural relation between flexion and certain combinations of octopole/higher-multipole moments which we call the Higher Order Lensing Image's Characteristics (HOLICs). This will allow one to measure directly flexion from observable octopole and higher-multipole moments of background images. We show based on simulated observations how the use of HOLICs can improve the accuracy and resolution of a reconstructed mass map, in which we assume Gaussian uncertainties in the shape measurements estimated using deep $i'$-band data of blank fields observed with Suprime-Cam on the Subaru telescope. ", "introduction": "It is now widely recognized that weak gravitational lensing is a unique and valuable tool to study the mass distribution of clusters of galaxies as well as large scale structure in the universe since it directly measures the projected mass distribution of the lens regardless of the physical state of the system and the nature of matter content (Bartelmann \\& Schneider 2001). In the usual treatment of the weak lensing analysis, the quadrupole moment of background galaxy images is used to quantify the image ellipticity. Then the lensing properties are extracted from the image ellipticities by assuming that source galaxies are randomly oriented in the absence of gravitational lensing. In practice we average over a local ensemble of image ellipticities to estimate the lensing properties. The local ensemble should contain a sufficient number of background galaxies to increase the signal-to-noise ratio of local shear measurements, whereas the region that contains the galaxies should be small enough to guarantee the constancy of the lensing properties over the region. The latter condition is necessary in the usual prescription for weak lensing because it is based on the locally linearized lens equation. On the other hand, the former limits the resolution of mass maps reconstructed via weak lensing techniques, which is of the order 1 arcmin in ground-based observations. Space-based high-resolution imaging surveys, such as the Cosmic Evolution Survey (Scoville et al. 2007) with the {\\it Hubble Space Telescope} (HST) and the proposed {\\it Supernova/Acceleration Probe} (SNAP) wide weak lensing survey (Massey et al. 2004), will provide significant gains with a higher surface number density of well-resolved galaxies due to the small, stable Point Spread Function (PSF), which will enable high-resolution mapping of the lensing mass distribution down to an angular resolution of $\\sim 0\\farcm 1$. On the other hand, such a small PSF will allow us to resolve not only the elliptical component, described by the quadrupole moment, but also higher-order shape properties of background galaxy images, which could also carry some sort of information of lensing properties. It may be therefore interesting to see if such higher multipole moments of the shape are useful for the weak lensing analysis. There have been some attempts to generalize the weak lensing analysis to include higher order moments of the light distribution. Goldberg and Natarajan (2002) suggested that higher order effects in gravitational lensing, described by the third order derivatives of the lensing potential, can give rise to octopole moments of the light distribution for background galaxies. Goldberg and Bacon (2005) have further developed their approach and proposed a new inversion technique based on the Shapelets formalism (Refregier 2003; Refreger \\& Bacon 2003; Massey and Refregier 2005), and labeled this third order effect as the {\\it flexion} of background images. Irwin and Shmokova (2006) developed a similar analysis method for measuring the higher order lensing effects and applied this method to the HST Deep Field North. Recently Irwin, Shmokova, \\& Anderson 2006 reported on the detection of lensing signals in the UDF due to small scale structure using their \"cardioid\" and \"displacement\" techniques. Recently Goldberg \\& Leonard (2006) has extended our HOLICs approach to developed a method to correct HOLICs for the effect of isotropic Point Spread Function (PSF). In the present paper, following the flexion formalism by Bacon et al. (2006), we study a possibility to use higher multipole moments of background source images for the weak lensing analysis, and demonstrate via simulations how such higher order moments can improve the accuracy and resolution of a weak lensing mass reconstruction. The paper is organized as follows. After briefly summarizing the basis of weak lensing and the flexion formalism in section 2, we introduce higher multipole moments of galaxy images in section 3. We define certain combinations of higher multipole moments as HOLICs (Higher Order Lensing Image's Characteristics) and establish an explicit relation between flexion and HOLICs. In section 4 we present simulations of a weak lensing mass reconstruction using mock observational data of image ellipticities and HOLICs. Finally some discussions and comments are given in section 5. ", "conclusions": "In the present paper, we have studied the possibility to improve the weak lensing analysis by utilizing the octopole and higher-multipole moments of lensed images that carries the third-order weak lensing effect. By defining proper combinations of octopole moments as HOLICs, we have derived explicit relations between the flexion fields $(F,G)$ and observable HOLICs $(\\zeta,\\delta)$. In the weak lensing limit, the first flexion $F$ excites in lensed images the first HOLICs $\\zeta$ with spin-1, while the second flexion $G$ excites the second HOLICs $\\delta$ with spin-3. One can employ the assumption of random orientation for intrinsic HOLICs of background sources to obtain an unbiased, direct estimator for flexion, in a similar manner to the usual prescription for weak lensing. We have also shown by using simulated observations how the use of HOLICs can improve the accuracy and resolution of a reconstructed mass map, in which we assumed Gaussian uncertainties in the shape measurements estimated using deep $i'$-band data of blank fields observed with Subaru/Suprime-Cam. The gravitational shear and flexion have different scale-dependence in mass reconstruction errors. The mass maps reconstructed using HOLICs recover substructures better than the shear-based reconstruction, allowing a high-resolution mass reconstruction. It is shown that an optimal linear combination of individual mass reconstructions can be formed using the statistical weight in Fourier space, which can improve the statistical significance of weak lensing mass reconstructions. In actual observations, on the other hand, we must apply various shape corrections (e.g., isotropic/anisotropic PSF corrections) to the higher order shape quantities in order to measure flexion to high precision. In particular, the first HOLICs $\\zeta$ with spin-1 is highly sensitive to the choice of the center-of-image. These issues will be discussed further in the forthcoming publications." }, "0607/astro-ph0607131_arXiv.txt": { "abstract": "The Friedman equation is solved for a universe containing hot-dark matter and cold dark matter. In this scenario, hot-dark matter drives an accelerating universe no cold dark matter. ", "introduction": "Supernovae observations suggest that our universe is accelerating and the dark energy contributes $\\Omega_{\\rm DE}\\simeq 0.60-0.70$ to the critical energy density of the present universe. In addition, cosmic microwave background observations imply that the standard cosmology is given by the inflation and FRW universe. On the one hand, a non-zero cosmological constant satisfies the existence of dark energy in the universe. On the other hand, our measurements are claimed to apply the existence of a dark energy in the background gravity with non-zero cosmological constant. Therefore, a typical candidate for the dark energy is the cosmological constant. In quantum field theory it is shown that, a short distance cutoff (UV cutoff: $\\Lambda$) is related to a long distance cutoff (IR cutoff: $L_{\\rm \\Lambda}$) due to the limit set by forming a black hole. Taking ${L_\\Lambda }$ as the size of the present universe, the resulting energy is comparable to the present dark energy. Even though when this approach leads to the data, the description is incomplete because it fails to explain the dark energy dominated present universe. To resolve this situation, one is forced to introduce another candidate for IR cutoff. One is the particle horizon $R_{H }$that used by Fishler and Susskind. This gives $\\rho_{\\rm \\Lambda} \\sim a^{-2(1+1/c)}$ which implies $\\omega_{\\rm H}>-1/3$ however; this result corresponds to a decelerating universe. A recent study of the x-ray emission of hot gas in a massive cluster of galaxy has allowed astronomers to determine the distribution of its dark matter content. The density of dark matter appears to increase towards the center of the cluster in agreement with cold-dark matter production. The x-ray data show that the dark matter density increases smoothly all the way into the central galaxy of the cluster. In the past five years, there has been growing evidence in favor of the cold-dark matter model. The Wilkinson Microwave Anisotropy Probe showed that normal baryon matter only accounts for 17 percent of the matter content of the universe; the rest being cold dark matter of unknown nature. Dark matter particles must have the property of increasing with each other and with normal baryon matter only through gravity. These so-called weakly interacting massive particles are difficult to detect and have been elusive until now. Massive neutrinos are a possible dark matter candidate, usually referred to them as hot dark matter because they travel at close the speed of light. Due to this high speed, hot dark matter models of the early universe create big structure of the size of galaxy clusters, which then fragment to form galaxies. By contrast, the slower cold dark matter particles cannot travel as far and so form small galaxies, which then merge to form the bigger structures such as clusters of galaxies. Our problem is related to answer the question: hot dark matter drives an accelerating universe or cold dark matter. The Friedman equation is solved for a universe contains hot-dark matter and could dark matter. It is shown that, hot-dark matter drives an accelerating universe no cold dark matter. ", "conclusions": "Dark particles must have the property of increasing with each other and with normal baryon matter only through gravity. So-called weakly gravity interacting massive dark particle are difficult to detect and have been elusive until now. Massive neutrinos are a possible dark particle candidate, usually referred to them as hot dark matter because they travel at close to the speed of light, due to this light speed, hot dark matter models of the early universe create big structure of the size of galaxy clusters, which then fragment to form galaxies. By contrast, the slower cold dark matter particles cannot travel as far and forms the small galaxies first, which then merge to form the bigger structures such as clusters of galaxy. We showed that, the hot dark matter (neutrinos) drives an accelerating universe." }, "0607/astro-ph0607307_arXiv.txt": { "abstract": "We report here X-ray imaging spectroscopy observations of the northeastern shell of the supernova remnant \\rcw\\ with \\chandra\\ and \\xmm. Along this part of the shell the dominant X-ray radiation mechanism changes from thermal to synchrotron emission. We argue that both the presence of X-ray synchrotron radiation and the width of the synchrotron emitting region suggest a locally higher shock velocity of $V_s \\approx 2700$~\\kms\\ and a magnetic field of $B \\approx 24\\pm 5~\\mu$G. Moreover, we also show that a simple power law cosmic ray electron spectrum with an exponential cut-off cannot explain the broad band synchrotron emission. Instead a concave electron spectrum is needed, as predicted by non-linear shock acceleration models. Finally, we show that the derived shock velocity strengthens the case that RCW 86 is the remnant of SN 185. ", "introduction": "Since the discovery of X-ray synchrotron radiation from SN1006 \\citep{koyama95} it has been found that most young, shell-type supernova remnants (SNRs) emit X-ray synchrotron radiation \\citep[][for a review]{vink06b}. For the youngest SNRs Cas~A, Kepler (SN1604) and Tycho (SN1572) this radiation is confined to a narrow region close to the shock front. This does not seem to be the case for somewhat older, but physically much larger objects like RCW 86, G266.2-1.2 \\citep{slane01a} and G347.3-0.5 \\citep{cassam04c}. The X-ray emission from the latter two SNRs does in fact only reveal synchrotron radiation, and both have been detected in TeV gamma-rays \\citep[][]{aharonian05,aharonian04}. It has been shown that the thickness of the X-ray synchrotron emitting region is directly related to the post-shock magnetic field strength \\citep[e.g][]{vink03a}. The relatively strong fields found are probably a result of magnetic field amplification by cosmic ray streaming \\citep[e.g.][]{bell04,bykov05}. The presence of X-ray synchrotron radiation is in itself an indication of efficient acceleration, and if the cut-off energy of the photon spectrum is determined by synchrotron losses, it is independent of the magnetic field strength, but scales with shock velocity, as $\\propto V_s^2$ \\citep{aharonian99}. Here we report on the analysis of \\chandra\\ and \\xmm\\ data of the northeastern (NE) part of \\rcw\\ (G315.4-2.1). \\rcw\\ is an interesting SNR. The non-thermal X-ray emitting regions are broader than those of the historical SNRs, and not confined to the forward shock region, possibly as a result of projection effects \\citep{vink97}. In that respect \\rcw\\ resembles G266.2-1.2 and G347.3-0.5. However, unlike those SNRs \\rcw\\ also emits noticeable thermal X-ray emission. This allows the determination of plasma properties, which can help to determine the conditions that may potentially lead to X-ray synchrotron emission in older SNRs. In particular the NE part seems best suited for that purpose, since the X-ray synchrotron radiation is confined to the region directly behind the shock front, and, as a result, the geometry of the emitting region is easier to asses. \\begin{table*} \\begin{center} \\caption{Best fit models for the \\xmm\\ MOS1\\&2 spectra.} {\\footnotesize \\begin{tabular}{lcccccc} \\hline\\hline\\noalign{\\smallskip} & NE & E bright & E faint & N bright & N faint\\\\ \\noalign{\\smallskip}\\hline \\noalign{\\smallskip} $EM{_1}$ ($10^{53}$cm$^{-3}$kpc$^{-2}$) & $12.6\\pm1.9$ & $311\\pm45$ & $70.6\\pm 8.0$ & $219\\pm48$ &$69.5$\\\\ $kT_{\\rm e 1}$ (keV) & $6.7\\pm2.6$ & $0.57\\pm0.05$ & $0.93\\pm0.07$ & $0.96\\pm0.13$&$6.3\\pm1.6$\\\\ $n_{\\rm e}t_1$ ($10^9$~\\netunit) & $2.25\\pm0.15$& $6.7\\pm0.6$ & $5.67\\pm0.03$ & $4.0\\pm0.3$ &$2.27\\pm0.12$\\\\ \\noalign{\\smallskip} $EM{_2}$ ($10^{53}$cm$^{-3}$kpc[$^{-2}$) & - & $43.8\\pm5.3$ & - & $40.9\\pm5.6$ & -\\\\ $kT_{\\rm e 2}$(keV) & - & $3.0\\pm0.3$ & - & $3.2\\pm0.6$ & -\\\\ $n_{\\rm e}t_2$ ($10^9~$\\netunit)& - & $17.0\\pm0.5$ & - & $19.7\\pm0.9$ & -\\\\ \\noalign{\\smallskip} O & 1 &\\multicolumn{2}{c}{$0.56\\pm0.06$} & $0.68\\pm0.09$ & 1\\\\ Ne & 1 &\\multicolumn{2}{c}{$0.59\\pm0.04$} & $0.72\\pm0.07$ & 1\\\\ Mg & 1 &\\multicolumn{2}{c}{$0.44\\pm0.04$} & $0.39\\pm0.05$ & 1\\\\ Si & 1 &\\multicolumn{2}{c}{$0.27\\pm0.04$} & $0.38\\pm0.07$ & 1\\\\ Fe & 1 &\\multicolumn{2}{c}{$0.64\\pm0.07$} & $0.79\\pm0.14$ & 1\\\\ \\noalign{\\smallskip} PL Norm ($10^{-3}$s$^{-1}$keV$^{-1}$ cm$^{-2}$)% & $1.22\\pm0.03$& - & - & - &$1.8\\pm0.1$\\\\ $\\Gamma$ & $2.82\\pm0.04$& - & - & - &$3.5\\pm0.2$\\\\ $N_{\\rm H}$ ($10^{21}$~cm$^{-2}$) & $4.1\\pm0.1$ &\\multicolumn{2}{c}{$3.6\\pm0.2$} & $3.6\\pm0.2$ &$3.9\\pm0.2$\\\\ \\noalign{\\smallskip} C-stat/d.o.f. & 190.6/169 &\\multicolumn{2}{c}{(302+260)/230}& 182/85 & 242/103 \\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\tablecomments{The models consist of either one or two NEI components, or of an NEI component with an additional power law component. Abudances are given with respect to the solar abunances of \\citet{anders89}. An entry ``1'' means the abundance as fixed to the solar value. Entries covering two columns were obtained by fitting two spectra simultaneously, and forcing the parameters to be identical. The emission measure, $EM$, is defined as $\\int n_{\\rm e}n_{\\rm H} dV/d^2$. The power law normalization is given at 1 keV. All errors are 1$\\sigma$ errors ($\\Delta\\chi^2=1$).\\label{tab_spectra} }} \\end{center} \\end{table*} \\begin{figure*}[t] \\centerline{ \\psfig{figure=f3.eps,width=0.95\\textwidth} } \\caption{\\xmm\\ EPIC-MOS spectra from the regions labeled in Fig.~\\ref{fig_maps}. Left: logarithmic plots of a thermal and a non-thermal spectrum. Right: a comparison of the line emission from various regions. From the northeastern spectrum (in red, left panel) the best fit power law model has been subtracted in order to emphasize the thermal emission. Dashed lines indicate (from left to right) the energies of OVII He$\\alpha$, OVIIILy$\\alpha$, and Fe XVIII line emission. \\label{fig_spectra} } \\end{figure*} ", "conclusions": "A recent development in the interpretation of X-ray synchrotron radiation is that the width of the region can be used to estimate the downstream magnetic field. However, different groups have used, at face value, different methods for estimating the magnetic field: \\citet{vink03a} assumed that the width, $l$, of the X-ray synchrotron emitting region is determined by a combination of plasma velocity relatively to the shock front, $u$, and the synchrotron loss time $\\tau_{loss} = 637/B^2E$~, i.e. $l_{adv} = u \\tau_{loss}$, with $E$ the particle energy in erg and $B$\\ the magnetic field strength in Gauss. The other method assumes that the width corresponds to the diffusion length scale \\citep{bamba05,voelk05}, given by $l_{diff} = D/u$, with $D = cE/3eB$ the diffusion coefficient. In essence $l_{diff}$\\ is the length scale at which advection starts to dominate over diffusion as means of transporting particles. Both methods give a combination of $E$ and $B$, which can be solved by using the fact that the observed photon energies peak around $\\epsilon = 7.4 B E^2~{\\rm keV}$.% Both methods rely on different assumptions. First of all, for standard shocks $u = V_s/4$, but for very efficient shock acceleration the compression factor may be stronger than a factor 4. Secondly, the diffusion length method assumes $D = cE/3eB$, which is the smallest diffusion coefficient possible, as the particle mean free path is then equal to the gyro-radius (the ``Bohm limit''). It turns out that with both methods very similar magnetic field estimates are obtained \\citep{ballet05,vink06b}. As shown by \\citet{vink06b}, this is to be expected if one observes the X-ray synchrotron spectrum near the spectral cut-off energies: The acceleration time for particles according to the first order Fermi acceleration theory is, within a factor of order one, $\\tau_{acc} \\approx D/u^2$\\ \\citep{malkov01}. For electrons the acceleration is limited by synchrotron losses. So there is only a net acceleration if $\\tau_{acc} < \\tau_{loss}$, and the maximum energy is reached when $\\tau_{acc} \\approx \\tau_{loss}$ \\citep{reynolds98}. So: \\begin{eqnarray} \\tau_{acc} \\approx \\tau_{loss} % \\Longleftrightarrow % D/u \\approx u \\tau_{loss} & \\Longleftrightarrow l_{diff} \\approx l_{adv} % \\end{eqnarray} Therefore, $l_{diff} \\approx l_{adv}$\\ is the geometrical equivalent of $\\tau_{loss} \\approx \\tau_{acc}$. Note that the {\\em observational} fact that the diffusion length scale and advection length scale methods give similar results is a justification for the assumption that the diffusion coefficient is close to the Bohm limit, and a compression ratio close to the standard value of 4 \\citep{vink06b}. The \\chandra\\ image reveals a width of the X-ray synchrotron shell of $\\sim100$\\arcsec, corresponding to $3.7\\times10^{18}$~cm for a distance of $d=2.5$~kpc \\citep{westerlund69,rosado94}. Fitting a projected shell model we estimate a physical width of $1.7\\times10^{18}$~cm. In addition we apply a factor 0.6, because the actual width is a convolution of advection and diffusion processes (i.e. we set $l_{diff} = l_{adv}$ in Eq. (1) of \\cite*{berezhko04a}), so $l_{adv} = 1.0\\times10^{18}$~cm. Assuming that the shock velocity is $V_s=600$~\\kms\\ \\citep{ghavamian01} gives inconsistent results for the two methods: $B\\sim90$~\\mug\\ employing the diffusion length method, and 6~\\mug\\ assuming the advection length method. However, if we no longer assume a shock velocity of 600~\\kms, which after all was based on optical observations of a different part of \\rcw\\ from which no X-ray synchrotron radiation is emitted, we can use the condition that $l_{adv} = l_{diff}$ in order to estimate $V_s$ and $B$, using an observed photon energy of $\\epsilon \\approx 1$~keV. We found that $B$ can be directly estimated from the diffusion/advection length alone, whereas the plasma velocity only depends on the photon energy: \\begin{eqnarray} B \\approx \\bigl(\\frac{c}{3e}\\bigr)^{1/3}l_{adv}^{-2/3} = 24 \\bigl(\\frac{l_{adv}}{1.0\\cdot 10^{18} {\\rm cm}}\\bigr)^{-2/3}\\ {\\rm \\mu G},\\\\ V_s = u \\chi \\approx \\chi \\sqrt{\\frac{\\epsilon}{7.4}% \\frac{c}{3e \\cdot 637} } = 2650 \\cdot \\frac{\\chi}{4} \\sqrt{\\frac{\\epsilon}{1\\, {\\rm keV}}}\\, {\\rm km/s}, \\label{eq_aharonian} \\end{eqnarray} where $\\chi$ the shock compression factor. Eq.~(\\ref{eq_aharonian}) was previously reported by \\citet{aharonian99}. To get more accurate estimations one has to solve the, model dependent, kinetic equations for the particle distribution. For the moment, we estimate the error in the physical width due to distance uncertainties and projection effects to be $\\sim 30$\\%, resulting in an error in $B$ of $\\sim 5$~\\mug. The downstream magnetic field therefore appears to be lower than for other SNRs \\citep{vink06b}. There is some uncertainty in the actual cut-off energy of the synchrotron spectrum (see below), but uncertainties about the assumptions - larger compression ratio, less efficient diffusion - make Eq.~(\\ref{eq_aharonian}) effectively a lower limit. One may wonder why the shock velocity in some regions may be so much higher, and why it has not resulted in a more distorted shell. Here the thermal spectra helps to answer this question. For all fitted regions the \\net\\ value of the primary component is low, being \\net$=6.7\\times 10^9$~\\netunit\\ even for the brightest region. From the size of our extraction box, we estimate that the emitting volume is $V= 10^{56}-10^{57}$~cm$^3$, together with the emission measure this implies an electron density of $n_{\\rm e}= 0.5-1.6$~cm$^{-3}$. If we use this to estimate how long ago the plasma was shock heated we find $t\\lesssim 425$~yr. This is surprisingly short for a large SNR as \\rcw. In such a time the difference in radius between regions with low and high shock velocity would be $\\sim 0.9$~pc on an average shock radius of 16~pc. The short interaction time supports the idea that \\rcw\\ is a SNR expanding in a wind blown bubble \\citep{vink97}. Such SNRs expand rapidly for a long time, but the shock velocity drops rather suddenly as soon as the shock starts interacting with the surrounding shell swept up by the stellar wind \\citep[e.g.][]{dwarkadas05}. This suggests that in \\rcw\\ the shock has reached in some regions the dense shell around the bubble some 400~yr ago, after which it rapidly decelerated. In other regions the shock velocity is still high, but, due to the low density, the thermal X-ray emission is weak. This explains the coexistence of relatively weak radio synchrotron emission with conspicuous X-ray synchrotron emission: due to the lower density fewer electrons are accelerated, but because of the high shock velocity they can be accelerated to higher energies. Interestingly, the X-ray synchrotron radiation is relatively bright, with respect to the radio emission, because we find that the simplest broad band synchrotron model, i.e. synchrotron radiation from a power law electron spectrum with an exponential cut-off does not fit the data (Fig.~\\ref{fig_broadband}). It can explain the X-ray flux, but not the spectral slope. Instead the electron spectrum needs to be concave, as predicted by non-linear shock acceleration models. However, we can not determine whether the electron power law index bends toward -2, predicted for an overall shock compression ratio of 4, or toward -1.5, predicted for strongly cosmic ray modified shocks \\citep{berezhko99}. The NE region of \\rcw\\ has % properties resembling those of the TeV emitting SNRs G347.3-0.5, and G266.2-1.2: weak radio emission, and X-ray emission (almost) entirely consisting of synchrotron radiation. For \\rcw\\ the (weak) thermal X-ray emission indicates that these properties are due to a low density combined with a, relatively, high shock velocity. We speculate that for G347.3-0.5 and G266.2-1.2 the shock also moves through a low density region, e.g. a stellar wind bubble, and % the shock velocity is similarly high. \\rcw\\ may be different in that some parts of the shock have reached the shell. Inside the bubble the shock evolution can be approximated by the Sedov self-similar model, but this breaks down as soon as the shell is reached. Finally, assuming a Sedov evolution and using the apparent radius of 22\\arcmin, the age of the remnant is estimated to be $t = \\frac{2r}{5V_s} \\approx 2250 (V_s/2700\\, {\\rm km\\,s^{-1}})$~yr. This would put the explosion date of \\rcw\\ closer to AD 185, the year a putative supernova was observed in China \\citep{stephenson02}. A shock velocity of $\\sim 600$~\\kms\\ would be more consistent with a 10,000~yr old SNR. Our results, therefore, strengthen the case that the event recorded by Chinese astronomers was indeed a supernova and that \\rcw\\ is its remnant." }, "0607/astro-ph0607461_arXiv.txt": { "abstract": "We study the first $\\sim$100\\,Myr of the evolution of isolated star clusters initially containing 144179 stars, including 13107 (10\\%) primordial hard binaries. Our calculations include the effects of both stellar and binary evolution. Gravitational interactions among the stars are computed by direct N-body integration using high precision GRAPE-6 hardware. The evolution of the core radii and central concentrations of our simulated clusters are compared with the observed sample of young ($\\aplt 100$\\,Myr) star clusters in the large Magellanic cloud. Even though our simulations start with a rich population of primordial binaries, core collapse during the early phase of the cluster evolution is not prevented. Throughout the simulations, the fraction of binaries remains roughly constant ($\\sim 10$\\,\\%). Due to the effects of mass segregation the mass function of intermediate-mass main-sequence stars becomes as flat as $\\alpha=-1.8$ in the central part of the cluster (where the initial Salpeter mass function had $\\alpha=-2.35$). About 6--12\\% of the neutron stars were retained in our simulations; the fraction of retained black holes is 40--70\\%. In each simulation about three neutron stars become members of close binaries with a main-sequence companion. Such a binary will eventually become an x-ray binary, when the main-sequence star starts to fill its Roche lobe. Black holes are found more frequently in binaries; in each simulated cluster we find $\\sim11$ potential x-ray binaries containing a black hole. Binaries consisting of two white dwarfs are quite common, but few (20--30) are sufficiently close that they will merge within a Hubble time due to the emission of gravitational radiation. Clusters with shorter relaxation times tend to produce fewer merging white dwarf binaries. The white dwarf binaries that do merge are all sufficiently massive to produce a type Ia supernova. The densest cluster produces about twice as many blue stragglers as a field population containing the same number of binaries, and these blue stragglers are more massive, bluer and brighter than in less dense clusters. ", "introduction": "High-quality ground- and space-based observations over the past two decades have revealed the existence of numerous young, dense star clusters in our Galaxy and beyond. Examples include (1) the Arches \\citep{2002ApJ...581..258F,2005ApJ...628L.113S} and Quintuplet \\citep{1999ApJ...514..202F} systems in the Galactic center \\citep{2003ApJ...594..812G,2003ApJ...586L.127G}; (2) the rich clusters NGC 3603 \\citep{1994ApJ...436..183M,2004AJ....128.2854M} and Westerlund 1 \\citep{1998A&AS..127..423P,2005A&A...434..949C} in the Galactic disk; (3) the R136 \\citep{1998ApJ...493..180M} system and other young clusters in the Large Magellanic Cloud (LMC) \\citep{2003MNRAS.338...85M}; (4) an increasing number of young star clusters in nearby starburst systems such as the Antennae and a newly discovered cluster \\citep{2003A&A...397..177B,2003A&A...404..223B,2006ApJ...643.1166F}. These systems are of great interest for a number of reasons. First, as both observations and simulation techniques continue to improve, evolutionary studies of model systems allow us to probe the dynamical state of observed clusters, and may offer key insights into the conditions under which star clusters are born. Second, these dense clusters are likely to be the sites of complex physical phenomena, such as stellar collisions and mergers \\citep{1999A&A...348..117P,2002ApJ...576..899P,2004ApJ...604..632G,2006MNRAS.368..141F}, placing them at the interface of stellar dynamics, stellar and binary evolution, and stellar hydrodynamics. Partly as a result of this overlap of traditionally distinct astrophysical disciplines, the past few years have seen an upsurge in interest in modeling dense stellar systems, which pose significant theoretical and technical challenges to researchers \\citep{2003NewA....8..337H,2003NewA....8..605S}.\\footnote{See for example {\\tt http://manybody.org/modest}.} Finally, since such clusters may plausibly be the progenitors of globular-cluster like systems, the studies presented here also offer valuable clues to the early evolution of the globular cluster systems observed in many galaxies. The early evolutionary conditions considered here may also have important consequences for the present-day content of globulars. In performing simulations of young star clusters we run into an immediate problem. The initial conditions of these systems have been actively debated for many years, but no consensus has been reached. Models of star formation are as yet insufficiently advanced to provide definitive predictions of initial structure for large systems \\citep{2001ApJ...556..837K,2002ApJ...576..870P,2005MNRAS.356.1201B}, and we cannot simply run an observed cluster backwards in time, even if its parameters were all known to arbitrary accuracy. Rather, we start with a poorly determined but plausible initial state, evolve it forward in time, then attempt to match observable properties of our model cluster with actual clusters in the universe to assess the reasonableness of our initial choice. In this study we simulate young (age $\\aplt 100$\\,Myr) star clusters by integrating the equations of motion of all stars and binaries. We use the Starlab environment (Portegies Zwart et al 2001),\\nocite{2001MNRAS.321..199P} which acquires it greatest speed on the GRAPE-6 special-purpose computer (GRAvity PipE, Makino et al 1997; 2003).\\nocite{1997ApJ...480..432M,2003PASJ...55.1163M} The calculations presented here were performed on the GRAPE hardware at the University of Tokyo, the MoDeStA\\footnote{See {\\tt http://modesta.science.uva.nl}} platform in Amsterdam, and the GRAPE-6 system at Drexel University. Both stellar and binary evolution are included self-consistently in our models. ", "conclusions": "We have simulated star clusters with highly concentrated initial density profiles and a wide range of initial relaxation times, from birth to an age of about 100\\,Myr. Our initial conditions include 10\\% hard primordial binaries, and the simulations incorporate the effects of stellar and binary evolution and binary dynamics. In this second paper on these simulations, we report on the time variation of the structural and internal composition parameters describing our model clusters, and compare our results directly to the sample of relatively young and isolated star clusters in the Large Magellanic Cloud. On the basis of this comparison, we conclude that the range of core radii and concentrations found in our simulated clusters is consistent with observations of the LMC clusters, and we argue that most of the LMC clusters are born with initial half-mass relaxation times of 200\\,Myr to 600\\,Myr and high central concentrations---$c \\simeq 2.7$ (King parameter $\\Wo \\simeq 12$). The only clear exception to this is the star cluster R\\,136 in the 30 Doradus region, which matches our simulation with an initial relaxation time of about 80\\,Myr. Due to mass segregation, the mass function of intermediate-mass main-sequence stars becomes as flat as $\\alpha=-1.8$ in the central part of the cluster (where the initial Salpeter mass function had $\\alpha=-2.35$). In the outer regions, the mass function exponent is as steep as $\\alpha = -2.6$. By the end of the simulations, at 100\\,Myr, the overall cluster binary fraction is still $\\sim 10$\\,\\%, but in the core the fraction of binaries is somewhat higher ($\\apgt 12$\\%). By this time about 7\\% of the single stars are remnants, and their number is increasing gradually at a rate of about 0.1\\% per Myr. In our simulations a large number of blue stragglers are formed. At any time, however, no more than 100--200 blue stragglers are visible in the cluster. The largest numbers of blue stragglers are formed in the densest clusters. The distribution of blue straggler masses depends quite sensitively on the initial cluster density. The densest clusters tend to produce more massive, brighter and bluer blue stragglers than less dense clusters. The trends visible in our simulations are consistent with observations of current globular clusters. A population of dormant blue stragglers is formed early in the evolution of the cluster. They remain hidden on the main-sequence until they emerge above the turn-off as the cluster ages. The fraction of high-velocity stars of spectral type O and B is considerably smaller than the fractions observed in the Galactic field. Our simulations, however, incorporate both of the effects thought to be responsible for the acceleration of the observed OB runaways: supernova in evolving binaries and gravitational slingshots from multi-body scattering encounters. The discrepancy with observations of the numbers of OB runaways might conceivably be explained by the initial binary fraction, which in our simulations is only 10\\%. Shortly after formation, the cores of our simulated clusters become quite rich in compact stars. Up to an age of about 40\\,Myr the remnant population in cluster cores is dominated by stellar-mass black holes; after that time white dwarfs take over. Neutron stars are easily ejected from the clusters and there are only a few present at any time. The neutron star retention fraction is about 6--12\\%, whereas 50--70\\% of black holes are retained. Clusters with longer relaxation times have smaller retention fractions. Binaries containing black holes with main-sequence companions outnumber those containing a neutron star and a stellar companion. We conclude that these clusters may be relatively rich in x-ray binaries with a black hole as accreting object, at least up to ages of a few hundred Myr. Binaries containing two white dwarfs are quite common in our simulations, and 20--30 have sufficiently small orbital periods that gravitational radiation will bring the two white dwarfs into contact within a Hubble time. Interestingly, clusters with shorter relaxation time produce systematically fewer white-dwarf binaries that will merge within a Hubble time." }, "0607/astro-ph0607182_arXiv.txt": { "abstract": "A method of the calculation of optical parameters of the nonisothermal giant planet atmospheres was developed using detailed intensity data of Raman scattering. We have used the model of Morozhenko (A.V. Morozhenko, 1997) as a baseline. In such a way, using observational data of Uranus and Neptune (E.Karkoschka, 1994), the spectral values of ratio of optical depth components: aerosol and gas components $\\tau_a/\\tau_R$, absorbing and scattering components $\\tau_\\kappa/\\tau_R$, and also single scattering albedo of aerosol component corrected for Raman scattering $\\omega'$ were obtained (where $\\tau_a, \\tau_R$ are aerosol and gas components, and $\\tau_\\kappa$ is absorbing components of effective optical depths of the formation of diffusely reflected irradiation). The averaged value of ratio $\\tau_a/\\tau_R$ is 0.96 but it slowly decreases in the spectral range of 350-450nm for Uranus and $\\tau_a/\\tau_R$ is 1.35 for Neptune. ", "introduction": "The atmospheres of Uranus and Neptune are known to be composed predominantly of molecular hydrogen. Since there is so much $H_2$ in the atmospheres of outer planets and $H_2$ has a reasonably strong Raman spectrum, it is very important to attempt to understand the physics of planetary Raman scattering. Raman scattering is the incoherent non-resonance scattering of photons by a molecule. During molecular scattering process, the photon may loose energy according to certain molecular transitions. If the incident solar photon of frequency $\\nu_0$ is scattered, it will emerge at frequency $\\nu_0 \\pm \\bigtriangleup\\nu$, where $\\bigtriangleup\\nu$ is the frequency of the Raman transition of the molecule. In recent years the observational data of detailed intensity of Raman scattering in the giant planet spectra was proposed to use to determine the relative contribution of the aerosol component of atmosphere. In such a way, we can determine the values of aerosol to gas ratio of optical depth components $\\tau_a/\\tau_R$, and absorbing to scattering ratio $\\tau_\\kappa/\\tau_R$ (M.S.Dementiev,1992; A.V.Morozhenko, 1997). In these papers, the model of atmosphere was taken to be isothermal, while the real giant planet atmospheres have complex temperature profiles (G.F.Lindal, et al., 1987; G.F.Lindal, et al., 1990). The relative number of hydrogen molecules in the ortho- and para- state depends on the depth in the nonisothermal atmosphere, while it doesn't depend on the depth in the isothermal one. So, the detailed intensity of Raman scattering will depend on the effective optical depth of the formation of diffusely reflected irradiation. The method of accounting of the real temperature profile in computing of Raman scattering effects was developed by Morozhenko and Kostogryz (A.V.Morozhenko and N.Kostogryz, 2005). This paper presents a method of computation of optical parameters of the Uranus's and Neptune's atmospheres such as $\\tau_a/\\tau_R$ and $\\tau_\\kappa/\\tau_R$ considering detailed intensity of Raman scattering and using observational data of Uranus's and Neptune's atmospheres (E.Karkoschka, 1994) and experimental temperature profiles (G.F.Lindal, et al., 1987; G.F.Lindal, et al., 1990). Section 2 contains reviews of the model of atmosphere. Section 3 is devoted to the method of computation and sections 4 and 5 describe results of computation and some conclusions of this work. ", "conclusions": "In this paper we determined the values of aerosol to gas ratio of optical depth component for Uranus ($\\tau_a/\\tau_R=0.96$) and for Neptune ($\\tau_a/\\tau_R=1.35$), and spectral dependence of absorbing to scattering ratio $\\tau_\\kappa/\\tau_S$. We confirmed that ignoring of real temperature profile leads to 50 $\\%$ errors in determination of ($\\tau_a/\\tau_R$). Real spectral values of single scattering albedo corrected for Raman scattering, were obtained for spectral region 350-450 nm." }, "0607/astro-ph0607657_arXiv.txt": { "abstract": "We constrain the post-Newtonian gravity parameter $\\gamma$ on kiloparsec scales by comparing the masses of 15 elliptical lensing galaxies from the Sloan Lens ACS Survey as determined in two independent ways. The first method assumes only that Newtonian gravity is correct and is independent of $\\gamma$, while the second uses gravitational lensing which depends on $\\gamma$. More specifically, we combine Einstein radii and radial surface-brightness gradient measurements of the lens galaxies with empirical distributions for the mass concentration and velocity anisotropy of elliptical galaxies in the local universe to predict $\\gamma$-dependent probability distributions for the lens-galaxy velocity dispersions. By comparing with observed velocity dispersions, we derive a maximum-likelihood value of $\\gamma = 0.98 \\pm 0.07$ (68\\% confidence). This result is in excellent agreement with the prediction of general relativity that has previously been verified to this accuracy only on solar-system length scales. ", "introduction": " ", "conclusions": "" }, "0607/astro-ph0607019_arXiv.txt": { "abstract": "We study the response of the gaseous component of a galactic disc to the time dependent potential generated by N-body simulations of a spiral galaxy. The results show significant variation of the spiral structure of the gas which might be expected to result in significant fluctuations in the Star Formation Rate (SFR). Pronounced {\\it local} variations of the SFR are anticipated in all cases. Bursty histories for the {\\it global} SFR, however, require that the mean surface density is much less (around an order of magnitude less) than the putative threshold for star formation. We thus suggest that bursty star formation histories, normally attributed to mergers and/or tidal interactions, may be a normal pattern for gas poor {\\it isolated} spiral galaxies. ", "introduction": "The local Universe furnishes many illustrations of the fact that vigorous star formation is highly localised in both space and time. The most spectacular examples, circumnuclear starbursts can, over timescales of $\\sim 10$ Myr, sustain star formation rates comparable to entire galaxies from a region less than a kiloparsec across (Lehnert and Heckman 1996). Equally intense, but more spatially distributed, is the current burst of star formation (and Super Star Cluster formation) in the Antennae Galaxy (Whitmore et al 1999). Star formation in the Antennaa is self-evidently triggered by a major merger; there is also ample evidence, however, that bursts of star and cluster formation can be triggered by interactions with small satellite galaxies (i.e. by minor mergers; see, for example, Homeier and Gallagher 2002). In isolated galaxies, however, it is usually assumed that the global star formation rate (henceforth SFR) shows little time variation although, in the case of disc galaxies, star formation is clearly spatially localised along spiral arms. This assumption (of a roughly constant quiescent SFR) is partly based on the expectations of spiral density wave theory in which the disc is susceptible to non-axisymmetric instabilities that generate a {\\it long lived} pattern of spiral arms (Bertin et al 1989). Thus whereas the SFR may vary locally as material is compressed by the passage of spiral density waves, the pattern is invariant in a frame corotating with the spiral disturbance and therefore the global SFR is expected to be roughly constant. N-body simulations of self-gravitating discs however paint a very different picture of the development and longevity of spiral structures. For example, Sellwood and Carlberg 1984 performed N-body simulations of a stellar disc subject to both self-gravity and to a rigid halo potential. They found that the spiral structures that developed in their simulations were {\\it transient} features which formed and re-formed on a timescale comparable with the galactic rotation period. (See also Huber and Pfenniger 2001, Gerritsen and Icke 1997 and Bottema 2003 for further simulations producing transient, regenerative spiral structures). The development and dissolution of such structures means that at a given location the gas is no longer subject to periodic variations in the potential, as in classical spiral density wave theory, and thus one would expect irregular variations in the resulting SFR. Likewise, the pattern of evanescent structures may not even give rise to SFRs that are {\\it globally} constant, an impression that is strengthened by inspecting snapshots of the N-body simulations (see Figure 5 of Sellwood and Carlberg 1984), which suggest that potential troughs are much more pronounced at some epochs than others. In this paper we undertake a preliminary investigation of the response of a galaxy's gaseous component to such irregular potential variations. In this simple approach, we neglect the self-gravity of the gas, and its potential role in amplifying the structures in the N-body simulations, and instead study the response of non-self gravitating isothermal gas. In order to make qualitative statements about the spatial and temporal variations of the resulting SFR we have to make some assumption about how SFR depends on the surface density distribution. Here we simply associate the SFR with the fraction of the gas mass that is instantaneously compressed to a column density greater than some adjustable threshold value, $\\Sigma\\sub{crit}$ (see Kennicutt 1989 and Martin and Kennicutt 2001 for a discussion of such a star formation threshold in disc galaxies). Note that we do not actually deplete the gas surface density in response to this nominal star formation rate but that since the simulations extend only over a few galactic rotation periods, this is not a major shortcoming. The structure of the paper is as follows. In Section 2 we describe the N-body and gas dynamical simulations on which the study is based, in Section 3 we describe the resulting gas structures and in Section 4 interpret these time dependent structures in terms of a nominal star formation rate. Section 5 summarises our conclusions. ", "conclusions": "\\label{sec:conclusions} We have presented the response of an isothermal gas disc to the gravitational potential of an N-body simulation of an isolated spiral galaxy. Complex spiral structures are formed, which break and re-join continually on a timescale comparable with the local orbital period. This highly varying structure leads to the formation of many dense features, which are generally short-lived. This would be expected to correspond to variations in the global star formation rate, as events such as the merging of spiral features occur sporadically on timescales of a few galactic rotation periods. Since the amplitude of such bursts varies from burst to burst, it is possible that rarer, larger scale bursts (traditionally attributed to galactic interactions) may also occur in {\\it isolated} disc galaxies, due purely to the regenerative nature of spiral structures in stellar discs. We however stress that large scale variations in {\\it global} SFR are probably confined to galaxies with mean gas surface density considerably below the threshold value. The simulations also imply spatial and temporal variations in star formation rates that are considerably more complex than those generated by a long lived spiral mode, since now the lifetime of individual spiral features is comparable with the timescale on which they cross the disc. Mapping of resolved stellar populations in nearby galaxies offers a potential tool to examine the observational evidence for such effects (Kodaira et al 1999, Williams 2003). We stress that our simulations omit many processes that {\\it must} be important in real galaxies (e.g. interconversion between the star and gas phase, the effects of stellar feedback, self-gravity and equation of state of the gas and back reaction of the gas on the stellar dynamics). For this reason, the results in this paper are more illustrative than quantitative (and indeed, as we discuss in 4.1, it is not obvious how to quantify star formation variations in {\\it any} simulation where the scales on which star formation and feedback occur are so far below the resolution of the simulations). Finally, we note that N-body simulations (such as those of Sellwood and Carlberg 1984), which find spiral structures to be recurrent, short-lived patterns, have attracted much interest from those studying the origin of spiral arms in galaxies. It is only rather recently that some of the {\\it observational} signatures of such transient structures have begun to be explored (Sellwood and Preto 2002, Sellwood and Binney 2002). Irregular variations in star formation rate are an obvious corollary of such a picture. The study of temporal and spatial variations in star formation rates in nearby galaxies may thus shed some light on the nature and origin of spiral structure in disc galaxies." }, "0607/astro-ph0607533_arXiv.txt": { "abstract": "The equations describing a two-component cosmological fluid with linearized density perturbations are investigated in the small wavelength or large $k$ limit. The equations are formulated to include a baryonic component, as well as either a hot dark matter (HDM) or cold dark matter (CDM) component. Previous work done on such a system in static spacetime is extended to reveal some interesting physical properties, such as the Jeans wavenumber of the mixture, and resonant mode amplitudes. A WKB technique is then developed to study the expanding universe equations in detail, and to see whether such physical properties are also of relevance in this more realistic scenario. The Jeans wavenumber of the mixture is re-interpreted for the case of an expanding background spacetime. The various modes are obtained to leading order, and the amplitudes of the modes are examined in detail to compare to the resonances observed in the static spacetime results. It is found that some conclusions made in the literature about static spacetime results cannot be carried over to an expanding cosmology. ", "introduction": "The analysis of cosmological perturbations in the Newtonian limit is a well studied problem in theories of structure formation, and it may be supposed that there is little left to learn from this theory. Most of the effort has gone into the study of the one-component cosmological fluid equations, and the results have been well expounded in many standard texts \\cite{padman,peebles,weinberg,kolb,zeld}. There is, however, still a wealth of problems remaining in the detailed analysis of two-component cosmological fluids and their linearized gravitational perturbation modes. In particular, if pressure effects are included so that the Jeans instability becomes an issue, the equations present a considerable analytic challenge, and a range of new physical effects become apparent. Some of these effects have been studied in the contrived case of a static spacetime background \\cite{russians1,carvalho}. In this scenario there is no expansion, so that the mathematics is considerably simplified, and solutions can easily be found. This is useful to gain some qualitative idea about physical phenomena observable, but to gain a true picture in a cosmological context, the expanding background spacetime given by the Friedmann-Robertson-Walker cosmologies is required. There have been a variety of studies of the multi-component cosmological fluid equations, ranging from some relatively specific applications under certain cosmological scenarios \\cite{russians2,fargion}, to a broad mathematical study and classification \\cite{haubold}. A discussion of the application and validity of some of the equations mentioned in these previous studies, together with the solution of an unsolved set of two-component post-recombination equations, has recently been undertaken by the authors \\cite{paper1}. The system of equations described the interaction between a dark matter and baryonic component in the Newtonian regime (density fluctuations on scales well within the Hubble radius). A series expansion of the solutions for small wavenumber $k$ (large scales) was presented. This allowed comparison with some of the previous work, in particular with the Meijer G-function classifications given by \\cite{haubold}. This region of $k$-space is also interesting because it is the region in which the Jeans instability is known to occur. In this paper we wish to complete this study by examining the large $k$ asymptotic region of the solutions. Such a study is worthwhile, in order to make contact with the static spacetime results of \\cite{carvalho}. Although not realistic as cosmological solutions, these results displayed a number of little known physical phenomena associated with the linearized modes, which we wish to expand on here. The techniques required to analyze the expanding universe solutions are also of interest in their own right mathematically, where a generalized WKB method will be expounded. It is possible to make a comparison with the work done in cosmological plasma physics in an Einstein-deSitter background \\cite{plasma,gailis}. This is interesting because of the mathematically very similar form of fluid equations for both type of systems, which is due to the similarity of the electromagnetic and gravitational forces. Thus mathematical techniques employed in the analysis of plasma equations will be useful in this paper, and give clues as to how to proceed with some challenging mathematical analysis of gravitational density perturbation modes. The paper is to be organized as follows. The relevant equations will be introduced in Section~2. The discussion will then be focused in Section~3, by reconsidering the two-component modes in a static spacetime. This investigation is by necessity of a qualitative nature, but gives a useful introduction to the concepts and interesting physical effects not found in the standard one-component analysis. The work of \\cite{carvalho} will also be extended. The expanding universe baryonic and dark matter equations will then be considered in Section~4. The short wavelength (WKB) approximation will be utilized to complete the study of these equations initiated in \\cite{paper1}. The relevance of the previous work on static spacetime systems is revealed through this analysis. This will allow meaningful conclusions to be drawn about this whole area of study, and point to where promising future work may lie. These aspects are discussed in Section~5. ", "conclusions": "The structure and behavior of the eigenvalues and eigenvectors of two-component cosmological density perturbations have been studied in great detail in this paper. We have reviewed the previous work done in a static spacetime background, and produced further results in this simple context. This has enabled the far more difficult expanding universe problem to be tackled. The WKB method employed has produced the full leading order behavior of all the modes in the Einstein-deSitter expanding Universe scenario. These solutions represent acoustic oscillations for wavelengths much smaller than the Jeans scale. The Jeans scale of the mixture has arisen in a natural way out of the analysis of the eigenvalues obtained through the WKB method, with some interesting interpretation. It is now a straightforward task to adapt the methods developed here to study a variety of further cosmological plasma modes. The ion-sound and two-component Langmuir oscillations would follow directly from the results presented here, and more complicated modes involving magnetic fields could also be obtained by similar procedures. We have also obtained the time- and $k$-dependent amplitudes of the modes in a fairly general setting (the one restriction being initial perturbations beginning from rest). These results have shown that the amplitudes are very constant in the region of interest. The existence of resonances in the amplitudes found for static spacetime results do not apply here, as all resonances occurred for wavenumbers far smaller than $k_M$. Thus a resonant amplitude cannot be viewed as a mechanism for producing structures of a preferred scale in a two-component model. The eigenvalues derived in this paper do also not have any direct physical interpretation around the Jeans scale, or for small $k$ expansions of the solutions of Eqs.(\\ref{canonicalB}) and (\\ref{canonicalD}). Thus the results obtained in this paper must be considered to be restricted to the parameter regions considered here. It may be interesting to investigate different models such as a three-component HDM+\\-CDM+\\-baryon fluid, or models involving a cosmological constant (especially given the weight of current observations \\cite{perlmutter}, \\cite{bennett}, \\cite{perlmutter2}). The analytics would become considerably more complicated, but some other interesting resonant scales may be found with a direct implication for structure formation." }, "0607/astro-ph0607643_arXiv.txt": { "abstract": "We consider the contribution of the Urca-type processes to the bulk viscosity of several spin-one color-superconducting phases of dense two-flavor quark matter. In the so-called transverse phases which are suggested to be energetically favorable at asymptotic densities, the presence of ungapped quasiparticle modes prevents that spin-one color superconductivity has a large effect on the bulk viscosity. When all modes are gapped, as for one particular color-spin-locked phase, the effect on the viscosity can be quite large, which may have important phenomenological implications. ", "introduction": "\\label{sec0} Dissipative processes play an important role in the evolution of neutron stars. These processes are governed by transport coefficients, such as the heat and electrical conductivities, the neutrino diffusion coefficient, as well as the shear and bulk viscosities. The conductivities are important for stellar cooling as well as for the magnetic field decay. Shear viscosity dampens differential rotation in a star and, thus, leads to a uniform rigid-body rotation. Bulk viscosity, on the other hand, dampens density oscillations inside the star. Both differential rotation and oscillations could be excited in newly formed (hot) neutron stars, or could develop in old (cold) stars due to external perturbations, e.g., such as matter accreted from a companion star. It is interesting to note that, in the absence of viscosity, all rotating stars would be unstable. The reason is that such stars spontaneously develop instabilities as a result of the emission of gravitational waves \\cite{Chandra1,Chandra2,Frid1,Andersson,Fried2} (for reviews on this topic see, e.g., Refs.~\\cite{Anderssonreview, Lind-lect}). The so-called r-mode (or rotation-dominated) instabilities might be the most important ones. They can develop at a relatively low angular velocity \\cite{Lind1}, and therefore may be relevant for a large number of compact stars. The main theoretical uncertainty in predicting whether the r-mode instabilities develop in a star lies in the poor understanding of the viscosity of dense baryon matter, as well as in the limited knowledge of the stellar composition. There is hope, however, that a systematic approach, based on a broad understanding of various properties of dense baryonic matter, can eventually result in a clear picture regarding the neutron star composition. The viscosity of nuclear and mixed phases of dense baryonic matter has been calculated under various conditions and assumptions over the last three decades \\cite{FlowersItoh1, FlowersItoh2,Sawyer,Jones1,Lindblom1,Lindblom2,Drago1,Haensel1,Haensel2, Chat}. The bulk viscosity of normal conducting strange quark matter was also calculated \\cite{Sawyer2,Madsen}. The latter might be relevant if the baryon density in the central regions of neutron stars is so high that matter becomes deconfined. The physical conditions in the interior of such stars are quite unique: this is the only place in the Universe where a deconfined state of cold and dense baryonic matter can naturally exist. This possibility has attracted a lot of attention since the notion of quarks was introduced \\cite{Ivanenko1965,Ivanenko1969, Itoh1970,Iachello1974,Collins1975}. If deconfined quark matter does exist inside stars, it is most likely color-superconducting. (For reviews on color superconductivity, see Refs.~\\cite{RajWil,Alfordreview,Reddy2002, Rischke2003,Buballa2003,Huangreview,Shovkovy2004}.) It is therefore of great interest to study various transport properties of color-superconducting phases of quark matter. First attempts have already been made to estimate the heat and electrical conductivity \\cite{ShovEllis1,ShovEllis2}, as well as the bulk and shear viscosities \\cite{Madsenprl2,Manuel} in the color-flavor-locked (CFL) phase of quark matter. Also, in the case of the two-flavor color-superconducting (2SC) phase, one can argue that most transport coefficients are dominated by the two ungapped (blue) quasiparticles \\cite{Shovkovy2004}. (For a recent detailed study of the bulk viscosity in the 2SC phase, see Ref.~\\cite{Alford:2006gy}.) In this paper, we calculate the bulk viscosity of the four most popular spin-one color-superconducting phases of two-flavor (non-strange) quark matter: the color-spin-locked (CSL), planar, polar, and the {\\it A} phase \\cite{Iwasaki,PD1,SchaferSpin1,Schmitt:2002sc,BuballaSpin1,AlfordSpin1,SchmittSpin1}. One of these is likely to be the ground state of dense baryon matter if the spin-zero Cooper pairing of quarks is prevented by the constraints of charge neutrality and $\\beta$-equilibrium. Moreover, cooling calculations for neutron stars favor small gaps of the order of $1$~MeV \\cite{Grigorian:2004jq} which is the typical size of the gap in spin-one color superconductors \\cite{Iwasaki,PD1,SchaferSpin1,Schmitt:2002sc,BuballaSpin1,AlfordSpin1,SchmittSpin1}. The absence of strange quarks in the system may be natural if the medium-modified constituent value of the strange quark mass is larger than the corresponding value of the chemical potential. The generalization of this study to the case of spin-one color-superconducting strange quark matter will be reported elsewhere \\cite{SadShR}. As is well known, in fully gapped spin-zero color-superconducting phases, the thermal densities of the quark- and hole-type quasiparticles are suppressed exponentially by the energy gap, $n_{\\rm qp}\\propto\\exp(-\\phi_{0}/T)$ where $\\phi_{0}$ and $T$ are the values of the spin-zero gap and the temperature, respectively. This then translates into an exponential suppression of the quasiparticle contributions to the transport coefficients. In Ref.~\\cite{Madsenprl2} such an argument was used in order to get a simple estimate of the viscosity in the CFL phase. It should be noted, however, that many transport properties are not dominated by the quasiparticles when there exist Nambu-Goldstone excitations in the low-energy spectrum, as is the case in the CFL phase \\cite{ShovEllis1,ShovEllis2,Manuel}. The situation is expected to be different also in the case of spin-one phases which, in general, are not isotropic, and whose gap functions may have nodes for some directions of momenta. The effect of non-isotropic gaps and various topologies of nodes on the neutrino emission and the cooling rate of spin-one color-superconducting phases were recently discussed in detail \\cite{SSW2,Wang2006}. Following a similar approach, in this work we study the bulk viscosity. Since the order parameter in spin-one color superconductors breaks rotational invariance, dissipative hydrodynamics is more complicated than for isotropic media. For instance, we expect that the viscosity becomes a tensor \\cite{prb15Bhatta}. We shall avoid these complications by making the implicit assumption that the hydrodynamic equations are averaged over solid angle. In this way, only one (angular-averaged) bulk viscosity coefficient, $\\zeta$, will appear in the hydrodynamic equations and will have to be extracted from the (angular-averaged) neutrino emission rate. In addition, one of the spin-one superconducting phases studied here, namely the CSL phase, breaks baryon number, i.e., it is also a superfluid. Fortunately, it is not an anisotropic superfluid because the order parameter does not break rotational invariance. (Here we assume that the magnetic field of the star is not strong enough to align the spins of the Cooper pairs.) Nevertheless, isotropic superfluids still have a rather complicated hydrodynamic behavior, involving three (instead of only one) bulk viscosity coefficients \\cite{LL6,Son:2005tj}. However, it is not completely unrealistic to assume that the {\\em relative\\/} velocity between normal and superfluid components is negligible compared to the {\\em absolute\\/} velocity of the normal component. In this case, only one coefficient contributes to energy dissipation. In this sense, our treatment is completely analogous to that of Ref.\\ \\cite{Haensel1}, the difference being that here we consider quark matter instead of nucleonic matter. Let us finally note that the dissipative hydrodynamics of {\\em anisotropic\\/} superfluids is even more complicated, see for instance the case of superfluid He-3 \\cite{Graham1974,prb15Bhatta}. In the CSL phase the breaking of baryon number gives rise to a phonon as the corresponding Goldstone excitation. We neglect the contribution of the Goldstone mode to the bulk viscosity coefficient for the following reason. First, the effective theory for phonons is approximately scale-invariant at very low energies. For scale-invariant superfluids, however, two of the three bulk viscosity coefficients have been shown to vanish \\cite{Son:2005tj}. The remaining coefficient may be non-zero, but corresponds to dissipation due to relative motion of superfluid and normal component, which we have already assumed to vanish. The remainder of this paper is organized as follows. In the next section, we introduce the formalism for calculating the bulk viscosity in non-strange quark matter. In Sec.~\\ref{bulk-vis-normal} we calculate the bulk viscosity in the normal phase of two-flavor quark matter. Then, in Sec.~\\ref{bulk-vis-spin-1} we present our results for the bulk viscosity in spin-one color-superconducting phases. The discussion of the results is given in Sec.~\\ref{Conclusion}. ", "conclusions": "\\label{Conclusion} In this paper we have calculated the bulk viscosity for the normal phase as well as for four spin-one color-superconducting phases of two-flavor dense quark matter. The main contributions come from the Urca processes shown diagrammatically in Fig.~\\ref{fig-Urca_d_u_e}. Note that the results for the normal phase are also relevant for the 2SC phase. Indeed, after taking into account that there are two (blue) ungapped modes of quasiparticles in the low-energy spectrum of the 2SC phase, the low-temperature bulk viscosity is approximately given by the same expression (\\ref{general-form}), provided the following redefinitions are made: $\\zeta_{max}^{\\rm 2SC} =3\\zeta_{max}$ and $\\omega_0^{\\rm 2SC}=\\omega_0/3$, where the normal-phase quantities are given in Eqs.~(\\ref{MAXzeta}) and (\\ref{omega0}), respectively. The redefinitions account for the decrease of the weak rates by a factor of 3 at $T\\ll \\Delta_0$ where $\\Delta_0$ is the value of the 2SC gap. The microscopic calculations of the bulk viscosity in the spin-one color-superconducting phases suggests that quasiparticles with different types of gapless nodes (e.g., points or lines at the Fermi sphere) could potentially play a very important role. In the case of the transverse phases, however, the presence of a single ungapped quasiparticle mode washes out essentially all information about spin-one Cooper pairing, see Fig.~\\ref{bv-fig4}. The presence of non-zero quark masses may provide a gap for such a mode and the situation changes. In this paper, we briefly discussed such a possibility in connection with the CSL phase of Ref.~\\cite{ABBY}. The results are shown in Fig.~\\ref{bv-fig5}. In agreement with the general expectation, we find that the bulk viscosity often tends to decrease when there is Cooper pairing of quarks whose main effect is to suppress the rates of the weak processes. In some cases (e.g., at sufficiently low frequencies and/or at temperatures close to the critical value) the behavior may reverse because of the non-trivial dependence of the bulk viscosity on the suppression factor, see Eq.~(\\ref{zeta-sp1}). Such an increase of the viscosity in the color-superconducting CSL phase is seen, for example, in a range of temperatures below $T_c$ in Figs.~\\ref{bv-fig4} and \\ref{bv-fig5} in the case when $T_c=4$~MeV." }, "0607/astro-ph0607196_arXiv.txt": { "abstract": "The presumed Wolf-Rayet star progenitors of Type Ib/c supernovae have fast, low density winds and the shock waves generated by the supernova interaction with the wind are not expected to be radiative at typical times of observation. The injected energy spectrum of radio emitting electrons typically has an observed index $p=3$, which is suggestive of acceleration in cosmic ray dominated shocks. The early, absorbed part of the radio light curves can be attributed to synchrotron self-absorption, which leads to constraints on the magnetic field in the emitting region and on the circumstellar density. The range of circumstellar densities inferred from the radio emission is somewhat broader than that for Galactic Wolf-Rayet stars, if similar efficiencies of synchrotron emission are assumed in the extragalactic supernovae. For the observed and expected ranges of circumstellar densities to roughly overlap, a high efficiency of magnetic field production in the shocked region is required ($\\epsilon_B\\approx 0.1$). For the expected densities around a Wolf-Rayet star, a nonthermal mechanism is generally required to explain the observed X-ray luminosities of Type Ib/c supernovae. Although the inverse Compton mechanism can explain the observed X-ray emission from SN 2002ap if the wind parameters are taken from the radio model, the mechanism is not promising for other supernovae unless the postshock magnetic energy density is much smaller than the electron energy density. In some cases another mechanism is definitely needed and we suggest that it is X-ray synchrotron emission in a case where the shock wave is cosmic ray dominated so that the electron energy spectrum flattens at high energy. More comprehensive X-ray observations of a Type Ib/c supernova are needed to determine whether this suggestion is correct. ", "introduction": "SNe Ib/c (Type Ib and Ic supernovae) are hydrogen free, or nearly hydrogen free, and are thought to have stripped massive star core, i.e. Wolf-Rayet star, progenitors. The observed rate of SNe Ib/c indicates that they make up about 1/4 of massive star supernovae \\citep[e.g.,][]{DF99} and thus are an important mode of massive star death. Although single very massive stars may end their lives as Wolf-Rayet stars, they are probably not in sufficient numbers to account for all of the SNe Ib/c \\citep{WL99}; binary progenitors are likely to be important contributors to the rate. Considerable interest in SNe Ib/c has been generated by their relation to GRBs (gamma-ray bursts). The spatial and temporal correlation of the Type Ic SN 1998bw with the burst GRB 980425 led to the probability of a chance superposition of only $10^{-4}$ \\citep{Gal98}, and the extraordinary energy displayed by the optical \\citep{Iwa98} and radio \\citep{Kul98} emission from the supernova made the identification even more secure. The case was clinched by the finding of Type Ic supernovae similar to SN 1998bw in the cosmological GRBs 030329 and 031203 \\citep{Sta03,Mal04}. The broad line character of the optical spectrum of SN 1998bw, with velocities up to $60,000\\kms$, led to the view that SNe Ib/c with such broad lines were likely to be related to GRBs, % including SN 2002ap \\citep{Maz02}. However, SN 2002ap was a relatively weak radio \\citep{BKC02} and X-ray source. On the basis of the radio emission, \\cite{BKC02} argued that the energy in high velocity ejecta was low, so that a connection to GRBs was unlikely for this object. Although it is not known which SNe Ib/c have a GRB connection, it is clear at least some of them do and this led to the prediction that late radio observations of SNe Ib/c could yield the detection of an off-axis GRB jet that slowed down and radiated more isotropically \\citep{Pac01}. Searches for such emission were undertaken \\citep{Sto03,Ber03,Sod06a}, leading to the detection of luminous radio and X-ray emission from SN 2001em \\citep{Sto04}, SN 2003L \\citep{Sod05}, and SN2003bg \\citep{Sod06b}. In the case of SN 2003L, broad-lined optical emission is not present \\citep{Mat03}, and in no case is there direct evidence for a GRB connection. SNe Ib/c were generally not detected at radio wavelengths, showing that such luminous supernovae are rare \\citep{Sod06a}. Although there have been discussions of the radio and X-ray emission from individual SNeIb/c, there has been little discussion of their overall properties and how they relate to the mass loss properties of the progenitor star. We consider these properties and discuss the mechanisms for X-ray emission from the supernovae. In \\S~2, the observed properties of SNe Ib/c are briefly reviewed. The hydrodynamics of the interaction and the resulting velocities are estimated in \\S~3. The expected relativistic electron spectrum, allowing for loss processes, is treated in \\S~4 and the resulting synchrotron radio emission in \\S~5. The possible mechanisms for the X-ray emission and their application to the observations are discussed in \\S~6. A discussion of the results is in \\S~7. ", "conclusions": "We have found that the radio and X-ray emission from SNe Ib/c can be accounted for by supernova interaction with the wind from the Wolf-Rayet star progenitor. Following the arguments of \\cite{Che98}, synchrotron self-absorption is generally responsible for the early absorption of the radio emission. The radio synchrotron emission from most of the SNe Ib/c considered here imply an electron energy index $p\\approx 3$. The radio spectral index of SN 1998bw is distinctly flatter, with $p\\approx 2.5$. The only other SN Ib/c with a relatively flat spectrum is SN 2002ap, with $p\\approx 2$ in the model of \\citet{BF04}. Interestingly, SN 2002ap has the highest velocity of any of the supernova in our group of normal Ib/c events. It may be that as the shock velocity increases and relativistic effects become important, there are changes in the character of the particle acceleration. The SNe Ib/c provide the opportunity to examine particle acceleration at shock velocities ($\\sim0.1c$) intermediate between those present in Galactic supernova remnants and those in GRBs. Although the properties of the synchrotron radio emitting particles are clear ($p\\approx 3$ is typical), the properties of the X-ray synchrotron emission are uncertain. If the circumstellar densities are typical of Wolf-Rayet stars, thermal radiation cannot explain the observed X-ray emission and a nonthermal mechanism is required. Although inverse Compton scattering of photospheric photons might be able to explain X-ray emission close to maximum optical light, it generally fails to explain late emission. If this emission is synchrotron radiation, there must be a flattening of the particle spectrum at high energy, as can occur in diffusive acceleration in cosmic ray dominated shock waves. More detailed X-ray observations are needed to confirm this picture. We expect the early X-ray luminosity to be related to the optical luminosity of the supernova if it is inverse Compton emission; as the supernova fades, synchrotron emission may become the dominant mechanism, with a relatively slow rate of decline. An additional piece of information from our interpretation is the magnetic field in the shocked region. If the circumstellar densities are typical of Wolf-Rayet stars, we find that efficient amplification of magnetic fields is needed, with $\\epsilon_B\\sim0.1$. A possible reason for the field is a streaming instability that accompanies cosmic ray acceleration \\citep{B04}, which may explain the high amplification at the high shock velocities present in SNe Ib/c. With this mechanism, $\\epsilon_B\\propto v_s$, so the value of $\\epsilon_B$ would be smaller in Galactic supernova remnants, as is observed." }, "0607/astro-ph0607475_arXiv.txt": { "abstract": "We analyze Rossi X--ray Timing Explorer ({\\it RXTE}) Proportional Counter Array (PCA) data of the transient low mass X--ray binary (LMXB) system 1A 1744--361. We explore the X--ray intensity and spectral evolution of the source, perform timing analysis, and find that 1A 1744--361 shows `atoll' behavior during the outbursts. The color-color diagram indicates that this LMXB was observed in a low intensity spectrally hard (low-hard) state and in a high intensity `banana' state. The low-hard state shows a horizontal pattern in the color-color diagram, and the previously reported `dipper QPO' appears only during this state. We also perform energy spectral analyses, and report the first detection of broad iron emission line and iron absorption edge from 1A 1744--361. ", "introduction": " ", "conclusions": "" }, "0607/astro-ph0607313_arXiv.txt": { "abstract": "We present the results of a spectroscopic survey of the Ca II H \\& K core strengths in a sample of 60 solar-type stars that are members of the solar-age and solar-metallicity open cluster M67. We adopt the HK index, defined as the summed H+K core strengths in 1 {\\AA} bandpasses centered on the H and K lines, respectively, as a measure of the chromospheric activity that is present. We compare the distribution of mean HK index values for the M67 solar-type stars with the variation of this index as measured for the Sun during the contemporary solar cycle. We find that the stellar distribution in our HK index is broader than that for the solar cycle. Approximately 17\\% of the M67 sun-like stars exhibit average HK indices that are less than solar minimum. About 7\\%-12\\% are characterized by relatively high activity in excess of solar maximum values while 72\\%-80\\% of the solar analogs exhibit Ca II H+K strengths within the range of the modern solar cycle. The ranges given reflect uncertainties in the most representative value of the maximum in the HK index to adopt for the solar cycle variations observed during the period A.D. 1976---2004. Thus, $\\sim$ 20\\%-30\\% of our homogeneous sample of sun-like stars have mean chromospheric H+K strengths that are outside the range of the contemporary solar cycle. Any cycle-like variability that is present in the M67 solar-type stars appears to be characterized by periods greater than $\\sim$ 6 years. Finally, we estimate a mean chromospheric age for M67 in the range of 3.8---4.3 Gyr. ", "introduction": "A key objective of investigations in the field of solar-stellar astrophysics is to provide insight on the nature and possible range of solar magnetic field-related atmospheric activity through observations of stellar analogs of the Sun. In this way, we immediately obtain information on the long-term variability of the Sun that otherwise would not be possible (or practically feasible) with, for example, the modern solar data base of high-precision, chromospheric Ca II H and K observations of about three decades (e.g., White \\& Livingston 1978; Livingston 1994). This is especially important given that the amplitude of solar and stellar variations in brightness are usually highly correlated with cycle variations in chromospheric emission (Hudson 1988;Radick 1991; Fr\\\"{o}hlich 1994; de Toma et al. 2004). Thus, the range of cycle variability in a sample of solar-type stars can be indicative of the magnitude of total brightness changes (Zhang et al. 1994) as well as the nature of variability in the spectral irradiance or the particle fluence in the Sun. These forms of solar variability, in turn, may affect global climate change (e.g., Soon et al. 1996; Soon et al. 2000a,b; Foukal 2003; Stott et al. 2003; Haigh 2003, 2005; Soon \\& Yaskell 2004 and references therein; Lean et al. 2005). In addition to the impact on our understanding of the role of the past and future variability of the Sun in global climate, the characteristics of cycle-related variability are a manifestation of the nature of the dynamo mechanism that is operative in the Sun and sun-like stars (Saar \\& Baliunas 1992; Saar 2002). Furthermore, the calibration of the empirical correlation between stellar age and chromospheric activity (Walter \\& Barry 1991; Soderblom et al. 1991; Simon 1992; Donahue 1998) will be affected by the intrinsic dispersion introduced by the cycle-related variation in chromospheric emission. As an illustration of the magnitude of this effect, if we combine the 25\\% range in the Ca II index from minimum to maximum in the solar cycle, as measured by Livingston \\& Wallace (2003) from 1975 to 2003, with the exponential decay law proposed by Walter \\& Barry (1991) for Ca II then the apparent uncertainty in the age of the Sun arising from its cycle variability would be $\\sim$ 21\\%, or an age-spread of roughly 4 Gyr -- 6 Gyr (see also Baliunas et al. 1998, their Fig. 4). Recent studies of the comparative properties of the Sun and sun-like stars have emphasized field stars in the solar neighborhood. This approach yields samples that are accessible to detailed and frequent observation with small-to-moderate aperture telescopes, and includes the extensive archive of Ca II H and K observations from the Mt. Wilson program (Wilson 1978; Vaughan et al. 1978; Baliunas et al. 1995). Baliunas \\& Jastrow (1990) utilized data from the Mt. Wilson program to examine the nature and distribution of activity in 74 stars that were selected to be solar-type according to a set of criteria that included age, mean level of chromospheric activity, and mass as inferred from photometric $B - V$ color. They estimated the potential range of solar magnetic activity over time scales of centuries, and the approximate magnitude of the associated changes that would be implied in the solar irradiance, particularly in the regime of exceptionally low-activity that is presumed to have occurred during the Maunder Minimum (White 1978). Among the distinctive features of the tentative results discussed by Baliunas \\& Jastrow (1990) were a significant width of the distribution of mean chromospheric activity measured at random phases of centuries-long variability and a possible bimodal distribution of mean chromospheric activity. They inferred from their stellar sample, using a particular method of sampling the time series of observations, that exceptionally quiescent levels of activity occurred at a frequency of about one-third of the time. This estimate, in turn, was primarily associated with four stars in their sample that exhibited little variability (i.e., instrumentally flat time series in the Mt. Wilson S-index of relative Ca II H and K strength) over periods of approximately two decades. Baliunas \\& Jastrow (1990) suggested that these objects were in the stellar counterpart of the solar Maunder-minimum state. In a subsequent study of a subsample of 10 solar-type stars, Zhang et al. (1994) cautioned that the small number of nearby solar-type stars, and the heterogeneity that is inherent in such field star samples, needed to be addressed. Hall \\& Lockwood (2004) reexamined the conclusions by Baliunas \\& Jastrow (1990) using a refined sample of solar counterparts in the field as monitored since 1994 with their Solar-Stellar Spectrograph (SSS) at Lowell Observatory in Flagstaff, Arizona (Hall \\& Lockwood 1995). These investigators also find that one-third of the solar counterparts in their sample do not exhibit cyclic behavior, in agreement with the results of Baliunas \\& Jastrow (1990). However, in contrast to the earlier study, the distribution of H and K activity in the Hall \\& Lockwood (2004) sample is unimodal rather than bimodal. Moreover, these investigators note that the level of chromospheric H and K emission in the non-cyclic, flat stars spans a range from levels below that of solar minimum to greater than that seen at solar maximum. Hence, as Hall \\& Lockwood conclude, non-cyclic behavior is not necessarily associated with relatively reduced chromospheric activity. These kinds of studies illustrate the crucial discrepancies in conclusions that can arise from investigations of heterogeneous samples of field stars. For example, in the case of an ostensibly solar-type star that exhibits a low-level of activity, it is particularly critical to know its evolutionary status before concluding that it is a Maunder-minimum candidate. In fact, Wright (2004) claims, on the basis of an analysis of $Hipparcos$ data and a transformation of the Mt. Wilson S-index to a normalized chromospheric emission flux ($R^{\\prime}_{HK}$), that nearly all the stars thus far identified as Maunder minimum candidates are actually evolved or subgiant stars with intrinsically low (and constant) chromospheric activity rather than solar-type stars that are in a temporary state of quiescence. A further examination of this result based on high resolution H \\& K line spectra combined with an independent flux calibration technique (e.g., following the methods of Linsky et al. 1979) is merited. In this investigation, we extend these pioneering studies to a sample of solar-type cluster members that are homogeneous in age and chemical composition. The open cluster M67 is an especially appropriate target of observation since it is approximately the same age (about 5 Gyr $\\pm$ 1 Gyr; Demarque, Guenther \\& Green 1992) and of the same metallicity as the Sun (Barry \\& Cromwell 1974). Thus, the solar-type members represent analogs of the Sun at random phases in their activity cycles. Therefore, a survey of the ``Suns of M67\" should reveal the possible range of {\\it solar} chromospheric activity, thus immediately yielding information on the potential long-term variability of the Sun. We discuss in $\\S2$ the observations and reduction. In $\\S3$, we describe the methodology adopted for the calibration of the stellar spectra to relative intensity and discuss our estimates of the random and systematic errors. We present the results in $\\S4$ followed by discussion in $\\S5$. We summarize our findings and indicate the future directions of this program in $\\S6$. ", "conclusions": "The broader distribution in Ca II H+K core strengths in the M67 solar-type stars, compared to that of the Sun during the contemporary epoch, suggests that the potential excursion in the amplitude of the solar cycle is greater than what we have seen so far. The stars with values of the HK index noticeably smaller than that of solar minimum may be in a prolonged state of relative quiescence analogous to the Maunder Minimum of the Sun during 1645-1715 C.E. when visible manifestations of solar activity, e.g., sunspot number, substantially decreased for a period much longer than the sunspot cycle. That persistent period of reduced solar activity occurred during the extreme phases (ca. mid-16th through 17th century) of a climatic anomaly characterized by a general cooling and referred to as the Little Ice Age. That anomaly, like the Medieval Warm Period which preceded it, is climatically, spatially and temporally complex, with evidence of regional variations in conditions indicating a sensitivity to a variety of forcings, including that due to a variable Sun (Rind 2002). The plausible assumption of a multi-decadal period of decreased total solar irradiance accompanying the observationally established low levels of solar magnetic activity has been postulated as a contributing factor to the sharply deteriorated conditions of the 17th-century (see, e.g., Eddy 1976; Foukal \\& Lean 1990; Hoyt \\& Schatten 1997; Soon \\& Yaskell 2004). However, quantitative estimates of the presumed decrease in the solar irradiance during this time are uncertain; recent work (Wang, Lean \\& Sheeley 2005; Lean et al. 2005) concludes that the secular increase in the total irradiance from the Maunder Minimum to the current cycle minima is less than the initial estimates that appeared in the literature beginning over a decade ago. The $\\sim$ 17\\% fraction of M67 sun-like stars in our sample that is at this relatively quiescent level would represent the frequency at which the Sun enters a Maunder Minimum. This frequency is lower than the estimated frequency that the Sun exhibits especially low levels of magnetic field-related activity. In particular, Damon (1977) inferred from the proxy terrestrial record of solar magnetic activity based on the $^{14}$C radioisotope that the Sun has spent about one-third of the time in the past several millennia in magnetic activity minima. Our results for the fraction of M67 solar-type stars that exhibit especially quiescent levels of activity is also lower than the $\\sim$ one-third frequency seen in samples consisting of field solar-type stars (Baliunas \\& Jastrow 1990; Hall \\& Lockwood 2004). However, it is higher than the fraction Wright (2004) finds when subgiants are removed from the Mt. Wilson field star sample. Our M67 sample of solar-type stars, which is presumably homogeneous in age and metallicity, does not exhibit the bimodal distribution that appeared in the heterogeneous sample of field stars studied by Baliunas \\& Jastrow (1990). From a purely qualitative perspective, our results imply that exceptional quiescence, such as that during a Maunder Minimum-like episode, is simply a low-amplitude extension of the solar dynamo rather than a separate dynamo mode. In this regard, Ribes \\& Nismes-Ribes (1993) concluded from their analysis of historical observations that the solar cycle persisted through the Maunder Minimum, though at exceptionally low amplitudes in terms of sunspot number. Likewise, the occurrence of activity levels that exceed contemporary solar maximum values could be a manifestation of the excursion to enhanced levels of activity in a single mode of dynamo operation. Our results ($\\S$4) would suggest that exceptionally high levels of activity can occur at a frequency of $\\sim$ 7\\% - 12\\% of the time. We present in Fig. 8 another depiction of the comparison between the solar cycle and the M67 solar-type star distributions, respectively, in the form of cumulative probability functions of the HK index. The curve for the M67 sample reflects the broader distribution in HK index compared to the solar cycle. Inspection of Fig. 8 readily reveals that the median values for the M67 and the modern solar cycle distributions, respectively, are nearly identical. Finally, a Spearman rank correlation test applied to the two distributions in Fig. 8 yields a moderately high correlation at $\\rho$ = 0.881. However, if we apply this test after omitting the ``supersolar\" stars with HK $>$ 300 m{\\AA} then we find an extremely high correlation at $\\rho$ = 0.978. This result, while not conclusive, is consistent with an activity-cycle origin, analogous to the solar cycle, for the observed HK index distribution for the solar-type stars in M67. The HK index distribution in M67 according to color, as illustrated in Figs. 4-5, merits further consideration. The distribution in Fig. 5(a) appears skewed with a tail toward higher activity stars with HK $\\gtrsim$ 200 m{\\AA}. A comparison with Fig. 4 reveals that the peak in the distribution consists mainly of bluer stars with intrinsic $B-V$ colors $\\lesssim$ 0.60. These objects may indeed be relatively more chromospherically active in this color range. However, it is important to recognize that the non-chromospheric, radiative equilibrium contribution to the core flux in the Ca II H \\& K lines will be relatively higher in hotter stars and decline toward cooler objects (Linsky et al. 1979; Noyes et al. 1984). A reliable calibration of the radiative equilibrium flux for the 1 {\\AA} bandpasses that define the HK index has not yet been done. The distribution in Fig. 5(b) is especially interesting because this subsample is the most photometrically similar to the Sun. In particular, the range of colors that has been quoted in the literature for the Sun is 0.63 $\\leq~B-V~\\leq$ 0.67 (VandenBerg \\& Bridges 1984, see their Table 2). Hence, the 21 M67 sun-like stars we observed in this color range are photometric counterparts of the Sun and thus can be considered true solar analogs, following the definition of a solar analog as given by Cayrel de Strobel \\& Bentolilla (1989; also see Cayrel de Strobel 1996). In the case of these M67 solar analogs, 2 stars or 10\\% are in a state of enhanced activity with HK indices in excess of solar maximum values while 4 stars (about 19\\%) exhibit values of HK that are below that of the contemporary solar minimum HK index; 71\\%, or 15 stars, have an average HK index that is within the range of the modern solar cycle. Thus, 29\\% of the solar analogs are characterized by levels of chromospheric activity that are outside the range of the modern solar cycle. In brief summary, a significant fraction of this homogeneous sample of solar analogs exhibits activity that is, in amplitude, consistent with the solar cycle as seen in the chromospheric Ca II resonance lines. The fraction that is outside of the cycle range can be indicative of the frequency of excursions in the solar cycle to either enhanced levels of activity or unusually quiescent episodes (e.g., Maunder minima). Those stars that are slightly cooler and less massive than the Sun in the color range of 0.68 $\\leq~(B-V)_0~\\leq$ 0.72 also show substantial overlap with the solar cycle with 6/11 stars in the cycle range; 3-4 stars are more active than solar maximum and 1-2 objects have values of the HK index less than the contemporary solar minimum. Similarly, the coolest stars observed in our sample with intrinsic $B-V$ colors in the range of 0.73 -- 0.76 include 2 stars more quiescent than solar minimum, no objects more active than solar maximum and 4 stars in the range of the solar cycle. In brief summary, the solar-type stars in this approximately solar-age and solar-metallicity cluster exhibit levels of chromospheric activity, as represented by our HK index, that are substantially within the range of the solar cycle. Excursions outside of the solar cycle range of activity occur at a frequency of about 29\\% for the entire solar-type sample, and at frequencies broadly in the range of $\\sim$ 22\\%--45\\% for the subsets of the sample given in Fig. 5. \\subsection{Implications for Brightness Changes} Given the correlation between variations in chromospheric emission and changes in brightness or irradiance, as in the case of the Sun, it is of interest to examine the implications of the range of chromospheric activity we see in M67 for the possible range of brightness amplitude variations that may occur. In their long-term study of the photometric variability of field solar-type stars, Lockwood, Skiff \\& Radick (1997; their Fig. 16) present the mean amplitude of variation in photometric brightness as a function of log $R^{\\prime}_{HK}$, i.e., the net chromospheric radiative flux in the H + K lines, normalized by the stellar bolometric flux (= $\\sigma T_{eff}^4$). In order to estimate the potential range in amplitude of brightness changes in our sample of M67 sun-like stars, we must first estimate the net chromospheric radiative loss rates in the H \\& K lines from the HK index and the stellar effective temperature. We utilize the calibrations given by Hall \\& Lockwood (1996) as a function of $B-V$ color in order to estimate the effective temperature and to obtain the total surface flux in the H \\& K cores from the HK index. An empirical correction for the radiative equilibrium (photospheric) contribution to the H and K core flux as a function of $B-V$ color is given by Noyes et al. (1984). This approach yields an estimate of $R^{\\prime}_{HK}$. Using our estimates of this parameter from above, we find a range in log $R^{\\prime}_{HK}$ of $\\sim$ -4.35 to -5.1. From the results of Lockwood et al. (1997) this range in log $R^{\\prime}_{HK}$ corresponds to brightness variations extending from sun-like ($\\sim$ 0.1\\%) to nearly 3\\%, or considerably in excess of that of the contemporary Sun. Whether long-term brightness variability of these amplitudes actually occurs in the M67 solar-type stars will require verification through high-precision photometric monitoring. \\subsection{Variability} The acquisition of the data for this program over several observing seasons enables us to conduct a preliminary investigation of the nature of long-term variability in the HK index in the solar-type stars in M67. We display in Fig. 9(a)-(n) the variation of the HK index by season in those program objects for which we have measurements in at least two observing seasons. The panels are ordered by decreasing $B-V$ color and span the 1996-2002 seasons. We also include in Fig. 9, for comparative purposes, the annual mean values of the HK index for the Sun during this same period. Inspection of Fig. 9 reveals the presence of variability in the HK index that exceeds the formal errors in practically all the program objects. The time series are not sufficiently long to enable us to conclusively identify the presence of magnetic activity cycles, analogous to the solar cycle, and to measure actual cycle periods, though some cases are suggestive of a cycle-like variation. It does appear from the results displayed in Fig. 9 that if activity cycles are present then their periods are greater than $\\sim$ 6 years. We include in Fig. 10 further examples of comparable portions (in time) of the solar cycle variation based on mean annual HK index values for comparison with the panels in Fig. 9. In general, the variability seen in the M67 solar-type stars appears to be characterized by higher amplitudes. We summarize the seasonal variability in HK index in Table 2 for our program objects. The rms variation in HK index among the M67 solar-type stars ranges from 4.60 m{\\AA} to 92.1 m{\\AA} with a mean rms of 19.9 m{\\AA}. We also include in Table 2 the rms variations for the Sun for the examples given in Fig. 10 as well as for the period from 1976-2004. In the case of the latter, the rms deviation of the annual mean HK index for the Sun is 7.90 m{\\AA}, which is within the range seen for the M67 solar-type stars though significantly lower than the mean rms for the sample. Thus, we find that, while there are seasonal variations comparable to that seen in the Sun in some M67 solar-type stars, variability in the HK index substantially in excess to that of the Sun dominates our results. We do not see a clear correlation between root mean square deviation and the HK index values in the entire data-set (Fig. 11). However, inspection of Fig. 11 suggests a possible correlation of increasing rms deviation with HK index for values of rms $\\gtrsim$ 20 m{\\AA}. A more precise study of variability will require a long-term program, ideally with spectra obtained at higher spectral resolution in order to accurately assess the nature of the variability in the H and K line cores. We construct in Fig. 12 the histogram of all the HK index values measured seasonally for the solar-type stars. Not surprisingly, the distribution is broader than that in Fig. 3, which is based on the HK index for each star averaged over several seasons. We find in Fig. 12 that 20\\% of the HK index measurements are less than the modern solar minimum value while Ca II core strengths in excess of solar maximum occur at a frequency of 11\\%. Thus, about 31\\% of the HK index measurements are outside the range of the solar cycle in our sample of M67 solar-type stars when examined on a seasonal basis. \\subsection{Chromospheric Activity and Stellar Age} The calibration of the empirical correlation between chromospheric activity and stellar age (Skumanich 1972) can yield an additional tool for the determination of the ages of field stars. It is therefore important to investigate quantitatively the magnitude of the impact of intrinsic variations in chromospheric activity, such as magnetic cycle variability, on the calibration of this relation. Relations between the net radiative chromospheric flux parameter in the Ca II H \\& K lines, $R^{\\prime}_{HK}$, and age have been given by Soderblom et al. (1991) and subsequently refined and extended to younger ages $\\lesssim$ 1 Gyr by Donahue (1993; also see Donahue 1998). We adopt this age-activity calibration to examine the {\\it apparent} age spread implied by the dispersion in chromospheric activity in both M67 and the solar cycle. In the case of the solar cycle, we utilize the HK index values based on the actual measurements from the high resolution solar spectra in order to give a more realistic estimate of the solar chromospheric age. Inserting our estimates of $R^{\\prime}_{HK}$ from $\\S$5.1 into the age-activity calibration given by Donahue (1998) yields the results in Fig. 13, shown both for the solar-type stars in M67 and the contemporary solar cycle. The broad extent of the apparent age range due to the cycle variation in Ca II in the Sun is $\\sim$ 2.5 Gyr to about 6 Gyr. The mean chromospheric age of the Sun is approximately 4.26 Gyr. We note, parenthetically, that Baliunas et al. (1998) similarly estimated the apparent (chromospheric) age, or error in estimating the age of the Sun over the last four hundred years if chromospheric activity were sampled at a random phase of long-term variability, namely, 3 to 8 Gyr at the extremes with presumably a narrower range in apparent age for most of the period. The age range among the M67 solar-type stars implied by the range of observed chromospheric activity is even broader, extending from less than 1 Gyr to about 7.5 Gyr. The mean chromospheric age inferred from our sample of solar-type stars is approximately 3.75 Gyr while the median age for the M67 distribution is about 3.33 Gyr. Taking into account the possibility of a systematic bias in our stellar HK index values ($\\S3$) yields a mean chromospheric age of 4.32 Gyr and a median of 3.83 Gyr. In summary, the mean chromospheric age for the solar-type members of M67 considered herein is in the range of roughly 3.8 - 4.3 Gyr while that for the Sun is approximately 4.3 Gyr. Clearly, age determinations based on chromospheric emission can vary considerably, depending for each star on the phase of its magnetic activity cycle at the time of observation and, in the case of cluster members that are presumably homogeneous in age, on the dispersion in rotational velocities that may be present. \\subsection{The Range of Chromospheric Activity in the M67 Solar-Type Stars} A surprising result of our investigation is the occurrence of solar-type stars that exhibit levels of chromospheric activity in excess of that seen at the maximum of the contemporary solar-activity cycle. An extreme example is illustrated in Fig. 14 where Ca II H and K line core emission is clearly seen in this spectrum of S1452. This object is not known to be a binary at a precision of $\\pm$2 km s$^{-1}$ in the possible variation of its radial velocity (R. D. Mathieu, private communication). An interpretation of our results is that the origin of the dispersion in chromospheric activity among the solar-type stars in M67 is due to the range of amplitudes in their cycle properties. But we also know, in general, that magnetic field-related chromospheric activity increases with increasing equatorial rotational velocity in late-type stars. Thus, those M67 solar-type stars with relatively enhanced levels of activity, compared to the Sun, may be rotating more rapidly, thereby giving rise to stronger dynamo action. If the chromospheric Ca II H and K strength is proportional to the rotational velocity (following Skumanich 1972) then this would suggest for the high-activity M67 stars in Fig. 3--with roughly twice the HK index as the mean Sun--that their rotational velocities are in the 3-4 km s$^{-1}$ range, or slightly more. Obviously, this issue can only be addressed through measurements of (a) rotation period, via ultra-high precision photometric observations of spot modulation, (b) through spectrophotometric observations of the rotational modulation of the H and K lines, or (c) by spectroscopic measurements of the projected rotation velocity. If the results ultimately reveal a significant spread in rotation rates then this would argue against activity cycles as the principal origin of the distribution of activity in Fig. 3. Moreover, it would raise the important question of why the angular momentum histories of the Sun and the M67 solar counterparts differ. If, however, the M67 sun-like stars exhibit rotational velocities that are more solar-like ($\\sim$ 2-3 km s$^{-1}$) then this would suggest that activity cycles--similar to the modern solar cycle but characterized by significant differences in amplitude--are the origin of the dispersion in chromospheric activity we see in Fig. 3. However, this would then raise the question of why the contemporary Sun is in a relatively subdued state of activity compared to other sun-like stars of solar age and solar metallicity." }, "0607/astro-ph0607639_arXiv.txt": { "abstract": "We explain the effect of dark matter (flat rotation curve) using modified gravitational dynamics. We investigate in this context a low energy limit of generalized general relativity with a nonlinear Lagrangian ${\\cal L}\\propto R^n$, where $R$ is the (generalized) Ricci scalar and $n$ is parameter estimated from SNIa data. We estimate parameter $\\beta$ in modified gravitational potential $V(r) \\propto -\\frac{1}{r}(1+(\\frac{r}{r_c})^{\\beta})$. Then we compare value of $\\beta$ obtained from SNIa data with $\\beta$ parameter evaluated from the best fitted rotation curve. We find $\\beta \\simeq 0.7$ which becomes in good agreement with an observation of spiral galaxies rotation curve. We also find preferred value of $\\Omega_{m,0}$ from the combined analysis of supernovae data and baryon oscillation peak. We argue that although amount of \"dark energy\" (of non-substantial origin) is consistent with SNIa data and flat curves of spiral galaxies are reproduces in the framework of modified Einstein's equation we still need substantial dark matter. For comparison predictions of the model with predictions of the $\\Lambda$CDM concordance model we apply the Akaike and Bayesian information criteria of model selection. ", "introduction": "Different astronomical observations \\cite{Riess:1998cb,Perlmutter:1998np} are pointing out that our Universe becomes, at present time, in accelerating phase of expansion. In principle, there are two quite different approaches to explain this observational fact. In the first approach (which can be called substantial) it is assumed that universe is filled by mysterious perfect fluid violating the strong energy condition $\\rho_x+3p_x>0$, where $\\rho_x$ and $p_x$ are, respectively, the energy density and the pressure of this fluid. The nature as well as origin of this matter, called dark energy, is unknown until now. Among these approaches have appeared concordance $\\Lambda$CDM model, which predicts that baryons contribute only about 4\\% of the critical energy density, non-baryonic cold dark matter (CDM) about 25\\% and the cosmological constant $\\Lambda$ (vacuum energy) remaining 70\\%. Although $\\Lambda$CDM model fits well SNIa data \\cite{Riess:2004nr,Astier:2005} this model offers only description of the observations not their explanation. From the methodological view point the conception of mysterious dark energy seems to be effective physical theory only and motivates theorists for searching of alternative approaches in which nature of dark energy will be known at the very beginning. In the first approach it is assumed that Einstein's theory of general relativity is valid which reduces in practice (after assuming Robertson-Walker symmetry of space slices) to the case of Friedman - Robertson - Walker models. Nevertheless, theoretically it is not a'priori excluded the possibility of cosmology based on some extension of Einstein's general relativity. In this paper we consider such particular cases. On the other hand, there are alternative ideas of explanation, in which instead of dark energy some modifications of Friedmann's equation are proposed at the very beginning. In these approaches some effects arising from new physics like brane cosmologies, quantum effects, nonhomogeneities effects etc. can mimic dark energy by a modification of Friedmann equation. Freese \\& Lewis \\cite{Freese:2002sq} have shown that contributions of type $\\rho^n$ to Friedmann's equation $3H^2=\\rho_{\\mathrm{eff}}$, where $\\rho_{\\mathrm{eff}}$ is the effective energy density and $n$ is a constant, may describe such situations phenomenologically. These models (called by their authors called the Cardassian models) give rise to acceleration, although the universe is flat, contains the usual matter and radiation without any dark energy components. This models have been tested by many authors (see for example \\cite{Sen03,Dev03,Z1,Z2,Z3,Z4,Godl04}). What is still lacking is a fundamental theory (like general relativity) from which these models can be derived after postulating Robertson Walker symmetry. In this paper we shall consider such particular type of generalization of Einstein's general relativity in which Lagrangian is proportional to $R^n$, where $R$ is generalized Ricci skalar. In particular, Einstein's general relativity is recovered if we put $n=1$. This theory is part of the larger class of so-called f(R) gravity, i.e. theories derived from gravitational Lagrangians that are analytical (usually polynomial) functions of $R$. (see e.g. \\cite{Cap02,Carroll04,Nojiri03,Flanagan03,Allemandi04,Mota,Mena,Clifton,Cruz06}. In this approach (we called it non-substantial) instead of postulating mysterious dark energy it is assumed some extension of general relativity. Then effect of acceleration appeared naturally as a dynamical effect of the model. For modified gravity one can find Newtonian potential in non-relativistic limit and ask about possibility to explain flat rotation curve of spiral galaxies - major evidence for dark matter in the universe \\cite{Stelle78,Milgrom83,Cap06}. (However, see also \\cite{Salucci} for non flat rotational curves.) The main goal here is to explore power of this particular generalization of gravity in the context of dark energy and dark matter problems. We argue that although cosmology with modified Lagrangian ${\\cal L}\\propto R^n$ can explain \"dark energy problem\" but baryon oscillation test distinguishes value of density parameter of mater to be equal $\\Omega_{m,0} \\simeq 0.3$, i.e. problem of dark matter is yet not solved within the framework of f(R) theories. We also demonstrate that models under considerations can reproduce rotation curves of spiral galaxies. The structure of the paper is as follows. In section II we define class of cosmological models of essential theory of gravity with lagrangian proportional to the Ricci scalar. Section III is devoted to analyze constraints on model parameters from SNIa, baryon oscillation peak and CMB shift. In section IV we investigate problem of rotation curves of spiral galaxies. Section V summarizes our results and formulates general conclusion that models modified gravity which are based on generalized lagrangians ${\\cal L}\\propto R^n$ and Palatini formalism although solve the acceleration and flat rotation curves problems still favor $\\Omega_{m,0} \\simeq 0.3$. ", "conclusions": "In this paper we consider the simplest choice of $f(R)$ theories with $f(R) \\propto R^n$. The basic motivation is searching for fundamental theory of gravity capable to explain both dark energy and dark matter problems without referring to mysterious dark energy conception. For this aim we consider cosmology based on such a theory of gravity and then we use different observational constraints on independent model parameters. We consider simple flat FRW model. It is integrable in exact form after re-parametrization of time variable. From estimation based on SNIa and BOP we obtain $n>2$ which means that bouncing phase instead of big-bang singularity is generic features of such models. Because $n>2$ new $\\tau$ parameter is monotonous function of cosmic time and acceleration epoch is transitional only phenomenon. In the future the universe decelerate which distinguish our model from $\\Lambda$CDM one. Note that because for small value of scale factor a curvature effects are negligible in the comparison to other matter contribution, therefore, in the generic case big-bang singularity is replaced by bounce. Analysis of SNIa Astier data shows that values of $\\chi^2$ statistic are comparable for both $\\Lambda$CDM and best fitted non-linear gravity model. For deeper analysis we use Akaike and Bayesian information criteria of model comparison and selection. We find these criteria still to favor the $\\Lambda$CDM model over non-linear gravity, because (under the similar quality of the fit for both models) the $\\Lambda$CDM model contains one parameter less. Moreover, we find that the effect of dark matter can be kinematically explained as a effect of nonlinear gravity with Lagrangian ${\\cal L} \\propto R^n$. Parameter $\\beta$ required for explaining accelerated expansion of the universe give rise to correct peculiarities of observed rotation curve. However from baryon oscillation peak prior we still obtain $\\Omega_{m,0} \\simeq 0.3$ (instead of $\\Omega_{m,0} \\simeq 0.05$ as we expected). Moreover, we find a disagreement between results obtained from CMB shift parameter analysis and that from joint SNIa and baryon oscillation peak. Finally, the substantial form of dark matter is still required." }, "0607/astro-ph0607125_arXiv.txt": { "abstract": "Multi-object spectroscopy (MOS) instruments, such as the Two-degree Field (2dF) facility of the Anglo-Australian Observatory (AAO), have facilitated large-scale redshift surveys. Yet despite their acclaim, instrument design has been suspected of introducing subtle selection effects into surveys. Investigation into these selection effects has been overshadowed by instrument complexity. We identify the field configuration algorithm (FCA) used to select targets for observation as mainly responsible for such effects. A FCA can imprint artificial structure on observed target distributions, which may accrue over large angular scales, potentially to the detriment of statistical analyses applied to such surveys. We present here a new FCA developed for 2dF that is based on simulated annealing (SA), a generic method commonly used to solve constrained optimisation problems. We generate synthetic fields and utilise mock 2dF volumes to contrast the behaviour of previous strategies with the SA FCA. The angular two-point correlation function and other sensitive techniques reveal that the new FCA achieves unprecedented sampling uniformity and target yield with improved target priority handling and observational flexibility over current FCAs. The SA FCA is generic enough to be used by current 2dF-like and potentially next-generation MOS instruments with little modification. ", "introduction": "The scientific motivation for large-scale redshift surveys has driven the development of efficient multi-object spectroscopy (MOS) instrumentation such as the 2dF facility (Lewis et al. 2002). Their high multiplex advantage is typically achieved by the placement of optical fibres in the focal plane of a telescope to relay light from multiple astronomical targets to a spectrograph. However, this placement can be severely constrained by the physical design of the instrument. Such constraints pose significant challenges to a field configuration algorithm (FCA) if it is to uniformly sample targets for observation whilst upholding the multiplex advantage of the instrument. Uniform sampling is essential to minimise any artificial power imprinted by a FCA on survey target distributions that are later subjected to sensitive statistical analyses. Although robotic fibre placement systems offer greater efficiency than manually operated counterparts, they are undoubtedly the most complex (for an overview see Smith et al. (2004)). Even the most basic of operations, such as observing a field, consists of multiple stages. Before a field can be observed a FCA is used to create a mapping between fibres and targets. This mapping is then used to determine a set of fibre movements that the positioner must make before the field can be exposed. The significant engineering effort required to manipulate fibres (e.g. Wilcox 1993) has overshadowed the requirements of field configuration, leaving FCAs relatively underdeveloped. Preliminary FCAs made possible the instrument design studies of Donnelly et al. (1992; hereafter DON92) and Lewis et al. (1993; hereafter LEW93). These works helped shape the final design of their respective instruments by optimising parameters such as button shape, angular fibre deviation and field geometry. Additionally, DON92 investigated two disparate field configuration strategies. The ideal strategy was an `exhaustive' approach, whereby possible solutions were explored by extensive field randomisation guided by criteria that describe an optimal configuration. DON92 found that although an `exhaustive' strategy would generate very optimal configurations, its heavy use of computer resources meant that a more `intelligent' approach was more viable at the time. The `intelligent' strategy mimics a very clever and patient human who iteratively recognises optimal moves towards a final configuration. Such `intelligent' strategies were readily adopted as the basis for early FCAs because of their relative speed and their ability to generate solutions of comparable quality to `exhaustive' methods. A major disadvantage of `intelligent' FCAs was their strong dependence on understanding how field plate components interact. The emergence of MOS in the early nineties meant that this understanding was not sufficiently mature enough to enhance the development of early FCAs. Substantial progress was later made coinciding with the 2dF Galaxy Redshift Survey (2dFGRS; Colless et al. 2001, hereafter COL01), which utilised the `Oxford' FCA that was tailor-made to 2dFGRS fields (\\S5.1; COL01). The `Oxford' FCA used an insight into fibre availability to attempt uniform target sampling of 2dFGRS fields. Furthermore, its use of fibre swaps to steadily optimise a configuration resulted in much higher target yields. Indeed, much of the success of the 2dFGRS can be attributed to the `Oxford' FCA, which has long been the default FCA for the 2dF \\textsc{configure} program that has been used to prepare fields for observation. Despite such progress, the effect FCAs have on target sampling in large-scale redshift surveys has remained an open question. Such an influence has long been thought to exist, albeit at a negligibly small level. Unless a finished survey has very high completeness, the FCA is likely to have some measurable influence. Although great scrutiny of the source and propagation of selection effects in the construction of such surveys is standard practice (e.g. colour and magnitude selection; tiling algorithm), attempts to measure FCA influence are somewhat of an afterthought. Such attempts are often left until a survey is complete where the perceptible influence of a FCA appears relatively benign. This approach bypasses detailed understanding of FCA behaviour, independent of any tiling algorithm, before a survey is designed. This preference has been exacerbated by a lack of sensitive analytical tools to quantify these influences. The preliminary work of Outram (2004; hereafter OUT04) addressed these issues by implementing sufficiently sensitive analytical tools and using them in a systematic fashion. Under certain conditions, OUT04 discovered previously unknown artificial structure imprinted by the `Oxford' FCA, highlighting the need for precaution when designing large-scale surveys. The work of OUT04 was expanded upon by Miszalski (2005; hereafter MIS05), in addition to implementing a new FCA based on simulated annealing (SA), an `exhaustive' method, that exploits the considerable computer power now available to produce highly optimal field configurations. MIS05 used a wider variety of synthetic fields than OUT04 to contrast the `Oxford' and SA FCAs, concluding that the `Oxford' FCA was relatively unsuitable for the needs of 2dF after the AAOmega spectrograph upgrade (Saunders et al. 2004). This paper serves a dual purpose to (i) Review field configuration requirements and strategies within the context of 2dF-AAOmega (Section 2 and throughout) and (ii) Describe the SA FCA developed by MIS05 for 2dF-AAOmega in Section 3. Section 4 describes the synthetic fields generated to facilitate a comparative study of the `Oxford' and SA FCAs in Section 5. Section 6 concludes with a summary of the SA FCA and its performance. ", "conclusions": "A batch version of \\textsc{configure} was created to facilitate the development of a new FCA based on SA. We bypassed preexisting \\textsc{configure} collision detection strategies by pre-calculating all possible conflicts between all possible allocations and storing them in an indexed collision matrix. This allowed for field randomisation orders of magnitude greater than previously possible. Field randomisation is guided by the Metropolis algorithm with the aid of a simple objective function that can wield great influence over the final solution obtained by the FCA. Our results were obtained by configuring vast quantities of synthetic fields tailored to address specific optimality criteria. They affirm the superior quality of the new FCA that consistently satisfies the criteria to a greater extent than the previous default `Oxford' FCA for 2dF. The new SA FCA has the following valuable attributes compared to the `Oxford' FCA: \\begin{itemize} \\item Gains of up to 7 per cent for low target density fields. \\item Gains of up to 11 per cent for heavily clustered Gaussian fields. \\item Optimal target priority weighting scheme that achieves maximum yields for highest priority targets (gains of up to 30 per cent). \\item Integrated sky target allocation, eliminating the unnecessary sacrifice of higher priority targets to fulfil sky target quotas. \\item Elimination of previous artificial structure imprinted on target distributions, including the severe structure imprinted on different priority populations, to achieve highly uniform target sampling. \\item Quantification of sampling behaviour via the two-point angular correlation function and the completenesses $C(x,y)$ and $C(r)$. \\item Ability to maximise the yield of close target pairs for possible use in surveys and CBS observations. \\item Greatly improved capacity for algorithm maintenance and enhancement arising from its simple design involving the key role of the objective function. \\end{itemize} The outstanding performance and flexibility of the new FCA, coupled with its generic design, makes it perfectly suitable for existing (e.g. 2dF and 6dF) and future MOS instruments with potentially little modification. Algorithm design and components such as the allocation sub-system should scale well to support future instruments with thousands of fibres, provided sufficient growth in computer memory and processing power takes place. The analytical techniques used here to analyse the imprint of artificial structure on FCAs and their general performance is recommended for use in any future FCA development." }, "0607/astro-ph0607580_arXiv.txt": { "abstract": "We report the results of a systematic near-infrared spectroscopic survey using the Subaru, VLT and Keck Telescopes of a sample of high redshift Ultra-luminous Infrared Galaxies (ULIRGs) mainly composed of submillimeter-selected galaxies. Our observations span the restframe optical range containing nebular emission lines such as \\hb, \\OIIID, and \\OII, which are essential for making robust diagnostics of the physical properties of these ULIRGs. Using the \\hal/\\hb\\ emission line ratios, we derive internal extinction estimates for these galaxies similar to those of local ULIRGs: $A_V\\sim 2.9\\pm 0.5$. Correcting the H$\\alpha$ estimates of the star formation rate for dust extinction using the Balmer decrement, results in rates which are consistent with those estimated from the far-infrared luminosity. The majority ($>60$\\%) of our sample show spectral features characteristic of AGN (although we note this partially reflects an observational bias in our sample), with $\\sim65$\\% exhibiting broad Balmer emission lines. A proportion of these sources show relatively low \\OIIIF/\\hb\\ line ratios, which are similar to those of Narrow Line Seyfert 1 galaxies suggesting small mass black holes which are rapidly growing. In the subsample of our survey with both \\OIIIF\\ and hard X-ray coverage, at least $\\sim 60$\\% show an excess of \\OIIIF\\ emission, by a factor of 5--10$\\times$, relative to the hard X-ray luminosity compared to the correlation between these two properties seen in Seyferts and QSOs locally. From our spectral diagnostics, we propose that the strong \\OIIIF\\ emission in these galaxies arises from shocks in dense gaseous regions in this vigorously star-forming population. We caution that due to sensitivity and resolution limits, our sample is biased to strong line emitters and hence our results do not yet provide a complete view of the physical properties of the whole high-redshift ULIRG population. ", "introduction": "There is almost irrefutible evidence for an increase in the star formation density with redshift, as demonstrated by emission line and continuum star formations tracers in wavebands from the ultraviolet to the submillimeter and radio wavebands. This evolution appears to be stronger for tracers which are less sensitive to dust obscuration (e.g.\\ Ivison et al.\\ 2006), suggesting that an increasing proportion of the activity in more distant galaxies may be highly obscured (e.g.\\ Blain et al.\\ 1999, 2002). Indeed, recent results on the mid- to far-infrared emission of luminous but dust obscured galaxies at high redshift ($z\\sim1$--3) suggests that the origin of their large infrared luminosities is a mix of dust obscured vigorous star formation and/or dust enshrouded active galactic nucleus (AGN) (Yan et al.\\ 2005; Houck et al.\\ 2005; Lutz et al.\\ 2005; Desai et al.\\ 2006). In many sources it is likely that both AGN and star formation contribute to the emission as a result of the close link required between the growth of super-massive black holes and bulges in massive galaxies (e.g.\\ Borys et al.\\ 2005). One of the best-studied populations of high-redshift, far-infrared luminous galaxies is that identified in the submillimeter waveband using the SCUBA camera (Holland et al.\\ 1999) on the James Clerk Maxwell Telescope (JCMT). Although they span less than an order of magnitude in submillimeter flux, these galaxies are responsible for much of the energy density in the submillimeter background (Barger et al.\\ 1998; Hughes et al.\\ 1998; Smail et al.\\ 2002; Cowie, Barger \\& Kneib 2002; Scott et al.\\ 2002). The faintness of these obscured galaxies in the optical waveband has made it difficult to obtain precise redshifts (e.g.\\ Simpson et al.\\ 2004), although some progress has been made using ultraviolet/blue spectrographs (Chapman et al.\\ 2003a; 2005). The median redshift for submillimeter galaxies with 850\\,$\\mu$m fluxes of $\\gs 5$\\,mJy, (hereafter SMGs) is $<\\! z\\! >\\sim 2.2$ (Chapman et al.\\ 2003a, 2005). The submillimeter and radio fluxes of these systems indicate their bolometric luminosities are $\\gs 10^{12}$\\,L$_\\odot$ (Kovacs et al.\\ 2006), confirming that they are examples of high-redshift Ultraluminous Infrared Galaxies (ULIRGs). This population provides critical constraints on models of galaxy formation and evolution. In particular, if the bolometric emission from SMGs is powered solely by star formation, then these galaxies form about half of the stars seen in the local Universe (Lilly et al.\\ 1999). However, it appears likely that both AGN and star formation activity contribute to the immense far-infrared luminosities of these systems, although it has been difficult to disentangle the precise balance between these two energy sources. Recent sensitive X-ray analysis suggest that star formation is likely to be the dominant source of the bolometric luminosity in SMGs (Alexander et al.\\ 2005a,b). Further evidence suggest it is plausible to identify SMGs as the progenitor of massive elliptical galaxies at the present-day, based on their large gas, stellar and dynamical masses (Neri et al.\\ 2003; Greve et al.\\ 2005; Tacconi et al.\\ 2006; Smail et al.\\ 2004; Borys et al.\\ 2005; Swinbank et al.\\ 2004, 2006). Furthermore, combining the X-ray constraints on the AGN within this population with the typical mass estimates suggests that SMGs are the sites of coeval growth of stellar bulges and central black holes (Borys et al.\\ 2005). Rest-frame optical emission lines provide a powerful tool to investigate many fundamental properties of galaxies, such as star formation rates (SFRs), power sources, internal extinction and metallicity. Swinbank et al.\\ (2004) conducted a systematic near-infrared spectroscopic survey of thirty SMGs to investigate their SFRs and metallicities and the kinematics of the emission line gas. However, the wavelength coverage was limited to the region around H$\\alpha$ and so they did not include several emission lines at shorter wavelengths, such as \\hb\\ and \\OIIID, which are useful for evaluating internal extinction and metallicity or determining the power source. We present in this paper the results from a near-infrared spectroscopic survey of redshifted \\OIIID, \\hb\\ and \\OII\\ lines for a sample of far-infrared luminous galaxies. The sample is composed of SMGs and optically faint radio galaxies (OFRGs), at $z\\sim1$--3.5. Chapman et al.\\ (2004) and Blain et al.\\ (2004) claim that high-redshift OFRGs are ULIRGs, with similar bolometric luminosities to SMGs but warmer characteristic dust temperature, resulting in them being undetectable in the submillimeter waveband. We use \\hal/\\hb\\ emission line ratios to derive the dust extinction in these systems and then employ these estimates to derive extinction-corrected SFRs from the \\hal\\ luminosities. In addition, we also use X-ray observations of these objects to compare the strength of the \\OIIIF\\ emission to their X-ray emission, and so investigate the power of the AGN in these galaxies. We adopt cosmological parameters of H$_0=$72 km sec$^{-1}$ Mpc$^{-1}$, and $\\Omega_M=0.3$ and $\\Omega_{\\Lambda}=$0.7 throughout. ", "conclusions": "Using near-infrared spectroscopy we have observed the redshifted \\hb\\, the \\OIIID\\ and \\OII\\ emission lines in a sample of 22 Ultra-luminous Infrared Galaxies at high redshifts. Twenty of the sources in our sample are submillimeter galaxies at $z\\sim 1.0$--3.5. Combining our observations with previous studies of the \\hal\\ and the \\NII\\ emission from these galaxies and also with observations of their hard X-ray and far-infrared emission, we have placed constraints on the physical properties of this population. We conclude the following: \\begin{enumerate} \\item A majority of our sample (14/22) have spectra which are classified as ``AGN'' or ``QSO'' based on several restframe optical spectroscopic diagnostics. Specifically, for those sources with detections of the four emission lines necessary to construct a BPT diagram, 8/9 are classified as ``AGN''. It should be noted that there is no confirmed pure starburst galaxy in our sample, although several sources show intermediate spectral properties. This is likely to be caused by our sample selection, which is biased towards galaxies with bright near-infrared magnitudes and also to those exhibiting strong line emission. Thus we caution that our results should not be taken as representative of the whole SMG population. \\item Using the \\hal/\\hb\\ flux ratio we are able to estimate the internal extinction in our SMGs. We measure a median extinction of $A_V=2.9\\pm0.5$, which is similar to the extinction measured in local ULIRGs. This value is also consistent with the estimates from the SED fitting in the restframe UV/optical which are derived under the assumption of a dominant dust-reddened young starburst (Smail et al.\\ 2004). \\item We compare the SFRs derived from the dust-extinction-corrected \\hal\\ luminosities with those derived from the far-infrared luminosities, and find reasonable consistency between these for most of the SMGs in our sample. The fact that the corrected \\hal-derived SFRs correspond closely to those estimated from the far-infrared suggests that star-formation is the major contributor to the far-infrared luminosities in SMGs. \\item At least 11/19 of the SMGs in our sample show a clear excess in the ratio of their \\OIIIF\\ to X-ray luminosities relative to values for local AGNs. The five sources with the highest \\OIIIF/\\hb\\ ratios ($>10$), which are classified as ``AGN'' from our spectral diagnostics, show this \\OIIIF\\ excess. One possible explanation for the \\OIIIF\\ excess is that it is produced by ``Compton-Thick'' AGNs. However, this is inconsistent with the column density measurements (N$_H$) from fitting of the X-ray spectra for the sources in CDFN and we argue that this is unlikely in most SMGs. Instead, we suggest that the most plausible cause of the \\OIIIF\\ excess is shock-induced emission arising from vigorous star formation (``super-wind'' activity). This scenario is supported in several galaxies by spatially extended and/or distorted/multiple \\OIIIF\\ emission line profiles. Furthermore, using limits on the electron temperatures from \\OIIIA\\ and \\NII\\ emission line ratios, we can explain the excess \\OIIIF\\ emission as arising from shocks in dense regions within these systems. \\item The Balmer line widths in 9/22 sample galaxies exhibit broad emission components with relatively small FWHMs ($\\sim1500$--3700\\,km\\,sec$^{-1}$). Three of them are classified as ``QSO'', but have smaller \\hb\\ FWHM (2100--2600\\,\\kmsec) than are typical for QSOs. They also have lower \\OIIIF/\\hb\\ ratios and relatively strong Fe{\\sc ii} emission, both of which are characteristics of local Narrow Line Seyfert 1s. Among the other six sources, only one shows a low \\OIIIF/\\hb\\ ratio, and four show high \\OIIIF/\\hb\\ ratios (larger than seen in NLS1's). However, the high \\OIIIF/\\hb\\ ratios may arise from \\OIIIF\\ excesses due to shock excitation and hence removing this contribution would yield lower ratios more consistent with NLS1 classification. Several of these sources also have tentative evidence for Fe{\\sc ii} emission, again characteristic of NLS1s. Thus, once account is taken of the potential contribution from shocks to the excess \\OIIIF\\ emission, there appears to be close similarities between SMGs and NLS1s. The spectral classification of SMGs as NLS1s may then indicate (as has been claimed for local NLS1s) that SMGs have small mass black holes which are rapidly growing at high accretion rates (Alexander et al.\\ 2005ab; Borys et al.\\ 2005). Deeper spectroscopic observations are essential to search for any obscured broad Balmer lines which might indicate larger SMBH masses and confirm the presence of Fe{\\sc ii} lines which are common in the NLS1s. \\end{enumerate} Summarising our results: we conclude that our sample of SMGs contains a population of vigorously star-forming galaxies with high SFRs and strong extinction. The activity in these systems is driving shocks through the dense gas reservoirs they contain and some of this material is being expelled from the galaxies. In addition, many of our sources show evidence for low-mass, but rapidly growing, super-massive black holes. These results confirm the critical place of the submillimeter-bright phase in defining the properties of massive galaxies forming at high redshifts." }, "0607/astro-ph0607063_arXiv.txt": { "abstract": "The emission spectra of TeV blazars extend up to tens of TeV and the emission mechanism of the TeV $\\gamma$-rays is explained by synchrotron self-Compton scattering in leptonic models. In these models the time variabilities of X-rays and TeV $\\gamma$-rays are correlated. However, recent observations of 1ES 1959+650 and Mrk 421 have found the ``orphan'' TeV $\\gamma$-ray flares, i.e., TeV $\\gamma$-ray flares without simultaneous X-ray flares. In this paper we propose a model for the ``orphan'' TeV $\\gamma$-ray flares, employing an inhomogeneous leptonic jet model. After a primary flare that accompanies flare-up both in X-rays and TeV $\\gamma$-rays, radiation propagates in various directions in the comoving frame of the jet. When a dense region in the jet receives the radiation, X-rays are scattered by relativistic electrons/positrons to become TeV $\\gamma$-rays. These $\\gamma$-ray photons are observed as an ``orphan'' TeV $\\gamma$-ray flare. The observed delay time between the primary and ``orphan'' flares is about two weeks and this is accounted for in our model for parameters such as $\\Gamma = 20$, $d = 4 \\times 10^{17}$cm, $\\alpha = 3$, and $\\eta = 1$, where $\\Gamma$ is the bulk Lorentz factor of the jet, $d$ is the distance between the central black hole and the primary flare site, $\\alpha/\\Gamma$ is the angle between the jet axis and the direction of the motion of the dense region that scatters incoming X-rays produced by the primary flare, and $\\eta/\\Gamma$ is the angle between the jet axis and the line of sight. ", "introduction": "Blazars are a subclass of active galactic nuclei and their high energy emission is generated in regions near the central super massive black holes. Because of their intense and rapidly variable radiation, the emission regions are thought to be in relativistic jets closely aligned to the line of sight \\citep[see][for a review]{aha04,kra04}. The spectral energy distributions (SEDs) of blazars are characterized by double broad peaks in the $\\nu$-$\\nu F_\\nu$ representation. One peak is in the optical to X-ray regions and the other is in the $\\gamma$-ray region. Very high energy (100 MeV -- 1 GeV) $\\gamma$-rays from blazars were discovered by {\\it EGRET} (Energetic Gamma-Ray Experiment Telescope) on board Compton Gamma-Ray Observatory \\citep{har99}. Recent observations by Cerenkov telescopes have revealed that many blazars emit $\\gamma$-rays up to tens of TeV \\citep[][for a review]{aha05hess-all}. Flare activities in X-rays and $\\gamma$-rays from blazars are also known. Although the emission mechanisms of blazars are still under study, the SEDs of blazars are well explained by leptonic or hadronic models. In the framework of the leptonic models, the emission in the optical to X-ray region is attributed to synchrotron radiation by nonthermal electrons/positrons in the jet and the TeV $\\gamma$-rays are accounted for by the synchrotron self-Compton (SSC) model \\citep{mgc92}. The SSC model assumes that synchrotron photons are inverse-Compton scattered by the same nonthermal electrons/positrons that emit synchrotron photons. If the SSC model is applied, the time variabilities of the X-rays and $\\gamma$-rays should be correlated \\citep[e.g.,][]{mkir97,lk00,kus00}, and indeed most flares occur almost simultaneously in X-rays and $\\gamma$-rays. However, the recent observations of 1ES 1959+650 \\citep{kra-et04,dan05} found that there is a TeV $\\gamma$-ray flare that is not accompanied by a X-ray flare, which was observed on June 4, 2002 (MJD 52,429) by the Whipple telescope. This is called the ``orphan'' TeV $\\gamma$-ray (OTG) flare. The OTG flare occurred about 15 days after a usual flare in which X-ray and $\\gamma$-ray flares were observed contemporaneously. More recently another OTG flare might have been observed from Mrk 421 \\citep{bla05}, although the X-ray flux seems to have peaked about 1.5 days before this TeV flare and this may not be a true OTG flare. The existence of the OTG flares is challenging for the leptonic models. \\citet{bot05} has recently proposed a model for the OTG flare of ES 1959+650, i.e., the hadronic synchrotron mirror model. He assumed that a fraction of X-rays produced in the primary flare by synchrotron radiation of leptons is reflected by a plasma cloud with scattering depth $\\sim 0.1$ located in the direction of the jet propagation. The reflected X-rays collide with protons in the jet and pions are produced. Subsequently a large number of neutral pions decay into $\\gamma$-rays, which are expected to be observed as TeV $\\gamma$-rays. If this model is applied to 1ES 1959+650, it is found that the reflecting plasma is located at a distance $\\sim 3 (\\Gamma/10)^2 \\Delta t_{20}$ pc from the central black hole, where $\\Gamma$ is the bulk Lorentz factor of the jet and $\\Delta t_{20} = \\Delta t/(20 \\, \\mathrm{days})$ is the normalized time delay between the primary and OTG flares. This model requires unreasonably large values of proton density in the jet and the hadronic jet power (B\\\"{o}ttcher 2006: erratum). Thus the hadronic synchrotron mirror model may not be applicable to OTG flares, although it may be still viable if the effects that reflected synchrotron photons increase as the blob approaches to the mirror are taken into account. (More recently \\citet{rbp05} calculated the predicted neutrino spectrum due to the decay of charged pions in the framework of the hadronic synchrotron mirror model.) In this paper, we propose another model of the OTG flares in the framework of the SSC model, assuming that the emission region is not homogeneous, which is different from the conventional SSC model. We assume that the injection of nonthermal electrons/positrons triggers the primary flare where both X-rays and TeV $\\gamma$-rays are emitted. If the jet is not uniform but consists of a few patchy regions, X-rays that were produced in the primary flare impinge on another dense region of the jet. This sudden increase of X-ray photons results in strong TeV $\\gamma$-ray flux by inverse Compton scattering, which is expected to be observed as an OTG flare. Because there is a time delay between the primary flare in one region and the scattering in another region, TeV $\\gamma$-rays are observed as an OTG flare. Note that TeV $\\gamma$-rays from the primary flare are not scattered in the OTG flare site because of Klein-Nishina effects. The observed delay time between the OTG flare and the previous flare is about two weeks in 1ES 1959+650 \\citep{kra-et04}. The delay time between the primary and OTG flares in the comoving frame of the jet is much shorter than the synchrotron cooling time of nonthermal electrons/positrons (see \\S \\ref{sec:lumi} for detail), if the magnetic field in the jet is about 0.1 G. Although high energy electrons/positrons injected into the primary flare are rapidly cooled, nonthermal electrons/positrons that emit TeV $\\gamma$-rays in a quiescent state may be still continuously injected into the jet and they may contribute to the secondary flare. In this paper we do not specify the acceleration mechanisms of those electrons/positrons or calculate the emission spectrum of the flare but focus our work on the kinematics of the jet. The kinematics of the jet is described in \\S \\ref{sec:kinematics} and the jet properties are presented in \\S \\ref{sec:lumi}. A summary of our results and discussion are given in \\S \\ref{sec:sum}. ", "conclusions": "\\label{sec:sum} In this paper we have proposed a structured jet model for the OTG flares. The high energy emission spectra from TeV blazars have been well explained by the leptonic jet models. Because the leptonic jet models assume synchrotron self-Compton scattering for the emission mechanism of very high energy $\\gamma$-rays, it is expected that flares occur in both X-ray and $\\gamma$-ray energy bands almost simultaneously. However, the OTG flares were not accompanied by the simultaneous increase in the X-ray flux \\citep{kra-et04,bla05}, and this is thought to be a challenging issue for the leptonic jet models. Our model of a structured jet assumes that $\\gamma$-rays are emitted in different components in the jet. X-rays produced by a flare in a region, where flares in X-rays and TeV $\\gamma$-rays occur simultaneously, are scattered in another region of the jet and the scattered photons are observed as a secondary (OTG) flare. Note that TeV $\\gamma$-rays of the primary flare are not scattered by electrons/positrons because of Klein-Nishina effects. The emitted $\\gamma$-ray energy is estimated as $\\epsilon'_\\mathrm{SSC} = \\gamma' m_e c^2$, because the scattering occurs mainly in Klein-Nishina regime. Then the observed $\\gamma$-ray energy is roughly given by \\begin{equation} \\epsilon_\\mathrm{SSC} \\sim \\Gamma \\epsilon'_\\mathrm{SSC} \\sim 5.5 \\times 10^{12} \\left( \\frac{B'}{0.1 \\mathrm{G}} \\right)^{-1/2} \\left( \\frac{\\Gamma}{20} \\right)^{1/2} \\left( \\frac{\\epsilon_\\mathrm{syn}}{10 \\mathrm{keV}} \\right)^{1/2} \\quad \\mathrm{eV} , \\end{equation} using equation (\\ref{eq:gamma-value}). Because the X-rays emitted in the primary flare need a propagation time before they are scattered in different parts of the jet, this results in the time delay between the primary and secondary flares. In our model, the delay time between the primary flare and the secondary TeV $\\gamma$-ray flare is dependent on $\\Gamma$, $\\alpha$, $\\eta$, and $d$. The observed delay time is about 15 days for 1ES 1959+650 and this value is obtained for $\\Gamma = 20$, $\\alpha = 3$, $\\eta = 1$, and $d = 4 \\times 10^{17}$ cm in Case 1. As the value of $\\alpha$ decreases the value of $d$ increases for a given value of $\\delta t_\\mathrm{obs}$. If the value of $\\Gamma$ is known from the observations of the primary flare, the values of $\\alpha$ and $d$ might be reasonably guessed from the delay time. It is not clear from the numerical values estimated in this paper whether Case 1 is favored to Case 2 or vice versa. Rare observations of OTG flares might be due to that this phenomenon needs a specific structure of the jets as shown in Figures \\ref{fig:jet-case1} or \\ref{fig:jet-case2}. The duration of the secondary flare depends on the time interval of the primary flare and the angular size of the secondary flare site, $\\Delta \\theta$. For example, in Case 1, if $\\Delta \\theta$ is given by $\\Delta \\alpha /\\Gamma$ and the primary flare is impulsive, the observed duration of the second flare is given by the difference between the values of $\\delta t_\\mathrm{obs}$ for $\\alpha$ and $\\alpha + \\Delta \\alpha$: \\begin{equation} t_\\mathrm{dur} \\approx 0.29 \\left( \\frac{20}{\\Gamma} \\right)^2 \\left( \\frac{d}{3 \\times 10^{17} \\, \\mathrm{cm} } \\right) \\frac{d f(\\alpha)}{d \\alpha} \\Delta \\alpha \\quad \\mathrm{days} , \\end{equation} where $\\Delta \\alpha \\ll 1$ is assumed. The values of $g(\\alpha) = df/d\\alpha$ are shown in Figure \\ref{fig:g-factor} for Case 1. For example, $g \\approx 54.4$ for $\\eta = 1$ and $\\alpha = 3$. If the observed duration of the secondary flare is about 5 hours, $\\Delta \\alpha \\sim 0.01$ for $\\alpha = 3$, $\\Gamma = 20$, and $d = 4 \\times 10^{17}$ cm. The observed duration of the secondary flare may depend on not only the value of $\\Delta \\alpha$ but also the duration of the primary flare, and the effect of the latter elongates the duration of the secondary flare. We have shown that the time lag between the primary and secondary flares are well accounted for by our model. Next we argue that the density of soft photons (X-rays) from the primary flare is high enough to produce the OTG flare if the particle injection time in the primary flare is sufficiently long, and that the particle injection rate needed in the secondary flare is not unreasonably high. First, it should be noted that photons from the primary flare blob (PB) are not injected from the rear side of the secondary flare blob (SB). In the comoving frame of the blobs, the angle between the line of sight and the incident direction of photons from PB is about 90 degrees. Then the energy boost by inverse Compton scattering in SB should be effective, although the detailed calculations of radiative transfer are necessary to obtain the emission spectrum of TeV $\\gamma$-rays from SB. Next, we discuss the amount of radiation received by SB and the cooling time of leptons. In the following, primed symbols are quantities during the primary flare in the comoving frame and doubly primed symbols are quantities during the secondary flare in the comoving frame. We assume that PB and SB are spherical plasma clouds for simplicity. The radii of PB and SB in the comoving frame at the occurrence of the primary flare are denoted by $R'_\\mathrm{prim}$ and $R'_\\mathrm{sec}$, respectively. The radius of SB at the time when the secondary flare occurs is denoted by $R''_\\mathrm{sec}$. The solid angle extended by SB for a point in PB is given by $\\Omega'_\\mathrm{sec} \\sim \\pi R^{\\prime \\, 2} _\\mathrm{sec} /d^{\\prime \\, 2}_\\mathrm{sec}$, where $d'_\\mathrm{sec}$ is the distance between PB and SB: $d'_\\mathrm{sec} \\sim \\xi d = (\\alpha/\\Gamma) \\, d$. If SB expands proportionally with the distance from the central black hole, $\\Omega'_\\mathrm{sec} = \\Omega''_\\mathrm{sec}$. If radiation from PB is emitted isotropically in the comoving frame, the fraction of the radiation received by SB is given by \\begin{equation} r_\\mathrm{rad} \\equiv \\frac{\\Omega'_\\mathrm{sec}}{4 \\pi} = \\frac{1}{4} \\, \\frac{\\Gamma^2}{\\alpha^2} \\, \\frac{R^{\\prime \\, 2}_\\mathrm{sec}}{d^2} \\sim \\frac{1}{36} \\left(\\frac{\\Gamma}{20}\\right)^2 \\left(\\frac{3}{\\alpha}\\right)^2 \\left(\\frac{R'_\\mathrm{sec}}{2 \\times 10^{16} \\mathrm{cm}} \\right)^2 \\left(\\frac{4 \\times 10^{17} \\mathrm{cm}}{d}\\right)^2 . \\end{equation} We assumed $R'_\\mathrm{sec} \\sim 2 \\times 10^{16}$ cm; this value is close to the values assumed for the emission region of the flare models of 1ES 1959+650 in \\citet{kra-et04}. Thus the soft photon (X-ray) energy density in SB is roughly about 1/30 of that in PB and the Compton cooling time of nonthermal particles in SB is longer by 30 times than in PB. On the other hand, the flare duration in PB is determined by the injection duration of nonthermal particles, because the cooling time should be short owing to the large energy densities of soft photons and magnetic fields. Thus the observed OTG flare is explained, if the particle injection duration in PB is about 30 times longer than the cooling time in SB. Furthermore, the energy densities of magnetic fields and internal synchrotron photons in SB are smaller than in PB, so that emission other than TeV $\\gamma$-rays is too weak to be observed. In the following we estimate the required injection rate of nonthermal leptons in the secondary flare. The observed energy flux of the secondary flare at 600GeV is $\\nu F_\\nu \\sim 3 \\times 10^{-10}$ ergs s$^{-1}$ cm$^{-2}$. The luminosity distance to 1ES 1959+650 ($z = 0.047$) is $d_L = 210$ Mpc, assuming that $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_\\Lambda = 0.7$, and $\\Omega_m = 0.3$, where $H_0$, $\\Omega_\\Lambda$, and $\\Omega_m$ are the Hubble constant and the density parameters of dark energy and matter, respectively. Then the luminosity of the very high energy gamma-rays is $L_\\mathrm{VHE, sec} \\sim 1.6 \\times 10^{45}$ ergs s$^{-1}$. The luminosity in the comoving frame is given by \\begin{equation} L''_\\mathrm{VHE, sec} = \\frac{L_\\mathrm{VHE, sec}}{{\\cal D}^4_\\mathrm{sec}} \\sim 9.9 \\times 10^{39} \\left( \\frac{{\\cal D}_\\mathrm{sec}}{20} \\right)^{-4} \\quad \\mathrm{ergs} \\, \\mathrm{s}^{-1} , \\end{equation} where ${\\cal D}_\\mathrm{sec}$ is the beaming factor of SB and ${\\cal D}_\\mathrm{sec} = 8$ and 20 for Cases 1 and 2, respectively, if $\\Gamma = 20$, $\\alpha = 3$, and $\\eta = 1$. The distance between the central black hole and SB is $d_\\mathrm{sec} = d + \\beta c \\delta t \\sim 11 d$, where $\\alpha = 3$ is assumed. Assume that SB expands linearly with the distance from the central black hole. Then the radius of SB is given by \\begin{equation} R''_\\mathrm{sec} = \\frac{d_\\mathrm{sec}}{d} \\, R'_\\mathrm{sec} \\sim 11 R'_\\mathrm{sec} = 2.2 \\times 10^{17} \\left( \\frac{R'_\\mathrm{sec}} {2 \\times 10^{16} \\mathrm{cm}} \\right) \\quad \\mathrm{cm} . \\end{equation} We obtain the energy injection rate per unit volume during the secondary flare: \\begin{equation} q''_\\mathrm{sec} = \\frac{L''_\\mathrm{VHE, sec}}{\\frac{4 \\pi}{3} R^{\\prime \\prime \\, 3}_\\mathrm{sec} } = 2.2 \\times 10^{-13} \\left( \\frac{R'_\\mathrm{sec}}{2 \\times 10^{16} \\mathrm{cm}} \\right)^{-3} \\left( \\frac{{\\cal D}_\\mathrm{sec}}{20} \\right)^{-4} \\quad \\mathrm{ergs} \\, \\mathrm{s}^{-1} \\, \\mathrm{cm}^{-3} . \\end{equation} The injection rate of nonthermal leptons in SB is also given by \\begin{equation} N''_e = \\frac{L''_\\mathrm{VHE, sec}}{\\langle \\gamma \\rangle m_e c^2} = 2.4 \\times 10^{40} \\left( \\frac{{\\cal D}_\\mathrm{sec}}{20} \\right)^{-4} \\left(\\frac{\\langle \\gamma \\rangle}{5 \\times 10^5} \\right)^{-1} \\quad \\mathrm{s}^{-1} , \\end{equation} where $\\langle \\gamma \\rangle$ is the value of the characteristic Lorentz factor of injected leptons in SB. When the energy injection rate during the primary flare is denoted by $q'_\\mathrm{prim}$, we obtain \\begin{equation} \\frac{q''_\\mathrm{sec}}{q'_\\mathrm{prim}} = \\frac{L''_\\mathrm{VHE, sec} /\\left(\\frac{4 \\pi}{3} R^{\\prime \\prime \\, 3}_\\mathrm{sec}\\right)} {L'_\\mathrm{VHE, prim} /\\left(\\frac{4 \\pi}{3} R^{\\prime \\, 3}_\\mathrm{prim}\\right)} = \\frac{[\\nu F_\\nu]_\\mathrm{VHE, sec}}{[\\nu F_\\nu]_\\mathrm{VHE, prim}} \\left(\\frac{{\\cal D}_\\mathrm{prim}}{{\\cal D}_\\mathrm{sec}}\\right)^4 \\left(\\frac{R'_\\mathrm{prim}}{R''_\\mathrm{sec}}\\right)^3 , \\end{equation} where $[\\nu F_\\nu]_\\mathrm{VHE, prim}$ and $[\\nu F_\\nu]_\\mathrm{VHE, sec}$ are the observed quantities of $\\nu F_\\nu$ in the primary and secondary flares, respectively. If $R'_\\mathrm{sec} \\sim R'_\\mathrm{prim}$ and $R''_\\mathrm{sec} \\sim 11 R'_\\mathrm{sec}$ are assumed, $q''_\\mathrm{sec} / q'_\\mathrm{prim} \\sim c_1 \\, [\\nu F_\\nu]_\\mathrm{VHE, sec} / [\\nu F_\\nu]_\\mathrm{VHE, prim}$, where $c_1 \\sim 3.0 \\times 10^{-2}$ and $1.9\\times 10^{-5}$ for Cases 1 and 2, respectively. From the observed data \\citep{kra-et04}, we may set \\\\ \\noindent $[\\nu F_\\nu]_\\mathrm{VHE, sec}/[\\nu F_\\nu]_\\mathrm{VHE, prim} \\sim 1$, and $q''_\\mathrm{sec} \\sim c_1 q'_\\mathrm{prim}$ is obtained. Thus the required value of $q''_\\mathrm{sec}$ is not unreasonably large. This particle energy is transfered to X-rays by inverse Compton scattering to emit TeV $\\gamma$-rays in the secondary flare. Our model assumes a structured jet and the occurrence of OTG flares depends on the geometry of dense regions in the jet. Alternatively the dense regions might be widely distributed, e.g., the layer-spine model by \\citet{gtc05}, and the value of $\\alpha$ depends on where the region of particle acceleration at pc-scale exists. Finally we comment on the observed large scale structure of jets. The line of sight is inside the opening angle of jets in our model, and this may result in a core-halo structure of the emission region. However, since the emission regions of high energy radiation are located very close to the central black hole, the large scale structure of jets is affected by the environment of the jets. That is, the interaction between jets and gases surrounding the central region may cause the bending of the jets. It is also possible that the opening angle may change with the distance from the central region. These may account for the jet structure observed by VLBI." }, "0607/astro-ph0607549_arXiv.txt": { "abstract": "We continue our series of papers on open cluster distances with a critical assessment of the accuracy of main-sequence fitting using isochrones that employ empirical corrections to the color-temperature relations. We use four nearby open clusters with multicolor photometry and accurate metallicities and present a new metallicity for Praesepe (${\\rm [Fe/H]} = +0.11 \\pm 0.03$) from high-resolution spectra. The internal precision of distance estimates is about a factor of 5 better than the case without the color calibrations. After taking into account all major systematic errors, we obtain distances accurate to about 2\\% -- 3\\% when there exists a good metallicity estimate. Metallicities accurate to better than 0.1 dex may be obtained from $BVI_{C}K_{s}$ photometry alone. We also derive a helium abundance for the Pleiades of $Y = 0.279 \\pm 0.015$, which is equal within the errors to the Sun's initial helium abundance and that of the Hyades. Our best estimates of distances are $(m - M)_0 = 6.33 \\pm 0.04$, $8.03 \\pm 0.04$, and $9.61 \\pm 0.03$ to Praesepe, NGC~2516, and M67, respectively. Our Pleiades distance at the spectroscopic metallicity, $(m - M)_0 = 5.66 \\pm 0.01$ (internal) $\\pm 0.05$ (systematic), is in excellent agreement with several geometric distance measurements. We have made calibrated isochrones for $-0.3 \\leq {\\rm [Fe/H]} \\leq +0.2$ available online. ", "introduction": "The determination of accurate distances is the key to understanding how stars and the Galaxy have formed and evolved. From protostars in star-forming regions to ancient tracers of the halo, improved distances have refined stellar evolutionary theory and Galactic structure models \\citep[e.g.,][]{reid99}. The {\\it Hipparcos} mission \\citep{perryman97b} was especially valuable, providing trigonometric parallaxes for $\\sim10^5$ stars to precision of 1--2 mas \\citep{perryman97a}. These parallaxes, however, are only useful for individual stars within $\\sim 100$~pc. Most open clusters are much more distant than this ``horizon,'' but a half-dozen of the nearest clusters have 10 -- 50 or more {\\it Hipparcos} stars, yielding cluster parallaxes ostensibly accurate to 5\\% or better \\citep{mermilliod97,perryman98,robichon99,vanleeuwen99}. Main-sequence (MS) fitting, also known as the photometric parallax method \\citep[e.g.,][]{johnson57,siegel02}, has long been used to estimate distances to individual stars and star clusters beyond the limits of parallax studies, and is considered to be a robust and well-understood technique. It was therefore a big surprise when the {\\it Hipparcos} distances to the Pleiades and Coma Ber open clusters were in disagreement with distances from the MS-fitting at more than a 3 $\\sigma$ level \\citep{pinsono98}. It is difficult to reconcile a short Pleiades distance with stellar interior and spectroscopic abundance studies. A high helium abundance would make a cluster fainter than expected from its metallicity, and solutions of this type have been discussed in the literature for the Pleiades \\citep{belikov98}. However, this is difficult to understand since there do not seem to be nearby field stars of similar characteristics in the {\\it Hipparcos} catalog \\citep{soderblom98}, and the helium enhancement would have to be enormous ($Y \\approx 0.34$). In addition, it has been suggested that the metal abundance from spectroscopy may have been significantly overestimated \\citep{percival03}. An argument was also made that distance estimates from theoretical stellar models have been overestimated for young clusters due to unknown, age-related physics \\citep{vanleeuwen99}. However, the most likely explanation is related to the {\\it Hipparcos} parallaxes themselves. \\citet{pinsono98} showed that the 12 bright stars near the center of the Pleiades all had virtually the same parallax, $\\sim 9$~mas, more than 1~mas larger than the mean parallax for other cluster stars. They attributed this to a local zero-point error of the individual stellar parallaxes that are correlated over the {\\it Hipparcos}' $0{\\fdg}9$ field of view \\citep{vanleeuwen98}. These quasi-random errors were caused by the {\\it Hipparcos} great-circle data reductions, as \\citet{makarov02,makarov03} proved by re-reducing the Pleiades and Comar Ber cluster parallaxes in a different way that correctly obtains the absolute zero point of parallax. Additional effects may result from the way the {\\it Hipparcos} data were obtained and analyzed, and a more elaborate reduction of the {\\it Hipparcos} parallaxes promises to produce improved distances and better understood errors \\citep{vanleeuwen05a,vanleeuwen05b}. The discrepant {\\it Hipparcos} result for the Pleiades subsequently led to many efforts to determine the cluster's distance from binaries and independent parallax measurements \\citep[e.g.,][]{munari04,pan04,johnskrull05,soderblom05}. These results support the longer distance scale from MS fitting, verifying that the {\\it Hipparcos} result was in error. With a formal error of $\\sim 1\\%$ from these measurements, the Pleiades represents a second system (besides the Hyades) with a sufficiently accurate distance for a precision test of stellar evolutionary models. Even though the controversy over the Pleiades distance is now settled, a critical assessment of the MS fitting technique is still required to reliably estimate a distance. In fact, MS fitting using theoretical isochrones is a complex process that involves both physical and empirical considerations \\citep[e.g.,][]{stauffer01}. There are, however, many opportunities to check the construction of the isochrones. Stellar evolution models can be tested against the Sun and other stars, such as eclipsing binaries, that have accurate masses and radii. Furthermore, multicolor photometry in nearby clusters and field stars can be used to test the bolometric corrections and color-effective temperature ($T_{\\rm eff}$) relations to transform theoretical quantities (luminosity and $T_{\\rm eff}$) to magnitudes and colors \\citep[e.g.,][]{vandenberg03}. In our first two papers of this series \\citep[][hereafter Paper~I and Paper~II, respectively]{pinsono03,pinsono04}, we began a long-term effort to assess the accuracy of distances from MS fitting and to reduce or eliminate systematic errors in the process, particularly those arising from the transformation of theoretical to observational quantities. In Paper~I, we demonstrated that stellar models from the Yale Rotating Evolutionary Code \\citep[YREC;][]{sills00} are in good agreement with masses and luminosities for the well-studied Hyades eclipsing binary vB~22 \\citep{torres02}. These models also satisfy stringent tests from helioseismology, and predict solar neutrino fluxes in line with observations \\citep{basu00,bahcall01,bahcall04}. In Paper~II, we showed that the models provide a good match to the spectroscopically determined temperatures \\citep{paulson03} for individual Hyades members with good parallaxes \\citep{debruijne01}. However, we found that any of the widely-used color-$T_{\\rm eff}$ relations \\citep[e.g.,][]{alonso95,alonso96,lejeune97,lejeune98} fail to reproduce the observed shapes of the MS in the Hyades; differences in broadband colors were as large as $\\sim0.1$ mag. The existence of these systematic errors in the colors in the presence of agreement between the spectroscopic and theoretical $L-T_{\\rm eff}$ scales strongly implies that there are problems with the adopted color-$T_{\\rm eff}$ relations instead of errors in the theoretical $T_{\\rm eff}$ scale. Therefore, we proposed empirical corrections to the color-$T_{\\rm eff}$ relations from \\citet{lejeune97,lejeune98} that were adopted in the isochrone computations. In this study, we generate a set of isochrones over a wide range of age and metallicity, and test the validity of the Hyades-based color-$T_{\\rm eff}$ corrections using extensive multicolor photometry of four well-studied nearby open clusters. We show that isochrones employing the Hyades empirical corrections precisely match the observed MS shapes, except where anomalously blue colors in young open clusters have been previously noted \\citep{stauffer03}. Furthermore, we demonstrate that the empirical corrections improve the internal precision of the isochrones by examining the consistency of distances derived from several color-magnitude diagrams (CMDs). We also assess various sources of systematic errors in the MS-fitting technique. Previously, \\citet{pinsono98} considered the effects of age, metal abundance, helium, reddening, and systematic errors in the photometry, demonstrating that these could not explain the short distance to the Pleiades from {\\it Hipparcos}. \\citet{terndrup02} paid attention to the adopted reddening law in a discussion of the distance to NGC~2516. Here we extend the error analysis more quantitatively, emphasizing photometric calibration issues and the bias in distance estimates induced by the presence of unresolved cluster binaries or field foreground/background stars. This paper also explores the effect of metallicity on the luminosity of the MS. Metallicity changes isochrone luminosities more strongly than many other input parameters, and the degree of sensitivity depends on the color index used. This permits a purely photometric derivation of the metallicity \\citep[e.g.,][]{pinsono98,pinsono00,stello01,terndrup02}, which can be compared to metallicities derived from high-resolution spectra. An agreement between the photometric and spectroscopic metallicities, as we find in this paper, provides supporting evidence that the effects of metallicity on the theoretical quantities ($L$, $T_{\\rm eff}$) and on the color-$T_{\\rm eff}$ relations are correctly computed. The distances in this paper are tied to the Hyades distance at $(m - M)_0 = 3.33 \\pm 0.01$ ($d = 46.34 \\pm 0.27$ pc), the cluster's center-of-mass inferred from the {\\it Hipparcos} catalog \\citep{perryman98}. Unlike the controversial {\\it Hipparcos} distance to the Pleiades, the large angular diameter of the Hyades on the sky makes the cluster parallax less vulnerable to the spatial correlation of the {\\it Hipparcos} parallax \\citep{narayanan99a,narayanan99b,debruijne01}. In \\S~2 we compile cluster photometry, metallicities, reddening estimates, and information on binarity and membership, and present a metallicity for Praesepe from new high-resolution spectra. In \\S~3 we briefly describe the construction of the isochrones. In \\S~4 we compute the distances to the sample clusters with the reddening fixed at previously known values and demonstrate that the empirical corrections improve the internal precision of the isochrones. In \\S~5 we simultaneously solve for the cluster metallicity, reddening, and distance from the $\\chi^2$ minimization. In \\S~6 we evaluate the effects of several systematic error sources, including those from cluster binaries and field star contamination. In \\S~7 we discuss several implications of our results. In the Appendix we address issues on the photometric zero points of the empirical Hyades isochrone. ", "conclusions": "" }, "0607/gr-qc0607134_arXiv.txt": { "abstract": "Because no closed timelike curve (CTC) on a Lorentzian manifold can be deformed to a point, any such manifold containing a CTC must have a topological feature, to be called a timelike wormhole, that prevents the CTC from being deformed to a point. If all wormholes have horizons, which typically seems to be the case in space-times without exotic matter, then each CTC must transit some timelike wormhole's horizon. Therefore, a Lorentzian manifold containing a CTC may nevertheless be causally well behaving once its horizon's are deleted. For instance, there may be a Cauchy-like surface through which every timelike curve passes one and only once before crossing a horizon. ", "introduction": "\\label{Censorshipofchronologicalviolations} An argument that CTCs should not be considered pathological is that all CTCs typically transit a timelike wormhole which prevent the CTC from being deformed to a point (see Monroe \\cite{Monroe:2006b}), and all wormholes in a vacuum space-time typically have an event horizon. Such a CTC is then ``censored\" by passing through the wormhole's event horizon and will be called a \\textit{c-CTC}; other CTCs are \\textit{uncensored-} or \\textit{u-CTCs}. The causal structure of a space-time ignoring c-CTCs can be examined by analyzing a space-time with points on event horizons excised. Modifying the definitions to this setting, most results Hawking and Ellis \\cite{Hawking-Ellis} Chapter 6 carry over in some form, and the excised space-time may be well-behaved by modified criteria. For instance, a Cauchy-like surface may exist through which all timelike curves pass one and only once before crossing an event horizon. A few results do not carry over: the proof of Proposition 6.4.2, stating that compact space-times contain CTCs, cannot be modified to show these are u-CTCs. The Appendix presents restatements of Chapter 6's definitions and results that are not needed for this section. Consider a space-time $(\\mathscr{M},\\textbf{g})$. Define $\\mathscr{M}_c=\\mathscr{M}-\\dot{J}^-(\\mathscr{I}^+,\\overline{\\mathscr{M}})-\\dot{J}^+(\\mathscr{I}^-,\\overline{\\mathscr{M}})$, that is, excise all points on event horizons. Below, the causal structure of a space-time $\\mathscr{M}$ containing c-CTCs is considered by analyzing $\\mathscr{M}_c$. Say a time-orientable space-time $\\mathscr{M}$ satisfying the Einstein vacuum equation is \\textit{safe} if all wormholes have event horizons; this is a weak assumption in a vacuum assumption in the absence of exotic matter (Morris and Thorne \\cite{Morris:1988cz}). This section and the next will establish that safe space-times may contain c-CTCs cannot contain no u-CTCs and are therefore well behaved by standard criteria with appropriate modifications. \\begin{theorem}\\label{trivialtopology} Suppose space-time $\\mathscr{M}$ is safe. Then, $\\mathscr{M}_c$ is simply connected. \\end{theorem} \\begin{proof} In the absence of topological defects, every CTC passes through a wormhole (see Monroe \\cite{Monroe:2006b}). Every wormhole of $\\mathscr{M}$ has an event horizon in a safe space-time, so the space-time $\\mathscr{M}$ becomes simply connected when its event horizons are excised to create $\\mathscr{M}_c$. \\end{proof} Suppose $\\mathscr{M}$ is time orientable. A \\textit{u-curve} is any curve of non-zero extent on $\\mathscr{M}_c$. For sets $\\mathscr{S}$ and $\\mathscr{U}$, the \\textit{c-chronological future} $I^+_c(\\mathscr{S},\\mathscr{U})$ \\textit{of} $\\mathscr{S}$ \\textit{relative to} $\\mathscr{U}$ is the set of all points in $\\mathscr{U}$ which can be reached from $\\mathscr{S}$ by a future-directed timelike u-curve. $I^+_c(\\mathscr{S},\\mathscr{M}_c)$ will be denoted by $I^+_c(\\mathscr{S})$. The \\textit{c-causal future of} $\\mathscr{S}$ \\textit{relative to} $\\mathscr{U}$, denoted by $J^+_c(\\mathscr{S},\\mathscr{U})$, is defined similarly. A point $p$ is a \\textit{c-future endpoint} of a future-directed non-spacelike u-curve $\\gamma:F\\rightarrow\\mathscr{M}_c$ if for every neighborhood $\\mathscr{V}$ of $p$ there is a $t\\in{}F$ such that $\\gamma(t_1)\\in\\mathscr{V}$ for every $t_1\\in{}F$ with $t_1\\geq{}t$. A non-spacelike u-curve is \\textit{c-future-inextendible} (respectively \\textit{c-future-inextendible in a set} $\\mathscr{S}$) if it has no future endpoint in $\\mathscr{M}_c$ (respectively \\textit{c-future-inextendible in a set} $\\mathscr{S}$). A set $\\mathscr{S}$ is \\textit{c-achronal} if $I^+_c(\\mathscr{S})\\cap\\mathscr{S}$ is empty, in other words, if there are no two points of $\\mathscr{S}$ that lie on a timelike u-curve. $\\mathscr{S}$ is a \\textit{c-future set} if $\\mathscr{S}\\supset{}I^+_c(\\mathscr{S})$. An example of a result from Hawking and Ellis that carries over with the modified causality definitions above is: \\prop{Proposition 6.3.1} If $\\mathscr{S}$ is a c-future set then $\\dot{\\mathscr{S}}$, the boundary of $\\mathscr{S}$ in $\\mathscr{M}_c$ is a closed, imbedded, c-achronal three-dimensional $C^{1-}$ manifold. A set with the properties of $\\dot{\\mathscr{S}}$ in Proposition 6.3.1$'$ a \\textit{c-achronal boundary}. An open set $\\mathscr{U}$ is \\textit{c-causally simple} if for every compact set $\\mathscr{K}\\subset\\mathscr{U}$, $\\dot{J}^+_c(\\mathscr{K})\\cap\\mathscr{U}=E^+_c(\\mathscr{K})\\cap\\mathscr{U}$ and $\\dot{J}^-_c(\\mathscr{K})\\cap\\mathscr{U}=E^-_c(\\mathscr{K})\\cap\\mathscr{U}$. A space-time satisfies the \\textit{c-chronology condition} if there are no u-CTCs. Trivially, the c-chronology condition is weaker than the chronology condition, because the latter rules out c-CTCs as well. The set of all points in $\\mathscr{M}_c$ that lie on a u-CTC is called the \\textit{c-chronology violating} set of $\\mathscr{M}_c$. For a safe space-time $\\mathscr{M}$, the c-chronology violating set of $\\mathscr{M}_c$ is empty, by definition. Proposition 6.4.2 is stated in its original form, because it cannot be modified to apply to $\\mathscr{M}_c$: \\textit{\\textbf{Proposition 6.4.2}} If $\\mathscr{M}$ is compact, the chronology violating set of $\\mathscr{M}$ is non-empty. Hawking and Ellis point out that any compact, four-dimensional manifold on which there is a Lorentzian metric cannot be simply connected. They conclude from this property and Proposition 6.4.2 that it is reasonable to assume space-time is compact. By contrast, if c-CTCs are considered reasonable, it will be argued below that compactness is a desirable feature. If $\\mathscr{M}$ is compact and therefore contains CTCs, the reasoning in Hawking and Ellis' proof of Proposition 6.4.2 does not necessarily imply that $\\mathscr{M}_c$ contains CTCs if $\\mathscr{M}$ contains a wormhole with event horizon---in that case $\\mathscr{M}_c$ is not compact at the excised points on event horizons. Let $\\mathscr{M}_c'$ be the compactification of $\\mathscr{M}_c$; in other words, event horizons of $\\mathscr{M}$ are unidentified, and $\\mathscr{M}_c'$ remains simply connected (see Theorem \\ref{trivialtopology}). Although $\\mathscr{M}_c'$ is compact, the proof cannot assume $\\mathscr{M}_c'$ can be covered by open sets of the form $I^+_c(q)$. For a given wormhole $w$, label the unidentified horizons $\\mathscr{I}^-_w$ and $\\mathscr{I}^+_w$ respectively; this notation will be justified by Theorem \\ref{causalboundarytheorem}. \\begin{theorem} $\\mathscr{M}_c'$ cannot be covered by open sets of the form $I^+_c(q)$ if $\\mathscr{M}$ contains a wormhole with event horizon. If $\\mathscr{M}_c'$ is covered by open c-future sets, these have a finite subcover. \\end{theorem} \\begin{proof} For any point $q\\in\\mathscr{I}^-_w$, $q\\notin{}I^+_w(q)$. If $\\mathscr{M}_c'$ is covered by open c-future sets, some sets must include a non-zero measure of $\\mathscr{I}^-_w$, so there is a finite subcover. \\end{proof} The \\textit{strong c-causality condition} is said to hold at $p$ if every neighborhood of $p$ contains a neighborhood of $p$ which no non-spacelike u-curve intersects more than once. The region $D^+_c(\\mathscr{S})$ to the future of $\\mathscr{S}$ is called the \\textit{c-future Cauchy development} or \\textit{c-domain of dependence} of $\\mathscr{S}$, defined as the set of all points $p\\in\\mathscr{M}_c$ such that every c-past-inextendible non-spacelike u-curve through $p$ intersects $\\mathscr{S}$. Appropriate data on a closed set $\\mathscr{S}$ would determine events in not only in $\\mathscr{M}_c$ but also in $\\mathscr{M}$, and there is an the additional consistency condition in the latter case. The future boundary of $D^+_c(\\mathscr{S})$, that is $\\overline{D^+_c(\\mathscr{S})}-I^-(D^+_c\\mathscr{S})$, marks the limit of the region that can be predicted from knowledge of data on $\\mathscr{S}$. Call this closed c-achronal set the \\textit{c-future Cauchy horizon} of $\\mathscr{S}$ and denote it by $H^+_c(\\mathscr{S})$. Define the c-edge$(\\mathscr{S})$ for a c-achronal set $\\mathscr{S}$ as the set of all points $q\\in\\overline{\\mathscr{S}}$ such that in every neighborhood $\\mathscr{U}$ of $q$ there are points $p\\in{}I^-(q,\\mathscr{U})$ and $r\\in{}I^+(q,\\mathscr{U})$ which can be joined by a timelike u-curve in $\\mathscr{U}$ which does not intersect $\\mathscr{S}$. \\prop{Corollary to Proposition 6.5.3} If c-edge$(\\mathscr{S})$ vanishes, then $H^+_c(\\mathscr{S})$, if non-empty, is a c-achronal three-dimensional imbedded $C^{1-}$ manifold which is generated by null geodesic segments which have no past endpoint in $\\mathscr{M}_c$. Call such a c-acausal set $\\mathscr{S}$ with no c-edge a \\textit{partial c-Cauchy surface}. That is, a partial c-Cauchy surface is a spacelike hypersurface which no non-spacelike u-curve intersects more than once. Define $D_c(\\mathscr{S})=D^+_c(\\mathscr{S})\\cup{}D^-_c(\\mathscr{S})$. A partial c-Cauchy surface $\\mathscr{S}$ is a global c-Cauchy surface (or simply \\textit{c-Cauchy surface}) if $D_c(\\mathscr{S})$ equals $\\mathscr{M}_c$. That is, a c-Cauchy surface is a spacelike hypersurface which every non-spacelike u-curve intersects once. Friedman \\cite{Friedman:2004jr} notes that the initial value problem is well-defined on a class of space-times broader than those which are globally hyperbolic. It is shown below that this larger class includes c-globally hyperbolic space-times. Suppose one is given a three-dimensional manifold $\\mathscr{S}$ with certain initial data $\\omega$ on it. The \\textit{c-Cauchy problem} requires one to find a four-dimensional manifold $\\mathscr{M}$ (without subscript), an imbedding $\\theta:\\mathscr{S}\\rightarrow\\mathscr{M}$ and a metric \\textbf{g} on $\\mathscr{M}$ which satisfies the Einstein equations, agrees with the initial values on $\\theta(\\mathscr{S})$ and is such that $\\theta(\\mathscr{S})$ is an c-Cauchy surface for $\\mathscr{M}$. Note that the c-Cauchy problems for $\\mathscr{M}$ and $\\mathscr{M}_c$ are distinct; in the former case, there are additional consistency constraints between $\\mathscr{I}^+_w$ and $\\mathscr{I}^-_w$. \\begin{theorem} For a safe space-time $\\mathscr{M}$, initial data on a c-Cauchy surface $\\mathscr{S}_c$ is sufficient to predict not only $\\mathscr{M}_c$ but also $\\mathscr{M}_c'$. \\end{theorem} \\begin{proof} By the definition of c-Cauchy surface, $D_c(\\mathscr{S}_c)$ equals $\\mathscr{M}_c$. The event horizons of $\\mathscr{M}_c$ are limit points of $D_c(\\mathscr{S}_c)$. Therefore, the Cauchy development includes the event horizons. \\end{proof} Say a safe space-time $\\mathscr{M}$ is \\textit{Novikov consistent} with respect to c-Cauchy surface $\\mathscr{S}$ if the c-Cauchy development of $\\mathscr{S}$ in $\\mathscr{M}_c$ coincides with the Cauchy development of $\\mathscr{S}$ in $\\mathscr{M}$. It is not obvious that Novikov consistency holds between $\\mathscr{I}^-_w$ and $\\mathscr{I}^+_w$. A set $\\mathscr{N}$ is \\textit{c-globally hyperbolic} if the strong c-causality assumption holds on $\\mathscr{N}$ and if for any two points $p,q\\in\\mathscr{N}$, $J_c^+(p)\\cap{}J_c^-(q)$ is compact and contained in $\\mathscr{N}$. \\prop{Proposition 6.6.8} If an open set $\\mathscr{N}$ is c-globally hyperbolic, then $\\mathscr{N}$, regarded as a manifold, is homeomorphic to $R^1\\times\\mathscr{S}$ where $\\mathscr{S}$ is a three-dimensional manifold, and for each $a\\in{}R^1$, $a\\times\\mathscr{S}$ is an c-Cauchy surface for $\\mathscr{N}$. Theorem \\ref{trivialtopology} shows that $\\mathscr{M}_c$ is simply connected, so there is no topological obstacle to $\\mathscr{M}_c$ being c-globally hyperbolic. ", "conclusions": "" }, "0607/astro-ph0607490_arXiv.txt": { "abstract": "Multi-band (0.9 to 1.6 $\\mu$m) images of the TW\\,Hydrae Association (TWA) brown dwarf, 2MASSWJ\\,1207334$-$393254 (also known as 2M1207), and its candidate planetary mass companion (2M1207b) were obtained on 2004 Aug 28 and 2005 Apr 26 with HST/NICMOS. The images from these two epochs unequivocally confirm the two objects as a common proper motion pair (16.0\\,$\\sigma$ confidence). A new measurement of the proper motion of 2M1207 implies a distance to the system of $59\\pm7$\\,pc and a projected separation of $46\\pm5$\\,AU. The NICMOS and previously published VLT photometry of 2M1207b, extending overall from 0.9 to 3.8 $\\mu$m, are fully consistent with an object of a few Jupiter masses at the canonical age of a TWA member ($\\sim8$\\,Myr) based on evolutionary models of young giant planets. These observations provide information on the physical nature of 2M1207b and unambiguously establish that the first direct image of a planetary mass companion in orbit around a self-luminous body, other than our Sun, has been secured. ", "introduction": "Beginning in 2004 July, with HST/NICMOS, we initiated a systematic imaging search for extra-solar gas giant planets around 116 young nearby stars and brown dwarfs. These targets are $\\lesssim50$\\,Myr old and located within 60\\,pc of Earth, making them among the best known targets for such a survey \\citep{ARAA}. A brown dwarf in the $\\sim8$\\,Myr old TW~Hydrae Association, 2MASSWJ\\,1207334$-$393254 (\\citealt{Gizis02}; hereafter 2M1207), was included in our HST target list and its observation was planned for 2005 April. However, on 2004 Apr 27 (UT), 2M1207 was observed with VLT/NACO and a faint companion candidate was discovered $\\sim0\\farcs78$ from the brown dwarf \\citep{Chauvin1}. NICMOS observations of 2M1207 were replanned and brought forward to 2004 August. The resulting NICMOS photometric data, shorter in wavelength than could be obtained with adaptive optics on the VLT, support the conjecture that 2M1207b is of mid- to late-L type based upon its color indices. With the limited precision of the proper motion data then available for 2M1207 and the short time between the VLT and 2004 August HST observations, common proper motion with 2M1207b was established at the $2.6\\,\\sigma$ level \\citep{Schneider}. Additional observations were obtained with the VLT during 2005 February and March that much more precisely demonstrated common proper motion between 2M1207 and its companion \\citep{Chauvin2}. As described in Section 4.1, the proper motion value of 2M1207 \\citep{Scholz} used in \\cite{Chauvin2} was not well measured, causing the analysis to be somewhat over-optimistic. With the higher accuracy HST astrometry and a new, more accurate, proper motion measurement of 2M1207 in the present paper, we report a more definitive common proper motion between 2M1207 and 2M1207b. We also present short near-IR wavelength diagnostic photometry which cannot currently be obtained from the ground given the performance limitations of adaptive optics imaging. ", "conclusions": "The common proper motion confirmation of the first imaged planetary mass companion to a celestial object other than our Sun enables the onset of a new era in extra-solar planet characterization -- direct spectroscopic analysis. Absorption spectroscopy of stellar light reprocessed through atmospheres of planets detected through radial velocity surveys has been demonstrated (e.g., HD~209458b; \\citealt{Brown}). 2M1207b provides the first opportunity to collect and spectroscopically analyze photons from an extra-solar planetary mass companion. Exploiting the superb stability of the HST, we will attempt to obtain a near-IR grism spectrum of 2M1207b. Relative to clear atmospheric models, dusty models \\citep[for example]{Dusty} predict more flux suppression in the $1.0-1.3 \\mu$m range, which can be readily compared to the anticipated S/N$\\sim$10 grism spectrum." }, "0607/astro-ph0607459_arXiv.txt": { "abstract": "We report high angular resolution, multi-epoch radio observations of the young pulsar PSR B1800$-$21. Using two pairs of data sets, each pair spanning approximately a 10 year period, we calculate the proper motion of the pulsar. We obtain a proper motion of $\\mu_{\\alpha}=11.6 \\pm 1.8$~mas\\,yr$^{-1}$, $\\mu_{\\delta}=14.8 \\pm 2.3$~mas\\,yr$^{-1}$, which clearly indicates a birth position at the extreme edge of the W30 supernova remnant. Although this does not definitively rule out an association of W30 and PSR B1800$-$21, it does not support an association. ", "introduction": "Pulsar-supernova remnant associations are important for various reasons. When an association is confirmed, knowledge of the supernova remnant (SNR) can be used to constrain pulsar (PSR) parameters such as birth magnetic fields, spin periods, luminosities, and beaming fractions. Conversely, knowledge of the pulsar can constrain SNR ages and distances, and illuminate remnant evolution and uncommon morphologies (Kaspi 1996). Although over 30 candidate pulsar-SNR associations are purported, fewer than a dozen have been confirmed. Although confirmation of an association is desirable, the opposite --- disproving an association --- can also be useful. Pulsar proper motions, which are often cited to support associations, can be a definitive means to disprove candidate associations. Even if an association is disproven, the pulsar transverse velocity is useful information. It can constrain stellar collapse models and contribute to studies of pulsar velocity distributions and galactic electron density distributions (e.g., Chatterjee et al.\\ 2001). ", "conclusions": "The newly measured proper motion, together with an adopted distance, can be used to determine the transverse velocity of the pulsar. The NE2001 distance, $D = 3.8 \\pm 0.4$~kpc, was used to compute a transverse velocity of $347^{+57}_{-48}$~km~s$^{-1}$. This calculation includes a correction of 0.64~mas~yr$^{-1}$ for the combined effects of the Sun's peculiar motion and the differential Galactic rotation. The net effect of this correction is that the LSR transverse speed is 8~km~s$^{-1}$ greater than its apparent speed. This speed is well within the normal range of young pulsar speeds (see, e.g. Arzoumanian et al. 2002 and Brisken et al. 2003) which suggests that the DM distance estimate is reasonable. Independently of the distance estimation, an assumed pulsar age, together with the measured apparent proper motion, allows the calculation of the pulsar birthplace coordinates. In particular, assuming 15.8 kyr for B1800$-$21, and a proper motion of $\\mu_{\\mathrm{total}} = 18.7^{+1.4}_{-1.5}$~mas~yr$^{-1}$ $38.1^\\circ$ east of north, its birthplace in J2000.0 is $\\alpha_0 =$ 18$^{\\mathrm h}$ 03$^{\\mathrm m}$ 38.0$^{\\mathrm s}$ $\\pm$ 2.3$^{\\mathrm s}$, $\\delta_0 = -21^\\circ 41' 18.2'' \\pm 41.6''$. The characteristic age for the pulsar is based purely on the observed pulse period and its derivative and is a good proxy for the true age given two conditions: (1) spindown is due to the magnetic dipole radiation leading to braking index $n=3$, and (2) the initial spin period was much less than the current spin period. Livingstone et al. (2005) note that all braking index measurements made give values less than the nominal $n=3$. Only very young pulsars have such measurements so it might be safe to assume that B1800$-$21 has $n < 3$. In fact, an index as low as 2 is reasonable, which would double its true age. If so, the birth site would be even further from the W30 remnant, and would nearly coincide with the HII region G8.14+0.23 (IRAS 17599$-$2148) and the dark cloud, traced by mid-IR emission (see Fig.~1). We note that this H\\,II region is probably unrelated to the pulsar-SNR complex; radio recombination line velocities suggest that it is kinematically distinct from the other HII regions found near the remnant \\citep{lockman89}. Faucher-Giguere \\& Kaspi (2006) have tabulated initial spin periods that have been estimated for nine young pulsars. Four have values less than or about 30~ms, four more are between 50 and 90~ms, and one exceptional case (PSR~J0538+2817) is quite long, at 140~ms. Based solely on these values it seems unlikely that the true age is less than half the characteristic age. However, serious selection effects have likely biased this sample. Through population modeling Faucher-Giguere \\& Kaspi (2006) conclude that a wider range of birth spin periods is actually likely. Assuming a normally-distributed birth spin period, they report a distribution with $\\mu=300$~ms and $\\sigma=150$~ms. Given that the current spin period of B1800$-$21 is 134~ms, this distribution is clearly not useful in estimating $P_0$ for this pulsar. Thus the characteristic age is likely correct to a factor of 2; claiming an age range much smaller than this is not well-justified. It is clear that the pulsar did {\\it not} originate at the geometric center of W30, as seen in the \\citet{kw90} image. First, the birthplace coordinates are far-removed from the SNR center for any assumed age. Second, the pulsar is currently at a position approximately 106$^\\circ$ west of north (in equatorial coordinates) with respect to the center of W30. The proper motion, at 38$^\\circ$ east of north, is more nearly {\\it toward} the center of W30, rather than away from it. A more sensitive image of the W30 remnant has been reported by \\citet{bgg06} (see Fig.~1). This image shows W30 to be the bright, eastern part of a larger, non-thermal nebula. Despite the more extended radio continuum emission seen in the Brogan et al. image, the pulsar birth position lies outside of the radio remnant (see also the three color image of Brogan et al.). \\citet{finley94} suggested that confinement by inhomogeneous, dense, star-forming gas may have shaped a complicated, asymmetric remnant, with the pulsar lying near one edge. The ambient interstellar medium, as traced by thermal 8~$\\mu$m radiation, clearly suggests an extensive interaction with the remnant. Moreover, as conjectured by Finley \\& \\\"Ogelman, the interstellar gas at lower Galactic longitudes probably does absorb X-ray emission from the remnant, resulting in the morphology that they reported. However, the {\\it MSX} image shows relatively little emission at Galactic latitudes north of the pulsar position, implying less (or at least cooler) cloud material in this direction. Hence, one would naively expect the SNR to expand toward northern latitudes in addition to or instead of toward higher Galactic longitude, as seen in the 90~cm image. Although we cannot definitively rule out some pathological SNR morphology, the inferred birth position of the pulsar does not lend support to an association. The Brogan et al. (2006) 90~cm image also shows the newly-discovered SNR G8.31$-$0.09 (seen in Fig.~1 as the cyan and yellow contours approximately 0.2$^\\circ$ below the pulsar positions). The proper motion we measure rules out this remnant as the B1800$-$21 birth site. In addition, the range of possible birth locations do not coincide with any interesting portion of the TeV source HESS J1804$-$216 \\citep{ahar06}. Associations between young pulsars and supernova remnants are attractive as it is widely accepted that a single supernova can produce such a pair. Some SNR / PSR associations are quite secure, such as PSR J0538+2817 and SNR S147 \\citep{klh03}. However, for both PSRs B1800$-$21 and B1757$-$24 the geometric coincidence and lack of contradictory data initially led to speculation of association with SNRs W30 and W28, respectively. In the case of B1757$-$24, a proper motion upper limit suggests insufficient velocity to allow for an association \\citep{tbg02}. In general, the lack of an associated radio pulsar with a known SNR is not surprising for several reasons: not all supernovae result in pulsars; not all pulsars are oriented for favorable detection at Earth, and some pulsars are intrinsically too faint to be seen. Young pulsars without shell-type SNRs are more mysterious. The Crab nebula, created in year 1054 with PSR B0531+21, is a pulsar-powered nebula without a well-defined SNR, possibly indicating a very low energy supernova \\citep{fkcg95}. The time since the supernova, the density structure of the ambient ISM, and the mass of the progenitor star greatly affect the visibility of SNRs. The lack of associations with either PSR B1800$-$21 or SNR W30 should not be alarming." }, "0607/astro-ph0607203_arXiv.txt": { "abstract": "A residual planetesimal disk of mass 10--$100M_\\earth$ remained in the outer solar system following the birth of the giant planets, as implied by the existence of the Oort cloud, coagulation requirements for Pluto, and inefficiencies in planet formation. Upon gravitationally scattering planetesimal debris, planets migrate. Orbital migration can lead to resonance capture, as evidenced here in the Kuiper and asteroid belts, and abroad in extra-solar systems. Finite sizes of planetesimals render migration stochastic (``noisy''). At fixed disk mass, larger (fewer) planetesimals generate more noise. Extreme noise defeats resonance capture. We employ order-of-magnitude physics to construct an analytic theory for how a planet's orbital semi-major axis fluctuates in response to random planetesimal scatterings. The degree of stochasticity depends not only on the sizes of planetesimals, but also on their orbital elements. We identify the conditions under which the planet's migration is maximally noisy. To retain a body in resonance, the planet's semi-major axis must not random walk a distance greater than the resonant libration width. We translate this criterion into an analytic formula for the retention efficiency of the resonance as a function of system parameters, including planetesimal size. We verify our results with tailored numerical simulations. Application of our theory reveals that capture of Resonant Kuiper belt objects by a migrating Neptune remains effective if the bulk of the primordial disk was locked in bodies having sizes $< \\mathcal{O}(100)$ km and if the fraction of disk mass in objects with sizes $\\gtrsim 1000$ km was less than a few percent. Coagulation simulations produce a size distribution of primordial planetesimals that easily satisfies these constraints. We conclude that stochasticity did not interfere with nor modify in any substantive way Neptune's ability to capture and retain Resonant Kuiper belt objects during its migration. ", "introduction": "Planet formation by coagulation of planetesimals is not perfectly efficient---it leaves behind a residual disk of solids. Upon their coalescence, the outer planets of our solar system were likely embedded in a 10--$100 M_\\earth$ disk of rock and ice containing the precursors of the Oort cloud (Dones et al.~2004) and the Kuiper belt (see the reviews by Chiang et al.~2006; Cruikshank et al.~2006; Levison et al.~2006). The gravitational back-reaction felt by planets as they scatter and scour planetesimals causes the planets to migrate (Fern\\'andez \\& Ip 1984; Murray et al.~1998; Hahn \\& Malhotra 1999; Gomes, Morbidelli, \\& Levison 2004). Neptune is thought to have migrated outward and thereby trapped Kuiper belt objects (KBOs) into its exterior mean-motion resonances, both of low-order such as the 3:2 (Malhotra 1995) and of high-order such as the 5:2 (Chiang et al.~2003; Hahn \\& Malhotra 2005). Likewise, Jupiter's inward migration may explain the existence of Hilda asteroids in 2:3 resonance with the gas giant (Franklin et al.~2004). A few pairs of extra-solar planets, locked today in 2:1 resonance (Vogt et al.~2005; Lee et al.~2006), may have migrated to their current locations within parent disks composed of gas and/or planetesimals. Orbital migration and resonant trapping of dust grains may also be required to explain non-axisymmetric structures observed in debris disks surrounding stars 10--100 Myr old (e.g., Wyatt 2003; Meyer et al.~2006). Only when orbital migration is sufficiently smooth and slow can resonances trap bodies. The slowness criterion requires migration to be adiabatic: Over the time the planet takes to migrate across the width of the resonance, its resonant partner must complete at least a few librations. Otherwise the bodies speed past resonance (e.g., Dermott, Malhotra, \\& Murray 1988; Chiang 2003; Quillen 2006). Smoothness requires that changes in the planet's orbit which are incoherent over timescales shorter than the libration time do not accumulate unduly. Orbital migration driven by gravitational scattering of discrete planetesimals is intrinsically not smooth. A longstanding concern has been whether Neptune's migration was too ``noisy'' to permit resonance capture and retention (see, e.g., Morbidelli, Brown, \\& Levison 2003). In N-body simulations of migration within planetesimal disks (Hahn \\& Malhotra 1999; Gomes et al.~2004; Tsiganis et al.~2005), N $\\sim \\mathcal{O}(10^4)$ is still too small to produce the large, order-unity capture efficiencies seemingly demanded by the current census of Resonant KBOs. At the same time, the impediment against resonance capture introduced by inherent stochasticity has been exploited to explain certain puzzling features of the Kuiper belt, most notably the Classical (non-Resonant) belt's outer truncation radius, assumed to lie at a heliocentric distance of $\\sim$48 AU (Trujillo \\& Brown 2001; Levison \\& Morbidelli 2003). If Neptune's 2:1 resonance captured KBOs and released them en route, Classical KBOs could have been transported (``combed'') outwards to populate the space interior to the final position of the 2:1 resonance, at a semi-major axis of 47.8 AU (Levison \\& Morbidelli 2003). As originally envisioned, this scenario requires that $\\sim$$3 M_\\earth$ be trapped inside the 2:1 resonance so that an attendant secular resonance suppresses growth of eccentricity during transport. It further requires that the degree of stochasticity be such that the migration is neither too smooth nor too noisy. Whether these requirements were actually met remain open questions.\\footnote{While Classical KBOs do have semi-major axes that extend up to 48 AU, the distribution of their perihelion distances cuts off sharply at distances closer to 45 AU (see, e.g, Figure 2 of Chiang et al.~2006). Interpreted naively (i.e., without statistics), the absence of bodies having perihelion distances of 45--48 AU and eccentricities less than $\\sim$0.1 smacks of observational bias and motivates us to re-visit the problem of whether an edge actually exists, or at least whether the edge bears any relation to the 2:1 resonance.} Stochastic migration has also been studied in gas disks, in which noise is driven by density fluctuations in turbulent gas. Laughlin, Steinacker, \\& Adams (2004) and Nelson (2005) propose that stochasticity arising from gas that is unstable to the magneto-rotational instability (MRI) can significantly prolong a planet's survival time against accretion onto the parent star. The spectrum of density fluctuations is computed by numerical simulations of assumed turbulent gas. In this work, we study stochastic changes to a planet's orbit due to planetesimal scatterings. The planet's Brownian motion arises from both Poisson variations in the rate at which a planet encounters planetesimals, and from random fluctuations in the mix of planet-planetesimal encounter geometries. How does the vigor of a planet's random walk depend on the masses and orbital properties of surrounding planetesimals? We answer this question in \\S\\ref{sec-oom} by constructing an analytic theory for how a migrating planet's semi-major axis fluctuates about its mean value. We employ order-of-magnitude physics, verifying our assertions whenever feasible by tailored numerical integrations. Because the properties of planetesimal disks during the era of planetary migration are so uncertain, we consider a wide variety of possibilities for how planetesimal semi-major axes and eccentricities are distributed. One of the fruits of our labors will be identification of the conditions under which a planet's migration is maximally stochastic. Apportioning a fixed disk mass to fewer, larger planetesimals renders migration more noisy. How noisy is too noisy for resonance capture? What limits can we place on the sizes of planetesimals that would keep capture of Resonant KBOs by a migrating Neptune a viable hypothesis? These questions are answered in \\S\\ref{sec-sizes}, where we write down a simple analytic formula for the retention efficiency of a resonance as a function of disk properties, including planetesimal size. Quantifying the size spectrum of planetesimals is crucial for deciphering the history of planetary systems. Many scenarios for the evolution of the Kuiper belt implicitly assume that most of the mass of the primordial outer solar system was locked in planetesimals having sizes of $\\mathcal{O}(100)$ km, like those observed today (see, e.g., Chiang et al.~2006 for a critique of these scenarios). By contrast, coagulation simulations place the bulk of the mass in bodies having sizes of $\\mathcal{O}(1)$ km (Kenyon \\& Luu 1999). For ice giant formation to proceed {\\it in situ} in a timely manner in the outer solar system, most of the primordial disk may have to reside in small, sub-km bodies (Goldreich, Lithwick, \\& Sari~2004). In \\S\\ref{sec-sum}, in addition to summarizing our findings, we extend them in a few directions. The main thrust of this paper is to analyze how numerous, small perturbations to a planet's orbit accumulate. We extend our analysis in \\S\\ref{sec-sum} to quantify the circumstances under which a single kick to the planet from an extremely large planetesimal can disrupt the resonance. We also examine perturbations exerted directly on Resonant KBOs by ambient planetesimals. ", "conclusions": "\\label{sec-sum} We summarize our findings in \\S\\ref{sec-sum1} and discuss quantitatively some remaining issues in \\S\\ref{sec-extend}. \\subsection{Summary}\\label{sec-sum1} Newly formed planets likely occupy remnant planetesimal disks. Planets migrate as they exchange energy and angular momentum with planetesimals. Driven by discrete scattering events, migration is stochastic. In our solar system, Neptune may have migrated outward by several AU and thereby captured the many Kuiper belt objects (KBOs) found today in mean-motion resonance with the planet. While resonance capture is efficient when migration is smooth, a longstanding issue has been whether Neptune's actual migration was too noisy to permit capture. Our work addresses---and dispels---this concern by supplying a first-principles theory for how a planet's semi-major axis fluctuates in response to intrinsic granularity in the gravitational potential. We apply our theory to identify the environmental conditions under which resonance capture remains viable. Stochasticity results from random variations in the numbers and orbital properties of planetesimals encountering the planet. The degree of stochasticity (as measured, say, by $\\sigma_{a_{\\rm p},T}$, the typical distance that the planet's semi-major axis random walks away from its average value) depends on how planetesimal semi-major axes $a$ and random velocities $u$ are distributed. We have parameterized $a$ by its difference from the planet's semi-major axis: $x \\equiv a-a_{\\rm p} \\equiv \\mathcal{R} R_{\\rm H}$, where $R_{\\rm H}$ is the Hill sphere radius and $\\mathcal{R} \\gtrsim 1$. In the case of high dispersion when $u > \\mathcal{R} v_{\\rm H}$ (where $v_{\\rm H} \\equiv \\Omega_{\\rm p}R_{\\rm H}$ is the Hill velocity and $\\Omega_{\\rm p}$ is the planet's orbital angular velocity), planetesimal orbits cross that of the planet. Stochasticity increases with decreasing $u$ in the high-dispersion case because the cross-section for strong scatterings increases steeply with decreasing velocity dispersion (as $1/u^4$). In the intermediate-dispersion case when $v_{\\rm H}/\\mathcal{R}^2 < u < \\mathcal{R}v_{\\rm H}$, planetesimal and planet orbits do not cross, and stochasticity decreases with decreasing $u$. In the low-dispersion case when $u < v_{\\rm H}/\\mathcal{R}^2$, the amount of stochasticity is insensitive to $u$. The values of $u$ and $\\mathcal{R}$ which actually characterize disks are unknown. The random velocity $u$, for example, is expected to be set by a balance between excitation by gravitational scatterings and damping by inelastic collisions between planetesimals and/or gas drag. Damping depends, in turn, on the size distribution of planetesimals. These considerations are often absent from current N-body simulations of planetary migration in planetesimal disks. Despite such uncertainty, we can still identify the circumstances under which stochasticity is maximal. Maximum stochasticity obtains when $\\mathcal{R} \\sim 1$ and $u \\lesssim v_{\\rm H}$, that is, when planetesimals have semi-major axes within a Hill radius of the planet's and when their velocity dispersion is no greater than the Hill velocity. A stochastically migrating planet cannot retain objects in a given resonance if the planet's semi-major axis random walks away from its average value by a distance greater than the maximum libration width of the resonance. This simple criterion is validated by numerical experiments and enables analytic calculation of the resonance retention efficiency as a function of disk parameters. A disk of given surface density generates more noise when composed of fewer, larger planetesimals. In the context of Neptune's migration, we estimate that if the bulk of the minimum-mass disk resided in bodies having sizes smaller than $\\mathcal{O} (100)$ km and if the fraction of the disk mass in larger bodies was not too large ($\\lesssim$ a few percent for planetesimals having sizes of $1000 \\km$, for example), then the retention efficiency of Neptune's first-order resonances would have been of order unity ($\\gtrsim 0.1$). Such order-unity efficiencies seem required by observations, which {\\it prima facie} place 122/474 $\\approx$ 26\\% of well-observed KBOs (excluding Centaurs) inside mean-motion resonances (Chiang et al.~2006). Drawing conclusions based on a comparison between this observed percentage and our theoretical retention percentage $P_{\\rm keep}$ is a task fraught with caveats---a more fair comparison would require, e.g., disentangling the observational bias against discovering Resonant vs.~non-Resonant objects; account of the attrition of the Resonant population due to weak chaos over the four-billion-year age of the solar system; and knowledge of the initial eccentricity and semi-major axis distributions of objects prior to resonance sweeping, as these distributions impact capture probabilities in different ways for different resonances (Chiang et al.~2003; Hahn \\& Malhotra 2005; Chiang et al.~2006). But each of these caveats alters the relevant percentages only by factors of a few, and when combined, their effects tend to cancel. Therefore we feel comfortable in our assessment that $P_{\\rm keep}$ must have been of order unity to explain the current Resonant population. In that case, $\\mathcal{O} (100\\km)$ is a conservative estimate for the maximum allowed size of planetesimals comprising the bulk of the disk mass, derived for the case of maximum stochasticity. How does an upper limit of $\\mathcal{O}(100)\\km$ compare with the actual size distribution of the planetesimal disk? While today's Kuiper belt places most of its mass in objects having sizes of $\\sim$100 km, this total mass is tiny---only $\\sim$$0.1 M_{\\earth}$ (Bernstein et al.~2004; see Chiang et al.~2006 for a synopsis). The current belt is therefore 2--3 orders of magnitude too low in mass to have driven Neptune's migration. The current size distribution is such that bodies having radii $\\gtrsim 40\\km$ are collisionless over the age of the solar system and might therefore represent a direct remnant, unadulterated by erosive collisions, of the planetesimal disk during the era of migration (Pan \\& Sari 2005). If so, the bulk of the primordial disk mass must have resided in bodies having sizes $\\lesssim 40 \\km$. Theoretical calculations of the coagulation history of the Kuiper belt are so far consistent with this expectation. Kenyon \\& Luu (1999) find, for their primordial trans-Neptunian disk of $10 M_{\\earth}$, that 99\\% of the mass failed to coagulate into bodies larger than $\\mathcal{O}(1)\\km$, because the formation of several Pluto-sized objects (comprising $\\sim$0.1\\% of the total mass) excited velocity dispersions so much that planetesimal collisions became destructive rather than agglomerative. The average-mass planetesimals in their simulation have sizes $\\mathcal{O}(1) \\km$, much smaller than even our most conservative estimate of the maximum allowed size of $\\mathcal{O}(100)\\km$. For a given size distribution of planetesimals, most stochasticity is produced by the size bin having maximal $\\eta\\, m^2$, which need not be the size bin containing the majority of the mass. Here, $\\eta$ and $m$ are the number of planetesimals and the mass of an individual planetesimal in a logarithmic size bin. For power-law size distributions $d\\eta/ds \\propto s^{-q}$ such that $q < 7$, stochasticity is dominated by the largest planetesimals. For disks having as much mass as the minimum-mass disk of solids and whose largest members are Pluto-sized, size distributions with $q \\ge 4$ enjoy order-unity efficiencies for resonance retention. The size distributions of Kenyon \\& Luu (1999) resemble $q=4$ power laws, but with a large overabundance of planetesimals having sizes of $\\mathcal{O}(1)\\km$. This sequestration of mass dramatically reduces the stochasticity generated by the largest bodies, which have sizes of $\\mathcal{O}(1000) \\km$. We conclude that Neptune's Brownian motion did not impede in any substantive way the planet's capture and retention of Resonant KBOs. \\subsection{Extensions}\\label{sec-extend} \\subsubsection{Single Kick to Planet} Our focus thus far has been on the regime in which many stochastic kicks to the planet are required for resonant particles to escape. Of course, a single kick from a planetesimal having sufficiently large mass $m_1$ could flush particles from resonance. To estimate $m_1$, we equate the change in the planet's semi-major axis from a single encounter, $\\Delta a_{\\rm p}$, to the maximum half-width of the resonance, $\\delta a_{\\rm p,lib}/2$ (see Equation [\\ref{eqn-aplib}] and related discussion). In the likely event that the perturber's eccentricity $e$ is of order unity, then $\\max (\\Delta a_{\\rm p}) \\sim (m_1/M_{\\rm p}) a_{\\rm p} e$ (Equation [\\ref{eqn-delamax}]) and therefore $m_1 \\gtrsim 0.6 \\, (0.5/e) \\, M_{\\oplus}$ for Plutinos to escape resonance. Our estimate for $m_1$ agrees with that of Malhotra (1993). Why such enormous perturbers have not been observed today is unclear and casts doubt on their existence (Morbidelli, Jacob, \\& Petit 2002). If such an Earth-mass planetesimal were present over the duration $T$ of Neptune's migration, then the likelihood of a resonance-destabilizing encounter would be $P_1 \\sim \\dot{\\overline N}T \\sim 10^{-2} (0.5/e)^4$, where the encounter rate $\\dot{\\overline N}$ is given by Equation (\\ref{eqn-Ndot}) with $\\Sigma_{m} \\sim m_1 / (2\\pi a_{\\rm p}^2)$, and we have set $T\\sim 2.7\\times 10^7 \\yr$. \\subsubsection{Kicks to Resonant Planetesimals} Finally, we have ignored in this work how disk planetesimals directly perturb the semi-major axis of a resonant particle. This neglect does not significantly alter our conclusions. Take the resonant planetesimal to resemble a typical Resonant KBO observed today, having size $s_{\\rm res} \\sim 100 \\km$. Then its Hill velocity is $e_{\\rm H,res}\\Omega a \\sim 10^{-4} \\Omega a$. The relative velocity between the resonant planetesimal and an ambient, perturbing planetesimal greatly exceeds this Hill velocity, if only because migration in resonant lock quickly raises the eccentricity of the resonant planetesimal above $e_{\\rm H,res}$. Equation (\\ref{eqn-superrms}), appropriate for the super-Hill regime, implies that $\\langle \\dot{a}_{\\rm p,rnd}^2 \\rangle^{1/2} \\propto M_{\\rm p}^0$---the RMS random velocity does not depend on the mass of the object being perturbed! Therefore when both the resonant planetesimal and the planet are scattering planetesimals in the super-Hill regime, their random walks are comparable in vigor. The conservative limit of $\\mathcal{O}(100)$ km on the planetesimal size which allows resonance retention is derived, by contrast, for the sub-Hill, maximum stochasticity regime, and is therefore little affected by these considerations.\\footnote{The probability that a planetesimal will be ejected from resonance by a planetesimal of comparable mass in a single encounter is negligibly small.}" }, "0607/astro-ph0607035_arXiv.txt": { "abstract": "The effect of photon-beam-induced turbulence on propagation of radio emission in a pulsar magnetosphere is discussed. Beamed radio emission with a high brightness temperature can generate low-frequency plasma waves in the pulsar magnetosphere and these waves scatter the radio beam. We consider this effect on propagation of radio emission both in the open field line region and in the closed field line region. The former is applicable to most cases of pulsar radio emission where the propagation is confined to the polar region; it is shown that the induced process is not effective for radio emission of moderately high brightness temperature but can have a severe effect on giant pulses. For giant pulses not to be affected by this process, they must be emitted very close to the light cylinder. We show that the induced process is efficient in the closed field line region, inhibiting propagation of the radio emission in this region. ", "introduction": "It is generally thought that pulsar radio emission is produced in the open field line region well inside the light cylinder (LC) (the radius at which the corotation speed equals $c$). Radio emission generated inside the LC must propagate over a substantial part of the magnetosphere before escaping to the interstellar medium. The intense radio emission is subject to various interactions with the intervening plasmas within the pulsar magnetosphere. Two types of interactions have been considered: cyclotron absorption in which electrons (positrons) are in cyclotron resonance with the wave (Blandford \\& Scharlemann 1976; Fussell, Luo, \\& Melrose 2001), and induced Compton scattering in which the radio wave is scattered to a higher or lower frequency~\\citep{bs76,p04b}. Both processes can affect or even disrupt the propagation of the beam. So far, no clear observational evidence for the absorption feature associated with the cyclotron resonance has been identified. However, the conventional polar cap theory, which assumes that the radio emission is generated in an outflowing electron-positron pair plasma, does predict a cyclotron resonance region where absorption should occur. Induced Compton scattering can be important in principle for any coherent emission with a sufficiently high intensity~\\citep{wr78,sk92,cbr93,hm95}. While cyclotron absorption is important only in the region near the LC, induced scattering can occur throughout the magnetosphere. Apart from these two processes, there is a third process that involves induced wave-wave interactions and that can significantly affect propagation of the radio beam. A highly collimated beam of radio waves can generate low frequency waves through three-wave interactions in the pulsar magnetospheric plasma in a way similar to wave generation in a particle beam instability. These low frequency waves interact with the radio waves, leading to a decrease in its intensity and an increase in its angular spread. In some cases, such effects can completely disrupt the radio beam: The medium becomes opalescent to the radio waves. In this paper, we discuss the effect of induced three-wave interactions on propagation of radio emission in a pulsar magnetosphere. Induced three-wave processes have been considered for pulsar winds \\citep{lm94}, pulsar eclipse by the stellar wind of its companion~\\citep{t-etal94,m94,lm95} and radio emission from active galactic nuclei~\\citep{k88,gk93,lb95}. Apart from a discussion of Raman scattering (a particular case of three-wave interactions, cf. Sec. 2.5) in a pulsar magnetosphere~\\citep{gk93,l98}, we are aware of no detailed study of the three-wave effect which explicitly takes account of the beaming and broadband nature of the radio emission. To our knowledge, there is no detailed study of propagation of radio emission in the closed field line region. Here we consider the three wave effect both in the open field line region (OFLR) and in the closed field line region (CFLR). In most models, propagation of radio emission is confined to the OFLR where the relativistic electron positron pair plasma streams along magnetic field lines at a relativistic velocity. Because electrons and positrons contribute to three-wave interactions equally but with opposite sign, three-wave interactions are significant weakened in a quasineutral pair plasma. However, such processes can still be important for radio emission with very brightness temperature provided that there is a modest excess of electrons or positrons and that the bulk Lorentz factor is modest. Since for young pulsars, the conventional polar cap model generally predicts a plasma density too high for propagation of radio waves, we assumes that the plasma is highly inhomogeneous across the polar cap and that radio emission is generated at a frequency near the plasma frequency (in the plasma's rest frame) in the underdense region where the plasma density is relatively low (cf. Sec. 3.1). This assumption seems to be required for consistency if the radio emission is produced due to a beam instability \\citep{mg99}. Propagation of radio emission in the CFLR is relevant at least for three specific cases: (1) the recently discovered double pulsar system PSR J0737-3039B in which the radio emission from one pulsar interacts with the plasma in the CFLR of the other, (2) radio emission from outer magnetospheric regions where the emission of the trailing component may propagate to the CFLR as a result of the aberration effect, and (3) backward emission predicted by the oscillatory polar gap model~\\citep{lmjl05} and by recent models for the peculiar `notch-like' feature of pulse profiles~\\citep{d-etal05}. For case (1), eclipse of the radio emission from one pulsar by the magnetosphere of the other has been observed. A possible interpretation of the eclipse is due to processes in the CLFR such as cyclotron or synchrotron absorption or induced scattering. For case (2) the trailing component may have an absorption-like feature or may be completely destroyed as the result of strong induced processes. For the third case, induced processes can severely constrain visibility of the backward emission that propagates through the CFLR. In Sec. 2, we outline a general formalism of induced three-wave interactions in the pulsar magnetosphere. We discuss the effect on propagation of radio emission in the OFLR in Sec. 3 and in CFLR in Sec. 4. ", "conclusions": "We consider the effect of induced three-wave turbulence on propagation of the radio emission in a pulsar magnetosphere. (As remarked in Sec. 2.5, this process is sometimes referred to, inappropriately we believe, as `stimulated Raman' or `induced Raman' scattering.) The low frequency wave is either an electrostatic plasma wave propagating along magnetic field lines or an Alfv\\'en wave propagating nearly along the magnetic field. We discuss both cases of propagation in the OFLR and CFLR. In the polar (OFLR) region, the induced processes strongly depend on the bulk velocity of the plasma and on the pair multiplcity ($M$). Since the scattering rate (cf. Eq. \\ref{eq:Gamma3}) decreases for increasing $\\gamma$ and $M$, the induced three-wave scattering effect on the propagation is suppressed due to the relativistic bulk flow ($\\gamma\\gg1$) and $M\\gg1$. Assuming a modest bulk Lorentz factor $\\gamma\\sim10^2$ and a multiplicity $M\\sim10^2$ for a young pulsar like the Crab pulsar, and $\\gamma\\sim 10^5$ and $M\\sim 10$ for a fast millisecond pulsar like PSR 1957+21, we show that propagation of radio emission in the OFLR is little affected by the induced three-wave processes. An exception is the propagation of giant pulses. If the giant pulses are emitted relatively close to the star, they are subject to very strong induced processes that prevent them from escaping. The favored region where induced processes can be effective is the transition region located at $r_*/\\gamma$ from the source, where the photon beam propagates nearly perpendicular to the magnetic field in the rest frame. For giant pulses not to be completely destroyed by such processes, their emission region needs to be close to the LC. This conclusion is consistent with the conclusion reached by \\cite{rj01} from a different argument, based on the alignment of giant pulses with the high energy emission. Induced three-wave interactions are generally important for propagation of radio emission in the CFLR. Because the plasma in this region is stationary in the corotating frame, with $|\\eta|=1$, the processes are much more efficient than in the polar region. The induced three-wave processes are important throughout the CFLR except for pulsars with a relatively long period. We discuss three specific examples where the processes are relevant. (1) For the double pulsar J0737-3039B, we find that the processes discussed here are not efficient enough to explain the observed eclipse, due to the very low plasma density of the CFLR of pulsar B which has a relatively long period. However, the induced three-wave scattering can be effective if one assumes a density much higher than the GJ density (by a factor of $10^4$). It is worth noting that a much higher plasma density is needed in the synchrotron/cyclotron absorption model as well~\\citep{lt05,rg05}. (2) Our result imposes a strong constraint on the location of the emission region of pulsars with a pulse profile with widely separated leading and trailing components. There is a particular radial range where the trailing component may be swept into the CLFR due to aberration. (3) Our result limits the visibility of backward emission. The backward emission must propagate through the CFLR where it can be dispersed due to three wave interactions. For young or fast millisecond pulsars, the radius of the opaque region (where induced three-wave interactions are important) is comparable with or larger than the LC radius. In this case, backward emission may not be visible. For long period pulsars, the plasma density in the CFLR is low and the radius of the opaque region is smaller than the LC. The backward emission can propagate through the magnetosphere provided that it is produced sufficiently close to the LC and that its propagation path is outside the opaque sphere. An important approximation made in our discussion is the strong magnetic field limit. Although such an approximation excludes the possibility of cyclotron resonance, our result should be valid for a finite, strong magnetic field. This is because the relevant frequency of the low frequency waves (Langmuir and Alfv\\'en modes) in the induced three-wave interactions considered here is much lower than the cyclotron frequency throughout the pulsar magnetosphere. \\appendix" }, "0607/astro-ph0607279_arXiv.txt": { "abstract": "Submilliarcsecond astrometry and imaging of the black hole Sgr A* at the Galactic Center may become possible in the near future at infrared and submillimetre wavelengths. This resolution is sufficient to observe the silhouette the supermassive black hole in the Galactic center casts upon background emission. However, more exciting is the prospect of observing ``hot spots'' in the accretion flow. Here we discuss how such measurements may be used to test not only the consistency of General Relativity, but also the validity of the Kerr metric in particular. ", "introduction": "Testing strong field gravity remains one of the primary objectives of observational astronomy. Due to their compact nature, black holes provide an ideal environment to do this. Nevertheless, an unambiguous confirmation of strong field relativity has been elusive thus far. There have been a number of attempts to probe the strong gravity regime, including observations of the relativistically broadened Fe K$\\alpha$ line (see, \\eg, \\cite{Pari-Brom-Mill:01,Reyn-Nowa:03}), interpretations of quasi-periodic oscillations (QPOs) (see, \\eg, \\cite{Remi:05,Genz_etal:03}), and multiwavelength spectropolarimetric observations (see, \\eg, \\cite{Conn-Star-Pira:80,Laor-Netz-Pira:90,Brod-Loeb:06}). However, the interpretation of each of these are dependent upon unknown accretion physics, making the implications for general relativity ambiguous. For example, the failure to find an expected correlation between the variability in the Fe K$\\alpha$ line emission and the soft X-ray continuum implies that the simplest emission models for the Fe K$\\alpha$ observations are incomplete (see, \\eg, \\cite{Wang_etal:99,Chia_etal:00,Lee_etal:00,Wang-Wang-Zhou:01,Weav-Gelb-Yaqo:01}), though attempts to rectify this with the inclusion of strong gravitational lensing have been made \\cite{Matt-Fabi-Reyn:97}. In addition, alternative explanations for the formation of the broad iron lines exist (see, \\eg, \\cite{Elvi:00,You-Liu-Chen-Chen-Zhan:03}), further complicating their interpretation. The most commonly discussed black hole QPO's are those observed in the X-ray spectra of stellar-mass black hole candidates. These typically have $Q$'s of $10$--$100$ and are at kilohertz frequencies. Unfortunately, in the absence of a definitive theory for how these are produced, the identification of these with the epicycles of black hole spacetimes is tenuous at best. A second class of QPO's are those observed in supermassive black holes. In the context of Sgr A*, these are observed in the near-infrared (NIR) and X-ray bands, have $Q$'s of $3$, (though see \\cite{Bela_etal:06}), and periods on the order of 20 minutes. If these are interpreted as the Keplerian orbital periods of the innermost stable circular orbit (ISCO), they imply black hole spins as high as $0.5$. More recently, arguments based upon the lack of a thermal peak in the spectra have been used to infer the absence of a surface in stellar-mass black hole candidates (see, \\eg, \\cite{Garcia_etal:01}) and Sgr A* \\cite{Brod-Nara:06}. In the latter case, where the putative thermal emission due to the small accretion rate peaks in the near infrared (NIR), this result appears especially robust. However, these arguments present only a qualitative confirmation of the observational consistency of general relativity. In contrast, it is now technologically feasible to image a black hole directly. A background illuminated black hole will appear in silhouette, with an angular size of roughly twice that of the horizon \\cite{Bard:73}, and may be directly observed. With an expected resolution of $\\sim20\\,\\muas$, submillimeter very-long baseline interferometry (VLBI) would be able to image the silhouette cast upon the accretion flow by Sgr A* (with an angular scale of $\\sim50\\,\\muas$), and M87 ($\\sim25\\muas$) \\cite{Falc-Meli-Agol:00,Brod-Loeb:05b}. In principle, detailed measurements of the size and shape of the silhouette could yield information about the mass and spin of the central black hole. In practice, the interpretation of such an image will likely depend upon the accretion flow model employed (this is discussed in more detail in \\S\\ref{S}). Sgr A* has exhibited strong flares in the NIR and X-ray \\cite{Ghez_etal:04,Ecka_etal:04,Genz_etal:03,Baga_etal:01}, and more recently in the submillimeter \\cite{Marr:06}, implying that the innermost portions of the emitting region are strongly variable. A simple model for the flares, motivated by the evidence of periodicity, is that of transient orbiting bright regions, hot spots, which dominate the flaring luminosity. Such hot spots appear inevitable, the product of shocks and magnetic reconnection events within the accretion flow. Due to its dynamical and compact nature, images of such a spot will contain significantly more information about the spacetime. Each imaged spot will allow the measurement of both, the mass and spin of the black hole. Combining observations of many hot spots orbiting at different radii provides a way to test not only the consistency of general relativity, but the validity of the Kerr metric generally, and the no-hair theorems specifically. In \\S\\ref{S} and \\S\\ref{HS} we discuss the expected images for a radiatively inefficient accretion flow (RIAF) and hot spots, respectively, in the context of the Galactic center. Some concluding remarks are in \\S\\ref{C}. ", "conclusions": "\\label{C} Despite their putative size, imaging the black hole horizons at submillimeter wavelengths is now technologically feasible. In addition, phase-referenced astrometry will be available in the NIR in the coming decade. These capabilities will allow, for the first time, direct observations of gravity in the strongly non-linear regime. Observations of the size, shape and location of the silhouette cast by the black hole on the surrounding accretion flow may be used in principle to determine its mass and spin. However, in practice this is unlikely to be simple due to uncertainties in accretion physics. Nevertheless, at the very least it will provide a means to learn about accretion onto compact objects. Of more interest will likely be observations of hot spots. Due to their compact and dynamical nature their images contain more information than those of the underlying quiescent accretion flow. It appears to be possible, even with the simplest models, to constrain the hot spot parameters (\\eg, orbit, size and spectral index). If this is indeed the case, then hot-spot observations will provide a method for quantitatively testing general relativity. \\ack A.E.B. gratefully acknowledges the support of an ITC Fellowship from Harvard College Observatory. A. L. was supported in part by NASA grants NAG 5-1329 and NNG05GH54G and by the Clark/Cooke fund of Harvard University." }, "0607/astro-ph0607565_arXiv.txt": { "abstract": "{} {We compute the afterglow of gamma-ray bursts produced by purely electromagnetic outflows to see if it shows characteristic signatures differing from those obtained with the standard internal/external shock model.} {Using a simple approach for the injection of electromagnetic energy to the forward shock we obtain the afterglow evolution both during the period of activity of the central source and after. Our method equally applies to a variable source.} {Afterglow light curves in the visible and X-ray bands are computed both for a uniform medium and a stellar wind environment. They are brighter at early times than afterglows obtained with the internal/external shock model but relying only on these differences to discriminate between models is not sufficient.} {} ", "introduction": "Lyutikov and Blandford (2003) proposed an alternative to the standard fireball model where the central engine produces a purely electromagnetic outflow instead of a re\\-la\\-tivistic baryonic wind. Observationally this electro\\-ma\\-gne\\-tic model (hereafter EMM) differs from the standard internal/external shock model by the absence of any reverse shock contribution, a different early afterglow evolution and a high polarization of the prompt emission (Lyutikov, 2004). In this paper we concentrate on the early afterglow (while the central source is still active) and compare the EMM to the standard model in X-rays and the visible for a uniform external medium or a wind environment. In Sect.2 we obtain simple equations that go\\-vern the evolution of the forward shock propagating in the burst environment. Their solutions are used in Sect.3 to compute afterglow light curves which are compared to those obtained in the standard model for the same total injected energy. We discuss our results in Sect.4 and conclude that it likely will be difficult to decide between models from afterglow observations only. ", "conclusions": "The lightcurves in Fig.3 show that the EMM and the standard model notably differ at early times (during the period of source acti\\-vi\\-ty). However relying on these differences alone to identify the physical origin of GRBs will be a difficult task requiring a very early follow-up of the afterglow. In X-rays, SWIFT should be able to do that (at least in some cases) but the problem here will come from the mixing of the afterglow contribution with the brighter prompt emission component. This mixing will also probably prevent an unambiguous detection of the imprint of source variability on the X-ray afterglow (Fig.4). In the visible, where the burst prompt emission is weak and probably negligible (see however the recent RAPTOR observations of GRB 041219a and GRB 050820a (Vestrand et al., 2005, 2006)), the EMM predicts a brighter afterglow for a given set of parameters $\\epsilon_e$, $\\epsilon_B$, $n$ or $A_*$. But in real afterglows these parameters are not known a priori and deciding between models will be tricky. Polarization properties of the burst prompt emission (Lyutikov, 2004) when they become more easily accessible may provide clearer evidence. A last interesting point concerning the EMM is related to the shallow part observed by SWIFT in many X-ray afterglows. The light curves in Fig. 3 show that the EMM indeed predicts an initially flat region in the early X-ray afterglow. However this flat region does not last more than the period of source activity (150 s in observer time in Fig.3). It would then extend to $10^4$ s (or more) only if the source can remain active for that duration, as was also suggested for the standard model (Zhang et al., 2005)." }, "0607/astro-ph0607086_arXiv.txt": { "abstract": " ", "introduction": "Thanks to recent data about the Cosmic Microwave Background (CMB), Large Scale Structures (LSS) and also Type Ia Supernov\\ae\\ (SNe), cosmology has become the most sensitive probe of some neutrino properties (e.g.\\ of neutrino masses: oscillation experiments test squared-mass differences, and other means of probing the absolute neutrino mass are currently less sensitive) and a very sensitive probe of other neutrino properties, including non standard ones~\\cite{review,LesgPasReview,BS}. In this paper we study how present cosmological data determine standard and non-standard `neutrino cosmology'. This includes three different issues. \\begin{itemize} \\item[i)] testing neutrinos: their masses, abundances, \\ldots \\item[ii)] do photons, neutrinos and gravitons make up the complete list of light particles? Data from particle physics allow extra light particles that are neutral under the Standard Model (SM) gauge group, and such extra light particles appear in many speculative extensions of the SM, one interesting example being simpler string models.\\footnote{Within the string scenario light particles can be avoided at the price of assuming that strings vibrate on complicated enough higher dimensional geographies~\\cite{AntroString}, such that predictivity seems lost.} \\item[iii)] The two above issues can be connected, because neutrinos are the least tested light particles and can easily interact with new light neutral particles, in a way that affects the evolution of cosmological inhomegeneities. \\end{itemize} In section~\\ref{th} we characterize the fundamental theories and describe the cosmological parameters that we want to extract from present data. Since our implementation of the cosmological computational tools needed for this analysis somewhat differs from the standard one, we describe it in section~\\ref{tool}. Section~\\ref{res} describes our results (table \\ref{tab:navigator} might help in navigating the paper), summarized in the conclusions. \\begin{table}[t] \\begin{center} \\begin{tabular}{r|lr|lr|lr|lr} & \\multicolumn{2}{c|}{Neutrino} & \\multicolumn{6}{c}{Cosmology with extra light particles} \\\\ & \\multicolumn{2}{c|}{cosmology} & \\multicolumn{2}{c|}{freely-streaming} & \\multicolumn{2}{c|}{self-interacting} & \\multicolumn{2}{c}{interacting with $\\nu$} \\\\ \\hline & & & & & & & & \\\\[-1.9mm] massless &\\parbox{8ex}{$A_s,n_s,h$,\\\\ $\\Omega_b,\\Omega_{\\rm DM},\\tau$} & \\S\\ref{tool} & $\\Delta N_\\nu$ & \\S\\ref{sec:Nnu} & $\\Delta N_\\nu$ & \\S\\ref{NnuI} & $N_\\nu,R\\equiv{\\displaystyle \\frac{N_\\nu^{\\rm normal}}{N_\\nu }}$ & \\S\\ref{00fluid} \\\\[8mm] \\ massive&$\\sum m_\\nu$ &\\S\\ref{sec:mNu}&$\\Delta N_\\nu, m_{\\rm s}$ & \\S\\ref{sec:nus}&$\\Delta N_\\nu, m_{\\rm s}$&\\S\\ref{sec:NnuIm}&$R=0, m_\\nu$ or $m_\\phi$ &\\S\\ref{m0fluid},\\ref{0mfluid} \\\\ \\end{tabular} \\caption{\\label{tab:navigator}\\em Schematization of the cases considered in this paper. For each one we list the notation of the relevant parameters probed by cosmology and refer to the relevant part of the text.} \\end{center} \\end{table} ", "conclusions": "We compared a non exhaustive but representative casistics of how cosmology is affected by extra light particles (with sub-keV masses), or by standard and non-standard properties of neutrinos, using CMB, LSS, Lyman-$\\alpha$, BAO, SN data. \\begin{itemize} \\item First, we considered ordinary massive neutrinos. We obtain the cosmological bound on neutrino masses, $\\sum m_\\nu \\circa{<}0.40\\eV$ at $99.9\\%$ C.L.\\ and fig.\\fig{Nm}a shows that the relatively less safe observations play a crucial r\\^ole. \\item The density of initially relativistic particles can be parameterized in terms of the usual number $N_\\nu$ of equivalent neutrinos. Assuming that all the $N_\\nu$ relativistic particles freely stream, we find that their density is constrained to be $N_\\nu=5\\pm 1$. The $2\\sigma$ preference for $N_\\nu>3$ is mainly due to the $2\\sigma$ anomaly in the Lyman-$\\alpha$ measurement of the matter power spectrum. \\item Assuming ordinary neutrinos plus an extra component of interacting particles, we find $\\Delta N_\\nu = 0\\pm1.3$. Fig.\\fig{NRnu} shows how data constrain the intermediate case where both kinds of relativistic particles are present. It is interesting that the uncertainty on $\\Delta N_\\nu$ is decreasing below 1. \\item The extra light particles might have a mass $m$ and an abundance $\\Delta N_\\nu$. Fig.\\fig{mN} shows how data constrain these parameters in the two limiting cases that these extra particles freely stream (fig.\\fig{mN}a) or interact among themselves (fig.\\fig{mN}b). \\item Finally, we considered one extra scalar of mass $m_\\phi$ that interacts with neutrinos of mass $m_\\nu$. We find that this scenario is strongly disfavored by the global fit, at about $4\\sigma$. \\end{itemize} All these results are based on assumptions and subject to caveats, that we discussed in the text. Technically, our analysis somewhat differs from typical analyses because we used a code developed by us and dealt with statistics using Gaussian analytical techniques, that become adequate nowadays that observations are rich and precise enough. Eq.s \\eq{means} and\\eq{corr} allow to check how well we reproduce the standard results for standard cosmology. \\small \\paragraph" }, "0607/astro-ph0607615_arXiv.txt": { "abstract": "{}{Despite extensive observations over the last decades, the central questions regarding the power source of the large IR luminosity of Ultra Luminous Infra Red Galaxies (ULIRGs), and their evolution, are still not fully answered. In this paper we will focus on massive star formation as a central engine and present an evolutionary model for these dust-enshrouded star formation regions.}{An evolutionary model was created using existing star formation and radiative transfer codes (STARBURST99, RADMC and RADICAL) as building blocks. The results of the simulations are compared to data from two IRAS catalogs.} {From the simulations it is found that the dust surrounding the starburst region is made up from two components. There is a low optical depth ($\\tau=0.1$, which corresponds to 0.1 \\% of the total dust mass), hot (T$\\sim$400K) non-grey component close to the starburst (scale size 10pc) and a large scale, colder grey component (100pc, 75K) with a much larger column ($\\tau=10$). The simulations also show that starburst galaxies can be powered by massive star formation. The parameters for this star forming region are difficult to determine, since the IR continuum luminosity is only sensitive to the total UV input. Therefore, there is a degeneracy between the total starburst mass and the initial mass function (IMF) slope. A less massive star formation with a shallower IMF will produce the same amount of OB stars and therefore the same amount of irradiating UV flux. Assuming the stars are formed according to a Salpeter IMF ($\\Psi(M) \\propto M^{-2.35}$), the star formation region should produce $10^9$~M$_{\\sun}$ ~of stars (either in one instantaneous burst, or in a continuous process) in order to produce enough IR radiation.} {Our models confirm that massive star formation is a valid power source for ULIRGs. In order to remove degeneracies and further determine the parameters of the physical environment also IR spectral features and molecular emissions need to be included.} ", "introduction": "\\label{sec:introduction} In 1983 the Infra-Red Astronomical Satellite (IRAS) surveyed 96\\% of the sky in four broad-band filters at 12$\\mu$m, 25$\\mu$m, 60$\\mu$m, and 100$\\mu$m. IRAS detected infrared (IR) emission from about 25,000 galaxies, primarily from spirals, but also from quasars (QSOs), Seyfert galaxies and early type galaxies. Among these galaxies IRAS discovered a new class of galaxies that radiate most of their energy in the infrared. The most luminous of these infrared galaxies, the (ultra-)luminous infrared galaxies [(U)LIRGs], have QSO-like luminosities of L $\\ge 10^{11}$ L$_{\\sun}$ ~(LIRGs) or even L $\\ge 10^{12}$ L$_{\\sun}$ ~(ULIRGs) \\citep{2000ARA&A..38..761G}. Despite extensive observations over the last decades, the central questions regarding the power source of the large IR luminosity of ULIRGs, and their evolution, are still not fully answered. \\cite{1988ApJ...325...74S} proposed that most ULIRGs are powered by dust-enshrouded QSOs in the late phases of a merger. The final state of such a merger would be a large elliptical galaxy with a massive quiescent black hole at its center \\citep{1992ApJ...390L..53K}. A significant fraction of the ULIRG population seems to confirm this assumption, since they exhibit nuclear optical emission line spectra similar to those of Seyfert galaxies \\citep{1988ApJ...325...74S}. Some also contain compact central radio sources and highly absorbed, hard X-ray sources, all indicative of an active nucleus (AGN). On the other hand, the (Far)IR, mm, and radio characteristics of ULIRGs are similar to those of starburst galaxies. A centrally condensed burst of star formation activity (called a starburst, hereafter denoted as SB), for instance fueled by gas driven into the center of the potential well of a pair of interacting galaxies by a bar instability, provides an equally plausible power source \\citep[e.g.][]{2002PhR...369..111B}. Observational evidence for the starburst nature of ULIRGs was found with the detection of a large number of compact radio hypernovae in each of the two nuclei of \\object{Arp220} by \\cite{1998ApJ...493L..17S}. \\begin{figure} \\centering \\includegraphics[angle=0,width=\\hsize]{figure1} \\caption{Schematic overview of the ``physical'' environment created in the simulations. In the center star formation is going on, which irradiates the dusty surroundings (here represented as a toroidal shape). This dust then re-radiates its given energy in the IR regime. Also the parameter space is shown. In the starburst, the total stellar mass ${\\rm M_{SB}}$, the star formation rate (SFR) and the IMF-slope $\\alpha$ are varied. In the dust the geometry is varied by changing the closing angle $\\eta$. Also the dust column density (expressed as the optical depth $\\tau$) is varied. A last variable is the observational inclination $\\delta$. Also some parameters are shown, which initially were not varied: the radius $r_0$ and the radial scale size $dr_0$. More information on the parameter space can be found in Sect. \\ref{sec:simulations}.} \\label{fig:setup} \\end{figure} Since both options are observationally supported, it is natural to think that there might be an evolutionary relation between the two. Several authors have suggested such schemes. In \\citeyear{1988ApJ...330..743B}, \\citeauthor{1988ApJ...330..743B} explained the evolution of the FIR properties of active nuclei with a model which incorporated both a relatively rapid decreasing thermal SB component and a slower evolving non-thermal component. Recently, more papers were published suggesting evolutionary schemes which incorporate both SB and AGN. \\cite{2006ApJ...637..104K} presented a sample of ULIRGs with type I Seyfert nuclei and showed with both observations and modeling that this type of galaxies can be powered by both a SB and a black hole (BH) with super-Eddington accretion. They suggest a scheme in which SB-powered ULIRGs and QSOs are two stages in the evolution of the host galaxy. The host will start as a SB-driven ULIRG, with only a small BH. Over time, the SB fades and the BH will grow into a super massive BH (SMBH), which will dominate the energy output of the host. By that time the host has become a QSO. Similar arguments are made by \\cite{2006ApJS..163....1H} and \\cite{2005astro.ph.11157H}. On the other hand, \\cite{2005ApJ...635L.121K}, \\cite{2005MNRAS.364.1337S} and \\cite{2005astro.ph.11157H} present schemes where the outflow of a super-Eddington accreting SMBH drives into the surrounding ISM, creating bubbles in which the gas cools and stars are formed. It is clear that detailed modeling is necessary to determine the evolution of ULIRGs and their engine(s). \\begin{figure} \\centering \\includegraphics[angle=0,width=\\hsize]{figure2} \\caption{This flowchart shows the computational setup of the simulation. First the starburst properties are calculated using STARBURST99 (Sect. \\ref{sec:starburst99}) and the results of these calculations are used as input for the dust calculations, which are performed by RADMC and RADICAL (Sects. \\ref{sec:radmc} and \\ref{sec:radical}).} \\label{fig:code-flow} \\end{figure} In this paper, we focus on massive star formation as a central engine. In SB-powered ULIRGs, that are at sufficiently low redshift for their internal structure to be resolved, the great majority of the IR emission is found to originate from sub-kpc dusty regions within merging systems of galaxies \\citep[e.g.][]{1998ApJ...507..615D}. These dusty SB galaxies are an important class of objects. About 25\\% of the high-mass star formation within 10 Mpc distance from us occurs in just four SB galaxies \\cite[\\object{M82}, \\object{NGC253}, \\object{M83}, \\object{NGC4945};][]{1998ASPC..148..127H}. Even though these galaxies create vast amounts of stars, the time scale of this formation is short. Near-IR imaging spectroscopy in \\object{M82}, \\object{IC342}, and \\object{NGC253} indicates that in the evolution of these galaxies there are several episodes of star formation activity, with timescales of around $10^7$ to $10^8$ years \\citep{2000ARA&A..38..761G}. The relatively low efficiency of the energy production of stars ($E\\sim 10^{-3} {\\rm M}_{\\star}{\\rm c}^2$) and the large energy output (up to $\\sim 10^{61}$ ergs for an ultra-luminous starburst like \\object{Arp220}) yield a production $10^8$ to $10^{10}$ M$_{\\sun}$ ~of stars per burst. Combined with the short timescales, this leads to star formation rates (SFRs) ranging from 10 up to 1000 M$_{\\sun}$ per year \\citep{1998ASPC..148..127H}. Numerical simulations confirm this picture. \\cite{2006ApJ...637..255J} show that during a typical merger event, there are several short periods of star formation at a very high rate. The goal of this paper is to present a model for these dust-enshrouded star formation regions and to study and explain the behavior of the broad band IR continuum properties of starburst ULIRGs during their evolution. The model consists of a number of existing codes, which are combined into one. In future work, we will also investigate the spectral features in the IR regime, which will provide more diagnostics to determine the source of the activity in ULIRGs \\citep[e.g.][]{1999RvMP...71..173H,2000ARA&A..38..761G,2003PhDT........18S}. We also intend to extend the model with a molecular environment, in order to further constrain the physical parameters of the cores of active galaxies \\citep{2005A&A...436..397M}. We will also investigate the similarities and differences between SB dominated and AGN dominated ULIRGs and will investigate the possibility of an evolutionary connection between the two. The structure of this paper is as follows: in Sect. \\ref{sec:model} the model is discussed. The parameter space and the simulations are presented in Sect. \\ref{sec:simulations} and the results of these simulations in Sect. \\ref{sec:results}. In the last section the results are discussed and suggestions for future work are made. ", "conclusions": "Our simulation model is able to reproduce the IR continuum properties of starburst galaxies, by using two dust components: a large scale component containing the bulk of the mass and a less massive component close to the starburst. The long wavelength radiation is influenced by macro-physics like the star formation activity and the large scale dust distribution. The short wavelength emission at 10 $\\mu$m comes from hot dust and is influenced by the stellar dust and micro-physics like the optical properties and optical depth of the dust. Not all parameters have a profound effect on the results. The stellar dust optical depth ($\\tau_2$) and the parameters controlling the star formation (the total stellar mass ${\\rm M_{SB}}$ and the IMF slope $\\alpha$) have a large effect on the final results, whereas the influence of the large scale dust geometry ($\\tau$ and $\\eta$) on the final IR properties is smaller.\\\\ \\\\ {\\bf Star formation region}\\\\ Because of their large influence on the results, the star formation parameters could be determined well. Increasing the stellar mass (${\\rm M_{SB}}$) has two effects. First of all, the shape of the spectrum changes, the short wavelength fluxes increase as compared to the long wavelengths. A second effect is an increase in ${\\rm L_{IR}}$, which changes a factor of 10 when changing the mass by a factor of 10. The effects of varying the IMF slope ($\\alpha$) are similar to those of the stellar mass, but are less pronounced when decreasing $\\alpha$ than when increasing it. Making the IMF slope more shallow increases the luminosity by a factor of 4.0, whereas steepening the slope decreases the luminosity by a factor of 7.9. There is a degeneracy between ${\\rm M_{SB}}$ and $\\alpha$. A less massive starburst with a shallower IMF will produce roughly the same amount of OB stars as a more massive starburst with a steeper IMF and therefore the same amount of massive stars and therefore the same irradiating UV flux. Assuming the stars are formed according to a Salpeter IMF ($\\Psi(M) \\propto M^{-2.35}$), the star formation region should produce $10^9$~M$_{\\sun}$ ~of stars (either in one instantaneous burst or in a continuous process) in order to produce enough IR radiation. \\\\ \\\\ {\\bf Stellar dust}\\\\ For the stellar dust component, the grey approximation of the optical dust properties is not valid. A more realistic dust model, including graphite, silicates and amorphous carbon, is necessary to produce the right results. In addition, the V-band optical depth ($\\tau_2$) is found to be an important factor. Using a $\\tau_2$ of 10 (or 1) overproduces the amount of 10$\\mu$m radiation by almost a factor of 4 (or 2.5). We require an optical depth of 0.1 to get values that agree with observations.\\\\ \\\\ {\\bf Torus dust}\\\\ The influence of the large scale dust geometry on the final IR properties was far less and therefore these parameters could not be determined to great precision. Increasing the dust column density ($\\tau$) has two effects. Like for ${\\rm M_{SB}}$, the short wavelengths are enhanced compared to the long wavelengths and the IR luminosity increases. These effects are, however, much smaller: the luminosity changes by a factor of 2.0 when going from $\\tau$=1 to 10 and by a factor of 3.2 when increasing it further to 100. A second effect is that the inclination dependence increases with increasing optical depth. The difference in the edge-on and face-on values of ${\\rm L_{IR}}$ changes by a factor of 1.6 for $\\tau=1$ to a factor of 16 for $\\tau=100$. The best fit to the data was obtained with $\\tau=10$, but the other simulations were also reasonable. The closing angle did not seem to have an optimum value at all. Varying it only affects the inclination dependence of the results. The variation of ${\\rm L_{IR}}$ for a flat disk-like structure ($\\eta=0.3$) is about a factor of 6.3, compared to a factor of 2.0 for a more shell-like dust geometry with a closing angle of 0.7. Even though the inclination effects are reduced for higher values, the edge-on results still do not match the data.\\\\ Considering the values determined for the parameters investigated, it seems that, although observationally starburst galaxies appear to have a very violent nature, the star formation environment does not need to be as ``exotic'' as one might expect. The only exceptional parameter needed to explain the high IR output of starburst galaxies is the large amount of massive stars (high ${\\rm M_{SB}}$ or SFR), but there is no need for an adjusted IMF to increase the number of heavy stars. The torus dust surrounding the stars is no exception to this trend. Both the size (100 pc) and the V-band optical depth of the torus dust are moderate (10, comparable to values found for photon dominated regions). Furthermore, most of the dust has a low temperature ($\\sim 75$K at 100pc from the center), while only a very small amount of hot dust ($\\sim 400$K at 10pc) is needed.\\\\ In all simulations, including the final ones, the data were only fit well by the face-on evolutionary tracks. All the completely edge-on results were a poor fit. This effect is stronger in the short wavelength results than in the long wavelength result. This indicates that inclination dependence is mostly caused by the obscuration of the hot dust in the center by the outer dust distribution. This can have two implications. On the one hand, it could be that the inclination dependence resulting from our model is too large, and that extra parameters are needed to address this problem. A likely parameter is the clumpiness of the dust. A clumpy medium with the same average density (i.e. the same mass) would have a lower apparent optical depth then a smooth medium \\citep[e.g.][]{1984ApJ...287..228N,2005Ap&SS.295..319C}, which makes it easier for the short wavelength radiation to travel in the edge-on direction. On the other hand, if the predictions of our model are correct, the implication is that there is a significant number of starburst ULIRGs, which are currently not classified as such based on their IRAS colors. A specific shortcoming of our models is that, although the IR properties are well explained, almost all parameters move the evolutionary tracks more or less along the same line and in the direction of the evolution. The result is that not all parameters of the physical environment in a given starburst galaxy can be uniquely inferred from observations, using this model. The UV input can be inferred from the total IR output, but this does not constrain the IMF or the SFR. Similarly, the optical properties of the dust can be determined, but the geometry is hard to infer. To address these problems, more information than just the IR continuum is needed and therefore we intend to extend the current model. First of all, the IR part of the code will be modified to include specific spectral characteristics (e.g. PAHs and high ionization lines). Also the molecular environment that surrounds the current dust region will be added to the model. The chemistry of such a region will give a better handle on the radiation field, as well as the densities and temperatures of the gas \\citep[e.g.][]{1999RvMP...71..173H, 2005A&A...436..397M, 2002A&A...381..783A, 2004A&A...419..897U, 2004ApJS..152...63G,2005ApJ...629..767O, 2006ApJ...640L.135G,2006A&A...000..000B} . Also other wavelengths will be studied, since optical and UV data will give more information about the star forming region, whereas (sub)millimeter and radio observations will reveal more about the outer dust regions and the molecular environment." }, "0607/astro-ph0607109_arXiv.txt": { "abstract": "{A survey of progress in recent years suggests we are moving towards a quantitative understanding of the whole cosmic ray spectrum, and that many bumps due to different components can hide beneath a smooth total flux. The knee is much better understood: the KASCADE observations indicate that the spectrum does have a rather sharp rigidity cut-off, while theoretical developments (strong magnetic field generation) indicate that supernova remnants (SNR) of different types should indeed accelerate particles to practically this same maximum rigidity. X-ray and TeV observations of shell-type supernova remnants produce evidence in favour of cosmic-ray origin in diffusive shock acceleration at the outer boundaries of SNR. There is some still disputed evidence that the transition to extragalactic cosmic rays has already occurred just above $10^{17}$ eV, in which case the shape of the whole spectrum may possibly be well described by adding a single power-law source spectrum from many extragalactic sources (that are capable of photodistintegrating all nuclei) to the flux from SNRs. At the very highest energy, the experiments using fluorescence light to calibrate energy do not yet show any conflict with an expected GZK ``termination''. (And, in ``version 2'',) Sources related to GRBs do not appear likely to play an important role. } ", "introduction": "} Because cosmic rays span such a huge range of energy, it is natural to start from a very deceptive broad view of the cosmic ray spectrum, such as that shown in figure 1, due to Gaisser (\\cite{gaisserfig}), which shows the flux reaching the Earth, in the form of the energy carried by particles per unit interval of $ln(E)$, or $E^2 J(E)$, where $J(E)$ is the number of particles arriving per unit interval of time, area, solid angle and kinetic energy, E. \\begin{figure} \\centering \\vspace{340pt} \\special{psfile=AMHillas_3_fig1.ps hscale=57 vscale=55 hoffset=-50 voffset=-30} \\caption {Many measurements of the cosmic ray flux over a wide energy range, assembled by Gaisser \\label{gaisfig}} \\end{figure} At the lowest energies, the fluxes of different nuclei can be measured, protons being the most numerous, and other common nuclei having practically the same shape of spectrum as a function of rigidity (momentum/charge $\\propto$ energy/charge at these relativistic energies). To identify the particles clearly, they have to be detected before they are broken up in the atmosphere, in detectors carried by balloons or satellites, and the flux is too low for this above about $10^5$ GeV ($10^{14}$ eV): beyond here the total flux of all particle types can be recorded by air shower experiments. The well-known power-law spectrum, $J(E) \\propto E^{-2.7}$ holds to a good approximation before the ``knee'', the downward bend near $10^{15.5}$ eV, the fall-off below 10 GeV being a very local effect within the solar system. For 3 decades of energy above the knee the flux continues to fall somewhat more steeply, to the ``ankle'', where the rate of fall briefly becomes less steep again, until statistics and possibly flux peter out near $10^{11}$ Gev ($10^{20}$ eV). At energies of several GeV there is good evidence from gamma rays produced in nuclear collisions (e.g. Hunter et al. \\cite{huntergam}) that the cosmic rays originate in the Galaxy, and diffuse out; and the belief that the major source is acceleration at the outer shock boundaries of expanding supernova remnants (SNR) has strengthened recently in several ways, outlined below. It now seems likely that this bland shape masks a superposition of bumps and variations which each tell their own story, though few of them can yet be disentangled clearly, so this field of diagnosing the components is still very active. Recent experimental work at Karlsruhe (discussed in section \\ref{galcompsec}) has made it seem very probable that the individual nuclear components each fall off rather steeply at a magnetic rigidity near $3\\times 10^{15}$ V (i.e. at energies $3\\times 10^{15}$ eV for protons, $6\\times 10^{15}$ eV for helium nuclei, extending to $8\\times 10^{16}$ eV for iron, the heaviest common nucleus). Assuming this to be right, even though the point of turn-down for iron is just beyond the range of this experiment, the main Galactic component is made up of elemental components each extending to an energy near $Z\\times 3\\times 10^{15}$ eV (where $Z$e is nuclear charge), beyond where the fluxes turn down much more sharply than does the total flux that is plotted in figure 1. Despite the separate bends, the total flux looks deceptively smooth after steepening a little at the knee, as shown in figure \\ref{gaisfig}, at least as far as $10^{17}$ eV. Beyond $10^{17}$ eV questions arise. The more extended gradual fall-off between the knee and the ankle has long been puzzling. Does some Galactic source (magnetars?) extend the spectrum of local particles well beyond $10^{17}$ eV (presumably highly-charged particles to allow them to be disoriented by Galactic magnetic fields)? A widespread view had been that some such additional component partly trapped within our galaxy eventually fell below the level of cosmic rays circulating throughout the universe, and originating in rarer far more energetic sources. The ``ankle'' might then mark the point where such extragalactic cosmic rays became dominant, but this is not necessarily so, as will appear from the discussion in section \\ref{extragalsec}. One much older view of the knee-to-ankle region had been that the Galactic sources might accelerate particles to much higher energies than $10^{16}$ eV, and the extended slightly steeper slope beyond the knee marked an increasingly rapid escape of particles from Galactic magnetic fields at higher energy, a view the present author mistakenly used when discussing anisotropies in a review 22 years ago (\\cite{araa}), when it appeared that there was an increasing anisotropy reflecting such a decreasing residence time. However, with much greater counting statistics, we now see no clear anisotropy apart from a small one in the region $10^{14}$ to $10^{15}$ eV. Unless there is indeed a high-Z flux generated by Galactic magnetars, it now seems that the extragalactic component becomes very important at a much lower energy than previously thought. If, then, there are no low-charge galactic particles above $10^{17}$ eV, the failure to find convincing anisotropies would be explained. These topics, and the particles of extreme energy, are discussed below. Cosmic-ray electrons will be mentioned only briefly. Figure \\ref{gaisfig} shows that at a given energy they are much less numerous than protons --- 1--2\\% around a GeV and even less at higher energies --- though their strong synchrotron radiation makes their presence in distant regions much easier to detect than that of protons and nuclei. The electrons may originate in SNR as we believe do the hadrons, or in plerions (e.g. the Crab Nebula), but if termination shocks of ultrarelativistic winds produce the acceleration in the latter, they probably accelerate an electron-positron medium, and the low relative abundance of cosmic-ray positrons indicates that plerions do not form a major source. ", "conclusions": "Cosmic ray physics is perhaps becoming less exciting for seekers of the exotic. There is as yet no necessity for new physics at the highest energies; and we may be approaching the situation where the mysteriously bland spectrum between the knee and the ankle is resolved into the sum of perhaps only two major kinds of source. If this interpretation is correct (notably the proton source model of Berezinsky et al.), the transition from Galactic to extragalactic cosmic rays has occurred at a much lower energy than was usually believed, and has left virtually no obvious sign in the flux level at the join point. Nevertheless, a rapid change in shower characteristics should occur here, and measurements of double-peaked distributions may be possible. The detection of a class of air showers of very uniform structure near $10^{17}$ eV, attributale to the most energetic Galactic cosmic rays, and rapidly diminishing in proportion as energy rises, would help to establish clearly whether proton-only acceleration occurs at the highest energies. Such an interpretation, and the demonstration of sharp sub-knees in the Galactic spectrum by the KASCADE experiment, have presented a lesson that the simplicity of a smooth spectrum, close to a power law, can be very deceptive. Cosmic-ray physicists should resist the temptation to read much into the position of a ``join point'' between two straight lines drawn through flux data points. The absence of obvious concavity in cosmic-ray proton spectra before the knee, despite its prediction in diffusive shock acceleration, may be another instance of many sub-spectra adding to give the appearance of a close approximation to a power law in the total. A much better understanding of the knee as a consequence of SNR development brings together theoretical and observational work very fruitfully. The central part played in astronomy by detailed images (as well as spectra) is exemplified by wonderful CHANDRA X-ray images of SNRs, and TeV detectors are at the threshold of this capability.." }, "0607/astro-ph0607423_arXiv.txt": { "abstract": "A linear modulation of the primordial perturbations is proposed as an explanation for the observed asymmetry between the northern and southern hemispheres of the Wilkinson Microwave Anisotropy Probe ({\\em WMAP\\/}) data. A cut sky, reduced resolution third year ``Internal Linear Combination'' ({\\em ILC\\/}) map was used to estimate the modulation parameters. A foreground template and a modulated plus unmodulated monopole and dipole were projected out of the likelihood. The effective chi squared was reduced by nine for three extra parameters. The mean galactic colatitude and longitude, of the modulation, with 68\\%, 95\\% and 99.7\\% confidence intervals were $56^{+17 +36 +65}_{-17 -35 -51}$ and $63^{+28 +59 +105}_{-26 -58 -213}$. The mean percentage change of the variance, across the pole's of the modulation, was $62^{+18 +35 +57}_{-18 -35 -47} $. Implications of these results and possible generating mechanisms are discussed. ", "introduction": "A fundamental assumption of cosmology is that the Universe is isotropic. This was confirmed, for the mean temperature of the cosmic microwave background ({\\em CMB\\/}), by the {\\em FIRAS\\/} experiment on the {\\em COBE\\/} satellite \\citep{wright92,bennett96}. However, the higher precision, measurements from the Wilkinson Microwave Anisotropy Probe ({\\em WMAP}) satellite \\citep{bennett03,hinshaw06,jarosik06,page06,spergel06}, have an anomalously asymmetric distribution, in the temperature fluctuation statistics, between the northern and southern hemispheres of the sky \\citep{erihanbangorlil03,hanbangor04,% vielva03,park03,cophutsta03,hansen04,larwan04,crumarvie04,lanmag04,hansen04a,bernui05,bernui06}. On scales greater than about $5^\\circ$, the variance of the {\\em CMB\\/} temperature fluctuations is anomalously higher in the southern hemisphere, in both galactic and ecliptic coordinates, compared to the northern hemisphere \\citep{erihanbangorlil03,hanbangor04}. This asymmetry also appears in higher order statistics \\citep{vielva03,park03,cophutsta03,hansen04,larwan04,crumarvie04,lanmag04,hansen04a,bernui05,bernui06}. In a spherical harmonic representation, scales ranging from $\\ell=2$ to $\\ell=40$ were found to be asymmetric. When optimized over direction, only 0.3\\% of isotropic simulations were found to produce higher levels of asymmetry \\citep{hanbangor04}. The result is not sensitive to the frequency band of the {\\em CMB\\/} \\citep{hanbangor04} and a similar pattern (at lower significance) is seen in {\\em COBE\\/} \\citep{hanbangor04}. This argues against a foreground or systematics explanation. Although a simple single field inflation model would give isotropically distributed perturbations, this is not necessarily the case in multi-field models \\citep{linmuk05}. Thus, if it can be shown that the {\\em CMB\\/} fluctuations are not isotropic, it may be an indication that inflation was a multi-field process. The layout of the paper is as follows: In Sec.~\\ref{sec:modulation}, a linearly modulated primordial power spectrum is proposed as the source of the observed isotropy breaking. Then, in Sec.~\\ref{sec:likana}, a method of evaluating the linear modulation parameters is outlined. The constraints are given in Sec.~\\ref{sec:results} and their implications and relation to other results are discussed in Sec.~\\ref{sec:discussion}. ", "conclusions": "\\label{sec:discussion} In this article the modulation model investigated by \\citet{spergel06} has been extended by including a marginalization over the unmodulated monopole and dipole. This additional feature is required if the apparent isotropy breaking had a primordial origin. Including this marginalization improved the $\\Delta \\chi_{\\rm eff}^2$ value from -3 to -9. As seen from the confidence intervals in Table~1 and \\fig{\\ref{fig:pdfs}}, the marginalized posterior probability of $\\Delta$ has its maximum more than three sigma away from the unmodulated case ($\\Delta=0$). The modulated model is also preferred by the Akaike Information Criteria (AIC) \\citep{akaike74,magsor06}. It is not preferred by the Bayesian Information Criteria (BIC) \\citep{schwarz78,magsor06}. However, the BIC is an approximation of the Bayesian evidence and assumes a prior for the parameters which is equivalent to one observation \\citep{raftery95}. The Bayesian evidence will be inversely proportional to the volume of the prior probability distribution of the modulation parameters. It may be hard to produce a modulation larger than one without effecting the observed dipole. A reevaluation of the Bayesian evidence is needed to see how it depends on the assumed prior. This could be done using a nested sampling algorithm \\citep{mukparlid05} which, unlike the BIC, does not require a Gaussian approximation to be made for the posterior distribution. \\citet{spergel06} also evaluated whether there was an additional quadrupolar component to the modulation. This component could potentially be useful in explaining the alignment and planarity of the quadrupole ($\\ell=2$) and octopole ($\\ell=3$) seen in the WMAP temperature data \\citep{oliveira03,schwarz04}. The normal direction of the plane of alignment is $(30^\\circ,-100^\\circ)$. Also, when the coordinate system is rotated in the direction of the normal of the $\\ell=2,3$ planarity there is anomalous power in the $m=3$ component of the $\\ell=5$ multipole \\citep{lanmag05}. \\citet{spergel06} found that including a dipolar and quadrupolar component to the modulation improved $\\Delta \\chi_{\\rm eff}^2$ by only 8 for a total of 8 extra parameters. Higher order terms in a spatial modulation could be implemented as terms quadratic in the spatial coordinates. Whether these additional terms will become significant when an unmodulated monopole and dipole are marginalized over will be part of a future investigation. The effect of marginalization over foregrounds was checked and found not to play a big role. Similar improvements are obtained when the foreground corrected V band is used instead of marginalization. Also, the results are not sensitive to the exact method of degrading and applying the mask. A Kp2 extended mask \\citep{eriksen06} did not make a significant difference. Including additional $C_{\\ell}$, with $\\ell>10$, as parameters to be estimated (rather than set to their unmodulated {\\em ML\\/} values), also does not significantly effect the results. It is interesting to compare the estimated modulation found in this article to that of \\citet{hanbangor04}. The 10 most effective axes of symmetry breaking, for a range of scales, are plotted in their \\fig{24}. A similar area, to the two dimensional confidence intervals in \\fig{\\ref{fig:results}}, is covered. Also, their \\fig{19} compares the power spectra in different hemispheres. The range of values is consistent with the confidence intervals for $\\Delta$ in Table~1 and \\fig{\\ref{fig:pdfs}}. \\citet{pruuzaberbru04} tested for a dipolar modulation. However, the largest scale they looked at was $\\ell=20$ to 100 binned. They did not get significant results in that range. As the observed modulation only occurs for $\\ell\\lesssim 40$ \\citep{hanbangor04}, the $\\ell=20$ to $100$ range would not be expected to show significant modulation when binned. \\citet{freeman05} propose that the modulation of $\\ell=2$ to $\\ell=7$ may be sensitive to any residual unmodulated dipole component. This is not a concern for the approach taken in this article as an unmodulated dipole is projected out of the likelihood, see \\eq{Cmarg}. Searches for lack of isotropy using a method based on a bipolar expansion of the two point correlation function do not detect the north south asymmetry in the $\\ell=2$ to $\\ell=40$ range \\citep{hajsou06,picon05}. The linear modulation model could be used to understand why the bipolar estimator is insensitive to this type of isotropy breaking. A small scale cut off in the modulation implies that a linear modulation of the primordial power spectrum would only apply to wave numbers larger than about $4\\times 10^{-3}h$~Mpc$^{-1}$. It would be interesting to evaluate whether this modulation would be detectable in future large scale galaxy surveys. However, at a redshift of one the change in the variance at opposite poles would only be about four percent, due to the smaller comoving distance. A number of attempts have been made to explain the asymmetry in terms of local nonlinear inhomogeneities \\citep{% moffat05,tomita05a,tomita05,inosil06}. It would be interesting to see if the polarization maps of the {\\em CMB\\/} \\citep{page06} could be used to distinguish local effects from a modulation of the primordial perturbations. Primordial magnetic fields \\citep{durkahyat98,chen04,naselsky04}, global topology \\citep{oliveira03,kunz06}, and anisotropic expansion \\citep{berbubkep03,bunberkep05,gumconpel06} can also lead to isotropy breaking. However, in these cases the modulating function is of higher order than dipolar and so these mechanisms are better suited for explaining the alignment between $\\ell=2$ and $\\ell=3$ and the high $(\\ell,m)=(5,3)$ mode \\citep{gorhuhut05}. An additive template based on a Bianchi $VII_h$ model has been shown to provide a good fit to the asymmetry \\citep{jaffe05}. However, the model is only empirical as it would require a very open Universe which is in conflict with many other observations. It is harder for additive templates to explain the alignment between $\\ell=2$ and $\\ell=3$ as this requires a chance cancellation between an underlying Gaussian field and a deterministic template \\citep{gorhuhut05, lanmag05a}. As seen in \\fig{\\ref{fig:results}}, the maximum likelihood direction of modulation was found to be about $44^\\circ$ from the ecliptic north pole. Only about 9\\% of the time would two randomly chosen directions be as close, or closer, together. This may be an indication that the modulation is caused by some systematic effect or foreground. However, as can be seen from \\fig{\\ref{fig:results}}, the confidence intervals, for the direction of modulation, cover just under half the northern hemisphere. Therefore, the actual direction, of modulation may be significantly further away from the ecliptic north pole. Standard single field inflation would produce isotropic perturbations. However, multi-field models, such as in the curvaton scenario \\citep{% lytwan01, mollerach90,linmuk96,enqslo01,mortak01}, can produce, what to a particular observer appear to be, non-isotropic perturbations \\citep{linmuk05}. The curvaton mechanism produces a web like structure in which relatively stable domains are separated by walls of large nonlinear fluctuations. If the mass of the curvaton field is sufficiently small, our observable Universe could be enclosed within a stable domain. If we happen to live near one of the walls, of a domain, then the amplitude of the perturbations will be larger on the side of the observed Universe closer to the wall \\citep{linmuk05}. However, if our observed Universe was far enough away from the web walls, the very large scale fluctuations would be linear and so isotropy would be unlikely to appear to be broken \\citep{lyth06}. As the non-isotropic nature only extends to about $\\ell=40$ \\citep{hanbangor04}, it would be necessary for the inflaton perturbations to dominate over the curvaton ones for wave numbers larger than about $4\\times 10^{-3}h$~Mpc$^{-1}$. The curvaton produces curvature perturbations proportional to $V^{1/2}$ \\citep{lytwan01}, where $V$ denotes the inflaton potential. While the inflaton produces curvature perturbations proportional to $V^{3/2}/V'$, where $V'$ denotes the slope of the potential. So if there is a sudden drop in $V$ and $V'$, it is possible for the non-isotropic curvaton perturbations to dominate for wave numbers smaller than $4\\times 10^{-3}h$~Mpc$^{-1}$ and inflaton perturbations to dominate for larger wave numbers. There are oscillations in the {\\em WMAP\\/} power spectrum, at around $\\ell=40$, which may be caused by a change of slope in the inflaton potential \\citep{covi06}. Whether all these elements can be put together to make a working curvaton model, that fits the data as well as a linear modulation, is still being investigated. The results presented here provide a parameterization for the observed asymmetry between different hemispheres of the {\\em WMAP\\/} data. Having a specific model for the primordial fluctuations will make it easier to develop new tests for this asymmetry and help determine if it is a genuine window into new physics at the largest observable scales." }, "0607/astro-ph0607437_arXiv.txt": { "abstract": "We compute the pair annihilation cross section of light (spin-0) dark matter particles into two photons and discuss the detectability of the monochromatic line associated with these annihilations. ", "introduction": "The precise determination by INTEGRAL/SPI \\cite{Jean:2003ci} of the characteristics of the 511 keV line emitted in our galaxy \\cite{Dixon:1997nn,Sreekumar:1997yg,Churazov05,Milne:2001zs,Dermer:1997cq,Pohl:1998bs,Dermer:2001wc,Milne:2001dt} has shed new light on the physics of the inner part of the Milky Way. This line has now been identified with a high level of confidence as originating from electron-positron annihilations. Although this recent detection probes unambiguously the existence of anti-matter inside our galaxy, its origin remains unknown. The observation of a relatively high fraction of low energy positrons in the bulge and a low fraction in the disk certainly constitutes the most puzzling aspect of this emission. Most of the astrophysical sources that have been proposed in the literature (e.g. Wolf-Rayet stars, Hypernovae, cosmic rays, pulsars, black holes) are associated with a low value of the bulge-to-disk (B/D) ratio or cannot explain why the 511 keV radiation seems to follow the stellar morphology of the galactic bulge. The remaining plausible sources are old galactic populations, such as Low Mass X-ray Binaries (LMXB) and Type 1a Supernovae (SN1A) \\cite{Knodlseder05}. However, to explain the observed flux and positron distribution, they both rely on strong hypothesis. LMXB require that the positrons emitted in the disk escape into the bulge while SN1A need a positron escape fraction and an explosion rate that are large enough to maintain a steady flux. At least eight point sources could explain the diffuse emission as observed by SPI \\cite{Knodlseder05}. However Ref.~\\cite{Weidenspointner:2006nu} did not find any evidence for significant emission from point sources in the galactic centre as yet. Another candidate could be light dark matter (LDM) particles \\cite{bens,bf,511} annihilating into electrons--positrons, neutrinos and photons. The positrons thus emitted lose their energy by ionization and eventually form para-positronium atoms with the thermal electrons present in the bulge of the galaxy \\cite{511}. The 511 keV line emission is expected to be strongly correlated with the dark matter energy density distribution. The latter is maximal in the inner part of the galaxy so the positrons should mostly be produced in the centre of the Milky Way and should naturally stop on the electrons present in the bulge. Depending on the cuspyness of the profile, this would explain why the emission is well described by a sphere of only $\\sim$ $8-10^\\circ$ of diameter. If dark matter is light enough, the amount of low energy gamma rays produced by the dark matter (DM) annihilations remains compatible with observations. Below the muon mass threshold ($m_{dm} \\leq 100$ MeV), the gamma ray production channels are: e.g. the DM pair annihilation into electron-positron plus a photon \\cite{bens,beacom,bu}, positronium formation, inflight $e^+ e^-$ annihilations, initial and final state radiation associated with the electron-positron annihilations (although they have not been included as yet in previous studies). All these processes generate a continuum. In addition, lines are produced at an energy corresponding to either the dark matter mass or at 511 keV. The former is the subject of the present paper. If the corresponding flux is large enough, this could be an unique tracer to answer the question of the low energy positrons. After summarizing recent progress on this topic and their implication for the dark matter characteristics, we will estimate the annihilation cross section of light dark matter particles into two photons and discuss the observability of the line at $E_\\gamma=m_{dm}$. We base our analysis on the model proposed in Ref.~\\cite{bens,bf}, which has been studied in detail in Ref.~\\cite{boehm,boehmS,ascasibar,boehmascasibar}. ", "conclusions": "In this paper, we computed the pair annihilation of LDM particles into two photons $\\sigma_{\\gamma \\gamma} v_r$ and determined the flux $\\phi_{\\gamma \\gamma}$ associated with the monochromatic line $E=m_{dm}$. To obtain a conservative estimate, we considered only $e-F_e-\\rm{dm}$ interactions and ignored all other possible interactions. With this simplistic assumption we could relate our estimate of $\\phi_{\\gamma \\gamma}$ to the 511 keV flux that has been measured by INTEGRAL/SPI. We made the reasonable assumption that the particle $F_e$ was much heavier than the dark matter and the electron. We also assumed that the couplings were small enough ($(c_l^2+c_r^2) m_e \\ll 2 c_l c_r m_{F_e}$) so that the contribution associated with the electron mass in the cross section could be neglected. We found that $\\phi_{\\gamma \\gamma}$ was ranging from $10^{-6}$ to $10^{-10} \\ {\\rm ph \\; cm^{-2} \\; s^{-1}}$ for dark matter masses from $m_e$ to 100 MeV. These values are well below the present SPI sensitivity. Next generation instruments such as AGILE/(super AGILE) or GLAST, which in principle could be more promising, will probably be limited by the energy range that they are able to investigate. Future instruments might nevertheless be able to see this line if their energy resolution and sensitivity are improved by a large factor with respect to SPI present characteristics. Maybe a better chance to detect this line is to do observations at a high latitude and a longitude slightly off the galactic centre. In this case, indeed, the background should drop significantly (the density of dark clouds has been measured recently \\cite{grenier}) but the line flux may decrease by a smaller factor. In dwarf galaxies, where the dark matter content dominates over baryons, the gamma ray background is also expected to be quite suppressed. The line $E=m_{dm}$ might be easier to detect. Among the closest dwarfs to us, Sagittarius Dwarf Galaxy (located at a distance of 24 kpc from us and with a size of about $10^8 \\ M_{\\odot}$ \\cite{Ibata:1995fz}), is particularly interesting. The amount of intrastellar gas is very low and the dwarf contains a large amount of popII stars. The problem is that it is somehow hidden by the galactic centre although it is a bit off. In principle, dwarf spheroidals are a powerful tool for testing the LDM hypothesis. E.g. the detection of a bright 511 keV line within INTEGRAL's sensitivity would provide a strong confirmation of this scenario \\cite{hooper-dwarf}. However, its detectability relies on the hypothesis that there are enough electrons to thermalize and to stop the positrons. No gas has ever been detected in any of the local group dwarf galaxies so it is hard to make reliable estimates. Also one needs to know the spatial distribution of the gas to make accurate predictions and determine whether the 511 keV emission will be extended or not. SDG is being disrupted by the tidal forces of our galaxy so the approximation of a spherical DM halo profile probably leads to incorrect predictions. At last, depending on the fraction of gas, inflight $e^+ e^-$ annihilations may happen before the positrons have time to thermalize. In this case, there should be a broad line at an energy $m_e < E < m_{dm}$\\footnote{Note that, for the Milky Way, the detection of both the 511 keV line and the continuum indicates that the interstellar medium is in a warm partially ionized phase. However the temperature of the gas in the dwarf may be larger. Even if the positrons can annihilate after being thermalized, this would imply a suppression of the positronium formation rate and would favour inflight $e^+ e^-$ annihilations.} instead of a 511 keV line. The study of the $\\gamma \\gamma$ channel has therefore two advantages compared to the 511 keV emission: it does not rely on the estimate of the electron number density nor the gas fraction inside the dwarf. It is also independent of the gas spatial distribution and can probe directly the DM halo profile. It has one major drawback: it is a higher order process and it is therefore suppressed compared to the the $e^+ e^-$ production. Since Ref.~\\cite{cordier} looked for the 511 keV line in SDG and did not find it, the chance to detect the monochromatic line is probably very small although Ref.~\\cite{cordier} was unable to probe the upper limit of the predicted range and the number density of electrons inside the dwarfs may have been overestimated." }, "0607/astro-ph0607601_arXiv.txt": { "abstract": "We present a multiwavelength study of GRB~060108 - the 100th Gamma Ray Burst discovered by \\textit{Swift}. The X-ray flux and light curve (3-segments plus a flare) detected with the XRT are typical of \\textit{Swift} long bursts. We report the discovery of a faint optical afterglow detected in deep $BVRi'$ band imaging obtained with the Faulkes Telescope North (FTN) beginning 2.75 minutes after the burst. The afterglow is below the detection limit of the UVOT within 100s of the burst, while is evident in $K$-band images taken with the United Kingdom Infrared Telescope (UKIRT) 45 minutes after the burst. The optical light curve is sparsely sampled. Observations taken in the R and i$'$ bands can either be fit with a single power law decay in flux, F(t)$ \\propto~t^{-\\alpha}$ where $\\alpha=0.43\\pm0.08$, or a 2-segment light curve with an initial steep decay $\\alpha_1$$<$0.88$\\pm$0.2, flattening to a slope $\\alpha_2$$\\sim$0.31$\\pm$0.12. A marginal evidence for rebrightening is seen in the i$'$ band. Deep $R$-band imaging obtained $\\sim 12$ days post burst with the VLT reveals a faint, extended object ($R \\sim 23.5$ mag) at the location of the afterglow. Although the brightness is compatible with the extrapolation of the slow decay with index $\\alpha_2$, significant flux is likely due to a host galaxy. This implies that the optical light curve had a break before 12 days, akin to what observed in the X-rays. We derive the maximum photometric redshift $z<3.2$ for GRB~060108. We find that the Spectral Energy Distribution at 1000~s after the burst, from the optical to the X-ray range, is best fit by a simple power law, F$_{\\nu}\\,\\propto\\,\\nu^{-\\beta}$, with $\\beta_{OX}\\,=\\,0.54$ and a small amount of extinction. The optical to X-ray spectral index ($\\beta_{OX}$) confirm GRB~060108 to be one of the optically darkest bursts detected. Our observations rule out a high redshift as the reason for the optical faintness of GRB~060108. We conclude that a more likely explanation is a combination of an intrinsic optical faintness of the burst, an hard optical to X-ray spectrum and a moderate amount of extinction in the host galaxy. ", "introduction": "Gamma ray bursts (GRBs) are brief, intense and totally unpredictable flashes of gamma rays on the sky that are thought to be produced during the core collapse of massive stars (long-duration bursts) or the merger of two compact objects such as two neutron stars or a neutron star and stellar-mass black hole. Until the recent launch of the {\\it Swift} satellite in November 2004, it was notoriously difficult to observe GRBs at other wavelengths within seconds or minutes after the burst. Nevertheless, the successful identification with BeppoSAX of bright, long-lived X-ray afterglow emission for long bursts \\citep{cos97} and that of the corresponding optical and infrared counterparts \\citep{vp97}, established GRBs as cosmological, and therefore the most instantaneously luminous objects in the Universe. Similar breakthroughs for short bursts have recently occurred, showing them also to be extragalactic, but less luminous and less distant than long bursts \\citep{geh05,vil05,fox05,cov06,bart05}. With the availability of \\textit{Swift}'s promptly-disseminated arcsec localizations and the on-board rapid-slew X-ray and ultraviolet/optical telescopes (XRT, UVOT; \\citealt{geh04}), multi-wavelength monitoring of GRBs from the earliest possible times is now being performed for a significant number of bursts. Additionally, large aperture ground-based robotic telescopes such as the 2-m Liverpool \\citep{steele04} and Faulkes telescopes respond rapidly to GRB alerts and begin automatically imaging the target field within minutes of receipt of an alert, providing early deep upper limits or multi-colour follow-ups of optical counterparts as faint as R\\,$\\sim$\\,18\\,-\\,22\\,mag (e.g. Guidorzi et al. 2005a, 2006b; Monfardini et al. 2006b). Despite increasingly rapid responses that provide sensitive limits within minutes of the burst, the absence of long-wavelength emission afterglows for a significant number of GRBs (so-called ``dark bursts'') remains a key unsolved problem. In the pre-\\textit{Swift} era, as many as 50\\% of BeppoSax bursts were lacking an optical detection \\citep{dp03,lcg02}. The discovery rate of optical afterglows was even higher for HETE2 than for SAX (\\citealt{lamb04}), and it was expected to increase significantly in the \\textit{Swift} era of rapid followup. Instead, a substantial fraction of \\textit{Swift} bursts remain undetected in the optical band \\citep{rom05}. Possible scenarios to explain the ``observed'' optical darkness of these bursts, apart from fast-fading transients lacking sufficiently deep, early-time observations \\citep{groot98}, include intrinsically-faint optical afterglows \\citep{fyn01, lcg02, dp03, ber05}, highly obscured afterglows whose optical light is absorbed by the circumburst or interstellar material \\citep{lcg02,dp06}, high redshifts \\citep{fru99,lr00,bl02,ber02,tag05}, and radiative suppression in a sub-class of bursts with unusually high $\\gamma$-ray efficiency producing intrinsically low X-ray and optical fluxes \\citep{rom05,ped06}. Here we present a multi-wavelength X-ray, optical and infrared study of the optically-faint GRB~060108. The Burst Alert Telescope (BAT, \\citealt{bart05}) was triggered by this GRB - {\\it Swift}'s 100th burst - at 14:39:11.76 UT on January $\\rm 8{^t}{^h} $ 2006. The $\\gamma$-ray light curve has a single peaked structure with a FRED time profile \\citep{oat06}, a duration T$_{90}$\\,=\\,14.4\\,$\\pm$\\,1.6\\,s (in the 15\\,$-$\\,350\\,keV band), and a 15\\,$-$\\,150\\,keV fluence $S_\\gamma = (3.7 \\pm 0.4) \\times 10^{-7} {\\rm erg\\,cm}^{-2}$. Here and in the following errors are at 90\\% confidence level, unless specified otherwise. The \\textit{Swift} X-ray and ultraviolet/optical telescopes (XRT/UVOT) began observing at 91\\,s and 76\\,s after the BAT trigger respectively, followed shortly after by the Faulkes Telescope North (FTN) at 2.75 minutes post trigger. Infrared observations with the United Kingdom Infrared Telescope (UKIRT) were acquired during the optical imaging period beginning at 45 min postburst. The Very Large Telescope (VLT) was used to obtain further near infrared imaging at 16.1 hours and 2.7 days. Following an initial estimate of the position of the X-ray afterglow \\citep{pag06}, a revised location was derived \\citep{bb06}; the location of the optical \\citep{mon06a} and infrared counterparts \\citep{dav06,lev06} was found to be consistent with this revised XRT position. Deep $R$-band imaging and spectroscopy were subsequently performed with the VLT at 12.7 days and $\\sim$\\,21 days respectively. ", "conclusions": "We presented and discussed the gamma-ray, X-ray and UV/optical/infrared observations of the \\textit{Swift} GRB~060108, performed with the instruments on-board the spacecraft as well as the with the ground-based Faulkes Telescope North, United Kingdom Infrared Telescope and Very Large Telescope. GRB~060108 has a moderately faint X-ray afterglow, but very low optical emission, making it one of the ``darkest'' GRBs ever observed. We suggest that faintness of the optical emission may be due to an intrinsic weakness of the burst and a hard optical to X-ray spectrum, accompanied by some degree of extinction which occurred in the GRB surroundings. The X-ray light curve shows the typical template discovered by {\\it Swift}, characterized by a rapid decay in the first 300~s (commonly interpreted as the tail of the prompt emission on the basis of the combined temporal/spectral properties) followed by a flat decay slope (typical of ``refreshed'' afterglows). An X-ray flare is observed at $\\sim300$~s, although the statistic is low. The flat decay phase lasts for about $\\sim$10~ks, after that the light curve breaks into a steeper segment with a power-law decay index of $1.0$, typical of a standard afterglow phase.\\\\ A comparison with the optical light curve is quite interesting. While the optical flux after 800 seconds has a slope of $\\alpha\\,=\\,0.4$, similar to that of the X-ray in the same interval, the decay index before this time is likely to be steeper. This behavior may be explained if we interpret the fast decay as the reverse shock emission, while the later, flatter emission, has the same origin of the X-ray one. A possible rebrightening is visible only in the infrared, although a monotonic decay is not ruled out. A similar feature has been seen in the case of GRB~060206 (\\citealt{monf06c}). If real, it might indicate a second reverse shock emission, initiated by the 300~s flare. Alternative models able to explain an optical rebrightening include an increase in the density of the medium where the forward shock is produced (\\citealt{laz02}), or energy injection by late shells (\\citealt{bjo04}, \\citealt{joh06}). This last scenario is less likely, because of our finding that the late energy injection is taking place at a steady rate. The coincidence of the rebrightening with the X-ray flare, however, favours the hypothesis of reverse shock. Another intriguing feature of the optical light curve is the absence of a break in correspondence to the X-ray one at $\\sim$10\\,ks, which may be either intrinsic or due to a significant contribution to the optical flux at late times by the host galaxy. However, the poor sampling of the late optical light curve does not allow us to better constrain this behaviour. The analysis of the optical spectrum, obtained from data gathered 45 minutes after the burst, has allowed us to determine an upper limit of $z < 3.2$ at 90\\% confidence level, by using a $\\chi^2$ minimization of the observed spectral energy distribution. The optical afterglow is below the detection limit of the UVOT within 100\\,s of the burst. This event has shown how observations taken promptly and deeply enough may reveal the interesting behaviour of the early optical emission. In the \\textit{Swift} era, further similar observations are a reality due to the prompt response of both the spacecraft and ground robotic telescopes, as well as due to the possibility to perform deep observations with large telescopes at early times after the trigger." }, "0607/astro-ph0607092_arXiv.txt": { "abstract": "The satellite systems of M31 and the Galaxy are compared. It is noted that all five of the suspected stripped dSph cores of M31 companions are located within a projected distance of 40 kpc of from the nucleus of this galaxy, whereas the normal dSph companions to this object have distances $>$ 40 kpc from the center of M31. All companions within 200 kpc $<$ D(M31) $<$ 600 kpc are late-type objects. In one respect The companions to the Galaxy appear to exhibit different systematics with the irregular LMC and SMC being located at small $R_{gc}$. It is speculated that this difference might be accounted for by assuming that the Magellanic Clouds are interlopers that were originally formed in the outer reaches of the Local Group. The radial distribution of the total sample of 40 companions of M31 and the Galaxy, which is shown in Figure 1, may hint at the possibility that these objects contain distinct populations of core (R $<$ 25 kpc) and halo (R $>$ 25 kpc) satellites.) ", "introduction": "In the present investigation the data on the companions to M31 and the Galaxy are extended by including a number of recently discovered satellites. Furthermore, following Koch \\& Grebel (2006), some compact objects that are widely believed to be the stripped cores of now defunct dwarf spheroidal galaxies, have been added to the list of satellites to M31 and the Galaxy. This enlarged database is then used to investigate some of the systematics of the M31 and Milky Way satellite systems. In particular we seek to answer three questions: (1) How does the morphological type of a satellite depend on its distance from the center of its parent galaxy? (2) Do inner and outer dwarf satellites belong to separate core and halo populations, and (3) were the Magellanic Clouds formed as satellites of the Galaxy, or might they have been captured from the outer reaches of the Local Group? ", "conclusions": "Inspection of the data in Table 1 and Table 2 shows that the Galactic satellite system differs from that of M31 in three important ways: (1) All inner satellites of M31 are early-type objects. On the other hand the LMC and the SMC are presently situated at small Galactocentric distances. This perhaps encourages the speculation (Byrd et al. 1994) that the Magellanic Clouds might be interlopers that were initially formed in a more remote region of the Local Group. [A recent paper on the orbit of the LMC (Pedreros et al. 2006) {\\it assumes} that the LMC is gravitationally bound to, and in an elliptical orbit around, the Galaxy.] (2) All of the suspected stripped cores in M31 occur at small (R $<$ 40 kpc) distances from the nucleus of M31. However, among companions to the Galaxy the putative stripped core NGC 2419 is located quite far ($R_{gc}$ = 92 kpc) from the Galactic center. This suggests that this object may have had a different evolutionary history from those of NGC 5139 = $\\omega$ Centauri $(R_{gc}$ = 6 kpc) and NGC 6715 = M54 $(R_{gc}$ = 19 kpc). (3) McConnachie \\& Irwin (2006) have drawn attention to the fact that the dwarf spheroidals associated with the Galaxy have half-light radii that are two or three times larger than those of the dSph galaxies surrounding M31. The recently discovered Galactic satellite in Ursa Major (Willman et al. 2005) strengthens and confirms this result. It is not yet clear if the observed systematic differences between the dSph satellites of M31 and the Galaxy are due to stronger tidal striping of Galactic companions, or if the presently available data sample might have been more strongly biased against the discovery of Galactic companions of low surface brightness. It should of course be emphasized that it is very likely that many additional very low luminosity satellites of both M31 and the Galaxy remain to be discovered. With the discovery of the UMa system one finds that the satellites of M31 and of the Galaxy, that are located within 150 kpc of their parents, now have a spread in surface brightness in excess of 5 mag arcmin$^{-2}$ . On the other hand the more distant satellites Leo I, Leo II, And II, And VI and And VII appear to have a much smaller luminosity dispersion and all have a surface brightness higher than 25 mag arcmin$^{-2}$. This difference might be due (McConnachie \\& Irwin 2006) to a radial surface density gradient, or perhaps more plausibly, to observational selection effects that have biased the sample of dwarf spheroidal satellites against the discovery of distant low surface brightness objects. In their comparison of (fossil) satellites with detailed numerical simulations of galaxy survival Gnedin \\& Kratsov (2006) note a discrepancy between theory and observation, in the sense that the observed radial distribution of fossils shows an excess of satellites at small radial distances. The present data increase the size of this discrepancy because of the inclusion of the putative stripped cores of dwarf spheroidals which (with the sole exception of NGC 2419) are all located at quite small radial distances from the nuclei of M31 and the Galaxy. A plot of the cumulative radial distribution of all of the satellites of M31 and the Galaxy is shown in Figure 1. This figure appears to show an abrupt break at R $\\sim$25 kpc. The existence of this sharp discontinity suggests that the six innermost satellites (B327 - 3 kpc, M32 - 6 kpc, NGC 5139 = $\\omega$ Centauri -6 kpc, Hux C1 -13 kpc, Hux 3 -14 kpc, and Sgr -19 kpc) might, in some way that is presently not understood, differ from the other satellites of the Galaxy and M31. The observed excess of satellites at small galactocentric distances is surprising because one would actually have expected disruptive tidal forces to have produced a deficiency of satellites with pericentric radii $<$ 30 kpc (Gauthier, Dubinski \\& Widrow 2006). It would be interesting to know if the apparent existence of an excess population of dwarfs at small radial distances is related to a result of recent N-body simulations (Lu et al. 2006) which appear to show that the assembly of cold dark matter halos occurs in two phases: (1) a fast-accretion stage with a rapidly deepening potential well, and (2) a slow-accretion stage characterized by a gentle addition of mass to the outer halo with little change to the inner potential well. A Kolmogorov-Smirnov test shows no statistically significant differences between the distributions of the galactocentric distances of the companions of M31 and of the Galaxy. This conclusion is consistent with that of McConnachie \\& Irwin (2006) which was, however, based on a smaller data sample. Within the, admittedly limited, accuracy of published metallicity values there is no obvious systematic difference between the Mv versus [Fe/H] relationships for the late-type satellites of M31 and of the Galaxy. In summary it appears that the M31 and Milky Way satellite systems are broadly similar, except for the presence of the LMC and the SMC, which might be interlopers that originated in distant reaches of the Local Group. To check on this possibility by detailed orbit computations one would have to have a much improved knowledge of the three-dimensional shape and radial profile of the gravitational potential of the Milky Way dark halo. It is a pleasure to thank Ken Freeman, Eva Grebel, Nitya Kallivayalil and Mario Pedreros for helpful exchanges of of correspondence. I am also indebted to an unusually helpful anonymous referee." }, "0607/astro-ph0607571_arXiv.txt": { "abstract": "We present $RIz$ photometry of four consecutive transits of the newly discovered exoplanet XO-1b. We improve upon the estimates of the transit parameters, finding the planetary radius to be $R_{\\rm P} = 1.184_{-0.018}^{+0.028}~R_{\\rm Jup}$, and the stellar radius to be $R_{\\rm S} = 0.928_{-0.013}^{+0.018}~R_\\odot$, assuming a stellar mass of $M_{\\rm S}=1.00 \\pm 0.03~M_\\odot$. The uncertainties in the planetary and stellar radii are dominated by the uncertainty in the stellar mass. These uncertainties increase by a factor of 2--3 if a more conservative uncertainty of $0.10~M_\\odot$ is assumed for the stellar mass. Our estimate of the planetary radius is smaller than that reported by \\citet{McCullough.2006} and yields a mean density that is comparable to that of TrES-1 and HD~189733b. The timings of the transits have an accuracy ranging from 0.2 to 2.5~minutes, and are marginally consistent with a uniform period. ", "introduction": "An exoplanetary transit is a rare opportunity to learn a great deal about both the planet and the star. With precise measurements of the amount of light blocked by the planet as a function of time, it is possible to infer the relative sizes of the star and planet, the orbital inclination, and the stellar limb-darkening function. Coupled with measurements of the time-variable Doppler shift of the star and an estimate of the stellar mass, one learns the planetary mass and the stellar radius. These fundamental measurements set the stage for a host of more subtle measurements of effects such as planetary atmospheric absorption lines, thermal emission, spin-orbit alignment, and timing anomalies, as reviewed recently by \\citet{Charbonneau.2006c}. For these reasons, newly-discovered transiting exoplanets are welcomed with open arms. The tenth such object was recently reported by \\citet{McCullough.2006}. The parent star, XO-1, is bright ($V=11$, G1~V), making it a favorable target for precise observations. The planet has an orbital period of $\\sim$4~days and a mass and radius comparable to Jupiter, although \\citet{McCullough.2006} point out that their photometry actually implies a mean density that is somewhat smaller than theoretical expectations for ``hot Jupiters.'' If confirmed, this would put XO-1b in the same category as the anomalously large planet HD~209458b, and may have implications for the various theories that have been espoused for that object. Two of us (M.J.H.\\ and J.N.W.) have initiated the Transit Light Curve (TLC) Project, a long-term campaign to build a library of high-precision transit photometry, with the dual goals of (1) refining the estimates of the physical and orbital parameters of the target systems, and (2) searching for secular and short-term variations in the transit times (and light curves) that would be indicative of perturbations from additional bodies~\\citep{Agol.2005,Holman.2005a}. Here, we present results for XO-1b that were obtained as part of this program. We describe the observations and the data reduction procedures in \\S~2. In \\S~3 we describe the model and techniques we used to estimate the physical and orbital parameters of the XO-1 system, and in \\S~4 we summarize our results. ", "conclusions": "Through observations of four consecutive transits, we have significantly improved upon the estimates of the system parameters of XO-1. The most interesting parameters are the radius of the star, the radius of the planet, and the mid-transit times, which will be discussed shortly. The results for the other parameters are not especially interesting but they do seem reasonable. The results for the orbital inclination are best described as bounds on the impact parameter $b$, which is the minimum projected star-planet distance, in units of the stellar radius. It is given by $b=a \\cos i / R_{\\rm S}$, where $a$ is the orbital distance. The data favor a central transit, with $b< 0.27$ and $i<88\\fdg53$ at the 95\\% confidence level. Although the survey and follow-up photometry of \\citet{McCullough.2006} were impressive, and built a convincing case for an exoplanet, those authors did not attempt to fit for the stellar radius when modeling the transit light curve. Instead, they used the value $R_{\\rm S}/R_\\odot=1.00\\pm 0.08$, based on an interpretation of the stellar spectrum. This is because the inference of $R_{\\rm S}$ from a transit light curve requires that the ingress and egress are well sampled and measured with a high signal-to-noise ratio. This type of data was not available. The higher precision and finer time sampling of our data, and of the $z$ band data in particular, allow for the determination of $R_{\\rm S}$ from the light curve, without relying on spectral modeling and theoretical isochrones. The resulting ``photometric'' value of $R_{\\rm S}$ is still subject to a systematic error due to the covariance with $M_{\\rm S}$, but the dependence is fairly weak, $R_{\\rm S} \\propto M_{\\rm S}^{1/3}$, generally leading to a smaller uncertainty in $R_{\\rm S}$ than can be achieved from spectral modeling and theoretical isochrones. Our result is $R_{\\rm S}/R_\\odot = 0.928_{-0.013}^{+0.018}$, which is consistent with (but more precise than) the value determined by \\citet{McCullough.2006}. Here we have incorporated the $0.03~M_\\odot$ uncertainty in $M_{\\rm S}$ determined by \\citet{McCullough.2006}. We note that this radius is somewhat small for the G1~V spectral type of XO-1, but it is still consistent, given the stated uncertainties. We remind the reader again that the quoted result assumes $M_{\\rm S}=1.0~M_\\odot$, and that the inferred $R_{\\rm S}$ scales as $(M_{\\rm S}/M_\\odot)^{1/3}$. From the Yonsei-Yale isochrones, a stellar mass of $M_{\\rm S}= 0.96~M_\\odot$ corresponds to a radius of $R_{\\rm S}=0.91~R_\\odot$, for solar metallicity and an arbitrary age of 3.6~Gyr, ~\\citep{Yi.2001}. Thus, a $1.3~\\sigma$ change in the estimated stellar mass yields an estimated stellar radius that is precisely in line with theoretical expectations. We also note that the stellar radius uncertainty is a factor of 2--3 larger if a more conservative uncertainty of $0.10~M_\\odot$ is assumed for the stellar mass, as shown in Table~2. Our derived radius of XO-1b is $R_{\\rm P}/R_{\\rm Jup} = 1.184_{-0.018}^{+0.028}$ (again assuming the uncertainty in the stellar mass to be $0.03~M_\\odot$). Previously, \\citet{McCullough.2006} found $R_{\\rm P}/R_{\\rm Jup} =1.30\\pm 0.11$. These figures are also in agreement right within their respective 68\\% confidence limits. Interestingly, we obtain very precise agreement with \\citet{McCullough.2006} for all parameters if we first time-average our data into 5-minute bins (i.e., by a factor of 8, for the $z$ band data). The \\citet{McCullough.2006} data were averaged into bins ranging in width from 3 to 9 minutes, depending on the telescope used. We suggest that it is possible that some of the previous results were slightly biased by the coarser time sampling of the photometry. Resolving the degeneracy among the stellar radius, planetary radius, and orbital inclination requires adequate sampling of ingress and egress. The uncertainties in the limb darkening parameters $u_1$ and $u_2$ are highly correlated, with the linear combination $2u_1 + u_2$ being well constrained by the data, and the orthogonal combination $u_1 - 2u_2$ being weakly constrained by the data (see the lower left panel of Fig.~3). We find $2 u_1 + u_2 = 0.86 \\pm 0.05$. This is $2~\\sigma$ larger than the value based on the theoretical calculations of \\citet{Claret.2004}, which predict $u_1 = 0.21$, $u_2 = 0.33$, and $2u_1 + u_2 = 0.75$ (for the standard $z$ band, $T=5750~\\rm K$, $\\log g = 4.5$, $[M/H] = 0.05$, and microturbulent velocity $v_t = 2.0~\\mathrm{km/s}$). The theoretical values are shown in Fig.~3 as the solid symbol in the $u_1$--$u_2$ plot, and as dotted lines in the other two limb-darkening plots. One might consider reducing the number of degrees of freedom and adopting the \\citet{Claret.2004} values as fixed quantities. When we do so, we find $R_{\\rm S}/R_\\odot = 0.94$, $R_{\\rm P}/R_{\\rm Jup} = 1.22$, and $b = 0.26$ ($i = 88.65$~deg), with the minimum $\\chi^2$ increased by 9. However, given the quality of the $z$ band data, the unknown level of uncertainty in the theoretical values, and the possible differences between the FLWO48/KeplerCam $z$ band and the standard SDSS $z$ band, we believe fitting for the limb-darkening coefficients is more appropriate. Some of the probability distributions shown in Fig.~3 are asymmetric. This is typical of all fits to transit light curve data. The fact that the orbital inclination has a maximum value (namely, 90\\arcdeg), combined with the measured durations of the ingress, egress, and the full transit, imposes this asymmetry among the covariant parameters $R_{\\rm S}$, $R_{\\rm P}$ and $b$. Our downward revision of the planetary radius translates into an increased value for the mean density, $ 0.67 \\pm 0.07 \\ {\\rm g \\, cm^{-3}}$. This value is $45$--$56$\\% that of Jupiter. This is comparable to, but slightly less than, the mean densities of TrES-1 ($0.84 \\ {\\rm g \\, cm^{-3}}$; \\citealt{Sozzetti.2004}) and HD~189733b ($0.93 \\ {\\rm g \\, cm^{-3}}$; \\citealt{Bakos.2006}). For XO-1b's estimated equilibrium temperature $T_{\\rm eq} = 1100~{\\rm K}$ (assuming Bond albedo $A_B = 0.4$ and our derived value of the stellar radius $R_{\\rm S} = 0.928~R_\\odot$) and its mass $M_{\\rm P} = 0.9~M_J$, the models of \\citet{Bodenheimer.2003} predict planetary radii of $R_{\\rm P} = 1.04~R_J$ and $1.11~R_J$ for models with and without a 20~$M_\\oplus$ core, respectively. Our estimate of the radius of XO-1b is $2~\\sigma$ larger than its predicted value, even for a planet without a core. HD~189733b's measured radius $R_{\\rm P} = 1.154 \\pm 0.032~R_J$ \\citep{Bakos.2006} is also larger than its theoretical value ($R_{\\rm P} = 1.03~R_J$ with a core, $R_{\\rm P} = 1.11~R_J$ without a core), given its equilibrium temperature $T_{\\rm eq} = 1050~{\\rm K}$ and mass $M_{\\rm P} = 0.82 \\pm 0.03$~\\citep{Bouchy.2005}. In contrast, \\citet{Laughlin.2005} showed that TrES-1's measured radius $R_{\\rm P} = 1.08 \\pm 0.05~R_J$ \\citep{Laughlin.2005} is consistent with its theoretically predicted value ($R_{\\rm P} = 1.05~R_J$ with a core, $R_{\\rm P} = 1.09~R_J$ without a core). The measured radii of both XO-1b and HD~189733 are consistent with the predictions of \\citet{Bodenheimer.2003} if ``kinetic heating'' is included. In these models $\\sim 2$\\% of the stellar insolation is deposited at depth, following the work of \\citet{Guillot.2002}. The age of XO-1 is also uncertain; the planet would be larger if the system were younger~\\citep{Burrows.2003}. Whether or not the radius of XO-1b requires an additional energy source, as is the case for HD~209458b ($0.36 \\ {\\rm g \\, cm^{-3}}$; \\citealt{Knutson.2006}), is an important topic for future theoretical work. The kinetic heating model, proposed to explain the apparent inflation of HD~209458b, would naturally predict that many other ``hot Jupiters'' should be inflated. Other explanations, such as ongoing tidal circularization due to an eccentricity exchange with a third body \\citep{Bodenheimer.2001}, or the trapping in a Cassini state with nonzero obliquity \\citep{Winn.2005a}, would seemingly have difficulty accounting for a large population of inflated objects. The accuracy of our transit times ranges from 0.2~minutes for the FLWO $z$ band observations to 2.5~minutes for the transit observed solely by TopHAT\\@. Fig.~5 shows the differences between the observed and predicted times of mid-transit, as a function of transit epoch. The predicted times assume the average orbital period determined by \\citet{McCullough.2006} and a reference time based on our observations. So far, all the times are marginally consistent with a constant period. These observations provide accurate anchors for future searches for transit time variations. \\clearpage \\begin{figure}[p] \\epsscale{1.0} \\plotone{f5.eps} \\caption{ The timing residuals for the 4 observed transits, according to the ephemeris of \\citet{McCullough.2006}. (see Eq.~1). The first point corresponds to the $T_c$ of \\citet{McCullough.2006}. The points lie on a horizontal line, and therefore the data are marginally consistent with a constant period. \\label{fig:mass_radius}} \\end{figure} \\clearpage" }, "0607/astro-ph0607021_arXiv.txt": { "abstract": "Shocks and blastwaves are expected to be driven driven into the intracluster medium filling galaxy groups and clusters by powerful outbursts of active galactic nuclei or quasars in the member galaxies; the first footprints of shock fronts have been tentatively traced out with X-ray imaging. We show how overpressures in the blasts behind the shock can prove the case and also provide specific marks of the nuclear activity: its strength, its current stage, and the nature of its prevailing output. We propose to detect these marks with the aimed pressure probe constituted by the resolved Sunyaev-Zel'dovich effect. We compute and discuss the outcomes to be expected in nearby and distant sources at different stages of their activity. ", "introduction": "Density jumps have been recently pinpointed by X-ray imaging of the hot intracluster medium (ICM) that pervades galaxy groups and clusters with average densities around $n \\sim 10^{-3}$ cm$^{-3}$ and temperatures $T$ in the keV range. These jumps have been interpreted in terms of shock fronts propagating into the ICM out to radial distances $r \\approx 0.2$ Mpc, with Mach numbers around $\\mathcal{M}\\approx 1.5$, and involving energies up to $\\Delta E \\approx 3\\times 10^{61}$ ergs (Mazzotta et al. 2004; Mc Namara et al. 2005; Forman et al. 2005; Nulsen et al. 2005a, b). Shocks over large scales with such \\emph{intermediate} strengths were specifically expected by Cavaliere, Lapi \\& Menci (2002) as marks of the energy being fed back into the ICM by active galactic nuclei (AGNs) when they flare up in member galaxies of a group or cluster. Such events occur when a central supermassive black hole (BH) accretes an additional mass $M_{\\bullet} \\sim 10^9 \\, M_{\\odot}$; with standard efficiency $\\eta\\sim 10^{-1}$ for mass-energy conversion, this yields over times $\\Delta t\\sim 10^8$ yr energies of order $2\\times 10^{62}\\, (M_{\\bullet}/10^9\\, M_{\\odot})$ ergs. If these outputs couple at levels $f\\sim$ a few percents to the surrounding ICM, they constitute impulsive, considerable additions $\\Delta E \\approx 10^{61}\\, (f/5\\%)\\, (M_{\\bullet}/10^9 M_{\\odot})$ ergs to its binding energy. In fact, the latter comes to $E \\approx G\\,M\\, m/4\\,r \\approx 3 \\times 10^{61}$ ergs in the central $0.2$ Mpcs of a cluster, that encompass a dark matter (DM) mass $M\\sim 5\\times 10^{13}\\, M_{\\odot}$ and an ICM mass fraction $m/M \\approx 0.15$. With such appreciable ratios $\\Delta E/E$, we expect in the ICM a large-scale \\textit{blastwave} bounded by a leading \\textit{shock} that starts from the host galaxy and moves into the surrounding ICM out to several $10^{-1}$ Mpc. In fact, the shock Mach numbers are provided in the simple form $\\mathcal{M} \\approx (1 + \\Delta E/E)^{1/2}$ by the hydrodynamics of the ICM (see Lapi, Cavaliere \\& Menci 2005), a good electron-proton plasma which in thermal equilibrium (see \\S~4) constitutes a single ``monoatomic\" fluid with pressure $P\\approx 2\\, n\\, kT$. For BH masses bounded by $M_{\\bullet} \\la 5 \\times 10^9\\, M_{\\odot}$ (Ferrarese 2002; Tremaine et al. 2002) relative energy inputs $\\Delta E/E\\la 1$ obtain and yield just $\\mathcal{M}\\approx 1.5$; they also yield standard Rankine-Hugoniot jumps of the post- to the pre-shock density $n_2/n_1 = 4\\, \\mathcal{M}^2 /(\\mathcal{M}^2 + 3)\\approx 1.7$, consistent with the X-ray analyses. What other marks will establish such shocks and blasts? One is constituted by the temperature. This rises sharply across a shock but then falls down in the blast, if nothing else by adiabatic cooling. In any case, resolved spectroscopic measurements of the electron $T$ require many X-ray photons, more than currently available from distant clusters or groups. Pressure provides another, independent mark; the electron pressure is \\emph{directly} sensed with the SZ effect (Sunyaev \\& Zel'dovich 1972). The pressure at the shock is to jump up from the unperturbed level $P_1$ by the factor $P_2/P_1 = (5\\, \\mathcal{M}^2 -1)/4$, which marks a shock from a cold front; moreover, $P$ must retain sufficiently high levels throughout the blast as to propel forward the ICM it sweeps up. We will see that the radial pressure run $P(r)$ actually \\emph{rises} from the leading shock to an inner ``cavity'', as long as the blast is driven on by a central AGN. Power must be transmitted from it to the surrounding ICM blast by means of an intervening medium. This may be constituted either by relativistic particles filling up a radiovolume energized by jets (Scheuer 1974; Heinz, Reynolds, \\& Begelman 1998); or by another and hotter plasma heated up by the impact of radiation-driven superwinds (see Lamers \\& Cassinelli 1999). Thus SZ pressure probing will have direct implications for the kind and the time-cycle of the AGN outputs, in particular of the radio-loud components. ", "conclusions": "SZ signals are widely measured in clusters at levels $y\\approx 10^{-4}$, see Reese et al. (2002) and Birkinshaw (2004); Lapi, Cavaliere \\& De Zotti (2003) discuss sub-arcmin resolutions to detect the integrated effects of AGN outbursts on the ICM of clusters and groups. Here we have focused on SZ signals \\emph{resolved} at levels of $10''$ or better, to probe the structure and the dynamics of the shocks and blasts so produced. The SZ probe is best used in scanning the ICM around density jumps selected in X-rays. In fact, the bremsstrahlung surface brightness proportional to $n^2$ is well suited for pinpointing density jumps and providing positions and Mach numbers of candidate shocks. But measuring $T$ from X-ray spectroscopy must contend with paucity of photons and the narrow post-shock range where $T$ exceeds the unperturbed value. So in conditions of low surface brightness (outskirts or distant structures) the SZ effect will lend a strong hand by unveiling the other key observable, namely, the overpressures behind the shocks; this is due to three circumstances. First, pressures are sensed \\emph{directly} through the parameter $y \\propto P$. Second, $y$ is independent of $z$ for sources wider than the instrumental beamwidth. Third, we expect (see Fig.~1) thermal pressures to rise throughout a blast continuously \\emph{driven} over a crossing time; from shock to piston at radii $40 - 30\\%$ smaller, $P(r)$ rises up to values $P_p\\approx 8 - 9\\, P_1$, considerably larger than the shock jump $P_2/P_1\\approx 2.6 - 4.8$ for $\\mathcal{M} \\approx 1.5 - 2$; this is due to the dynamical stress in running blasts. This rising behavior of $P(r)$ is just opposite to the run down from shock to center expected for \\emph{free} blasts launched by short-lived AGN activity. We compare in Fig.~1 the results we expect; they constitute an aim particularly interesting for shocks of intermediate $\\mathcal{M}$ pinpointed in close proximity to a radiovolume or around a currently shining AGN. Clear study cases will be provided by clusters or groups in quiet conditions, with no sign of outer merger induced dynamics. What is needed for SZ probing at $z\\approx 0.1$ is a resolution around $10''$ in the upper $\\mu$wave band, already approached with the Nobeyama radiotelescope (Kitayama et al. 2004, also \\url{http://www.nro.nao.ac.jp/index-e.html/}). Upcoming instruments such as CARMA (\\url{http://www.mmarray.org/}) will do better; at resolutions of a few arcsecs, blasts in clusters at $z\\approx 0.5$ may be probed at levels comparable to the \\textit{top right} panel of Fig.~1. ALMA (\\url{http://www.alma.nrao.edu/}) with its planned resolutions around $1''$ and sensitivities down to $1\\,\\mu$K, will do better yet both in the $\\mu$wave and in the submm band. Following the noise analysis by Pfrommer et al. (2005), a $5\\sigma$ detection with ALMA of the thermal SZ signal we focus on will require scanning a limited area in the vicinity of an X-ray preselected position for a few hours per cluster. We add that once kpc scales will be resolved, two-fluid effects will be interesting as may arise from disequilibrium between the electrons and ions (see Ettori \\& Fabian 1998; Fabian et al. 2006). They may cause a gradual rise of the electron pressure in front of the shock, to converge behind it with the declining ion pressure toward the equilibrium values considered above. Here we stress three issues. First, how common are shocks in clusters? The frequent occurrence of quenched cooling flows argues for widespread shocks driven over some Gyrs by central AGNs injecting energies $\\Delta E$ in excess of the cooling losses (see B\\^{i}rzan et al. 2004; Nulsen et al. 2005a, b). Second, in what prevailing mode does an AGN release the driving energy $\\Delta E$, in mechanical form and radio jets or in radiation-driven superwinds? To a first approximation the mode little affects the total energy injected, based on the simple rule $\\Delta E \\sim \\alpha\\, f\\,\\eta\\, c^2 \\, \\dot M_{\\bullet}\\, \\Delta t \\sim $ const that we extract from Churazov et al. (2005); that is, mechanical energy and jets gain on grounds of coupling efficiency ($\\alpha f\\approx 1/2$ vs. a few $\\%$) what they lose to radiation on grounds of $\\eta\\dot M_{\\bullet}$. Conversely, the shock energetics alone will not distinguish the mode. Third, what other probe may help? The specific SZ probing we propose can trace the distinguishing features of the injection: its timescale (past or ongoing), and its prevailing content (relativistic particles or photons). The SZ effect resolved at levels of $10''$ or better will directly detect the mark of a blast launched or driven by a powerful AGN, namely, the hydro overpressure \\emph{jumping} up at the shock and \\emph{sustained} throughout the blast. But when the radiosource drive persists over the transit time the pressure actually \\emph{rises} in the blast, and the SZ signal is enhanced; it is boosted up by plasma in the cavity if the drive is helped by AGN superwinds. These outcomes are independent of model details, rather they depend on a few overall parameters: the relative injection energy $\\Delta E/E$, evaluated from $\\mathcal{M}$; the active time of the AGN compared with the blast crossing time; the injection mode, whether dominantly in mechanical or radiative energy. We conclude that SZ measurements concurring with the X-ray imaging can effectively probe the injection mode, the dynamics of the blasts, and the history of the driving AGN sources." }, "0607/astro-ph0607217_arXiv.txt": { "abstract": "Most population I Wolf-Rayet (WR) stars are the He-rich descendants of the most massive stars ($M_{i} = 25 - 100 M_{\\odot}$). Evidence has been accumulating over the years that among all pop I WR stars, those of the relatively cool, N-rich subtype \"WN8\" are among the most peculiar: \\\\ 1. They tend to be runaways, with large space velocity and/or avoid clusters. 2. Unlike their equally luminous WN6,7 cousins, only a very small number of WN8 stars are known to belong to a close binary with an OB companion. 3. They are the systematically most highly stochastically variable among all (single) WR stars. \\\\ Taken together, these suggest that many WN8 stars may originally have been in close binaries (like half of all stars), in which the original primary exploded as a supernova, leaving behind a very close binary containing a massive star with a neutron star/black hole companion (like Cyg X-3). When the massive remaining star evolved in turn, it engulfed and eventually swallowed the compact companion, leading to the presently puffed-up, variable WN8 star. Such stars could fall in the realm of the exotic Thorne-$\\dot{\\mbox{Z}}$ytkow objects. ", "introduction": "Thorne-$\\dot{\\mbox{Z}}$ytkow Objects (TZOs) have been proposed for the first time by Thorne and $\\dot{\\mbox{Z}}$ytkow (1977). They are stars with a degenerate neutron core which provides peculiar conditions to ensure the presence of a nuclear burning region. This supports the envelope of the star, which appears like a Red Supergiant (see e.g. Biehle 1991, 1994; Cannon et al. 1992; Cannon 1993). Their lifetime in the RSG phase is also expected to last as long as for a normal RSG (i.e. with nuclear burning right to the center of the core). Being RSGs (or appearing so), TZOs probably have strong winds. We develop here the idea that if TZOs evolve like normal RSGs do, the most massive of them will evolve further to the WR stage, following the usually adopted scenario that RSG become WR stars because of strong mass loss by stellar wind, stripping the outer envelope (see e.g. Garcia-Segura et al. 1996). Among the population I Wolf-Rayet stars, those of subtype WN8 are peculiar. We present here why they are good candidates to be \"evolved\" TZOs, formed via a binary scenario. This paper is organized as follows: in section 2 we give the details of two binary scenarii following which TZOs can be formed. These two scenarii provide the same paradigm to explain TZOs and objects like Cyg X-3. In section 3 we discuss how these WN8 stars are really different and how they can really come from a binary evolution scenario. In section 4 we discuss observations which could be undertaken to confirm such exotic objects. Section 5 gives our conclusions. ", "conclusions": "We have presented here the two known scenarii in which TZOs form following close binary evolution and a common envelope phase. The hydrodynamical simulations of Terman et al. (1995) dealing with the entry of the NS inside the envelope of a more-or-less evolved RSG allowed us to construct a paradigm within which WN8 stars, TZOs and objects like Cyg X-3, find an explanation. We also drew attention to the peculiarities of WN8 stars which fit well in our \"model\". Finally we believe that spectroscopic observations might be feasible to reveal the presence of a degenerate neutron core inside single WN8 stars, which would then confirm the existence of TZOs. To the question: \"Are peculiar Wolf-Rayet stars of type WN8 Thorne-$\\dot{\\mbox{Z}}$ytkow Objects?\" we should now respond \"Why not?\"." }, "0607/astro-ph0607108.txt": { "abstract": "{The highly obscured radio-bright galaxy PKS\\,1343\\,--\\,601 at Galactic coordinates of $(l,b) = (309\\fdg7, +1\\fdg8$) has been suspected to mark the centre of a hitherto unknown cluster in the wider Great Attractor region. As such it presents an ideal region for a search of galaxies in the near-infrared (NIR) and an in-depth study of their colours as a function of extinction. A visual search of a $\\sim\\!30$ square-degree area centered on this radio galaxy on images of the NIR DENIS survey (\\IJK ) revealed 83 galaxies (including two AGNs) and 39 possible candidates. Of these, 49 are also listed in the 2MASS Extended Source Catalog 2MASX. Taking the IRAS/DIRBE extinction values (Schlegel \\etal 1998) at face value, the absorption in the optical ($A_B$) ranges from $\\sim\\!2^{\\rm m}$ to over $100^{\\rm m}$ across the Galactic Plane. Comparing the detections with other systematic surveys, we conclude that this search is highly complete up to the detection limits of the DENIS survey and certainly surpasses any automatic galaxy finding algorithm applied to crowded areas. The NIR galaxy colours from the $7\\arcsec$ aperture were used as a probe to measure total Galactic extinction. A comparison with the IRAS/DIRBE Galactic reddening maps suggests that the IRAS/DIRBE values result in a slight overestimate of the true extinction at such low Galactic latitudes; the inferred extinction from the galaxy colours corresponds to about 87\\% of the IRAS/DIRBE extinctions. Although this determination still shows quite some scatter, it proves the usefulness of NIR surveys for calibrating the IRAS/DIRBE maps in the extinction range of about $2^{\\rm{m}} \\la A_B \\la 12^{\\rm{m}}$. % ", "introduction": "\\label{intro} Various extragalactic large-scale structures are hidden behind the dust and stars of the Milky Way, the so-called Zone of Avoidance (ZoA), resulting in a poor understanding of the dynamics of the nearby Universe; for a detailed overview see Kraan-Korteweg \\& Lahav (2000), Kraan-Korteweg (2005), and the conference proceedings ``Nearby Large-Scale Structures and the Zone of Avoidance'' (Fairall \\& Woudt 2005). The Great Attractor (GA), an extended mass overdensity in the nearby Universe, lies for instance close to the crossing of the Supergalactic plane and the Galactic plane. Its presence was inferred by the systematic large-scale flow of elliptical galaxies (Lynden-Bell \\etal 1988). Applying the potential reconstruction method of the mass density field POTENT (Dekel 1994), Kolatt \\etal (1995) found its centre at $(l,b,v) = (320\\deg, 0\\deg, 4000$\\kms ). Close to the potential well of the GA lies the cluster ACO 3627 $(l,b,v) = (325\\deg, -7\\deg, 4848$\\kms ). A deep optical galaxy search (Woudt \\& Kraan-Korteweg 2001) revealed this cluster to be as massive and rich a cluster as the Coma cluster (Kraan-Korteweg \\etal 1996, Woudt \\etal 2005). It therefore most likely marks the centre of the potential well of the GA. However, the GA is an extended region of high galaxy density (about $40\\deg \\times 40\\deg$ on the sky, see Kolatt \\etal 1995), and other clusters (rich and poor) may well contribute substantially to this mass overdensity. Identifying them is a challenge as the central part of the wider GA area lies behind the thickest dust layer of the Milky Way. About $10\\deg$ from the Norma cluster, at $(l,b) = (309\\fdg7, +1\\fdg8$), lies the galaxy PKS\\,1343\\,--\\,601 with a recession velocity of 3872\\kms\\ (West \\& Tarenghi 1989). Near-infrared (NIR) observations revealed PKS\\,1343\\,--\\,601 to be a giant elliptical galaxy, which often reside at the centre of galaxy clusters. It is also one of the brightest radio sources in the sky (McAdam 1991): its flux density is only surpassed by Cygnus A, Centaurus A, Virgo A, and Fornax A. Two of these four radio sources are situated at the centre of a rich cluster, one in a smaller cluster, and one in a group of galaxies (Jones \\etal 2001). This evidence motivated Kraan-Korteweg \\& Woudt (1999) to investigate by different means whether PKS\\,1343\\,--\\,601 points to another cluster in the Great Attractor region. Such a cluster would have a considerable impact on the local velocity field calculations. Results are still controversial. A preliminary analysis of the systematic deep \\HI\\ search for galaxies with the Parkes Multibeam receiver found a concentration of galaxies in redshift space around this radio galaxy (Kraan-Korteweg \\etal 2005b). A deep NIR search (\\JHK ) of half a degree radius and a deep \\II -band survey of 2 degrees around PKS\\,1343\\,--\\,601 (Nagayama \\etal 2004; Kraan-Korteweg \\etal 2005a, respectively) have revealed a distribution of galaxies consistent with a (medium-sized) cluster around PKS\\,1343\\,--\\,601. X-ray observations with ASCA have only revealed diffuse emission from PKS\\,1343\\,--\\,601 itself (Tashiro \\etal 1998, see also the discussion in Ebeling \\etal 2002), which would rule out a rich cluster. This paper presents the results of a search for galaxies based on the NIR DENIS survey (Epchtein \\etal 1997) in a much larger but shallower area than the above ones. The advantages of using the NIR to search for galaxies in the ZoA are manifold: (i) the NIR is less affected by the foreground extinction than the optical (the extinction in the \\K -band is about 10\\% of the extinction in the \\B -band); (ii) the NIR is sensitive to early-type galaxies, tracers of massive groups and clusters (contrarily to far-infrared and blind \\HI\\ surveys); (iii) the NIR shows little confusion with Galactic objects such as young stellar objects and cool cirrus sources. In pilot studies, we have assessed the performance of the DENIS survey at low Galactic latitudes (Schr\\\"oder \\etal 1997; Kraan-Korteweg \\etal 1998; Schr\\\"oder \\etal 1999; Mamon et al. 2001). We tested the potential of the DENIS survey to detect galaxies where optical and far-infrared surveys fail, \\ie at high foreground extinctions and in crowded regions; we established that the NIR colours of galaxies lead to values for the foreground extinction; and we cross-identified highly obscured galaxies detected in a blind \\HI\\ search at $|b| < 5\\deg$. Overall, both systematic NIR surveys DENIS (\\IJK ; Paturel \\etal 2003, Vauglin \\etal 1999) and 2MASS (\\JHK ; Skrutskie \\etal 2006, Jarrett \\etal 2000a) have proven their effectiveness in penetrating the ZOA (Jarrett \\etal 2000b, Rousseau \\etal 2000, Schr\\\"oder \\etal 2000) -- as long as the star density does not exceed a certain value (Kraan-Korteweg \\& Jarrett 2005). In the following, we will introduce the DENIS survey and the implication of extinction on galaxy counts in general (Sections~\\ref{denis} and~\\ref{counts}, respectively). We will then describe the search area and the quality of the DENIS data (Sect.~\\ref{strips}), and the methods of galaxy and parameter extraction (Sect.~\\ref{param}). In Sect.~\\ref{cat} the catalogue data are described, Sect.~\\ref{lit} gives a detailed comparison with the data of other searches and catalogues in this area, and in Sect.~\\ref{ext} we investigate the extinction in this area using the derived NIR colours. Conclusions are presented in the final Sect.~\\ref{conclusion}. Throughout the paper, we assume a Hubble constant of $H_0=70 \\rm km\\,s^{-1}\\,Mpc^{-1}$. A second paper will provide a detailed discussion of the local environment of PKS\\,1343\\,--\\,601 using the local galaxy density, the velocity distribution, and the X-ray luminosity to assess its mass and contribution to the GA overdensity (Schr\\\"oder \\& Mamon 2006, hereafter Paper II). ", "conclusions": "" }, "0607/astro-ph0607351_arXiv.txt": { "abstract": "The inner few hundred parsecs of our galaxy provide a laboratory for the study of the production and propagation of energetic particles. Very-high-energy $\\gamma$-rays provide an effective probe of these processes and, especially when combined with data from other wave-bands, $\\gamma$-rays observations are a powerful diagnostic tool. Within this central region, data from the H.E.S.S. instrument have revealed three discrete sources of very-high-energy $\\gamma$-rays and diffuse emission correlated with the distribution of molecular material. Here I provide an overview of these recent results from H.E.S.S. ", "introduction": "The central 200~pc of our Galaxy is a unique region that harbours many remarkable objects --- including several potential sites of effective particle acceleration. Non-thermal emission (particularly in the radio, X-ray and $\\gamma$-ray bands) can be used to trace the energetic particle populations of this region. X-ray and radio observations of synchrotron emission provide information on the product of the local magnetic field energy density and the density of relativistic electrons. The flux of inverse Compton $\\gamma$-rays, on the other hand, is proportional to the radiation field density (and the electron density). The combination of $\\gamma$-ray and X-ray measurements therefore provides a powerful tool for probing both magnetic field strength and energetic particle content. Moreover, $\\gamma$-rays provide an effective tracer for hadronic particles: proton-proton interactions in the interstellar medium lead to the production and decay of pions and hence $\\gamma$-ray production. The combination of $\\gamma$-ray measurements with tracers of atomic and molecular material may be the \\emph{only} way to effectively trace energetic hadrons in our galaxy. The usefulness of $\\gamma$-ray observations has traditionally been limited by poor angular resolution and modest sensitivity. A major step forward in the very-high-energy (VHE, $>100$ GeV) domain has recently been taken with the commissioning of H.E.S.S. (High Energy Stereoscopic System). H.E.S.S. is an array of four, 13~m diameter, imaging Cherenkov telescopes located in the Khomas highlands of Namibia~\\cite{HESS}, a southern hemisphere location ideal for observations of the Galactic Centre (GC). H.E.S.S. has an angular resolution of a few arc-minutes and a locational accuracy of $\\sim 30''$ for typical point sources. The instrument reaches an energy flux sensitivity of $10^{-12}$ erg cm$^{-2}$ s$^{-1}$, an order of magnitude lower than the previous generation of VHE instruments. The wide field of view of H.E.S.S. (5$^{\\circ}$ in diameter) enables us to simultaneously monitor the entire central 200 parsec region. Results from the first two years of H.E.S.S. GC observations are described here. \\begin{figure}[h] \\begin{center} \\includegraphics[width=36pc]{fig1_gc_fov.eps} \\caption{\\label{fig1} The H.E.S.S. view of the central 200 parsecs (reproduced from \\cite{HESS_gc}). Top panel: smoothed count map (without background subtraction) for data taken with H.E.S.S. in 2004. Bottom panel: the same data after subtraction of point-like excesses at the positions of Sgr A$^{\\star}$ and G0.9+0.1. The location of Sgr~A$^{\\star}$ is marked with a black star and the G\\,0.9+0.1 as a yellow circle. 95\\% confidence regions for the positions of EGRET sources are shown as dashed ellipses~\\cite{EGRETEllipses}. } \\end{center} \\hspace{2pc} \\end{figure} ", "conclusions": "Present and future $\\gamma$-ray observations will play a key role in our understanding of the physical processes at work in the Galactic Centre. Key outstanding issues, such as the spatial distribution and strength of magnetic fields and the energy density of relativistic particles, can be addressed with the help of such observations. H.E.S.S. has provided the first sensitive view of this region in very-high-energy $\\gamma$-rays. In a few years time the combination of the second phase of the H.E.S.S. project and the GLAST satellite will provide unbroken sensitive coverage of the $10^{8}$--$10^{13}$ eV $\\gamma$-ray domain, with important consequences for Galactic Centre research." }, "0607/astro-ph0607398_arXiv.txt": { "abstract": "We measured the angular clustering at $z\\sim6$ from a large sample of \\ip\\ dropout galaxies (293 with \\zp$\\le$27.5 from GOODS and 95 with \\zp$\\le$29.0 from the UDF). Our largest and most complete subsample (having $L\\gtrsim0.5L^*_{z=6}$) shows the presence of clustering at 94\\% significance. For this sample we derive a (co-moving) correlation length of $r_0=4.5^{+2.1}_{-3.2}$ $h_{72}^{-1}$ Mpc and bias $b=4.1^{+1.5}_{-2.6}$, using an accurate model for the redshift distribution. No clustering could be detected in the much deeper but significantly smaller UDF, yielding $b<4.4$ (1$\\sigma$). We compare our findings to Lyman break galaxies at $z\\sim3-5$ at a fixed luminosity. Our best estimate of the bias parameter implies that \\ip\\ dropouts are hosted by dark matter halos having masses of $\\sim10^{11}$ $M_\\odot$, similar to that of \\vp\\ dropouts at $z\\sim5$. We evaluate a recent claim that at $z\\gtrsim5$ star formation might have occurred more efficiently compared to that at $z=3-4$. This may provide an explanation for the very mild evolution observed in the UV luminosity density between $z=6$ and 3. Although our results are consistent with such a scenario, the errors are too large to find conclusive evidence for this. ", "introduction": "\\label{sec:intro} The Advanced Camera for Surveys \\citep[ACS;][]{ford98} aboard the {\\it Hubble Space Telescope} has made the detection of star-forming galaxies at $z\\sim6$ (\\ip\\ dropouts) relatively easy. The largest sample of \\ip\\ dropouts currently available \\citep{bouwens05_z6} comes from the Great Observatories Origins Deep Survey \\citep[GOODS;][]{giavalisco04_survey}, allowing the first quantitative analysis of galaxies only 0.9 Gyr after recombination \\citep[][see also \\citealt{shimasaku05,ouchi05}]{stanway03,bouwens03,yan04,dickinson04,malhotra05}. \\citet{bouwens05_z6} found evidence for strong evolution of the luminosity function between $z\\sim6$ and 3, while the (unextincted) luminosity density at $z\\sim6$ is only $\\sim0.8$ times lower than that at $z\\sim3$. Some \\ip\\ dropouts have significant Balmer breaks, indicative of stellar populations older than 100 Myr and masses comparable to those of $L^*$ galaxies at $z\\approx0$ \\citep{eyles05,yan05}. Through the study of the clustering we can address fundamental cosmological issues that cannot be answered from the study of galaxy light alone. The strength of clustering and its evolution with redshift allows us to relate galaxies with the underlying dark matter and study the bias. The two-point angular correlation function (ACF) has been used to measure the clustering of Lyman break galaxies (LBGs) at $z=3-5$ \\citep[e.g.,][]{adelberger98,adelberger05,arnouts99,arnouts02,magliocchetti99,giavalisco01,ouchi01,ouchi04_r0,porciani02,hildebrandt05,allen05,kashikawa06}. LBGs are highly biased ($b\\simeq2-8$), and this biasing depends strongly on rest frame UV luminosity and, to a lesser extent, on dust and redshift. The clustering statistics of LBGs have reached the level of sophistication that one can measure two physically different contributions. At small angular scales the ACF is dominated by the non-linear clustering of galaxies within single dark matter halos, whereas at large scales its amplitude tends to the ``classical'' clustering of galaxies residing in different halos \\citep{ouchi05_smallscale,lee05}, as explained within the framework of the halo occupation distribution \\citep[e.g.,][]{zehavi04,hamana04}. Understanding the clustering properties of galaxies at $z\\sim6$ is important for the interpretation of ``overdensities'' observed towards luminous quasars and in the field \\citep{ouchi05,stiavelli05,wang05,zheng06} that could demarcate structures that preceded present-day massive galaxies and clusters \\citep{springel05}. Our aim here is to ``complete'' the census of clustering by extending it to the highest redshift regime with sizeable samples. In \\S\\S\\ 2 and 3 we describe the sample, and present our measurements of the ACF. In \\S\\ 4 we discuss our findings. Throughout we use the cosmology ($\\Omega_M$, $\\Omega_\\Lambda$, $h_{72}$, $n$, $\\sigma_8$)$=$(0.27,0.73,1.0,1.0,0.9) with $H_0=72$ $h_{72}$ km s$^{-1}$ Mpc$^{-1}$. \\begin{figure}[t] \\begin{center} \\includegraphics[width=\\columnwidth]{f1.ps} \\end{center} \\caption{\\label{fig:nz}Redshift distributions of \\ip\\ dropouts in our GOODS ({\\it left}) and UDF ({\\it right}) selections \\citep[estimated by projecting a complete UDF \\bp\\ dropout sample scaled to the sizes and colors as found for the \\ip\\ dropout sample to $z\\sim5-7$; see][for details]{bouwens05_z6}. As a result of a more significant photometric scatter in \\ip--\\zp, the selection extends to lower redshifts in GOODS than it does for the UDF.} \\end{figure} ", "conclusions": "" }, "0607/astro-ph0607620_arXiv.txt": { "abstract": "We present a study of the early (days to weeks) X-ray and UV properties of eight Type Ia supernovae (SNe Ia) which have been extensively observed with the X-Ray Telescope (XRT) and UV/Optical Telescope (UVOT) onboard \\S, ranging from 5--132 days after the outburst. SN~2005ke is tentatively detected (at a 3--3.6$\\sigma$ level of significance) in X-rays based on deep monitoring with the XRT ranging from 8 to 120 days after the outburst. The inferred X-ray luminosity [$L_{0.3-2}=(2\\pm1)\\times10^{38}~{\\rm ergs~s}^{-1}$; 0.3--2~keV band] is likely caused by interaction of the SN shock with circumstellar material (CSM), deposited by a stellar wind from the progenitor's companion star with a mass-loss rate of $\\dot{M} \\approx 3 \\times 10^{-6}~M_{\\odot}~{\\rm yr}^{-1}~(v_{\\rm w}/10~{\\rm km~s}^{-1})$. Evidence of CSM interaction in X-rays is independently confirmed by an excess of UV emission as observed with the UVOT onboard \\S, starting around 35 days after the explosion. The non-detection of \\sn\\ with \\C\\ 105~days after the outburst implies a rate of decline steeper than $L_{\\rm x} \\propto t^{-0.75}$, consistent with the decline expected from the interaction of the SN shock with a spherically symmetric CSM ($t^{-1}$). None of the other seven SNe Ia is detected in X-rays or shows a UV excess, which allows us to put tight constraints on the mass-loss rates of the progenitor systems. ", "introduction": "\\label{introduction} Type Ia supernovae (SNe Ia) are a subclass of exploding stars defined observationally by the absence of hydrogen lines in their optical spectra and the presence of lines from elements such as silicon and sulfur (Leibundgut 2000). There is consensus that SNe Ia are explosions of white dwarfs which occur when accretion from a companion star drives the white dwarf mass close to the Chandrasekhar limit (Woosley \\& Weaver 1986, Nomoto et~al.\\ 2003). However, the details of the system are not fully understood, especially with regard to the type of companion (a main sequence star, a red giant). A useful indicator of the properties of the companion is through its mass loss, which depends strongly on the stellar type. The X-rays and the radio regimes are especially well suited to study the interaction of the SN ejecta with the surrounding CSM, which should be dominated by the companion's wind. However, no SN Ia has ever been detected in either regime. In this Letter, we present X-ray and UV data for a sample of eight SNe~Ia observed with \\S\\ between 5 and 132 days after outburst. The dates of outburst were estimated to be $18\\pm2$ days before the peak in the $B$-band. The highlight of this study is the tentative detection of SN~2005ke in X-rays with \\S, which was discovered on 2005-11-13.33~UT with the Katzman Automatic Imaging Telescope (KAIT; Baek, Prased \\& Li 2005) and later classified as an under-luminous SN Ia from the presence of the characteristic 420~nm Ti~II and 635~nm Si~II lines (Patat \\& Baade 2005). ", "conclusions": "" }, "0607/astro-ph0607416_arXiv.txt": { "abstract": "We present infrared images and spectra of comets 2P/Encke, 67P/Churyumov-Gerasimenko, and C/2001~HT50 (LINEAR-NEAT) as part of a larger program to observe comets inside of 5~AU from the sun with the \\textit{Spitzer Space Telescope}. The nucleus of comet 2P/Encke was observed at two vastly different phase angles (20\\degr{} and 63\\degr). Model fits to the spectral energy distributions of the nucleus suggest comet Encke's infrared beaming parameter derived from the near-Earth asteroid thermal model may have a phase angle dependence. The observed emission from comet Encke's dust coma is best-modeled using predominately amorphous carbon grains with a grain size distribution that peaks near 0.4~\\micron{}, and the silicate contribution by mass to the sub-micron dust coma is constrained to $<31$\\%. Comet 67P/Churyumov-Gerasimenko was observed with distinct coma emission in excess of a model nucleus at a heliocentric distance of 5.0~AU. The coma detection suggests that sublimation processes are still active or grains from recent activity remain near the nucleus. Comet C/2001~HT50 (LINEAR-NEAT) showed evidence for crystalline silicates in the spectrum obtained at 3.2~AU and we derive a silicate-to-carbon dust ratio of 0.6. The ratio is an order of magnitude lower than that derived for comets 9P/Tempel~1 during the \\textit{Deep Impact} encounter and C/1995~O1 (Hale-Bopp). ", "introduction": "Comets are frozen reservoirs of primitive solar dust grains and ices. Analysis of the composition and size distribution of cometary dust grains from infrared imaging and spectroscopic observations expedites an appraisal of the physical characteristics of the solid materials that constituted the primitive solar nebula \\citep{ahearn04, ehrenf04, wooden05}. The study of comets is an indirect probe of the origin of the constituents of the primitive solar system, their subsequent evolution into planetesimals, and their relationship to materials in other astrophysical environments \\citep{wooden05}. Although comets of all types have undergone some amount of post-formation processing, they remain the best preserved sources of material extant during our solar system's epoch of planet formation. In the current paradigm, the nearly isotropic comets (i.e., Oort cloud and Halley-type comets) formed amongst the giant planets and were scattered into large orbits, in a spherically symmetric manner \\citep{dones04}. The ecliptic comets (including Jupiter-family comets) originate from the Kuiper-belt and scattered disk populations and likely formed \\textit{in situ} or in the transneptunian region \\citep{duncan04, morbidelli04}. Comparisons between the nearly isotropic and ecliptic comets may reveal the differences in their post-formation processing or the structure and mineralogy of the proto-planetary disk. Both nearly isotropic and ecliptic comets have been exposed to bombardment by UV photons and cosmic rays, although to varying extents (the Jupiter-family comets have been exposed to $\\sim10^5$ times more UV and 100~KeV solar cosmic rays than the Oort cloud comets). The ecliptic comets have suffered frequent collisions during their residence in the Kuiper-belt and are likely to be fragments of larger Kuiper-belt bodies \\citep{stern03}. The number of comets studied by mid-infrared spectroscopic methods necessary to determine their detailed mineralogies is increasing \\citep[e.g., see][]{hanner96, harker06, harker05, harker02, harker99, honda04, kelley05b, lynch02, lynch00, sitko04, wooden04} and we may soon be able to compare comets to each other as groups, rather than individually. We present \\textit{Spitzer Space Telescope} \\citep{werner04} images and spectra of comets 2P/Encke, a 3.3~yr period ($P$), Jupiter-family comet with a perihelion distance, $q=0.3$~AU, known for an abundance of large dust particles \\citep{reach00} and an extensive debris trail \\citep{sykes92}; 67P/Churyumov-Gerasimenko (67P), $P=6.6$~yr, $q=1.3$~AU, a Jupiter-family comet and the primary mission target of the European Space Agency's \\textit{Rosetta} spacecraft; and C/2001~HT50 (LINEAR-NEAT) (HT50), a long period, Oort cloud comet, $P=40,250$~yr, $q=2.8$~AU. Comet Encke frequently approaches the Earth on it's 3.3~yr orbit and is one of the most studied of all comets \\citep{sekanina91}. Comet Encke was also one of the first comets discovered to have a dust trail \\citep{sykes92}. Dust trails are composed of large ($\\gtrsim0.1$~mm), slow moving particles and are precursors to meteor streams [Encke is associated with the Taurid meteor stream \\citep{whipple50}]. The comet also exhibits weak or non-existent 10~\\micron{} silicate emission \\citep{campins82, gehrz89, lisse04}. The existence of a dust trail, association with a meteor stream, and the lack of a strong silicate feature has led investigators to conclude that Encke's dust production is dominated by large particles. We present \\textit{Spitzer Space Telescope} observations of comet Encke in \\S\\ref{spectra-obs-text}, and derive the comet's dust coma mineralogy at 2.4~AU in \\S\\ref{encke-coma}. The mineralogy of comet Encke is discussed in \\S\\ref{mineral-discuss}. We derive the temperature and effective size of the nucleus of comet Encke in \\S\\ref{encke-results} and discuss the results in \\S\\ref{eta-discuss}. Comet 67P is the primary target of the \\textit{Rosetta} mission. The spacecraft is designed to intercept and orbit the comet at $r_h = 4.5$~AU (pre-perihelion) to study the development of coma activity as the comet approaches the sun\\footnote{\\url{http://www.esa.int/esaMI/Rosetta/}}. Information on the comet's dust environment is crucial to mission planning, which motivated our \\textit{Spitzer} observations of the comet at 5.0~AU (post-perihelion). The observation is presented in \\S\\ref{spectra-obs-text} and the results presented in \\S\\ref{cg-results}. Comet HT50 is an Oort cloud comet with an orbital period that suggests it has orbited the sun many times ($P=40,250$~yr). It was discovered to be cometary at the large heliocentric distance of 7.5~AU \\citep{pravdo01}. We observed comet HT50 twice with \\textit{Spitzer}. Both observations are presented in \\S\\ref{spectra-obs-text} and we derive dust mineralogies at both epochs in \\S\\ref{ht50-results}. We discuss HT50's mineral content and compare it to other Oort cloud comets in \\S\\ref{mineral-discuss}. ", "conclusions": "We present \\textit{Spitzer}/IRS spectra of comets 2P/Encke, 67P/Churyumov-Gerasimenko, and C/2001~HT50 (LINEAR-NEAT), and \\textit{Spitzer}/IRAC and MIPS images of comet 2P/Encke. Comet Encke exhibited a smooth continuum, best modeled by carbonaceous grains with a small peak grain size ($a_p=0.4$~\\micron). Previous investigations into comet Encke's dust coma revealed a weak silicate feature at perihelion ($r_h = 0.3$~AU). We conclude the weak silicate feature is due to the paucity of silicate grains and the preponderance of carbonaceous grains (or some other warm, deeply absorbing material). We constrain the sub-micron silicate fraction to $<31$\\% by mass. The nucleus of comet Encke is fit by the near-Earth asteroid thermal model with an effective radius $R = 2.34\\pm0.14$. The nucleus was observed at phase angles 20\\degr{} and 63\\degr{} and may be exhibiting a variation of the infrared beaming parameter with phase angle, which indicates of a rough nucleus surface or appreciable night side temperature. Comet 67P exhibited a significant coma at a heliocentric distance of 5~AU, $F_\\lambda = 2.01\\pm0.10 \\times 10^{-21}$~\\wcm{} at 27.9~\\micron{}. 67P's known dust trail comprises approximately 3\\% of the measured dust flux density. The remaining coma flux was due to 1) recently ejected dust (age of order hours to weeks), 2) large, slowly moving particles from the 2003 perihelion passage, or 3) some combination of the two. Comet HT50 displayed a significant silicate mineralogy with a silicate-to-carbon sub-micron mass ratio of 0.6. The derived ratio of 0.6 is an order of magnitude lower than the silicate-to-carbon ratios of post-\\textit{Deep Impact} comet 9P/Tempel~1 and other Oort cloud comets, C/1995~O1 (Hale-Bopp) and C/2001~Q4 (NEAT). The differences in silicate-to-carbon mass ratios in comet comae may be linked to strong jet activity in comets. Comet HT50's derived silicate-to-carbon sub-micron mass ratio is 0.6, but analysis of comet Hale-Bopp, which exhibited strong jet activity, derived a ratio of 8.1. At this time, the wide diversity in comet comae mineralogy likely has not been probed." }, "0607/astro-ph0607285_arXiv.txt": { "abstract": "It has been recognized that magnetic reconnection process is of great importance in high-energy astrophysics. We develop a new two-dimensional relativistic resistive magnetohydrodynamic (R$^2$MHD) code, and carry out numerical simulations of magnetic reconnection. We find that outflow velocity reaches Alfv\\'{e}n velocity in the inflow region, and that higher Alfv\\'{e}n velocity provides higher reconnection rate. We also find Lorentz contraction plays an important role in enhancement of reconnection rate. ", "introduction": "Magnetic reconnection is widely recognized as a very important phenomenon in astrophysics. Over the last decade, it has been recognized that magnetic reconnection processes are very important in high-energy astrophysics. Dissipation of such super strong magnetic fields may play an important role both in global dynamics of the system and as a way to produce high-energy emission. Relativistic magnetic reconnection was proposed as a source of the high-energy emission \\citep{1996A&A...311..172L,2002A&A...388L..29K} and as the solution to the $\\sigma$-problem \\citep{1990ApJ...349..538C,2001ApJ...547..437L,2003ApJ...591..366K,2003MNRAS.339..765L}. Similar models were also developed for the cosmological gamma-ray bursts \\citep{2002A&A...387..714D,2002A&A...391.1141D,2001MNRAS.321..177L}. Magnetic reconnection was evoked for explanation of the rapid variability observed in active galactic nuclei \\citep{1998MNRAS.299L..15D}. The particle acceleration in the reconnection process was proposed to operate in radio jets \\citep{1992A&A...262...26R,2001ApJ...559...96B}. Due to the extreme complexity and richness of the possible effects arising in relativistic plasma physics, there is a strong interest for developing computer codes for relativistic magnetohydrodynamics (hereafter RMHD). \\cite{1993JCoPh.105..339V} illustrated the implementation on the Riemann problem for MHD. \\cite{1996ApJ...463L..71K} then developed a RMHD code, which has been extensively used in relativistic two-dimensional and three-dimensional jet simulations. \\cite{1999MNRAS.303..343K} and others developed and tested a Godunov-type code which is a truly multidimensional scheme \\citep{2001ApJS..132...83B,2002MNRAS.333..932K}. Recently, \\cite{2003A&A...400..397D} presented a third order shock-capturing scheme for three-dimensional RMHD and validated it by several numerical tests. On the other hand, \\cite{1998ApJ...495L..63K,1999ApJ...522..727K} extended to general relativistic (GRMHD) effects, and applied it to the jet formation mechanism. \\cite{2003ApJ...589..444G} and \\cite{2003ApJ...589..458D} also developed GRMHD codes. Despite magnetic reconnection is recognized as an important process in high-energy astrophysics, there is not a lot of theoretical studies. \\cite{1994PhRvL..72..494B} considered kinematics of relativistic reconnection in the Sweet-Parker and Petschek configurations and concluded that due to the Lorentz contraction, the reconnection inflow is significantly enhanced and may approach the speed of light. \\cite{2003ApJ...589..893L} confirmed this conclusion for the Sweet-Parker case. \\cite{2005MNRAS.358..113L} presented generalization of Sweet-Parker and Petschek reconnection models to the relativistic case, and argued that the reconnection inflow does not approach the speed of light. Particle acceleration in relativistic current sheets was studied both in the test particle approximation \\citep{1992A&A...262...26R,2001ApJ...559...96B} and in two-dimensional PIC simulations \\citep{2001ApJ...562L..63Z,claus04}. Furthermore, \\cite{2005PhRvL..95i5001Z} studied three-dimensional PIC simulations, and suggested the importance of the current-aligned magnetic field for studying the energetics of relativistic current sheet. Meanwhile, there are several RMHD simulations as we write, all these codes, however, are applied to ideal MHD and take no account of resistivity. In this paper, we develop a new two-dimensional relativistic resistive MHD (R$^2$MHD) code, and carry out numerical simulations of two-dimensional relativistic magnetic reconnection. ", "conclusions": "Fig. \\ref{fig:2d} shows the density distribution of the typical case ($\\beta = 0.1$). Because of the enhanced resistivity around the origin, magnetic reconnection starts at this point. This point evolves to become an X-type neutral point. The reconnected field lines together with the frozen-in plasma are ejected from this X-point to the positive and negative $y$-directions because of the tension force of the reconnected field lines. The velocity of the reconnection outflow $V_{\\rm out}\\simeq 0.9$ is approximately the \\alfven speed of the inflow region ($C_{A0}=0.894$). To complement these outflows, inflows take place from positive and negative $x$-directions of the current sheet. At the boundary between this inflow and the outflow, a shock is formed, emanating from the neutral point. Fig. \\ref{fig:1d} shows one-dimensional plots of various physical variables at $t=100$, when the distribution becomes nearly steady state, and along $y=10$, which is well upstream of the plasmoid ejected in the positive $y$-direction. At $x\\sim \\pm 0.5$, there are strong jumps for several variables. The value of current density $j_z$ becomes large, and $y$-component of magnetic field $B_y$ becomes weak at these jumps. Therefore, we can say these jumps are the slow-mode MHD shocks. We also checked these jump conditions using the arranged model of \\citet[shown by dotted and dashed lines]{2005MNRAS.358..113L}. From Fig. \\ref{fig:2d} and Fig. \\ref{fig:1d}, we obtained $\\tan\\theta\\sim 0.21$ where $\\theta$ is the angle between the magnetic field and the shock plane, and this value is close to the inflow velocity at the slow shock (e.g., Fig. \\ref{fig:1d}(e)). According to the model of \\cite{2005MNRAS.358..113L}, inflow velocity $v_{\\rm in}\\sim \\tan\\theta$ in the highly relativistic regime, and our results supports this model. We next studied the dependency on the initial plasma $\\beta$. Fig. \\ref{fig:para} shows (a) inflow velocity at $x=4$ and $y=0$, (b) maximum inflow velocity, (c) outflow velocity, and (d) outflow 4-velocity as function of time for $\\beta =0.1$, 0.2, 0.5, and 1.0 ($C_{A0}=0.894$, 0.816, 0.667, and 0.535, respectively). Velocities are normalized by $C_{A0}$ in (a), (b) and (c), and time is normalized by \\alfven transit time $\\tau_A=L/C_{A0}$ in all figures. The maximum inflow velocity shown in (b) is almost the same as the inflow velocity at the edge of anomalous resistivity spot ($x\\approx\\pm 0.8$ and $y=0$). Each line shows the case of a different value of $\\beta$. The outflow velocity reaches $C_{A0}$ in all the cases. However, we obtain higher inflow velocity with lower $\\beta$ (higher $C_{A0}$). This means higher $C_{A0}$ causes higher reconnection rate $v_{\\rm in}/C_{A0}$. In other words, reconnection rate is higher at the relativistic regime. For a steady state reconnection, we can also express the reconnection rate by using the conservation of the mass at the steady state, $\\nabla\\cdot (D{\\Vect{v}})=\\nabla\\cdot(\\gamma\\rho{\\Vect{v}})=0$, so that we obtain a following equation: \\begin{equation} \\fr{v_{\\rm in}}{v_{\\rm out}}\\approx\\fr{\\delta}{d} \\fr{\\rho_{\\rm out}}{\\rho_{\\rm in}}\\gamma_{\\rm out} \\label{eq:rate} \\end{equation} where $\\rho_{\\rm out}$ and $\\rho_{\\rm in}$ are proper density of the inflow and the outflow region, and $\\gamma_{\\rm out}$ is the Lorentz factor of the outflow velocity, respectively. $\\delta$ and $d$ are evaluated by $\\delta/d=\\tan\\theta$, where $\\theta$ is the angle between the $y$-axis and the slow-shock plane. We consider the Lorentz factor of the inflow velocity $\\gamma_{\\rm in}\\sim 1$. Fig. \\ref{fig:ratio} shows the dependency of $\\delta /d$, $\\rho_{\\rm out}/\\rho_{\\rm in}$, $(\\delta /d)(\\rho_{\\rm out}/\\rho_{\\rm in})$, and $\\gamma_{\\rm out}$ to initial plasma $\\beta$ under the relativistic regime ($P_0 = 1.0$), and the non-relativistic regime ($P_0 = 10^{-2}$). From these panels, we can see the similar behaviors of ratios $\\delta /d$ and $\\rho_{\\rm out}/\\rho_{\\rm in}$ in the both regimes. Furthermore, $(\\delta /d)(\\rho_{\\rm out}/\\rho_{\\rm in})\\sim 0.11 - 0.14$ under the relativistic regime, while $(\\delta /d)(\\rho_{\\rm out}/\\rho_{\\rm in})\\sim 0.14 - 0.18$ under the non-relativistic regime. Therefore the product of the two terms are roughly constant in both regimes, and we can say that the effect of the Lorentz factor, namely, the Lorentz contraction is an important factor to determine the reconnection rate under the relativistic regime. Let us summarize this paper. The motivation of this study is to investigate relativistic effects of magnetic reconnection to apply for high-energy phenomena. For this purpose, what we have done were; (i) to develop a new resistive relativistic MHD code, and (ii) to do numerical simulations for relativistic magnetic reconnection. From our study, we obtain that outflow velocity become close to the light speed, and due to the high inflow velocity, high reconnection rate is obtained. For the enhancement of the reconnection rate, we find that the effect of the Lorentz contraction is significant which is suggested by \\cite{1994PhRvL..72..494B}. However, our results also supports the suggestions of \\cite{2005MNRAS.358..113L}, specially in the physics of jump conditions at the slow shocks. We recognize that simulations at the ultra-relativistic regime are required, so we would like to report these results in future." }, "0607/astro-ph0607550_arXiv.txt": { "abstract": "Using adaptive optics at the Gemini North telescope we have obtained a K-band spectrum of the star near the center of the luminous Galactic center bowshock IRS~8, as well as a spectrum of the bowshock itself. The stellar spectrum contains emission and absorption lines characteristic of an O5-O6 giant or supergiant. The wind from such a star is fully capable of producing the observed bowshock. However, both the early spectral type and the apparently young age of the star, if it is single, mark it as unique among hot stars within one parsec of the center. ", "introduction": "The nature of the Galactic center source IRS~8 \\citep{bec75}, one of the brightest compact mid-infrared sources in the central infrared cluster, was unknown until adaptive optics H- and K-band imaging revealed that the bulk of its infrared emission originates in a classic bowshock \\citep{rig03,geb04}. \\citet{geb04} showed that the IRS~8 bowshock is a straightforward consequence of the interaction of a dense and high velocity wind from a hot star that is traversing moderately dense interstellar gas. Adaptive optics imaging on large telescopes easily resolves the central star of IRS~8 (hereafter IRS~8*) from the bowshock. That stars undergoing mass loss within the interstellar medium in the Galactic center produce bowshock-like structures of swept-up gas had already been suggested from observations of IRS~21 by \\citet{tan02} and, by analogy, for a number of other luminous mid-infrared sources in the Northern Arm. This has now been observationally verified \\citep{tan05} and additional examples of the phenomenon have been found \\citep{gen03,cle04a,cle04b}. Thus IRS~8 is the most graphic example of a common phenomenon in the Galactic center that may also include a large number of lower luminosity sources with less massive winds \\citep[e.g.,][]{eck04}. Because of its large angular dimensions, fortuitous orientation, and relatively isolated location, spectroscopy of IRS~8 at high angular resolution can provide information both on the nature of the stellar source of the wind and the properties of the wind and the interstellar medium in the Galactic center. Here we describe exploratory low resolution K-band spectroscopy of IRS~8* and its bowshock. ", "conclusions": "\\subsection{Radial velocities} At the location of IRS~8 \\citet{lac91} identified two velocity components in the Ne~II line at 12.8~$\\mu$m, one at -10~km~s$^{-1}$ which is localized at IRS~8 and the other centered near +110~km~s$^{-1}$ associated with the Northern Arm. The signal-to-noise ratios of the lines in the IRS~8* spectrum are not high and the resolving power is low; hence the radial velocities of the lines cannot be determined accurately. Nevertheless, we can draw some limited conclusions from our model fits (see section 3.3) to the observed spectrum. Fits to the 2.07-2.12~$\\mu$m complex of lines gives -10 ~km~s$^{-1}$. Fits to the He absorption lines give +30 ~km~s$^{-1}$. The mean value of these is +10~km~s$^{-1}$. The strong Br~$\\gamma$ line in the bowshock is centered at -40~km~s$^{-1}$. All of these values match the radial velocity of the more negative component of the Ne~II line to within the uncertainties. We thus conclude that the blueshifted Ne~II component seen by \\citet{lac91} is associated with IRS~8* and suspect that it arises in gas that has been swept up by the wind from IRS~8*. Apparently, the Northern Arm is not interacting with IRS~8 and thus IRS~8* lies either well in front of it or behind it. \\subsection{Classification of IRS~8*} The relatively small equivalent widths of the emission lines in the spectrum of IRS~8* indicate that the star is an OB type rather than a Wolf-Rayet type. A very few WC9 stars with weak 2~$\\mu$m lines have been found by \\citet{fig97}; however, the lines are considerably broader than observed here. The isolated C~IV lines in IRS~8* have widths that are only marginally broader than the velocity resolution of 350~km~s$^{-1}$, indicating that the intrinsic full widths at half maximum (FWHMs) of the lines are not larger than 200~km~s$^{-1}$, consistent with rotational velocities found for O stars with lines originating at the photosphere. The presence of the two helium lines in absorption suggests that the star is either an Of or WNL type \\citep{han96,fig97}, but the latter classification is unrealistic as the carbon lines at 2.07~$\\mu$m and 2.08~$\\mu$m are prominent. Thus we are confident that we are observing an O star. \\citet{han96}, and \\citet{han05} have developed an infrared classification scheme for OB stars, based on spectra in the H window and short wavelength half of the K window, which is useful for hot stars that suffer large extinctions. Figure~3 is a comparison of the IRS~8* K-band spectrum with online-available K-band spectra from the \\citet{han96} catalog for O stars ranging from O4 to O6.5 and different luminosity classes. The resolving powers for all template spectra have been degraded to 800 for direct comparison with the observed spectrum. From Fig.~3 we judge that IRS~8* falls within the O5-O6.5 and III-If ranges, with likely O5-O6~If spectral type and luminosity class. For stars earlier than O5-O6If the C~III/N~III lines become much weaker while the strengths of the C~IV lines are considerably reduced for both earlier and later spectral types. In addition, in cooler stars He I usually begins to develop a noticeable absorption at 2.113~$\\mu$m (see \\citet{han96} for further objects with later spectral type), but this line is not detected in IRS~8*. A similar trend is observed in the He~I line at 2.059~$\\mu$m line, as its absorption strength clearly increases toward later types. Two spectral features may be used to determine the luminosity class: the emission feature at 2.116~$\\mu$m and Br~$\\gamma$. From Fig.~3 we see that for a given spectral type the emission strength of the 2.116~$\\mu$m feature is larger for supergiants than for dwarfs. Also emission in the strong stellar winds present in the supergiants starts to fill the Br~$\\gamma$ photospheric absorption profile to drive this line into emission as opposed to the clean absorption profile observed for giants and dwarfs. Further constraints on the spectral type of IRS~8* might come from measurements of the ionization state of the surrounding gas using mid-infrared fine structure lines of ions with a diagnostic range of ionization potentials. \\citet{lac80} found from observations of the Ne~II, Ar~III, and S~IV lines that the overall ionization state in the central parsec of the Galaxy is consistent with excitation by stars with $T_{eff}$~$\\le$~35,000~K. If that constraint applies to IRS~8 the earlier O spectral subtypes in the above range could be ruled out. The Ne~II line intensity at IRS~8 has been measured and is strong, but searches for the Ar~III and S~IV lines at IRS~8 have not been reported. \\subsection{Stellar parameters and abundances} Given the strong spectral similarities of IRS~8* with the O5-6 supergiants in Cyg~OB2 (see Fig.~3), we computed model fits covering that parameter domain, drawing from our analysis of the Cyg~OB2 stars for which UV, optical and IR spectra are available \\citep[][Najarro et al. in preparation]{h02}. To perform quantitative analysis we utilized the iterative, non-LTE line blanketing method presented by \\citet{hillier98} and proceeded as described in \\citet{naj04}, also taking into account the effects of the Fe~IV lines close to the wavelength of the He~I 2.06~$\\mu$m line \\citep{naj06}. The reader is referred to \\citet{hillier98} and \\citet{hillier99} for a detailed discussion of the code. Our best-fitting model is displayed in Fig.~3 (dashed line) and reproduces the observed K-band spectrum of IRS~8* quite well. Although the observed K band flux and low resolution spectrum are insufficient to tightly constrain all of the stellar properties of a supergiant such as IRS~8* (higher resolution and other H lines would be required), it is possible to obtain accurate estimates of some crucial parameters such as temperature and luminosity and useful constraints on the wind density and metal abundances. To determine the effective temperature we use He~I/II ionization balance as the main constraint via the He~I 2.06 and 2.112/3~$\\mu$m lines and the He~II~2.189~$\\mu$m line. The C~IV lines are considered as secondary T$_{eff}$ indicators (see below). The strengths of the absorption components of neutral helium lines display the highest sensitivity to changes in temperature, while the He~II 2.189~$\\mu$m line also shows strong sensitivity to wind density ($\\dot M$). Within the parameter domain of interest, for reasonable He enrichment (e.g. He/H $<$ 0.25 by number), these lines are not highly sensitive to changes in He abundance if the mass loss rate is also adjusted to reproduce Br~$\\gamma$. This may be understood as follows. When the He abundance increases, the He absorption lines become somewhat stronger. However, to recover the observed Br~$\\gamma$ strength $\\dot M$ must be increased to compensate for the reduction in the hydrogen abunance. The higher mass loss rate refills the He absorption components and, therefore, the resulting He profiles do not change significantly. This situation is reversed when the He abundance is high enough (He/H~$>$~0.25 by number) as both effects, enhancement of He abundance and increased $\\dot M$, drive the He lines into emission. From the above, we obtain T$_{eff}$=36000 $\\pm$ 2000, log~$L$/L$_{\\odot}$ = 5.6 $\\pm$ 0.2 (set by the derived T$_{eff}$ and the observed K magnitude), and 0.10~$<$~He/H~$<$0.25 by number. The error estimate for the effective temperature may seem small, but we are confident of the robustness of the derived value given the strong sensitivity of the He~I absorption components and the He~II line together with the presence of the C~IV lines. For temperatures above 38000~K the He~I lines absorption components disappear while they clearly become too strong for effective temperatures below 34000~K. Likewise, the He~II line starts to fade for temperatures below our quoted range and gets too strong in absorption for temperatures above it. Is also important to note that, for the relevant parameter domain, the He~I/II lines react only minorly to changes on the surface gravity. The error in the stellar luminosity reflects those in the effective temperature, extinction, contribution of the bow-shock spectrum and the uncertainty in the flux calibration. The Br~$\\gamma$ line is strongly sensitive to wind density, clumping, and velocity field, and to a lesser extent to gravity. Since we do not have diagnostics for the terminal velocity to constrain the wind density, we adopt a typical value, 2500~km~s$^{-1}$, found for other galactic O5If stars \\citep{h02}. Further, the spectral resolution is not high enough to clearly constrain either the shape of the velocity field \\citep[e.g., the $\\beta$ parameter][]{naj04} or the clumping factor ($f$). Thus, depending on the adopted $\\beta$ and He abundance we obtained clumping scaled mass loss rates $\\dot M / \\sqrt{f}$ in the range4.5--6.2~$\\times~10^{-6}$~M$_\\odot$~yr$^{-1}$. For our models we assumed a clumping factor of $f$=0.15, based on our experience of modeling the Cyg~OB2 objects, and hence the resulting mass-loss rates ranged from 1.75 to 2.45~$\\times~10^{-6}~$M$_\\odot$~yr$^{-1}$ for the adopted $\\beta$ and terminal velocity. Likewise, the lack of cleaner diagnostic lines for gravity such as the Brackett series lines in the H-band, hampers the determination of log~g. Models with gravities above log~g=3.70 started to fail to reproduce the spectrum. Reasonable fits were obtained for models with log~g=3.40--3.60. The C~IV lines were utilized as a secondary temperature diagnostic \\citep[see also][]{len04}, as they not only depend quite strongly on temperature, but also on wind density, carbon abundance, gravity and the $\\beta$ parameter. If none of the above parameters are precisely determined, the uncertainty in the carbon abundance may be as high as a factor of three. In the case of IRS~8* currently acceptable values for the carbon abundance range from 0.25$~\\times$~solar to 0.6~$\\times$~solar. The N~III doublet at 2.25~$\\mu$m reacts mainly to nitrogen abundance and only slightly to effective temperature within the parameter domain of interest and thus constitutes an excellent diagnostic of the nitrogen abundance \\citep[see also][]{naj04}. However, the weakness of this feature relative to the noise level results in a relatively high uncertainty in the lower limit to the nitrogen abundance. We estimate an enrichment of roughly 5~$\\times$~solar with 7.5~$\\times$~solar and 2.5~$\\times$~solar as reasonable upper and lower limits respectively. The strong emission feature at 2.116~$\\mu$m in IRS~8* is of particular interest. In the past this feature, which is present over a very wide range of O spectral types and luminosities \\citep{han96} has been attributed to C~III and N~III n=8--7 transitions. We tried to reproduce this feature with the carbon and nitrogen abundances derived from the C~IV 2.07-2.08~$\\mu$m and N~III 2.25~$\\mu$m lines and missed more than half of the observed equivalent width of the feature. Increasing either the carbon or the nitrogen abundance to match the 2.116~$\\mu$m feature resulted in far too strong lines of C~IV at 2.07-2.08~$\\mu$m and/or N~III at 2.25~$\\mu$m. This result is confirmed by inspection of the dominant ions of these species within the atmospheric region where the above lines form. Carbon is largely divide between C~IV and C~V, so the C~III lines are too weak even with a large increase in the carbon abundance. The dominant ions of nitrogen and oxygen are N~IV and O~IV and hence the N~III and O~III recombination lines are reliable diagnostics for deriving abundances. Hence, we now believe that the 2.116~$\\mu$m feature in IRS~8* is dominated by O~III n=8--7 transitions. After extending our O~III model atom to account for these transitions we could satisfactorily reproduce the observations (see Fig.~3). Further, the O~III component of the 2.116~$\\mu$m feature largely depends on the oxygen abundance and only slightly on gravity, effective temperature, wind density, and velocity field. Thus, this feature may be a powerful diagnostic of oxygen abundance, and therefore an important metal abundance determiner, over a wide range of O spectral types (Najarro et al. in preparation). Using it we obtain an oxygen abundance of 0.8 to 1.1$~\\times$~solar in IRS~8*, which indicates solar metallicity for the cloud in which IRS~8* formed. To summarize this subsection, the stellar parameters T$_{eff}$, L, $\\dot M / \\sqrt{f}$, and log g are fully compatible with OIf supergiants \\citep[e.g.][]{h02} and confirm the nature of IRS~8* derived from classification schemes. The derived N/C and N/O ratios reveal partial CNO processing at the surface of IRS~8 and are consistent with the values expected for an O supergiant if rotational mixing has taken place. \\citet{geb04} used the measured standoff distance of the bowshock from the star together with a rough estimate of the particle density ($n$~=~10$^{3}$~cm$^{-3}$) of the ambient interstellar medium in the Galactic center and typical values of windspeed (10$^{3}$~km~s$^{-1}$ and mass loss rate in a hot star (10$^{-6}$M$_\\odot$~yr$^{-1}$) to estimate a space velocity $v_*$ for IRS~8* of 150~km~s$^{-1}$, which is a reasonable value in the Galactic center. The standoff distance is proportional to ($\\dot{M}$~$v_w$~/~$n$~~$v_*$$^{2}$)$^{0.5}$, where $\\dot{M}$ is the mass loss rate and $v_w$ is the wind speed. The values of $v_w$ and $\\dot{M}$ used in the above modelling suggest that in order to produce the observed standoff distance $n$ would need to approach 10$^{4}$~cm$^{-3}$, which would not be surprising. \\subsection{Evolutionary status and mass of IRS~8*: An outsider within the Galactic center context?} Recently, \\citet{paum06} have reported the spectroscopic identification of $\\sim$40 OB supergiants, giants and dwarfs in the central parsec of the galaxy. Interestingly they find no OB stars outside the inner 0.5~pc (radius) of the galaxy and the earliest spectral type in their OB sample lies around O8-9I. They derive a common age of $6~\\pm~2~$Myr for the cluster. Our analysis suggests that IRS~8*, although only 1~pc from the center, does not fit into this picture of the central cluster of hot stars. It is of much earlier spectral type than any of the stars classified by \\citet{paum06}. Currently it is the only known OB star outside the central 0.5~pc region of the cluster. Figure~4 shows the position of IRS~8* (solid cross) as estimated from our model fits in the HR diagram compared with different evolutionary scenarios. An estimate of the evolutionary status of IRS~8*, based on comparisons with evolutionary tracks of stars without rotation \\citep[e.g.][not displayed in Fig.~4]{madx} yields a star with a zero age main sequence (ZAMS) mass of 48~M$_\\odot$ and a current age of 2.8~Myr. Such a star would not show any processed material on its surface. This is clearly at odds both with the current estimate for the age of the Galactic center cluster and with the abundance pattern derived from our models. The situation improves when evolutionary models accounting for rotation (dashed-lines in Fig.~4) are considered \\citep{madx}. Then the current position in the HR diagram corresponds to a star with a ZAMS mass of 44.5~M$_\\odot$ and an age of 3.5~Myr. Such a star would show CNO-processed material on its surface. Except for the age, still well below the estimate obtained by \\citet{paum06} using non-rotating models for a single burst scenario, the stellar parameters, including the abundance pattern, are fully consistent with those derived from our modelling. The crucial question thus is whether this star is really much younger than the cluster and probes the existence of ongoing (or at least much more recent) star formation, or if on the contrary the star is either an impostor or a cluster member that underwent a rejuvenation cure. A possible way out is provided if the star originally was a member of a massive close binary system. In such a case, we could be looking now at the secondary star, with the primary either exploded as supernova or in an evolutionary phase when it is much dimmer at $K$ than the secondary. Models for massive close binaries have been developed by \\citet{well99} to explain the optical counterparts of massive X-ray binaries. Using these models we have found that for a massive close binary system with initial masses of 25~M$_\\odot$ and 24~M$_\\odot$ (their model 10a) the current position of IRS~8* may be elegantly explained without violating the age of the Galactic center cluster. Similar scenarios are a possible explanation for some of the overluminous He~I objects in the central parsec. The solid lines in Fig.~4 correspond to the evolution of the primary (grey) and secondary (black) components of the massive close binary system. The thick solid line displayed on the track of the secondary corresponds to the phase where the primary is at least 10 times less bright in $K$ than the secondary. The observed location of IRS~8* in the HR diagram is reached after 7.1~Myrs, which is consistent with the $6~\\pm~2$~Myr estimate from \\citet{paum06}. Furthermore, at this stage the surface enrichment displayed by the secondary shows excellent agreement with the values derived in our model. For this particular massive close binary model the primary has not exploded yet, as otherwise the secondary would be already halfway through core helium burning and would have a much lower effective temperature. However, models with a slightly more massive primary and a slightly less massive secondary (N. Langer, private communication) could also reproduce the elemental abundances and current position of IRS~8* in the HR diagram after the explosion of the primary (at $\\sim$6~Myr) and, in addition, could have provided a kick to the secondary, placing it at its current position outside the central 0.5~pc. The spectroscopic mass of IRS~8* is in the range 23--37~M$_\\odot$ (for log~g 3.40--3.60), although the range could be larger, bearing in mind the uncertainties in determining the stellar gravity, as discussed previously. From the evolutionary models of \\citep{madx}, the current mass of IRS~8* is 38~M$_\\odot$ for a single star evolution with rotation, and thus consistent with the highest values of spectroscopic mass. Likewise, in the close binary scenario, the current mass of IRS~8* (after consuming a significant fraction of its companion) is 36--40~M$_\\odot$. Future spectra at higher resolution than presented here should be able to test if IRS~8* is part of a close binary. If IRS~8* is single its origin is highly uncertain. It is then either an impostor in the central parsec or it is a loosely associated member of the central cluster of hot and massive stars. The current motion of IRS~8*, nearly directly away from the center (as judged by the orientation of the bowshock and low radial velocity), suggests that it passes close to Sgr~A* \\citep{geb04} and may once have been a member of central cluster. If so it might be that a small population of less evolved and less luminous mid-late O type stars with weak emission lines such as those of IRS~8* are still hidden within the cluster, which could challenge the current understanding of that cluster as having a common age \\citep{paum06}." }, "0607/astro-ph0607599_arXiv.txt": { "abstract": "We present observations of the H\\,91$\\alpha$ recombination line emission towards a sample of nine \\ionhy\\/ regions associated with 6.7-GHz methanol masers, and report arcsecond-scale emission around compact cores. We derive physical parameters for our sources, and find that although simple hydrostatic models of region evolution reproduce the observed region sizes, they significantly underestimate emission measures. We argue that these findings are consistent with young source ages in our sample, and can be explained by existence of density gradients in the ionised gas. ", "introduction": "\\label{sec:introduction} Ultracompact (UC) \\ionhy\\/ regions are pockets of ionised hydrogen that form around massive stars in the earliest stages of their evolution. Together with massive bipolar outflows, strong far infrared emission by dust, and the presence of molecular masers, they are indicative of massive star formation in its earliest stages. The simple model of \\ionhy\\/ region evolution \\cite{Spitzer78} does not explain many of their observed properties, such as the frequent occurrence of non-spherical morphologies and the lifetime paradox. This lifetime problem, first noted by Wood \\& Churchwell \\shortcite{WoodChurchwell89a}, is that the number of observed UC \\ionhy\\/ regions exceeds that predicted from their dynamical expansion timescales by two orders of magnitude, given the accepted massive star formation rate in the Galaxy. A number of modifications and enhancements to the basic model have been suggested, including the work of Dyson \\etal\\ \\shortcite{DysonEA95}, Hollenbach \\etal\\ \\shortcite{HollenbachEA94}, Tenorio-Tagle \\shortcite{Tenorio-Tagle79} and van~Buren \\etal\\ \\shortcite{vanBurenEA90}, Franco \\etal\\ \\shortcite{FrancoEA90}, Arthur \\& Lizano \\shortcite{Arthur97}, Keto~\\shortcite{Keto03}. The thermal \\cite{DePreeEA95a} and turbulent \\cite{XieEA96} pressure confinement models are appealing due to their dependence on the ambient conditions observed to commonly exist in molecular clouds. De~Pree \\etal\\/ \\shortcite{DePreeEA95a} suggested thermal pressure confinement as an explanation of the lifetime paradox. They noted that when Wood \\& Churchwell proposed the lifetime problem in 1989, the molecular medium surrounding the UC \\ionhy\\/ regions was thought to have temperatures $\\sim 25$~K and densities $\\sim 10^5$~cm$^{-3}$. More recent observations indicate T$\\sim 100$~K and n$\\sim 10^7$~cm$^{-3}$. The resulting 400$\\times$ increase in thermal pressure limits the expansion of the Str\\\"omgren sphere. A weakness of this model, noted by Xie \\etal\\/ \\shortcite{XieEA96}, is the exceedingly high emission measures that it predicts ($\\sim 10^{10}$~\\pccm), which are more than two orders of magnitude greater than the values typically observed. Hierarchical density and temperature structures are known to exist within star-forming regions, with hot cores embedded in larger, less dense molecular clumps which themselves are within still larger and less dense molecular clouds. The densities decrease by approximately an order of magnitude in going from core to clump and again from clump to cloud \\cite{CesaroniEA94}. In a seminal work, Franco et al. \\shortcite{FrancoEA90} showed that these density inhomogeneities are important for \\ionhy\\/ region evolution. That the hierarchical structure of molecular clouds plays an important role in \\ionhy\\/ evolution is supported by the marked similarity in the ionised and neutral gas density structures observed within molecular clouds \\cite{KimKoo96,KimKoo02,KooKim03}. The large thermal molecular line and recombination line widths observed towards many UC \\ionhy\\/ regions suggest that significant turbulent motions are present, probably of magnetic origin \\cite{GarciaSeguraFranco96}. This led Xie \\etal\\/ \\shortcite{XieEA96} to suggest that turbulent pressure is the dominant mechanism to restrict the expansion of an \\ionhy\\/ region. In contrast to thermal pressure confinement, the assumed densities are lower, resulting also in lower emission measures. During much of the expansion phase, however, the turbulent pressure is expected to play a lesser role than thermal pressure in confining the \\ionhy\\/ region. This is discussed briefly in section~\\ref{sec:turb}. Icke \\shortcite{Icke79} investigated the formation of \\ionhy\\/ regions in non-homogeneous media and was able to explain some non-spherical morphologies. Observational evidence obtained in the last decade, however, suggests that many \\ionhy\\/ regions have compact cores within diffuse, arcminute-scale extended emission \\cite{KurtzEA99,KimKoo01}. Other sources exhibit this to a smaller degree --- the so-called core-halo morphology; see Wood \\& Churchwell~\\shortcite{WoodChurchwell89a} and Kurtz \\etal\\/~\\shortcite{KurtzEA94}. A study of the compact and extended radio continuum emission and radio recombination lines (RRLs) from eight \\ionhy\\/ regions known to be associated with 6.7-GHz methanol masers has been undertaken. The RRL analysis is presented here, while details of the continuum observations can be found in Ellingsen, Shabala \\& Kurtz \\shortcite{EllingsenEA05} (hereafter~\\otherpaper). In section~\\ref{sec:observations} we briefly outline our observations. The results are presented in section~\\ref{sec:results}, and these are compared with a simple model in section~\\ref{sec:model}. A discussion of our findings is presented in section~\\ref{sec:discussion}. ", "conclusions": "We have detected arcsecond-scale emission around UC \\ionhy\\/ cores. Using region parameters derived from continuum and H\\,91$\\alpha$ recombination line data we show that although simple models of expansion in hydrostatic equilibrium reproduce the observed region sizes, their emission measures are significantly underestimated. This discrepancy can be explained by the presence of density gradients in the ionised gas, consistent with young source ages and observations of the diffuse emission." }, "0607/hep-th0607092_arXiv.txt": { "abstract": "We study classically unstable string type configurations and compute the renormalized vacuum polarization energies that arise from fermion fluctuations in a 2+1 dimensional analog of the standard model. We then search for a minimum of the total energy (classical plus vacuum polarization energies) by varying the profile functions that characterize the string. We find that typical string configurations bind numerous fermions and that populating these levels is beneficial to further decrease the total energy. Ultimately our goal is to explore the stabilization of string type configurations in the standard model through quantum effects. We compute the vacuum polarization energy within the phase shift formalism which identifies terms in the Born series for scattering data and Feynman diagrams. This approach allows us to implement standard renormalization conditions of perturbation theory and thus yields the unambiguous result for this non--perturbative contribution to the total energy. ", "introduction": "In the past decades the perturbative treatment of the electroweak standard model has proven to be a very powerful tool to describe the properties and interactions of elementary particles within a wide energy regime. On the other hand the role and even the existence of non--perturbative solutions in this model is still quite uncertain. While the standard model does not contain topological solitons, one can construct nontopological string solutions to the classical equations of motion, called $Z$-strings or electroweak strings \\cite{Vachaspati:1992fi,VachaspatiReview, Nambu}. In the absence of topological arguments, however, one is not guaranteed that these classical solutions actually correspond to true local minima of the full effective energy, or even the classical energy. Indeed, Naculich~\\cite{Naculich:1995cb} has shown that in the limit of weak coupling, fermion fluctuations destabilize the string solution. This analysis leaves open the possibility, however, that deformed string solutions could exist, which would be local minima of the full effective energy. Since fermions are tightly bound in the string background, one potential mechanism for restoring stability is for fermions to bind to the string, yielding a lower total energy than a corresponding density of free fermions. When including this effect, however, one must also take into account the shifts in the zero-point energies of all the unoccupied modes as well, computed consistently in a standard renormalization scheme. If stabilized by the fermion binding mechanism, electroweak strings could have significant cosmological consequences \\cite{Kibble,Hindmarsh:1994re}. A network of strings could contribute to the dark energy that is required to explain the recently observed cosmic acceleration~\\cite{Perlmutter:1998np,Riess:1998cb}. However, a complete dynamical description of that scenario is still missing~\\cite{Spergel:1996ai,McGraw:1997nx,Bucher:1998mh}. Also, as pointed out by Nambu~\\cite{Nambu}, Z--strings are expected to terminate in monopole--antimonopole pairs, which could give rise to a primordial magnetic field. Furthermore, a network of stable strings at the electroweak phase transition would provide a scenario for electroweak baryogenesis without requiring a first--order phase transition~\\cite{Brandenberger}. The strings provide out--of--equilibrium regions, and the core of the string has copious baryon number violation due to the suppressed Higgs condensate. They thus provide an alternative to the usual idea of bubble--nucleation baryogenesis, which requires a first order phase transition to go out of thermal equilibrium. From a theoretical point of view, the quantum properties of $Z$-strings have been connected to non-perturbative anomalies \\cite{Klinkhamer:2003hz}. Decoupling arguments \\cite{D'Hoker:1984ph,D'Hoker:1984pc} suggest that when a fermion's mass is made very large by increasing its Yukawa coupling, soliton configurations should appear in the low-energy spectrum to maintain cancellation of such anomalies. In these calculations, the analysis of the fermion determinant -- whose logarithm is the sum over zero-point energies -- is essential. A first attempt at a full calculation of the quantum corrections to the $Z$-string energy was carried out in \\cite{Groves:1999ks}. Those authors were only able to compare the energies of two string configurations, rather than comparing a single string configuration to the vacuum. Furthermore, they used a proper time prescription, in which it was not possible to explicitly split off the divergent parts from the fermion determinant. Thus the final answer was expressed at large but finite cut--off in terms of two large terms that could only be computed numerically, representing the fermion determinant and the counterterms. In this formulation, no numerically stable results could be extracted for the limit ``cut--off to infinity.'' In the present work, we will employ phase shift techniques \\cite{Leipzig}, which allow us to make this separation cleanly and unambiguously. In this paper we restrict ourselves to the case of $2+1$ dimensions. We make this restriction to simplify the renormalization procedure. However, from the experience gathered in the case of QED magnetic flux tubes~\\cite{Graham:2004jb}, it is reasonable to expect that the qualitative behavior might be very similar to the actual 3+1 dimensional theory. The techniques of \\cite{GrahamInterface} allow for a straightforward generalization of the calculational procedure used here to that case as well. We also consider only the weak interactions, neglecting electromagnetism. Although introducing electromagnetism into our calculation is not entirely straightforward because of the effects of Aharonov-Bohm phases discussed in ref.~\\cite{Graham:2004jb}, we expect the approach of that paper to provide a natural generalization to the full electroweak interactions. In particular, the symmetry operator we rely on for the partial wave decomposition of our phase shifts continues to hold in the full theory. The paper is organized as follows. In Section II we will consider the Higgs--gauge sector model, consisting of a doublet Higgs coupled to an $SU(2)$ gauge field, and discuss various aspects of classical string configurations. In Section III we discuss the coupling of the fermions to string configurations and compute the energy that arises from summing the fermion zero modes. We will present our numerical results for the corrected string energy in Section IV. Some preliminary results of the current investigation were reported in conference proceedings~\\cite{Schroeder:2006hk}. ", "conclusions": "In this paper we have studied the quantum energies of static and localized $W$--string configurations in a $D=2+1$ dimensional gauge theory. This is to be understood as a precursor to a full investigation of quantum energies of $W/Z$--string configurations in the standard model in $D=3+1$. As we know from studies on QED flux tubes, the limitation to $D=2+1$ dimensions is a reliable truncation for string--type configurations when the $D=3+1$ renormalization conditions are imposed and the vacuum expectation value of the Higgs field is suitably scaled. We have concentrated on the fermion contributions to the vacuum polarization energy as a first calculation that comprehensively compares that energy to the one associated with the the trivial vacuum configuration, where the Higgs field sits at the minimum of its potential and the gauge fields vanish. This approach is motivated by the observation that there exists a fermion zero--mode for field configurations that correspond to the $W$--string with the charged Higgs field being identically zero. The population of this zero mode could yield a significant energy gain. The incorporation of contributions due to a single fermion mode requires us to also include the full fermion vacuum polarization energy because it is of identical order in both the loop and large $N_C$ expansions, where $N_C$ counts the degeneracy of the fermions, {\\it e.g.\\@} the number of color degrees of freedom. We compute the effective energy as the sum of the classical energy, the bound state contributions and the (renormalized) fermion vacuum polarization energy. In the combined limit, $\\hbar\\to0$ and $N_C\\to\\infty$ this approximation becomes exact at next to leading order. We have employed techniques that use scattering data for the computation of the vacuum polarization energy. These techniques have the important and outstanding feature that standard renormalization conditions, formulated in terms of Green's functions in momentum space, can be straightforwardly imposed. This is essential for a sensible comparison to the energy of the translationally invariant vacuum configuration and to generalize the $D=2+1$ results to $D=3+1$. To make efficient use of this powerful scattering theory approach, however, we have to assume that the fermions in the doublet are degenerate in order to obtain a partial wave expansion. Our numerical results show that this contribution to the total energy is small for configurations of interest, and does not destroy stability of the string. We have found energetically favored string configurations within an ansatz of variational parameters characterizing the $W$--string configuration. Although the true minimum likely lies outside our ansatz, its energy can only be less. Though we find that the fermion vacuum polarization energy reduces the effective energy, our numerical results show that the large classical energy carried by the pure $W$--configuration cannot easily be compensated by fermion quantum contributions with smaller classical energies. Rather, in the search for energetically stable configurations we are led to Higgs field dominated scenarios. In turn, this result reduces the importance of the zero mode. This is a significant result, one that runs contrary to our expectations and demonstrates that the search for stable configurations requires one to fully account for the fermion spectrum in the string background, rather than merely concentrating on just the zero mode. Furthermore, the change in character of the fields when going from the gauge field dominated to the Higgs field dominated scenarios is quite interesting. In the gauge field dominated case the fields have non--zero winding at infinity and magnetic flux. In the other case we are essentially left with only a shallow ring in coordinate space in which a hole is dug into the Higgs vacuum expectation value. In the regime of interest, where the extension of the background fields is a few times the Compton wave--length of the fluctuating fermion, the number of bound states essentially grows quadratically with the widths of the background fields. The hole in the Higgs vacuum expectation value causes the corresponding binding. Although these states are generally only weakly bound, they are so numerous that they can cause an object with large fermion number to indeed be bound. Though we did not observe such a configuration for empirically motived model parameters, we have seen that a doubling of the fermion degeneracy or an increase of the fermion mass provides sufficient emphasis on the fermion contribution to the total energy such that stable objects emerge. Typically these stable objects carry fermion numbers of about one hundred or even more and the configurations are quite wide, a few times the Compton wave--length of the free fermion. We have not taken the effort to determine a minimal value for the fermion mass for which a stable configuration emerges. However, for the empirical value of the fermion mass the critical value $N_C=7$ is not too far away from the physical datum. It thus does not seem totally absurd to imagine that in $D=3+1$ a stable object exists when adopting the standard model parameters. This certainly motivates the corresponding extension of the present study, but some technical obstacles must be overcome first. In $D=2+1$ it was sufficient to compute Feynman diagrams (and Born terms) up to second order in the external fields. The expansion with respect to external fields is gauge variant, but only gauge invariant combinations of Higgs and vector fields have well defined Fourier transforms that enter the Feynman diagrams. (The full effective energy is, of course, gauge invariant as the analogous peculiarity is contained in the Born series for the scattering data.) Already in the present case we had to introduce a fake Higgs field to circumvent that problem. Once we go to higher order in the Feynman diagrams, additional manipulations of this kind will be needed. One possibility is to unwind the string configuration at some distant point in space by making the angle $\\xi_1$ in the parameterization, eq.~(\\ref{eq:SphaleronSquare}) space dependent, such that it vanishes at $r\\to\\infty$. Unlike in the case of the QED flux tubes the {\\it return flux} associated with 'unwinding' the strings will contribute to the energy even if the distant point is sent to infinity. Though that contribution is easy to estimate (and thus subtract) for the classical energy, it is not so for the vacuum polarization energy, which is a non--local functional of the background fields. We expect, however, that the artificial 'unwinding' part of the effective energy can be quantified by some approximate techniques such as the gradient expansion. Since the present study indicates that $\\xi_1\\to0$ for the energetically favored configuration, it is also worthwhile to explore pure Higgs background configurations in $D=3+1$. Studies in that direction are in progress." }, "0607/hep-ph0607229_arXiv.txt": { "abstract": "I investigate the possibility that the observed curvature perturbation is due to a massive vector field. To avoid generating a large scale anisotropy the vector field is not taken to be driving inflation. Instead it is assumed to become important after inflation when it may dominate the Universe and imprint its perturbation spectrum before its decay, as in the curvaton scenario. It is found that, to generate a scale invariant spectrum of perturbations, the mass-squared of the vector field has to be negative and comparable to the Hubble scale during inflation. After inflation the mass-squared must become positive so that the vector field engages into oscillations. It is shown that, such an oscillating vector field behaves as pressureless matter and does not lead to large scale anisotropy when it dominates the Universe. The possibility of realising this scenario in supergravity is also outlined. ", "introduction": "Observations of the curvature perturbation in the Universe strongly suggest that it is generated during inflation by the gravitational production of particles. The field, whose quantum fluctuations are responsible for the particle production is typically considered to be a scalar field; one of the many flat directions that are envisaged in theories beyond the standard model. Very little has ever been discussed about gravitational production of vector fields during inflation (see Refs.~\\cite{VI,Lidsey}). This is mostly because, to achieve particle production in a de-Sitter background, the field in question must be light enough for its Compton wavelength to extend beyond the horizon. However, a massless vector field is conformally invariant and, therefore, it does not couple to the inflating gravitational background, which means that it does not undergo particle production. Hence, vector field generation during inflation has been ignored. In this work I investigate the possibility that a vector field with non-zero mass undergoes indeed particle production during inflation. My motivation was originally the possibility that a small, albeit non-zero, mass may lead to something interesting. However, I have found that this is not a promising direction, as it is shown below. Nevertheless, I have discovered that a negative mass-squared comparable to the Hubble scale can indeed result to the desired scale-invariant superhorizon spectrum of perturbations. In contrast to previous work \\cite{VI,Lidsey} the vector field considered is not assigned to the task of driving inflation. This is so in order to avoid generating a large scale anisotropy, which is otherwise inevitable (see, however, Ref.~\\cite{triad}). The curvature perturbations are produced in the same spirit as in the curvaton scenario \\cite{curv}. Thus, it is assumed that the vector field is subdominant during inflation. Consequently, particle production gives rise to isocurvature perturbations, which turn adiabatic at some point after inflation if the vector field manages to dominate the Universe before its decay. What I find is that an oscillating massive vector field does not result in a large scale anisotropy even when it dominates the Universe. Therefore, a vector field can indeed realise the curvaton scenario. Throughout the paper I use natural units, where \\mbox{$c=\\hbar=1$}. The signature of the metric is (1,-1,-1,-1). ", "conclusions": "In summary we have shown that, in principle, a vector field can indeed be responsible for the observed curvature perturbation in the Universe. In order to be so the mass-squared of the vector field has to be \\mbox{$m^2\\approx -2H^2$} during inflation. The vector field must be subdominant during inflation to avoid generating a large-scale anisotropy. Hence, particle production generates an originally isocurvature perturbation, which can become adiabatic if, after inflation, the vector field dominates the Universe before its decay, in accordance to the curvaton scenario. Indeed, we have shown that, if after inflation the mass-squared of the vector field becomes positive, then the field begins oscillating. We have demonstrated that an oscillating vector field scales as pressureless matter with the Universe expansion and does not cause any large scale anisotropy even when it dominates the Universe. Hence, provided the mass of the vector field complies to the above requirements, the `vector curvaton' scenario can account successfully for the observations. Admittedly, the condition for the mass of the vector field during inflation is hard to achieve and, at best, amounts to a certain level of tunning. Still, using a `vector curvaton' may be additionally motivated by the fact that no scalar fields have been observed as yet in nature." }, "0607/astro-ph0607236_arXiv.txt": { "abstract": "{The Local Group galaxies constitute a fundamental step in the definition of cosmic distance scale. Therefore, obtaining accurate distance determinations to the galaxies in the Local Group, and notably to the Andromeda Galaxy (M31), is essential to determining the age and evolution of the Universe. With this ultimate goal in mind, we started a project to use eclipsing binaries as distance indicators to M31. Eclipsing binaries have been proved to yield direct and precise distances that are essentially assumption free. To do so, high-quality photometric and spectroscopic data are needed. As a first step in the project, broad band photometry (in Johnson $B$ and $V$) has been obtained in a region ($34^\\prime\\times 34^\\prime$) at the North--Eastern quadrant of the galaxy over 5 years. The data, containing more than 250 observations per filter, have been reduced by means of the so-called difference image analysis technique and the DAOPHOT program. A catalog with 236\\,238 objects with photometry in both $B$ and $V$ passbands has been obtained. The catalog is the deepest ($V<25.5$ mag) obtained so far in the studied region and contains 3\\,964 identified variable stars, with 437 eclipsing binaries and 416 Cepheids. The most suitable eclipsing binary candidates for distance determination have been selected according to their brightness and from the modelling of the obtained light curves. The resulting sample includes 24 targets with photometric errors around 0.01 mag. Detailed analysis (including spectroscopy) of some 5--10 of these eclipsing systems should result in a distance determination to M31 with a relative uncertainty of 2--3\\% and essentially free from systematic errors, thus representing the most accurate and reliable determination to date.} ", "introduction": "Extragalactic distance determinations depend largely on the calibration of several distance indicators (e.g., Cepheids, supernovae, star clusters, etc) on galaxies of the Local Group with known distances. The Large Magellanic Cloud (LMC) has traditionally been used as a f\\mbox{}irst rung for extragalactic distance determinations. However, even taking into account that some recent results seem to converge to an LMC distance modulus of $(m-M)_0=18.50\\pm0.02$ \\citep{A2005024}, its low metallicity and a possible line-of-sight extension have posed serious difficulties to the accurate calibration of several distance indicators. The Andromeda galaxy (M31), on the contrary, has a more simple geometry and a metallicity more similar to the Milky Way and other galaxies used for distance estimation \\citep[see, e.g.,][]{A2006005}. Therefore, although stars in M31 are about six magnitudes fainter than those in LMC, the particular characteristics of this galaxy make it a promising f\\mbox{}irst step of the cosmic distance scale \\citep{A2006008}. Given the importance of M31 as an anchor for the extragalactic distance scale, many studies have provided distance determinations to M31 using a wide range of methods. A comprehensive list of distance determinations to M31, with explicit errors, are shown in Table \\ref{distances}. As can be seen, the values listed are in the range $(m-M)_0=24.0-24.6$ mag. Although the weighted standard deviation is of $\\sim$4\\%, most of the distance determinations in Table \\ref{distances} rely on previous calibrations using stars in the Milky Way or the Magellanic Clouds. As a consequence of this, a large number of subsequent distance determinations, based only on new recalibrations, can be found in the literature. These are not included in Table \\ref{distances}. Therefore, a direct and precise distance determination to M31 is of central importance since this would permit the use of all the stellar populations in the galaxy as standard candles. \\begin{table}[!ht] \\caption{Distance determinations to M31 as presented in the references. Values resulting of posterior calibrations and distance moduli without extinction corrections are not included.} \\label{distances} \\centering \\begin{tabular}{l r@{$\\pm$}l r@{$\\pm$}l c} \\hline \\hline Method & \\multicolumn{2}{ c }{$(m-M)_0$} & \\multicolumn{2}{ c }{Distance} & Reference \\\\ & \\multicolumn{2}{ c }{[mag]} & \\multicolumn{2}{ c }{[kpc]} & \\\\ \\hline Cepheids & 24.20 & 0.14 & 690 & 40 & [1] \\\\ Tip of the RGB & 24.40 & 0.25 & 760 & 90 & [2] \\\\ Cepheids & 24.26 & 0.08 & 710 & 30 & [3] \\\\ RR Lyrae & 24.34 & 0.15 & 740 & 50 & [4] \\\\ Novae & 24.27 & 0.20 & 710 & 70 & [5] \\\\ Cepheids & 24.33 & 0.12 & 730 & 40 & [6] \\\\ Cepheids & 24.41 & 0.09 & 760 & 30 & [6] \\\\ Cepheids & 24.58 & 0.12 & 820 & 50 & [6] \\\\ Carbon--rich stars & 24.45 & 0.15 & 780 & 50 & [7] \\\\ Cepheids & 24.38 & 0.05 & 752 & 17 & [8] \\\\ Carbon--rich stars & 24.36 & 0.03 & 745 & 10 & [8] \\\\ Glob. Clus. Lum. Func. & 24.03 & 0.23 & 640 & 70 & [9] \\\\ Red Giant Branch & 24.47 & 0.07 & 780 & 30 & [10] \\\\ Red Clump & 24.47 & 0.06 & 780 & 20 & [11] \\\\ Red Giant Branch & 24.47 & 0.12 & 780 & 40 & [12] \\\\ Cepheids & 24.49 & 0.11 & 790 & 40 & [13] \\\\ RR Lyrae & 24.50 & 0.11 & 790 & 40 & [14] \\\\ Tip of the RGB & 24.47 & 0.07 & 785 & 25 & [15] \\\\ Eclipsing binary & 24.44 & 0.12 & 770 & 40 & [16] \\\\ \\hline Mean \\& std. deviation & 24.39 & 0.08 & 750 & 30 & \\\\ \\hline \\end{tabular}\\\\ \\flushleft [1]:~\\citet{A2004021}; [2]:~\\citet{Mould86}; [3]:~\\citet{A2004025}; [4]:~\\citet{Pritchet87}; [5]:~\\citet{Capaccioli89}; [6]:~\\citet{A2004026}; [7]:~\\citet{Richer90}; [8]:~\\citet{Brewer95}; [9]:~\\citet{Ostriker97}; [10]:~\\citet{Holland98}; [11]:~\\citet{Stanek98a}; [12]:~\\citet{Durrell01}; [13]:~\\citet{A2004006}; [14]:~\\citet{Brown04}; [15]:~\\citet{A2006007}; [16]:~\\citet{P2005002} \\end{table} Following the procedure already used in the LMC to obtain precise distance measurements to several eclipsing binary (EB) systems \\citep[see][ and references therein]{A2003004}, a project was started in 1999 to obtain a direct distance determination to M31 from EBs \\citep{A2003018,P2004001}. The methodology involves, at least, two types of observations: photometry, to obtain the light curve, and spectroscopy, to obtain the radial velocity curve for each component. By combining the two types of observations the individual masses and radii for each of the components can be obtained. The remaining needed parameters (luminosities, temperatures and line-of-sight extinction) can be obtained from the modelling of the spectral energy distribution \\citep[see, e.g.,][]{A2002007} or the modelling of spectral lines \\citep[][ hereafter Paper~I]{P2005002}. The combination of the analysis results yields an accurate determination of the distance to the EB system, and hence, to the host galaxy, which is the main objective of the project. But in addition to providing distance measurements, the resulting stellar physical properties can be used as powerful diagnostics for studying the structure and evolution of stars that have been born in a chemical environment different from that in the Milky Way. The process briefly described above requires high quality light curves to obtain precise fundamental properties. By the time this project started, the highest quality EB light curves were those obtained by the DIRECT group \\citep[see][ and references therein]{A2004003}, but the scatter was too large for a reliable determination of the elements. Therefore, we began a new photometric survey to obtain high quality EB light curves (see \\S\\ref{obs}). With the application of the difference image analysis algorithm in the data reduction process (see \\S\\ref{datared}), precise photometry was obtained not only for the EBs in the f\\mbox{}ield, but also for all the variable stars. Two photometric catalogs were compiled (see \\S\\ref{cats}). The reference catalog contains photometry, in both $B$ and $V$ passbands, of 236\\,238 stars from which 3\\,964 are also in the variable star catalog (containing 437 EBs and 416 Cepheids). For some of the best quality light curves, corresponding to the brightest EB systems -- 24 in total, -- a preliminary determination of the orbital and physical properties is presented (see \\S\\ref{ebs}). These systems can be considered potential targets for distance determination and, indeed, already one of them was used to derive the f\\mbox{}irst direct distance determination to M31 \\citepalias{P2005002}. ", "conclusions": "Deep and high-quality time-series photometry has been obtained for a region in the North--Eastern quadrant of M31. The present survey is the deepest photometric survey so far obtained in M31. The photometry has been checked for compatibility with that from other surveys and, therefore, our resulting catalog can be used as a reference to extend the time baseline of future stellar surveys in M31. The study of the variable star population has revealed about 4\\,000 variables. Further analysis is needed for most of them for a proper classif\\mbox{}ication but over 800 variable stars have already been identif\\mbox{}ied as EBs and Cepheids. The study of the Cepheid variables will be the subject of a forthcoming publication \\citep{P2006002}. In the current paper we have presented the analysis of the EB population, which has resulted in an extensive list of EBs suitable for accurate determinations of the components' physical properties and their distances. The catalog provided here constitutes an excellent masterlist to select systems for further analysis. For example, the resulting physical properties will be an important tool to study stellar evolution of massive stars in another galaxy and, therefore, in a completely independent chemical environment. But, in accordance with the main goal of the the present survey, the technique of determining accurate distances from EBs has been demonstrated in \\citetalias{P2005002} to harbor great potential. Full analysis of an additional 5--10 EBs from the sample provided here can result in a distance determination to M31 with a relative uncertainty of 2--3\\% and free from most systematic errors. This will represent the most accurate and reliable distance determination to this important Local Group galaxy." }, "0607/astro-ph0607486.txt": { "abstract": "In this second paper of the series, we pursue the analysis of the 180~ks \\xmm\\ campaign towards the young open cluster NGC~6231 and we focus on its rich OB star population. We present a literature-based census of the OB stars in the field of view with more than one hundred objects, among which 30\\% can be associated with an X-ray source. All the O-type stars are detected in the X-ray domain as soft and reasonably strong emitters. In the 0.5-10.0~keV band, their X-ray luminosities scale with their bolometric luminosities as $\\log L_\\mathrm{X} - \\log L_\\mathrm{bol}=-6.912\\pm0.153$. Such a scaling law holds in the soft (0.5-1.0~keV) and intermediate (1.0-2.5~keV) bands but breaks down in the hard band. While the two colliding wind binaries in our sample clearly deviate from this scheme, the remaining O-type objects show a very limited dispersion (40\\% or 20\\% according to whether `cool' dwarfs are included or not), much smaller than that obtained from previous studies. At our detection threshold and with our sample, the sole identified mechanism that produces significant modulations in the O star X-ray emission is related to wind interaction. We thus propose that the intrinsic X-ray emission of non-peculiar O-type stars can be considered as constant for a given star. In addition, the level of X-ray emission is accurately related to the star luminosity or, equivalently, to its wind properties.\\\\ Among B-type stars, the detection rate is only about 25\\% in the sub-type range B0-B4 and remains mostly uniform throughout the different sub-populations while it drops significantly at later sub-types. The associated X-ray spectra are harder than those of O-type stars. Our analysis points towards the detected emission being associated with a physical (in a multiple system) PMS companion. However, we still observe a correlation between the bolometric luminosity of the B stars and the measured X-ray luminosity. The best fit power law in the 0.5-10.0~keV band yields $\\log L_\\mathrm{X} = 0.22(\\pm0.06) \\log L_\\mathrm{bol}+22.8(\\pm2.4)$. %The transition zone between the two behaviours and the extent of the canonical O star relation towards lower luminosities is poorly mapped by the present sample. It should however occur at about $L_\\mathrm{bol}\\approx10^{38}$~\\ergs\\ as previously suggested.% by \\citet{BSD97}. ", "introduction": "\\label{sect: intro} X-ray emission from early-type stars of spectral type O was, in December 1978, one of the earliest findings of the \\einst\\ satellite \\citep{HBG79, SFG79}. It was soon realized that all O-type stars were X-ray emitters. Most of them were characterized as soft (k$T<1$~keV) and reasonably strong ($10^{31} \\la L_\\mathrm{X} \\la 10^{33}$~\\ergs) sources. \\citet{HBG79} already suggested that the X-ray luminosity was directly linked to the characteristic luminosity of the emitter, though at the time the authors proposed a scaling law between \\lx\\ and the visual luminosity. However as the stars of their sample had mostly the same colors, it is equivalent to assert that \\lx\\ scales with \\lbol. Using various samples of stars observed with \\einst, different authors \\citep{LoW80, PGR81, CWS81, VCF81, SeC82} confirmed the so-called canonical relation $L_\\mathrm{X} \\approx 10^{-7}\\ L_\\mathrm{bol}$. Based on the \\einst\\ X-ray Observatory catalog of O-type stars \\citep{CHS89}, \\citet{SVH90} performed a more comprehensive study of the relationship between the optical and X-ray properties of O-type stars. They confirmed the existence of a canonical \\lxlbol\\ relation, though with a rather large dispersion. They were unable to find any significant correlation with the rotation rate ($v \\sin i$), the wind terminal velocity ($v_\\infty$) or the mass-loss rate ($\\dot{M}$), but observed a strong correlation with the wind momentum ($\\dot{M} v_\\infty)$ and with the wind luminosity ($0.5 \\dot{M} v^2_\\infty$). More recently, \\citet{BSD97} investigated the properties of the bright OB-type stars detected in the \\rosat\\ all-sky survey \\citep{BSC96} and found that the canonical relation extends down to spectral type B1-B1.5. They established the separation line between the O star relation and a less constrained relation for B stars to lie at a bolometric luminosity $L_\\mathrm{bol}\\approx10^{38}$~\\ergs.\\\\ Historically, two physically different models have been proposed to explain the hot O-type star X-ray emission. While no standard dynamo-driven surface magnetic field is theoretically expected for any star hotter than A7, \\citet{CaO79} suggested a scaled-up version of a solar-type coronal emission model that yielded a roughly correct prediction for the X-ray flux but could not explain the softness of the observed spectrum. Indeed in the coronal emission model, the X-rays are produced near the photosphere and are thus expected to suffer absorption by the overlying dense wind. Such an absorption is however not seen in the X-ray spectrum of hot stars. This suggests that the X-ray emission is rather produced throughout a significant fraction of the wind volume, which has been a strong argument in favor of the embedded wind-shock model. In this second model, the X-ray emission is supposed to arise from shocks occurring in the denser layers of the winds and that grow from small-scale instabilities of the line-driven winds. Since the phenomenological model of \\citet{Luc82}, hydrodynamical simulations \\citep[e.g.][]{OCR88,FPP97} have brought further support to this interpretation.\\\\ However, recent high resolution spectral observations of early-type stars are bringing the wind shock model to its limits. The unprecedented spectral resolution reached by the \\xmm\\ and \\chandra\\ observatories (corresponding to about 300~\\kms\\ in velocity space) now allows us to probe the widths and profiles of the X-ray emission lines seen in the spectra of hot stars. These contribute to put constraints on the velocity (Doppler broadening) and location \\citep[{\\it fir} line ratio, ][]{PMD01} of the emitting plasma, yielding thus an unprecedented characterization of its localization in the expanding winds. \\citet{KCO03} reported that the O-type supergiant $\\zeta$ Puppis (O4~Ief) displays broad, blue-shifted and asymmetric line profiles that are generally consistent with the wind-shock model. However, observations of other early-type stars suggest different pictures and hybrid magnetic wind models have been proposed. For example, the X-ray emission lines in the spectrum of $\\tau$~Scorpii (B0.2~V) are significantly narrower than expected from the standard wind-shock model \\citep{CdMMF03}. These authors rather suggested magnetically confined wind shocks \\citep{uDO02}, eventually coupled with the clump infall model \\citep{HCB00}, as the prime origin for the observed X-ray emission. A similar model was also successfully applied to the magnetic rotator $\\theta^1$~Ori~C \\citep[O5.5V,][]{GOC05}. Finally, \\citet{SCH03} have suggested that young massive stars could enter the main sequence carrying a significant residual magnetic field. Therefore, massive ZAMS stars could generate their X-ray luminosities via the standard model and magnetic confinement may provide an additional source of X-rays.\\\\ Compared to single stars of the same spectral type, close massive binaries are also known to display an extra X-ray emission \\citep{ChG91} which is generally attributed to a wind-wind collision. Within the interaction region, the shocked gas is expected to be heated to temperatures of a few $10^7$~K and to generate a substantial amount of X-rays, which are thus produced in addition to the intrinsic emission by each of the components. This extra-emission can be further modulated because e.g.\\ of a variation of the optical depth along the line of sight towards the wind interaction zone due to the orbital motion. It could also reflect the changing properties of the shocks due, for example, to a variation of the distance between the two stars in an eccentric binary system. \\\\ While the earliest works based on \\einst\\ data suggested the canonical relation to extend through the B-type range and down to A5 stars \\citep[e.g.][]{PGR81}, \\citet{SGH85} showed that the latter law does not hold for A-type stars. Later works \\citep{RGV85, CCM94} further suggested that the X-ray emission from B stars of spectral type B2 or later was not following the same scheme as O-type stars. Because these B-type stars do not have the convective zones required to sustain a magnetic dynamo, coronal emission is not expected. Their stellar winds are also much weaker than those of their hotter sisters, and are therefore not supposed to provide a sizable amount of X-rays. Indeed, \\citet{BSD97} reported that, for stars of spectral type B2 or later, the detection rate drops below 10\\%. Actually the intrinsic emission from B-type stars could, comparatively, be much lower than from O-type stars, with $\\log \\left( L_\\mathrm{X}/L_\\mathrm{bol} \\right) \\sim-8.5$ \\citep{CCM97}. This suggests that most of the detected X-ray emission for B stars in distant clusters is actually associated with unresolved companions, either in a binary system or located by coincidence on the same line of sight. The question of the intrinsic B-type X-ray emission has however not yet received a satisfactory answer. One of the difficulties is the intrinsic lack of homogeneity of the B-type population that contains different kinds of objects (\\bcep, Be stars, shell stars, ...). An additional difficulty is their lower emission level (if any), which thus limits the number of detections on shorter duration exposures or in distant fields. It was hoped that the advent of the `large' X-ray observatories, which combine improved sensitivity and spatial resolution, would help to solve this question. A scan of the recent literature \\citep[e.g.][and references therein]{SHH03} indeed favours the `companion' scenario but provides by no mean a definitive answer to this question.\\\\ The young open cluster \\ngc\\ \\citep[age $\\sim$ 3 to 5~Myr,][]{BVF99} is considered as the core of the \\sco\\ association and contains a large OB star population. Located at about 1.6~kpc ($DM=11.07\\pm0.04$, see discussion in \\citealt{SGR06}, hereafter Paper~I), it offers an excellent opportunity to probe a homogeneous sample of early-type stars in terms of e.g.\\ distance, age, reddening, environment and chemical composition. Our \\xmm\\ campaign has been described in \\citetalias{SGR06} and has revealed hundreds of point-like sources in the 15\\arcmin\\ radius field of view (FOV) of the satellite. In the present paper, we focus on the O- and B-type star population in the FOV. The analysis of the sources associated with optically faint counterparts is postponed to a forthcoming paper in this series (Sana et al. -- Paper III, in preparation). Preliminary results of this work were presented in \\citet{SNG06} but should be considered as supplanted by the present analysis. %Since the early 1990's, several studies have provided an estimation of the distance to \\ngc. Based on optical photometry, different authors obtained a distance modulus (DM) of %$11.56 \\pm 0.32$ \\citep{PHC91}, %$11.08 \\pm 0.05$ \\citep{BL95}, %$11.2 \\pm 0.4 $ \\citep{RCB97}, %$11.00 \\pm 0.07$ \\citep{SBL98} and %$11.50 \\pm 0.25$ \\citep{BVF99}. %These values are in good agreement with the estimated distance $\\mathrm{DM}=10.92\\pm0.16$ \\citep{SAR05} of the O+B eclipsing binary \\cpd7742 located in the cluster core. On the other hand, as already mentioned in \\citetalias{SGR06}, Hipparcos parallaxes of \\ngc\\ were known to be problematic with a negative mean value of $-0.8\\pm0.4$~mas \\citep{AL99}. Even the recent revision by \\citet{Ma03} ($1.7\\pm0.4$~mas, corresponding to $\\mathrm{DM}=8.9\\pm0.5$) is still far from the distance obtained from the other methods. However, Hipparcos parallaxes of distant massive stars have since long been known to be inaccurate \\citep[see e.g.][]{SKL04}. In the present study, we therefore adopt the weighted mean of the photometric values $DM=11.07\\pm0.04$ as the cluster reference distance. This paper is organised as follows. The next section provides details of the data handling. Sect.~\\ref{sect: ET} makes a census of the OB star population in the core of \\ngc\\ and identifies the early-type X-ray emitters. Sect.~\\ref{sect: Xray} investigates the X-ray properties of the detected O and B-type stars and the appropriate \\lxlbol\\ relations are derived in Sect.~\\ref{ssect: lxlbol}. Sect.~\\ref{sect: indiv} presents the properties of the individual early-type sources while Sect.~\\ref{sect: discuss} discusses the results of our study. Finally, Sect.~\\ref{sect: ccl} summaries our main results. %_____________ OBSERVATIONS & DATA REDUCTION ___________________________ ", "conclusions": "\\label{sect: ccl} In this second paper of the series, we have pursued the analysis of the X-ray data concerning the young open cluster \\ngc. While \\citetalias{SGR06} focused mainly on the detection and identification of the numerous X-ray sources in the \\xmm\\ FOV, this paper was devoted to the properties of the rich early-type star population. A detailed census of the OB-type stars within the FOV resulted in more than one hundred objects identified. Using a limited cross-correlation radius of 2\\farcs5, about one third of them could be associated with an X-ray counterpart. Among these, the 15 O-type stars/binaries are all detected in the X-rays as soft and usually bright sources characterized by \\mek\\ temperatures of k$T=0.3$ and 0.7~keV. On the other hand, the B-type star detection rate only amounts to about 20\\%. Compared to the O-type stars, the B-type stars have a similar low energy component but their second temperature is well above 1~keV. The B-type X-ray emitters are thus significantly harder than the O-type sources. O- and B-type emitters clearly present different behaviours in the $\\log L_\\mathrm{X} - \\log L_\\mathrm{bol}$ diagram, though both types draw up a linear relation in the log--log plane. The separation between the two sub-sets is located at about $\\log L_\\mathrm{bol}=38$ (\\ergs), as previously suggested by \\citet{BSD97}. The dispersion around the expected linear relation is apparently quite small. In the O-type star sample, the two objects that show the largest deviations are known to display an extra-emission component due to a wind interaction (\\hda, \\citealt{SSG04}; \\cpd7742, \\citealt{SAR05}). These were thus excluded from the subsequent discussion. We showed that, for the O-type stars, the X-ray luminosities are scaling with the bolometric luminosities. In the 0.5-10.0~keV energy range, we obtained: $$\\log L_\\mathrm{X}-\\log L_\\mathrm{bol}= -6.912\\pm0.153 .$$ We also found that a power law relation did not provide any significant improvement to the quality of the fit. The obtained dispersion around this new canonical relation is very limited. It becomes even smaller when excluding the `cooler' (i.e. with a spectral type later than O9) O dwarfs from the fit. In this case, the typical dispersion drops to about 0.087 in the \\lg\\ plane, thus corresponding to only 20\\% on the X-ray luminosities. Within our sample, the only identified mechanism that provides a significant deviation from this relation is extra-emission produced in a wind interaction region. It is also the sole mechanism that, at our detection threshold, produces a significant variability of the observed fluxes in our O-type star sample (see Sect.~\\ref{sect: indiv}). Though relatively limited, the present sample suggests thus that the intrinsic X-ray emission from O-type stars is very tightly correlated with their bolometric luminosity. Beyond the two strong CWB systems, our sample is formed by single stars and binaries belonging to different luminosity classes. Though not extending towards sub-types earlier than O6, they all seem to follow the canonical relation, suggesting thus a common mechanism for X-ray production in these objects. We also provide a new analysis of recent flux measurements obtained in the Carina region \\citep{ACMM03}. This sample is mainly formed by main-sequence stars ranging from O3 to O8.5. We note that this new analysis confirms much of our present conclusions. The best \\lxlbol\\ relation is indeed in the form of a scaling law rather than a power law. The dispersion around the obtained relation is very limited and the difference with our own relation might possibly be accounted for by the different energy ranges considered.\\\\ We emphasize that this apparent scaling might indirectly result from a scaling of the X-ray luminosity with the wind properties, themselves scaling with the bolometric luminosities for these stars with radiatively driven winds. Clues for this are provided by the work of \\citet{VdKL00, VdKL01} which has allowed us to naturally convert the scaling of the X-ray luminosity with the bolometric luminosity into a scaling with the wind parameters. Nonetheless the \\lxlbol\\ relation probably remains the most accurate observational constraint to link the intrinsic X-ray emission of the O-type stars with their fundamental properties. Turning to B-type stars, the fact that only about one quarter of the stars of a given spectral sub-type are actually associated with an X-ray source argues strongly against X-ray emission being an intrinsic property of these stars. We however note that we still observed a linear relation that links $\\log L_\\mathrm{X}$ and $\\log L_\\mathrm{bol}$: $$\\log L_\\mathrm{X} = (0.22\\pm0.06) \\log L_\\mathrm{bol} + 22.8 (\\pm 2.4) $$ The dispersion around this relation is quite limited ($\\sim$0.14) and the linear pattern is also seen in the different energy sub-ranges considered. From our analysis, the most probable explanation points towards the X-ray emission originating from a low mass PMS physical companion. The observed relation between \\lx\\ and \\lbol\\ remains however a puzzle and could eventually result from an observational effect. Alternatively it could be linked to some particularities of the B-type stars in \\ngc\\ or to a putative interaction between the PMS object and its B-type companion, yielding an X-ray emission partly governed by the intrinsic properties of the B-type primary. \\\\ Finally, we note that the separation line between the O- and ``B-type'' behaviours (around $L_\\mathrm{bol}=10^{38}$~\\ergs) is poorly mapped by our present sample. The adopted criterion to preserve the homogeneity of the spectral analysis (requiring at least 2-T models) has led to the rejection of the objects in the transition zone. The extent of the canonical relation towards lower luminosities probably deserves a more particular attention. Dedicated observations, combined with the already observed fields, could help to increase the number of objects in this zone. This is probably a necessary condition to answer with more details this still open question. %_____________ ACKNOWLEDGMENTS __________________________________________" }, "0607/astro-ph0607370_arXiv.txt": { "abstract": "Isotherms clustering in cosmic microwave background (CMB) has been studied using the 3-year WMAP data on cosmic microwave background radiation. It is shown that the isotherms clustering could be produced by the baryon-photon fluid turbulence in the last scattering surface. The Taylor-microscale Reynolds number of the turbulence is estimated directly from the CMB data as $Re_{\\lambda} \\sim 10^2$. ", "introduction": "Luminous matter in the observable universe is strongly clustered (see, for a review \\cite{p1}). Main reasons for this apparent clustering are gravitational forces and an initial non-uniformity of the baryon-photon matter just before the recombination time. An information about this non-uniformity we can infer from the Cosmic Microwave Background (CMB) radiation maps. Although these maps indicate nearly Gaussian distribution of the CMB fluctuations, certain clustering may be observed already in these maps even after they have monopole and dipole set to zero. The clustering of the isotherms in the last scattering surface then can be used as an initial condition for the gravitational simulations. One of the purposes of this paper is to describe the clustering quantitatively. Another related purpose is to find out underlying primordial physical processes, which may cause the isotherms clustering in the CMB maps. It is known \\cite{sbar},\\cite{dky} that rotational velocity perturbations in the primordial baryon-photon fluid can produce angular scale anisotropies in CMB radiation through the Doppler effect. The conclusions are relevant to arcminute scales (for large scales see \\cite{jaf}). On the other hand, it is recently shown in \\cite{sb1} that clustering is an intrinsic property of the turbulent (rotational) velocity field. Therefore, one should check whether the CMB isotherms clustering can be related to turbulent motion of the baryon-photon fluid just before recombination time (an observational indication of turbulent motion in the baryon-photon fluid just before recombination time is given in \\cite{bs1}, see also \\cite{gibson}-\\cite{kah}). Moreover, we will estimate the Taylor-microscale Reynolds number ($Re_{\\lambda}$, see Appendix A) of this motion, that is $Re_{\\lambda} \\sim 10^2$ (let us recall that critical value is $Re_{\\lambda} \\simeq 40$ ). As far as we know this is first estimate of the Reynolds number obtained directly from a CMB-map (we have used a cleaned 3-year WMAP \\cite{tag}). \\\\ To find physical (dynamical) origin of the isotherms clustering in the CMB we will compare certain statistical properties of the CMB maps (and their Gaussian simulations) with corresponding properties of the fluid turbulence. It is well known that energetically the CMB fluctuations are dominated by acoustic perturbations of the primordial velocity field. The velocity field can be represented by following way $$ {\\bf u} = {\\bf u_a} + {\\bf u_r} $$ where the acoustic component is a potential one ${\\bf u_a} = \\nabla \\varphi $ and ${\\bf u_r}$ is rotational component of the primordial velocity field. To exclude the acoustic component one can take operation $curl ~{\\bf u}= curl ~{\\bf u_r}$, that is vorticity. Average energy dissipation in turbulent velocity field is determined just by its rotational component \\cite{my} $$ \\langle \\varepsilon \\rangle \\sim \\nu \\langle (curl~ {\\bf u})^2 \\rangle $$ where $\\langle ... \\rangle$ means statistical average and $\\nu$ is kinematic viscosity. Since it is expected that high frequency events in the velocity field should provide the most significant contribution to the turbulent dissipation and, especially, to its high order moments one can also expect that clustering in turbulent velocity field should be intimately related to fluctuations of the energy dissipation (so-called {\\it intermittency} phenomenon \\cite{my},\\cite{sa}). Taking into account that turbulent velocity is nearly Gaussian \\cite{my} (see also Appendix B) the clustering phenomenon in the turbulent velocity field should be also a Gaussian one. On the other hand, the intermittency phenomenon in fluid turbulence, usually associated with fine non-Gaussian properties of the velocity field \\cite{my},\\cite{sa}. Thus, simultaneous consideration of the cluster and intermittency characteristics can provide an additional valuable information about physical origin of the isotherms clustering in the CMB maps and about the fine non-Gaussianity of the maps. Therefore, the paper starts with an introduction to this subject. ", "conclusions": "" }, "0607/astro-ph0607146_arXiv.txt": { "abstract": "Sensitive measurements of the interstellar gas-phase oxygen abundance have revealed a slight oxygen deficiency ($\\sim$ 15\\%) toward stars within 500 pc of the Sun as compared to more distant sightlines. Recent $FUSE$ observations of the interstellar gas-phase nitrogen abundance indicate larger variations, but no trends with distance were reported due to the significant measurement uncertainties for many sightlines. By considering only the highest quality ($\\geq$ 5 $\\sigma$) N/O abundance measurements, we find an intriguing trend in the interstellar N/O ratio with distance. Toward the seven stars within $\\sim$ 500 pc of the Sun, the weighted mean N/O ratio is 0.217 $\\pm$ 0.011, while for the six stars further away the weighted mean value (N/O = 0.142 $\\pm$ 0.008) is curiously consistent with the current Solar value (N/O = 0.138$^{+0.20}_{-0.18}$). It is difficult to imagine a scenario invoking environmental (e.g., dust depletion, ionization, etc.) variations alone that explains this abundance anomaly. Is the enhanced nitrogen abundance localized to the Solar neighborhood or evidence of a more widespread phenomenon? If it is localized, then recent infall of low metallicity gas in the Solar neighborhood may be the best explanation. Otherwise, the N/O variations may be best explained by large-scale differences in the interstellar mixing processes for AGB stars and Type II supernovae. ", "introduction": "The CNO group constitutes the most abundant elements in the Galaxy after hydrogen and helium. As such, these elements are heavily involved in the life cycles of stars of all masses and compositions (Wheeler, Sneden, \\& Truran 1989) and strongly influence the character of the interstellar medium through which these stars recycle material. Consequently, the current epoch abundances and abundance spreads of the CNO elements are important to establish accurately for studies of Galactic chemical evolution (Timmes, Woosley, \\& Weaver 1995) and for comparisons with the abundances of the young galaxies sampled by QSO absorption-line systems (Timmes, Lauroesch, \\& Truran 1995). Carbon and oxygen are produced during helium shell burning and returned to the ISM through Type II supernovae, while nitrogen is primarily formed during the CNO cycle and expelled into the ISM through the massive winds of Asymptotic Giant Branch (AGB) stars. The abundance of interstellar carbon and oxygen have been well studied (e.g., Cartledge et al. 2004; Sofia et al. 2004). However, due to the paucity of accurate interstellar nitrogen abundance measurements, the situation for interstellar nitrogen is not as secure. The abundance of interstellar oxygen (Meyer, Jura, \\& Cardelli 1998) is relatively constant, O/H$_{\\rm{tot}}$ = (3.43 $\\pm$ 0.15) $\\times$ 10$^{-4}$, where $N$(H$_{\\rm{tot}}$)=$N$(\\ion{H}{1}) + 2 $N$(H$_2$), with any variability less than the measured 1-$\\sigma$ uncertainties for lines of sight reaching $\\sim$ 1 Kpc. Meyer et al. (1998) found no significant variation of O/H with the fractional abundance of H$_2$ [$f$(H$_2$)], a measure of the physical conditions in the gas. The lack of a change in abundance with change in $f$(H$_2$) supports the notion that depletion onto interstellar grains is not a major sink of \\ion{O}{1}. Their results showed that \\ion{O}{1} can be used to trace the \\ion{H}{1} column density in diffuse and translucent clouds. More recent studies of O/H$_{\\rm{tot}}$ (e.g., Cartledge et al. 2004) find evidence for slight depletion effects ($\\sim$ 0.1 dex) for a few lines of sight and a weak trend of increased O/H$_{\\rm{tot}}$ ($\\sim$ 0.1 dex) for lines of sight with $d$ $\\geq$ 1 Kpc with relatively small scatter about the mean for all measurements. Using high-quality {\\it Hubble Space Telescope (HST)} data on the weak interstellar \\ion{N}{1} $\\lambda\\lambda$1159.837, 1160.917 doublet toward seven moderately reddened stars ($d$ $\\leq$ 500 pc), Meyer, Cardelli, \\& Sofia (1997) suggested that the interstellar nitrogen abundance is constant with N/H$_{\\rm{tot}}$ = (7.5 $\\pm$ 0.4) $\\times$ 10$^{-5}$. In an effort to expand on this work, Knauth et al. (2003) used moderate to high signal-to-noise (S/N) {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)} data toward an additional 17 stars ($d$ $\\leq$ 2.5 Kpc). Knauth et al. (2003) revealed a systematic trend of lower N/H$_{\\rm{tot}}$ with increasing $N$(H$_{\\rm{tot}}$) above 10$^{21}$ cm$^{-2}$, see their Figure 2. They also compared their \\ion{N}{1} measurements to the \\ion{O}{1} $\\lambda$1356 abundance toward some of the same lines of sight and found a similar trend of decreasing N/O for increasing $N$(H$_{\\rm{tot}}$), but with slightly smaller measurement uncertainties. Focusing on these weak, optically thin transitions enable more precise relative determinations of $N$(\\ion{N}{1}) and $N$(\\ion{O}{1}) (Meyer et al. 1997; Knauth et al. 2003; Cartledge et al. 2004) by minimizing saturation effects and eliminating the oscillator strength uncertainties associated with column densities determined from different transitions. Thus, the large scatter and potential deficiency seen in the \\ion{N}{1} abundance (e.g., Knauth et al. 2003) is quite perplexing. Could the scatter be due to measurement uncertainties, enhanced depletion or ionization of interstellar \\ion{N}{1}, or does it represent a real cosmic variance? Here, we re-examine the highest quality \\ion{N}{1} and \\ion{O}{1} data in the $FUSE$ and $HST$ data archive and from the literature (Meyer et al. 1997; Knauth et al. 2003). ", "conclusions": "One possible mechanism for cosmic variance arises due to differences in interstellar mixing processes. Nitrogen is primarily formed during the CNO cycle and expelled into the interstellar medium (ISM) through the winds of intermediate mass (4-8 M$_{\\odot}$) AGB stars (Henry et al. 2000), while oxygen is produced during helium shell burning and returned to the ISM through Type II supernovae. A larger scatter in the N/H or N/O ratio compared to O/H, would suggest that the products of AGB stars are enhanced in the Solar neighborhood and/or are not as well mixed compared to products of supernovae. It is generally acknowledged that interstellar mixing processes produce a homogeneous ISM on timescales of $\\geq$ 10$^8$ years (e.g., Moos et al. 2002; de Avillez 2000), which suggests that the nitrogen enhancement in the Solar neighborhood is relatively recent. Additional evidence for enhanced s-process elements in the Solar neighborhood comes from the possible enrichment of Sn with respect to the Solar abundance (Sofia, Meyer, \\& Cardelli 1999; Lauroesch et al. 2006, submitted). Additionally, recent work on the $^{85}$Rb/$^{87}$Rb isotope ratio toward $\\rho$~Oph A (Federman, Knauth, \\& Lambert 2004) suggests a dearth in r-process Rb in the Solar neighborhood but their findings could represent an enhancement in s-process Rb. These results all point to a nucleosynthetic origin for the larger N/O abundance due to contributions from AGB stars in the Solar neighborhood. However, Cartledge, Meyer, \\& Lauroesch (2003) find enhanced krypton abundances toward two distant stars HD~116852 \\& HD~152590, which they attribute to additional contributions from low and intermediate mass stars (e.g., AGB stars), hinting at possibly a more wide spread phenomenon. If localized to the Solar neighborhood, another mechanism could be the infall of low-metallicity gas previously invoked to explain the slight \\ion{O}{1} deficiency within a 500-800 pc radius of the Sun (Meyer et al. 1994; Cartledge et al. 2004). Comer\\'{o}n \\& Torra (1994) investigated the infall of a low metallicity 10$^6$ M$_{\\odot}$ cloud, which explain the dynamics of Gould's Belt. A recent study on the effect of infalling low metallicity gas on a galaxy (K\\\"{o}ppen \\& Hensler 2005) showed that such an event would have a propensity to form intermediate mass stars. The infalling material would initally decrease the N/O ratio in a localized portion of the galaxy and then after a time delay of $\\sim$10$^8$ years, the timescale needed for AGB stars to begin enriching the ISM, the N/O ratio would increase. Ultimately, the N/O ratio would return to its canonical value for the host galaxy. In this scenario, K\\\"{o}ppen \\& Hensler (2005) predict that N/O will be high for low O/H$_{\\rm{tot}}$ and low for high O/H$_{\\rm{tot}}$. The right panel of Figure~1 shows that for distances closer than 500 pc, the N/O ratio (solid squares) is higher at lower O/H$_{\\rm{tot}}$ than the N/O ratio at further distances (solid circles) indicating that local infall of low-metallicity gas may be the likely explanation (K\\\"{o}ppen \\& Hensler 2005). Both plots in Figure~1 provide the best evidence that the N/O variability is primarily due to differences in the nucleosynthetic origin of nitrogen and oxygen and not to observational uncertainties. However, it is important to note that the stellar abundance yields (van den Hoek \\& Groenewegen 1997) that K\\\"{o}ppen \\& Hensler (2005) utilized in their models of galactic chemical evolution do not include stellar rotation, which may alter their N/O abundance ratio predictions. Is the enhanced N/O ratio in the Solar neighborhood evidence for infall of low-metallicity gas or for large-scale differences in interstellar mixing processes between stellar winds of AGB stars or Type II supernovae? Only through further precise observations of the interstellar N/O ratio toward stars at variety of distances from the Sun will this question be resolved. If infall is the solution, then all high N/O ratios should be clustered within a $\\sim$ 500 pc radius of the Sun, otherwise further observations will show significant scatter in the N/O ratio at all distances. If verified, the infall of near-primordial gas in the vicinity of the Sun has important implications for a wide variety of astrophysical problems, including the enhanced D/H ratio within the Local Bubble (Hoopes et al. 2003; Moos et al. 2002; Wood et al. 2004)." }, "0607/astro-ph0607478_arXiv.txt": { "abstract": "We report on the search for optical bursts from J1819--1458, a member of the recently discovered Rotating Radio Transients (RRATs). J1819--1458 exhibits 3 millisecond bursts with a peak flux of $f_{\\nu}^{1.4 GHz}$ = 3.6 Jy every $\\sim 3.4$ minutes, implying that it is visible for only $\\sim 1$ second per day at radio wavelengths. Assuming that the optical light behaves in a similar manner, the most sensitive way of detecting RRATs is hence not to take long exposures of the field, but instead to capture individual bursts using a high-speed camera mounted on a large aperture telescope. Using ULTRACAM on the 4.2-m William Herschel Telescope (WHT) we obtained 97\\,100 images of the field of J1819--1458, each of 18.1 milliseconds exposure time and with essentially no dead-time between the frames. We find no evidence for bursts in $u'$, $g'$ and $i'$ at magnitudes brighter than 15.1, 17.4 and 16.6 (5$\\sigma$), corresponding to fluxes of less than 3.3, 0.4 and 0.8 mJy at 3560\\AA, 4820\\AA\\ and 7610\\AA, respectively. ", "introduction": "The RRATs are a remarkable new class of variable star characterized by their radio bursts of duration 2--30 milliseconds which recur every 4--180 minutes \\citep{mclaughlin06}. The RRATs, of which 11 are currently known, exhibit periodicities of 0.4--7 seconds, inferred by dividing the intervals between bursts by the largest common denominator. Such periods are long in comparison with most radio pulsars and are instead reminiscent of the periods found in the radio-quiet Anomalous X-ray Pulsars (AXPs), Soft Gamma Repeaters (SGRs) and X-ray Dim Isolated Neutron Stars (see \\citealt{woods06} and \\citealt{haberl04}). The distances to the RRATs can be estimated from their dispersion measures and it is found that they lie 2--7 kpc away in the Galactic plane. In the three RRATs with the most frequent bursts it has also been possible to measure a period derivative, showing no evidence for binarity but instead that these objects spin down like other pulsars. In the case of J1819--1458, the inferred magnetic field strength is high ($B\\sim 5 \\times 10^{13}$ G), providing another link between the RRATs and the magnetars (i.e. the AXPs and SGRs). The sporadic nature of the bursts in RRATs makes localization to better than the 14 arcminute beam of the Parkes Telescope difficult. Fortunately, the positions of the three RRATs with period derivatives can be refined through radio timing and were quoted to an accuracy of arcseconds by \\citet{mclaughlin06}. This enabled \\cite{reynolds06} to identify the {\\em Chandra} source CXOU J181934.1--145804, lying within 2 arcseconds\\footnote{It should be noted that this assumes updated radio coordinates which differ from the radio position listed in \\citealt{mclaughlin06} (Steve Reynolds, private communication).} of the radio source, as the X-ray counterpart to J1819--1458. The X-ray properties of the source, which is point-like and shows no variability, are consistent with thermal emission from a cooling neutron star \\citep{reynolds06}, lending further weight to the hypothesis that RRATs are rotating neutron stars. To further constrain the nature of RRATs, it is desirable to observe them at different wavelengths. \\cite{reynolds06} made a first attempt at this by searching optical and infrared archives for counterparts to CXOU J181934.1--145804. They found none, but their limits are not particularly deep ($I=17.5, J=15.6, H=15.0, K=14.0$). Taking longer exposures to go deeper is not necessarily the best solution, however, as the RRATs may have very faint persistent optical/IR emission and only emit strongly at these wavelengths during bursts\\footnote{\\cite{reynolds06} found no evidence for X-ray bursts in CXOU J181934.1--145804, but this does not mean optical bursts will be undetectable -- the AXP 4U\\,0142+61, for example, exhibits a pulsed fraction 5--7 times greater in the optical than the X-ray \\citep{dhillon05}. In addition, by analogy with the Crab pulsar, searching for optical bursts (as opposed to any persistent emission) may also be the most sensitive method of detecting RRATs -- the main optical pulse of the Crab is 5 magnitudes brighter in the B and V-bands compared to its persistent light level and is coincident (to within 100 $\\mu$s) with the radio peak \\citep{golden00}.}. Given that the duration of the bursts in J1819--1458, for example, only total $\\sim 1$ second a day, the best strategy would then be to reduce the contribution of the sky and take a continuous sequence of extremely short exposures on a large-aperture telescope covering a number of burst cycles in order to catch a burst in one or two of the frames. With the high-speed, triple-beam CCD camera ULTRACAM at our disposal \\citep{dhillon01}, we have the ideal tool with which to search for optical bursts from RRATs. In this paper, we report on an attempt to detect such bursts from the RRAT J1819--1458. We selected this RRAT as it exhibits both the most frequent and most powerful radio bursts -- a 3 millisecond burst of 3.6 Jy (at 1400 MHz) occurring every $\\sim 3.4$ minutes, with the burst intervals showing a periodicity of 4.26 seconds. Thanks to the X-ray identification by \\cite{reynolds06}, J1819--1458 is also the RRAT with the best-determined position, accurate to 0.5 arcseconds, and is one of the closer RRATs to the Earth (at a distance of 3.6 kpc). ", "conclusions": "We find no evidence for optical analogues to the radio bursts seen in the Rotating Radio Transient J1819--1458. Using a frame rate of $\\sim50$~Hz, ULTRACAM has enabled us to place 5$\\sigma$ limits on the burst magnitudes of 15.1, 17.4 and 16.6 in $u'$, $g'$ and $i'$, respectively. In comparison with the AXP 4U\\,0142+61, for example, which has magnitudes of $i'=23.7$, $g'=27.2$ and $u'>25.8$ \\citep{dhillon05}, our limits on J1819--1458 do not appear to be particularly deep. To place our limits in some context, therefore, it should be noted that if we had taken a single 1 hour exposure of the field with the WHT under identical conditions, and assuming the object emitted 18 bursts, each of $i'=16.6$ and 18.1 milliseconds duration, we would have detected the object at only $\\sim 0.7\\sigma$. Using the high-speed photometry technique described in this paper, on the other hand, we would have detected the source at 5$\\sigma$. The difference in sensitivity between the two techniques is due to the fact that the long exposures would be sky limited, whereas the data presented in this paper are readout-noise limited. Due to the nature of the bursts in RRATs, therefore, the only way we can significantly improve on the magnitude limits is to observe at higher frame rates (in order to reduce the small contribution of sky noise still further) and/or use a larger aperture telescope (in order to increase the number of counts detected from each burst). The discussion above assumes, of course, that the optical and radio light behave in a similar manner. If, however, the optical light has only a low (or no) pulsed fraction, then deep, long-exposure imaging might prove fruitful, as might searches for pulsed light on the proposed spin period of the neutron star (e.g. \\citealt{dhillon05})." }, "0607/astro-ph0607152_arXiv.txt": { "abstract": "It is now well established that many young brown dwarfs exhibit characteristics similar to classical T Tauri stars, including infrared excess from disks and emission lines related to accretion. Whether the same holds true for even lower mass objects, namely those near and below the Deuterium-burning limit, is an important question. Here we present optical spectra of six isolated planetary mass candidates in Chamaeleon II, Lupus I and Ophiuchus star-forming regions, recently identified by Allers and collaborators to harbor substantial mid-infrared excesses. Our spectra, from ESO's Very Large Telescope and New Technology Telescope, show that four of the targets have spectral types in the $\\sim$M9-L1 range, and three of those also exhibit H$\\alpha$. Their luminosities are consistent with masses of $\\sim$5-15 M$_{Jupiter}$ according to models of Chabrier, Baraffe and co-workers, thus placing these four objects among the lowest mass brown dwarfs known to be surrounded by circum-sub-stellar disks. Our findings bolster the idea that free-floating planetary mass objects could have infancies remarkably similar to those of Sun-like stars and suggest the intriguing possibility of planet formation around primaries whose masses are comparable to those of extra-solar giant planets. Another target appears to be a brown dwarf ($\\sim$M8) with prominent H$\\alpha$ emission, possibly arising from accretion. The sixth candidate is likely a background source, underlining the need for spectroscopic confirmation. ", "introduction": "Thanks to extensive studies conducted in the past five years, it is now clear that many (if not most) young brown dwarfs undergo a classical T Tauri phase similar to that of solar mass stars. Evidence for dusty disks around young sub-stellar objects come from infrared and millimeter excess (e.g., Natta et al. 2002; Jayawardhana et al. 2003; Scholz et al. 2006). Broad and asymmetric emission lines, such as H$\\alpha$, indicate on-going, and often variable, accretion (e.g., Jayawardhana, Mohanty \\& Basri 2003; Natta et al. 2004; Mohanty et al. 2005; Scholz et al. 2005). Forbidden line emission, thought to arise in outflows or winds, is also seen in some cases (Fern\\'andez \\& Comer\\'on 2001; Barrado y Navascu\\'es \\& Jayawardhana 2004; Whelan et al. 2005). These striking similarities between young low-mass stars and objects near and below the sub-stellar boundary are sometimes taken as evidence for a common formation scenario. Meanwhile, a few dozen free-floating objects with inferred masses near and below the Deuterium-burning limit of $\\sim$12 Jupiter masses (Chabrier et al. 2000a) have been identified in the $\\sigma$ Orionis region (Zapatero Osorio et al. 2000) and the Orion Nebula Cluster (Lucas and Roche 2000; Lucas, Roche and Tamura 2005). Low-resolution spectra have confirmed cluster membership and ultra low masses for some of them according to evolutionary models (Barrado y Navascu\\'es et al. 2001; Mart\\`in et al. 2001; Lucas et al. 2001). These `isolated planetary mass objects' (`IPMOs', `planemos') or `sub-brown dwarfs', as they are sometimes called, represent the bottom end of the stellar initial mass function (IMF). Thus, any successful theory of star formation must account for their origin as well. However, little is known about the properties of `planemos' because they are very faint at the distances of $\\sigma$ Ori and ONC ($\\sim$350-450 pc), and only a few more have been identified in closer star forming regions. In particular, it is extremely interesting to investigate whether they also harbor accretion disks, just like many of the higher mass young brown dwarfs. At least in a few cases, there is already evidence of mid-infrared excess (e.g., Testi et al. 2002; Luhman et al. 2005a, 2005b) and strong H$\\alpha$ emission from objects near and below the Deuterium-burning limit (Barrado y Navascu\\'es et al. 2001, 2002). Now, by combining ground-based optical and near-infrared photometry with {\\it Spitzer} ``Cores to Disks'' Legacy Survey data, Allers et al. (2006) have identified six new candidate planemos in three nearby regions, at distances $\\sim$150 pc. What's more, based on their infrared excess, these objects appear to be surrounded by dusty disks. However, spectroscopy is essential to determine their true nature, in particular whether they indeed have late spectral types, and thus relatively cool temperatures and very low masses. That is the goal of the present {\\it Letter}. ", "conclusions": "We find that five of the six candidates have spectral types in the late M to early L range, as would be expected for low mass brown dwarfs and planetary mass objects, while a sixth (\\#12) is probably a background source. Given that the spectral types we derive and the effective temperatures found by Allers et al. from photometry are consistent within the errors, these five objects are likely to be at the distances of the Chamaeleon II, Lupus I and Ophiuchus clouds. Their {\\it Spitzer}-detected mid-infrared excesses, modeled as emission from dusty disks by Allers et al., provide strong evidence of youth. Here it is interesting to note that the only object in their sample for which Allers and co-workers were not able to fit the spectral energy distribution (SED) well with a flat or flared disk model is \\#12, which we now find to be a likely background source. The large minimum H$\\alpha$ EWs we find in at least three of the five late-type objects bolsters the case for their youth. H$\\alpha$ could originate from disk accretion or chromospheric activity in young low-mass objects (e.g., Jayawardhana, Mohanty \\& Basri 2002), and it is often difficult to distinguish definitively between these two possibilities with low-resolution spectra. However, according to the criterion developed by Barrado y Navascu\\'es \\& Mart\\`in (2003), three of these objects, if not four, could well be accreting. The lack of clear H$\\alpha$ emission in \\#17 does not necessarily rule out youth; young L-type objects may not show H$\\alpha$ unless they are accreting (Barrado y Navascu\\'es et al. 2001, 2002). There are many uncertainties involved in deriving the masses of young very low mass objects, given the difficulties of determining their ages and distances reliably as well as the uncertainties in the spectral type to T$_{eff}$ conversion and in the evolutionary models themselves (e.g., Baraffe et al. 2002; Mohanty, Jayawardhana \\& Basri 2004). Allers et al. (2006) estimate masses for these candidates by matching their source luminosities to the widely used isochrones of Baraffe et al. (2001; 2003), and assuming ages of 1 Myr for Ophiuchus and Lupus I and 3 Myr for Chamaeleon II. They also point out that three of the objects we have now confirmed as late type (\\#01, \\#05, \\#17) have luminosities equivalent (within the errors) to the lowest luminosity young brown dwarf with mid-infrared excess reported previously (Luhman et al. 2005b). While we do not think an exhaustive analysis is warranted, given the uncertainties involved, we can at least check whether the spectral types we have derived are roughly consistent with the temperatures derived by Allers et al. from source SEDs. Thus, in Table 2 we report effective temperatures for our targets, based on the spectral type to T$_{eff}$ conversion scale of Mart\\`in et al. (1999), along with those from Allers et al. The comparison shows that the two sets of derived T$_{eff}$ indeed agree within $\\pm$100 K for four out of five objects. For the fifth, \\#18, we find a somewhat earlier spectral type, and a correspondingly higher T$_{eff}$, than implied by the Allers et al. analysis. Fig. 3 shows a Hertzsprung-Russell diagram of the five late-type objects. The absolute $K$ magnitude was derived from the apparent $K$ magnitude, the distance modulus and the visual absorption ($A_v$) given in Allers et al., and adopting the reddening law from Rieke \\& Lebofsky (1985). The effective temperatures ($T_{eff}$) are those reported in Table 2. The isochrones for 1 Myr and 5 Myr DUSTY00 models (Chabrier et al. 2000b; Baraffe et al. 2001) are also plotted. This figure shows that four of the candidates have masses near or below the deuterium-burning limit, while \\#18 is a somewhat higher mass brown dwarf. Given the slope of the isochrones around the locus occupied by these objects and the $\\approx \\pm$100 K uncertainty in $T_{eff}$, this figure suggests an uncertainty in the mass estimates on the order of a few Jupiters for four objects and $^{+7}_{-5}$ for the fifth, \\#18. The ages of all five objects are consistent, within the uncertainties, to those assumed by Allers et al. The targets \\#1, \\#5, \\#11 and \\#17 are among the lowest mass objects with dusty disks known to date, and at least three of the four also show possible signs of accretion from those disks. Whatever their exact masses are, these objects represent the bottom end of the stellar IMF. It is worth noting that Mohanty et al. (2006) have recently found evidence for an edge-on disk surrounding the planetary mass companion to the nearby young brown dwarf 2MASSW J1207334-393254, which itself is known to harbor a disk. Our findings, combined with previous work, suggest that some planetary mass objects have characteristics usually seen in T Tauri stars and higher mass brown dwarfs, implying strikingly similar infancies for our Sun and objects that are some hundred times less massive. Thus, a successful theory for star formation should be able to account for these similarities in young objects with a wide range of masses. The shape of the IMF at these lowest masses is not yet well defined observationally. Deep, wide-field optical and infrared surveys with 8-meter class telescopes are needed to investigate this regime with larger samples (e.g., Lucas, Roche \\& Tamura 2005). With current observing facilities, spectroscopic confirmation of very low mass candidates is challenging, but not impossible as we have demostrated here (also see Lucas et al. 2001; Barrado y Navascu\\'es et al. 2001). As Allers et al. have shown, {\\it Spitzer} could play a pivotal role in determining the disk frequency in the planetary mass regime." }, "0607/astro-ph0607364_arXiv.txt": { "abstract": "The spinning down (up) of a superfluid is associated with a radial motion of its quantized vortices. In the presence of pinning barriers against the motion of the vortices, a spin-down may be still realized through ``random unpinning'' and ``vortex motion,'' as two physically separate processes, as suggested recently. The spin-down rate of a pinned superfluid is calculated, in this framework, by directly solving the equation of motion applicable to only the unpinned moving vortices, at any given time. The results indicate that the pinned superfluid in the crust of a neutron star may as well spin down at the same steady-state rate as the rest of the star, through random unpinning events, while pinning conditions prevail and the superfluid rotational lag is smaller than the critical lag value. ", "introduction": "Spinning down (up) of a superfluid at a given rate is associated with a corresponding rate of outward (inward) radial motion of its quantized vortices. If the vortices are subject to pinning, as is observed in the experiments on superfluid Helium (\\markcite{hed80}Hedge \\& Glaberson 1980; \\markcite{sch81}Schwarz 1981; \\markcite{adam85} Adams, Cieplak \\& Glaberson 1985; \\markcite{Ziv02}Zieve \\& Donev 2000) and also assumed for the superfluid in the crust of a neutron star (pinned to the lattice nuclei) (\\markcite{T75}Tsakadze \\& Tsakadze 1975; \\markcite{TT}Tsakadze \\& Tsakadze 1980; \\markcite{A87} Alpar 1987; \\markcite{titi90}Tilley \\& Tilley 1990), a spin-down would require also unpinning of the vortices, in order to become moveable. Unpinning may be realized by the combined effects due to the Magnus effect, quantum tunnelling and/or thermal activation. However, the subsequent {\\em radial} motion of the unpinned vortices (before repinning) is a separate {\\em dynamical} process, subject to their equation of motion, apart from the unpinning process. This is a view different than that adopted in the model of ``vortex creep\" (\\markcite{alet84} Alpar et.~al. 1984; \\markcite{jonAp91}Jones 1991b; \\markcite{elb92}Epstein, Link \\& Baym 1992), which envisages the spin-down to occur through quantum tunnelling {\\it between} adjacent pinning sites, at different {\\em radial} distances. A critical discussion of the model of vortex creep, as well as further justification of the presently adopted viewpoint, may be found elsewhere (\\markcite{MJ05}Jahan-Miri 2005a;\\markcite{MJ05b}Jahan-Miri 2005b). The derivation of the spin-down rate of a superfluid, in presence of random unpinning, as discussed here, aims to pay due attention to the dynamical role of the vortex {\\em radial} motion. That is, vortex radial motion accompanies a transfer of the spin-down (-up) torque between the ``container'' and the bulk superfluid, which has to be necessarily meditated by the {\\em moving} (not the stationary {\\em pinned}) vortices, as in the absence of any pinning (\\markcite{son87}Sonin 1987; \\markcite{titi90}Tilley \\& Tilley 1990). Nevertheless, there exist uncertainties in the (micro)physics of vortex motion, as opposed to the structure of a vortex lattice, as well as in the theoretical understanding of the pinning/unpinning mechanisms. Such issues are beyond the scope of the present discussion, and are dealt with by making justified assumptions. The predicted general relation here reduces to that reported previously (\\markcite{MJ05a}Jahan-Miri 2005a), as an approximate limiting case. Moreover, the present calculation is based on a direct solution of the equation of motion for the (temporarily unpinned movable) vortices, in contrast to the heuristic arguments used in \\markcite{MJ05a}Jahan-Miri 2005a. ", "conclusions": "" }, "0607/astro-ph0607014_arXiv.txt": { "abstract": "We have created a general methodology for calculating the wavelength-dependent light curves of close-in extrasolar giant planets (EGPs) as they traverse their orbits. Focussing on the transiting EGPs HD189733b, TrES-1, and HD209458b, we calculate planet/star flux ratios during secondary eclipse and compare them with the {\\it Spitzer} data points obtained so far in the mid-infrared. We introduce a simple parametrization for the redistribution of heat to the planet's nightside, derive constraints on this parameter (P$_n$), and provide a general set of predictions for planet/star contrast ratios as a function of wavelength, model, and phase. Moreover, we calculate average dayside and nightside atmospheric temperature/pressure profiles for each transiting planet/P$_n$ pair with which existing and anticipated {\\it Spitzer} data can be used to probe the atmospheric thermal structure of severely irradiated EGPs. We find that the baseline models do a good job of fitting the current secondary eclipse dataset, but that the {\\it Spitzer} error bars are not yet small enough to discriminate cleanly between all the various possibilities. ", "introduction": "\\label{intro} Probing the atmospheres of extrasolar giant planets (EGPs) by measuring their spectra is the paramount means to determine their physical and chemical character (Burrows 2005). Such direct measurements complement the kinematic and orbital information obtained through the radial-velocity (RV) technique by which the vast majority of the EGPs have to date been discovered and studied. However, an EGP's spectrum and phase-dependent light curve can in principle reveal or constrain the molecular and atomic compositions, atmospheric temperatures, cloud properties, albedos in the optical, and the degree to which the heat absorbed on the dayside is redistributed to the nightside before reradiation. The advection of heat and material from the dayside by jet streams and zonal winds will alter the dayside atmospheric temperatures and non-equilibrium compositions (Menou et al. 2002; Cho et al. 2003; Burkert et al. 2005; Iro, B\\'ezard, \\& Guillot 2005; Showman \\& Guillot 2002; Guillot \\& Showman 2002), and might measurably shift the light curve with respect to the orbital ephemeris (Cooper \\& Showman 2005; Williams et al. 2006). Dynamic meteorology could also introduce zonal banding, as seen in Jupiter and Saturn, and temporal fluctuations, and does influence the rate with which heat is lost from the inner core (Burrows, Sudarsky, \\& Hubbard 2003; Burrows et al. 2004), and thereby the radius of the planet and its evolution. Moreover, such redistribution affects the near- and mid-infrared emissions from the night side, and as a result will affect the interpretation of nightside data when they become available. The planet/star flux ratios of wide-separation EGPs ($>0.2$ AU) are quite low ($10^{-4}$ to 10$^{-14}$) and vary widely as a function of wavelength and orbital separation (with the concomitant non-monotonic variations in geometric and Bond albedos) (Sudarsky, Burrows, \\& Pinto 2000; Burrows, Sudarsky, \\& Hubeny 2004; Sudarsky et al. 2005; Burrows 2005). Nevertheless, space-based coronagraphic techniques can be designed with inner working angles and contrast capabilities that will eventually image such planetary systems in the optical and mid-IR and distinguish planet from star (Trauger et al. 2000; Trauger, Hull, \\& Redding 2001). However, the close-in EGPs with orbital semi-major axes less than $\\sim$0.1 AU will not be imaged separately any time soon. In these cases, distinguishing the planet's spectrum from that of the star requires different techniques that don't rely on imaging. Fortunately, it has been shown recently by Charbonneau et al. (2005) and Deming et al. (2005,2006) that the {\\it Spitzer} infrared space telescope can discern changes in the summed light of a transiting EGP and primary star due to the occultation of the planet by the star during secondary eclipse (phase angle, $\\alpha$, near 0$^{\\circ}$). The difference in the summed light just before and during planetary eclipse provides a measure of the irradiated planet's emissions in the {\\it Spitzer} IRAC bands at 3.6 \\mic, 4.5 \\mic, 5.8 \\mic, and 8.0 \\mic, in the MIPS band at 24 \\mic, and via the {\\it Spitzer}/IRS. To date, nine transiting EGPs have been discovered (Charbonneau, Brown, Burrows, \\& Laughlin 2006), four (HD209458b, TrES-1, HD189733b, and HD149026b) are close enough to attempt secondary eclipse measurements with adequate signals-to-noise, and, as of this writing, eclipses for three transiting EGPs have in fact been detected\\footnote{The planet/star flux contrast ratio for HD149026b may, however, be too low for a successful {\\it Sptizer} campaign.}. The corresponding planet/star flux ratios\\footnote{actually, detected electron ratios} at superior conjunction are for TrES-1 0.00066$\\pm$0.00013 and 0.00225$\\pm$0.00036 at 4.5 \\mic and 8.0 \\mic, respectively (Charbonneau et al. 2005), for HD209458b 0.0026$\\pm$0.00045 at 24 \\mic (Deming et al. 2005), and for HD189733b 0.0055 $\\pm$0.00017 at $\\sim$16 \\mic in the IRS peak-up band (Deming et al. 2006). Along with the inferences using HST/STIS of the presence of sodium in the atmosphere of HD209458b (Charbonneau et al. 2002; Fortney et al. 2003; Allard et al. 2003) and of photolytic atomic hydrogen in its wind (Vidal-Madjar et al. 2003), these data are the first direct ``spectral\" measurements of extrasolar planet atmospheres and are early harbingers of the numerous programs of EGP remote sensing from the ground and from space being planned and/or proposed. The secondary eclipse data for HD209458b and TrES-1 have been subjected to preliminary theoretical analysis by four groups. Burrows, Hubeny, \\& Sudarsky (2005) concluded that these data are best interpreted with atmospheres containing water and carbon monoxide for which redistribution to the night side is significant, but partial. They conclude that the metallicity dependence is very weak and predict that the flux at 3.6 \\mic is higher than that at 4.5 \\mic. They also predict a broad peak near 10 \\mic, not so obvious in the theoretical results of others. Seager et al. (2005) emphasize the potential effects of non-solar C/O ratios above 1.0, in particular the associated lowering of the water abundance and weakening of the water absorption features. They also suggest that the dayside reradiates most of the stellar heat absorbed and incorporate into their arguments the upper limit near 2.2 \\mic found for HD209458b by Richardson, Deming, \\& Seager (2003). Fortney et al. (2005) have trouble fitting the steep spectral slope seen in TrES-1 between 4.5 \\mic and 8.0 \\mic, without a significant enhancement in metallicity. With enhancements of 3 to 5, they fit the two TrES-1 data points to within 2-$\\sigma$ (8 \\mic) and 1-$\\sigma$ (4.5 \\mic). Furthermore, for both TrES-1 and HD209458b they prefer uniform reradiation of the absorbed stellar light over the entire planetary sphere and, hence, complete heat redistribution. Barman et al. (2005) calculate 2D planetary atmospheres and a set of light curves for TrES-1 and HD209458b and redistribute heat from the dayside with a redistribution factor $f$ (Burrows et al. 2000; Burrows, Sudarsky, \\& Hubbard 2003), also used by Burrows, Hubeny, \\& Sudarsky (2005) and Fortney et al. (2005). They assume nightside core fluxes consistent with fixed values of effective temperature (\\teff) of 225 K and 500 K for TrES-1 and HD209458b, respectively. Barman et al. (2005) conclude that some redistribution must be occurring in the atmospheres of both HD209458b and TrES-1, but have trouble simultaneously fitting the two TrES-1 data points. All groups find that the atmospheric temperatures are, as expected, hot and above $\\sim$1000 K, but the predicted planet/star contrast ratio spectra at superior conjunction vary perceptibly from group to group. With the IRS peak-up measurement of HD189733b at $\\sim$16 \\mic, and the secondary eclipse data anticipated in the near future for the remaining combinations of object and {\\it Spitzer} band \\footnote{bringing the total number of data points or constraints around secondary eclipse to 18 (!)}, the time is ripe for a new set of theoretical spectral models and predictions for HD209458b, TrES-1, and HD189733b at superior conjunction ($\\alpha = 0$), as well as for the corresponding light curves for the general phase angle, $\\alpha$. In this paper, we provide such models and compare to the extant data to extract physical information about the atmospheres of these three transiting EGPs. We also make predictions for the light curves as a function of wavelength and the degree of redistribution to the nightside, and explore the metallicity dependence of the secondary eclipse predictions. To calculate the phase light curves, we use the 2D photon transport code and technique described in Sudarsky et al. (2005), but introduce the redistribution parameter, P$_{n}$, which is the fraction of the stellar energy intercepted by the planet that is redistributed to the nightside \\footnote{Note that with this definition, if the Bond albedo were large (which is the case only for cloudy models we don't discuss in this paper), the P$_n$ = 0.5 model would result in slightly greater IR fluxes from the nightside than the dayside.}. P$_{n}$ = 0 means no redistribution. We calculate for a given star/planet system and P$_{n}$ both the dayside and nightside atmospheric temperature($T$)/pressure($P$) profiles and the associated spectra, and then for a given phase angle, $\\alpha$, combine the emissions from the two hemispheres to derive the total planet fluxes at the Earth for 300 wavelengths logarithmically spaced from 2.5 \\mic to 30 \\mic and the corresponding planet/star flux ratios. Limb darkening effects for the day and night sides and planetary Bond, geometric, and spherical albedos for the day side are automatically derived in the calculations and are not imposed artificially. We have opted in this paper for the P$_n$ parametrization, and not the $f$ parametrization mentioned above and introduced by Burrows et al. (2000), because it is better tied to the core issue of heat redistribution, and because $f$ is definitionally tied to stellar irradiation, which is in fact absent on the nightside. Modeling the heating on the nightside with a flux at the base of the atmosphere that accounts for the advection of heat by winds seemed a bit more physical than heating the nightside by the ensatz of external insolation. The results are predictions for the three transiting EGPs as a function of wavelength or {\\it Spitzer} band and six values of P$_{n}$ (\\{0,0.1,0.2,0.3,0.4,0.5\\}). We thereby derive the dependence of the planet/star flux ratio spectra upon redistribution fraction and phase angle, albeit in the context of a simplified meteorological model. The spectral model, however, is state-of-the-art. In addition, we determine for the three close-in EGPs the approximate P$_{n}$ dependence of the day and night side $T/P$ profiles and pay special attention to the temporal and phase dependence of the flux ratios in the IRAC bands and the 24-\\mic MIPS band during a full orbit ($0^{\\circ} < \\alpha < 180^{\\circ}$), not just at superior conjunction (secondary eclipse). In this way, we provide a complete set of theoretical models both for comparison with current data and for predicting future measurements. Though our baseline models are for solar metallicity, we find that the metallicity dependence, without clouds and for solar abundance ratios, is small (see \\S\\ref{comparison}). ", "conclusions": "\\label{comparison} We now turn to specific comparisons between the extant {\\it Spitzer} data and our theoretical results. Figure \\ref{fig:13} portrays the planet/star flux ratios versus wavelength at superior conjunction for P$_n$ = 0.5 ($\\sim$complete redistribution) and the three transiting EGPs: HD189733b (blue), TrES-1 (red), and HD209458b (green). Superposed as large squares with 1-$\\sigma$ flux error bars are the four secondary eclipse measurements to date (two for TrES-1, one for HD209458b, and one for HD189733b). Also included as round dots in the appropriate color are the band-integrated detected electron ratio predictions for these EGP models, with approximate band widths indicated and no error bars in the flux direction. A comparison between the measured points and corresponding theoretical points for this P$_n$ = 0.5 model is encouraging, particularly for the TrES-1 data at 4.5 \\mic and 8.0 \\mic, but also for the HD189733b IRS peak-up data point near 16 \\mic. The 24-\\mic point for HD209458b is within about 1-$\\sigma$, but slightly below the theory. Since these P$_n$=0.5 models yield the lowest theoretical values for the contrast ratios among the set from 0.0 to 0.5, for the HD209458b 24-\\mic point this may be the best we can do currently. For the HD189733b point at 16 \\mic, the entire P$_n$ range studied would still be consistent to within the 1-$\\sigma$ range quoted (see the top left panel of Fig. \\ref{fig:12}), with perhaps only the P$_n$ = 0.0 model mildly discounted, rendering problematic for this EGP the constraint on the degree of heat redistribution from this one point alone. As Fig. \\ref{fig:6} suggests, values of P$_n$ of 0.0, 0.1, 0.2, and 0.3 do not fit the TrES-1 data point at 4.5 \\mic to at least 2-$\\sigma$, while values of P$_n$ of 0.0 and 0.1 do not fit the TrES-1 data point at 8.0 \\mic to the 2-$\\sigma$ level. The other values of P$_n$ can not be excluded. Hence, we conclude that while some heat redistribution by winds to the nightside is definitely indicated for TrES-1 and HD209458b, the degree of redistribution is harder to constrain, with a slight bias towards the larger values of P$_n$. The steep slope from 4.5 \\mic to 8.0 \\mic is best explained by the rise to a peak near $\\sim$10 \\mic, which in our models is a natural consequence of the relative strength and positions of water bands longward of $\\sim$5.5 \\mic (Fig. \\ref{fig:11}). Note that without a large water abundance, none of the data nor their ratios would make collective sense. This was the conclusion of Burrows, Hubeny, \\& Sudarsky (2005), which we reconfirm here. Furthermore, the presence of the 4.67 \\mic band of CO is indicated by the depth of the 4.5-\\mic feature of TrES-1, but due to the fact that this IRAC-2 band measurement perforce sums over steeply rising fluxes in regions of the spectrum that bracket the 4.67 \\mic feature, and the fact that with reasonable abundances the band is saturated, almost nothing can be said about the CO abundance (Burrows, Hubeny, \\& Sudarsky 2005). We predict a rise from IRAC-2 to IRAC-1 for all our models, indicative of the peak we generically see just shortward of 4.0 \\mic. We also predict a slight peak around 10 \\mic, and a plateau from $\\sim$14 \\mic to 30 \\mic. The peak near 10 \\mic might be discernible for HD189733b using the full capability of {\\it Spitzer}/IRS. The predicted plateau seems suggested by the comparison between theory and data for the HD189733b 16-\\mic and HD209458b MIPS points, taken together, but mixing objects (as we have been forced to do with only four data points) is not very satisfying. We have calculated a P$_n$ = 0.5 model for HD209458b with $10\\times$solar metallicity and, contrary to the conclusion of Fortney et al. (2005), we find that the band contrast ratios are within $\\sim$5\\% of those with solar abundances. This is because, without clouds, the Bond albedos are very low ($\\sles$5\\%). Since changing the metallicity does not change the total stellar light intercepted by the planet for a given planet radius, the characteristic atmospheric temperatures are similar. What is more, we find that the $T/P$ profiles are also similar, with the result that the fluxes and contrast ratios are little altered. We have not been able to trace the origin of the difference between our results for higher metallicities and those of Fortney et al. (2005). However, we interpret the very weak metallicity dependence of the contrast ratios at secondary eclipse for EGP models without clouds that we find theoretically to indicate that the metallicity may well be supersolar and large. However, by the same token, we conclude that the metallicity can not easily be constrained nor measured by secondary eclipse data alone. Cloud models, which we expect may be relevant for HD209458b alone among the three EGPs (Fortney et al. 2003; Sudarsky, Burrows, \\& Hubeny 2003), may well change this conclusion and variations in the C/O ratio, while we do not see any need at this time to invoke them to fit the four {\\it Spitzer} data points, are still of interest (Seager et al. 2005). One way to significantly alter the planet/star contrast ratios is to introduce at altitude a strong absorber in the optical and near-UV, where the incident stellar flux can be large. In this way, the upper atmosphere is heated. The associated reradiated optical flux is also greater and the $T/P$ profile manifests a ``stratospheric\" inversion (Hubeny, Burrows, \\& Sudarsky 2003). Since the mid-IR fluxes originate higher up in the atmosphere than where $\\tau_{\\rm Rosseland}$$\\sim$1, the associated brightness temperatures from 4 \\mic to 30 \\mic are also enhanced. The increase in the emergent fluxes in the optical and mid-IR leads to a corresponding suppression in the near-IR ($\\sim$1-4 \\mic). This potential mechanism for altering our baseline predictions in the {\\it Spitzer} bands and for suppressing flux in the near-IR, particularly in the $Z$, $J$, and $K$ bands, should be borne in mind. Figure \\ref{fig:14}, constructed from a theoretical model found in Hubeny, Burrows, \\& Sudarsky (2003), depicts an extreme version of this effect for P$_n$ = 0 models of OGLE-TR56b with (red curve) and without (blue curve) TiO and VO in its upper atmosphere. In fact, we expect that TiO and VO are both flushed out of the upper atmosphere by the coldtrap effect, but suppressing this effect allows us to make the general point. Note that the bumps near 10 \\mic and 4 \\mic seen in our fiducial models (Figs. \\ref{fig:10} -- \\ref{fig:12} and Fig. \\ref{fig:13}) can be altered, shifted (4 \\mic), or muted (10 \\mic) by this upper-atmosphere absorption effect, so if we fail to see these features as predicted interesting stratospheric or upper-atmosphere chemistry might be implied. Conversely, the presence of these bumps will put useful limits on such upper-atmosphere absorbers. As mentioned in \\S\\ref{intro}, we employ the latest measurements of the transit radii of each of the three EGPs in determining the planet/star contrast levels. However, these radii, being transit radii that probe along the chord of the planet in the optical (Burrows, Sudarsky, \\& Hubbard 2003), are not strictly the appropriate radii to use in determining the planet/star contrast ratios. They are close, but the total stellar energy intercepted by the planet depends upon wavelength and is different in the near-IR water bands, where the transit radius should be slightly larger than in the optical (for HD209458b, by $\\sim$2-3\\%; Fortney et al. 2003). Furthermore, planetary emission is from a radius that is not corrected for by the ``transit radius effect.\" For HD20948b, this radius difference can be 8-10\\% (Burrows, Sudarsky, \\& Hubbard 2003), while for TrES-1 and HD189733b, due to the lower atmospheric temperatures and higher gravities the effect is smaller ($\\sles$5\\%). The upshot is that the predicted flux ratios for HD209458b could be smaller by as much as $\\sim$15\\%, bringing the 24-\\mic MIPS point better in line with our prediction, but introducing further ambiguities into predictions at all wavelengths until the radius issue is resolved. No one doing theoretical secondary eclipse calculations has yet corrected for these subtle radius effects. Moreover, even the measured transit radii retain a residual ambiguity due to the systematic uncertainty in the stellar radius, which could easily be 5-10\\%. Knutson et al. (2006) argue that measurements of the star, constraints of stellar evolution theory (Cody \\& Sasselov 2002), and the detailed fits to the HST/STIS transit measurements together yield a transit radius for HD209458b with an error of only $\\sim$2\\%. Perhaps, but our same concerns apply to TrES-1 and HD189733b. In sum, ambiguities in the appropriate radii to employ in comparing secondary eclipse data with theory remain and slightly compromise their interpretation. However, that the data and theory we have developed here correspond as well as they do is gratifying. In fact, the theory does a good job fitting the four secondary eclipse data points (Fig. \\ref{fig:13}), whatever the ambiguities. In addition, there is evidence for redistribtuion to the nightside, particularly for TrES-1 and HD209458b, (though its specific magnitude remains to be determined), and the presence of H$_2$O and of CO is strongly indicated. Moreover, we find that the metallicity dependence of cloud-free models is quite mild, but that ambiguities in the radii remain to slightly compromise the interpretation of the data. Due to the greater sensitivity in the IRAC bands to variations in P$_n$ (Figs. \\ref{fig:5} - \\ref{fig:8}), such data have greater potential to determine the degree(s) of redistribution. Data off secondary eclipse at other phase angles, and particularly in IRAC-4 (Figs. \\ref{fig:10} - \\ref{fig:12}), would further constrain the models, but the most propitious phase angles in this regard are larger than 90$^{\\circ}$ and, hence, will prove very difficult to measure. However, JWST, with its two to three orders-of-magnitude greater sensitivity in the mid-IR, will be able to measure a large fraction of the planetary light curves. In the shorter term, data, or even upper limits, at the 14 other anticipated {\\it Spitzer} band points are anxiously awaited." }, "0607/astro-ph0607222_arXiv.txt": { "abstract": "Most stars reside in binary/multiple star systems; however, previous models of planet formation have studied growth of bodies orbiting an isolated single star. Disk material has been observed around both components of some young close binary star systems. Additionally, it has been shown that if planets form at the right places within such disks, they can remain dynamically stable for very long times. Herein, we numerically simulate the late stages of terrestrial planet growth in circumbinary disks around `close' binary star systems with stellar separations 0.05 AU $\\leq a_B \\leq$ 0.4 AU and binary eccentricities 0 $\\leq e_B \\leq$ 0.8. In each simulation, the sum of the masses of the two stars is 1 M$_{\\odot}$, and giant planets are included. The initial disk of planetary embryos is the same as that used for simulating the late stages of terrestrial planet formation within our Solar System by Chambers (2001, Making more terrestrial planets, Icarus 152, 205-224), and around each individual component of the $\\alpha$ Centauri AB binary star system by Quintana et al. (2002, Terrestrial planet formation in the $\\alpha$ Centauri system, Astrophys. J. 576, 982-996). Multiple simulations are performed for each binary star system under study, and our results are statistically compared to a set of planet formation simulations in the Sun-Jupiter-Saturn system that begin with essentially the same initial disk of protoplanets. The planetary systems formed around binaries with apastron distances $Q_B$ $\\equiv$ $a_B$(1 + $e_B$) $\\lesssim$ 0.2 AU are very similar to those around single stars, whereas those with larger maximum separations tend to be sparcer, with fewer planets, especially interior to 1 AU. We also provide formulae that can be used to scale results of planetary accretion simulations to various systems with different total stellar mass, disk sizes, and planetesimal masses and densities. ", "introduction": "More than half of all main sequence stars, and an even larger fraction of pre-main sequence stars, are in binary/multiple star systems (Duquennoy and Mayor 1991; Mathieu \\etal\\ 2000). Virtually all previous models of planet formation, however, have assumed an isolated single star. Of the first 131 extrasolar planet systems that have been confirmed, at least 30 are on so-called S-type orbits that encircle one component of a binary star system, including at least 3 that orbit one member of a triple-star system (Raghavan \\etal\\ 2006). The effect of the stellar companion on the formation of these planets, however, remains unclear. One planet has been detected in a P-type orbit which encircles both members of a binary star system. This planet, which has a minimum mass of $\\sim$ 2.5 times the mass of Jupiter (M$_{\\jupiter}$), orbits $\\sim$ 23 AU from the center of mass of PSR 1620-26, a radio pulsar binary comprised of a neutron star and a white dwarf in a $\\sim$ 191 day stellar orbit (Lyne \\etal\\ 1988, Sigurdsson 1993, Sigurdsson \\etal\\ 2003). The most plausible model for its formation is accretion within a metal-rich disk produced by post-main sequence Roche lobe overflow (Lissauer 2004). Planets have not been detected in a P-type orbit around two main sequence stars, but short-period binaries are not included in precise Doppler radial velocity search programs because of their complex and varying spectra. Planets in P-type orbits around the eclipsing binary star system CM Draconis have been searched for using the eclipse timing variation method (Deeg \\etal\\ 2000), but results were not definitive. Two substellar companions have been detected around the G6V star HD 202206, with minimum masses of 17.4 M$_{\\jupiter}$ (at 0.83 AU) and 2.44 M$_{\\jupiter}$ (at 2.55 AU) (Udry \\etal\\ 2002). The inner companion is so massive that it is considered to be a brown dwarf, and it is likely that the outer companion formed from within a circumbinary (star-brown dwarf) disk (Correia \\etal\\ 2005). A more general discussion of the detectability of circumbinary planets is presented by Muterspaugh (2005). Note also that the observation of two small moons orbiting in nearly circular/planar orbits about Pluto-Charon (Weaver \\etal\\ 2006), a system which is like a binary with an 8/1 mass ratio, suggests that accretion can occur in P-type orbits about close binaries. The main objective of this article is to numerically examine the late stages of terrestrial planet formation around both members of a binary star system. The existence of Earth-like planets in orbit about one or both components of main sequence binary stars has yet to be determined, though ground- and space-based efforts to search for extrasolar terrestrial planets are currently in development. An additional benefit of understanding the differences between planet formation around single stars and that around close binaries is that for eclipsing binaries, the contrast ratio between brightness of the stars and that of the planet(s) is reduced during the eclipse. For a total eclipse of identical stars, this reduction is a factor of two; as lower mass main sequence stars can be just slightly smaller but significantly less luminous, the detectability of the planet can be enhanced by more than a factor of two when the fainter star transits the brighter one. In an evolved close binary having undergone mass transfer, the fainter star can actually completely eclipse its much brighter companion, leading to an even larger improvement in planetary detectability. In the conventional model of planet formation, terrestrial planets are believed to have formed by an accretion process from within a disk of gas and dust that has remained around a newly formed star (Safronov 1969, Lissauer 1993). The coexistence of disks of material with stars that possess a stellar companion support the idea that planet formation within binary star systems may be common. Circumbinary disk material has been detected through millimeter and mid-infrared excess emission around several spectroscopic pre-main sequence binary star systems with stellar semimajor axes $a_B$ $\\lesssim$ 1 AU. These systems include GW Ori (Mathieu \\etal\\ 1995), UZ Tau E (Jensen \\etal\\ 1996), and DQ Tau (Mathieu \\etal\\ 1997). The masses of these disks are each comparable to or exceed the minimum mass of the solar nebula, $\\sim$ 0.01 solar mass (M$_{\\odot}$) (Weidenshilling 1977), and are also comparable to the masses of disks found around single stars. Numerical models of circumbinary disks find that, for binary star systems with binary eccentricities ($e_B$) increasing from 0 -- 0.25, the inner edge of a gaseous disk is truncated to within $\\sim$ 1.8 -- 2.6 times the semimajor axis of the binary stars' mutual orbit ($a_B$) (Artymowicz and Lubow 1994, Lubow and Artymowicz 2000). A star with a giant planet orbiting interior to the terrestrial planet region is dynamically a binary system of extreme mass ratio. Raymond \\etal\\ (2005) showed that planetary embryos can accrete into terrestrial planets around a star that has a close-in (between 0.15 AU -- 0.5 AU) Jupiter-mass planet. Herein, we simulate terrestrial planetary accretion within a circumbinary disk of protoplanets around `close' ($a_B$ = 0.05 -- 0.4 AU) binary star systems that each have a combined stellar mass of 1 M$_{\\odot}$. Our numerical method and the initial states of the systems that we have simulated are given in Section 2. Section 3 examines the regions of stability for test particles orbiting about these binary star systems. The results of the close binary accretion simulations, including a quantitative analysis of the final planetary systems formed, are presented in Section 4. Our conclusions are discussed in Section 5. Appendix A presents new simulations of the late stages of terrestrial planet formation in the Sun-Jupiter-Saturn system that we have performed to facilitate comparisons between planet growth around single and close binary stars, and simulations using an initial disk of bodies whose eccentricities are forced by the binary stars are presented in Appendix B. In Appendix C, we discuss the scaling of our results to systems with different planetesimal densities, disk sizes, and stellar masses. ", "conclusions": "In the present work, we have examined the effect of 14 different short-period binary star configurations (each with a combined stellar mass of 1 M$_{\\odot}$) on the late stages of terrestrial planet formation within a circumbinary protoplanetary disk. Stellar mass ratios of 1:1 and 4:1 were examined, and the initial orbits of the stars were varied (with semimajor axes between 0.05 AU $\\leq a_B \\leq$ 0.4 AU and eccentricities $e_B \\leq$ 0.8) such that the stellar apastron ranged from 0.05 AU $\\leq Q_B \\leq$ 0.4 AU. The midplane of the disk began coplanar to the stellar orbit in all but one set of runs; in that exceptional set, the initial inclination of the disk started at 30$^{\\circ}$ relative to the binary orbital plane. Giant planets analogous to Jupiter (at $\\sim$ 5.2 AU) and, in all but one set of runs, Saturn (at $\\sim$ 9.5 AU) were included. The evolution of the protoplanets was followed using a symplectic `close binary' algorithm which was developed for this purpose (Chambers \\etal\\ 2002), and 5 or 6 simulations were performed for each binary star system under study (with small changes in the initial conditions of the disk) to account for the chaotic nature of these $N$-body systems. We statistically compared our results to a large set of simulations of the Sun-Jupiter-Saturn system that began with virtually the same initial disk mass distribution (initially performed by Chambers (2001), but also integrated herein (Appendix A)). The close binary stars with maximum separations $Q_B$ $\\equiv$ $a_B$(1 + $e_B$) $\\leq$ 0.2 AU and small $e_B$ had little effect on the accreting bodies, and in most of these simulations terrestrial planets formed over essentially the entire range of the initial disk mass distribution (and even beyond 2 AU in many cases). The stellar perturbations cause orbits to precess, thereby moving secular resonances out of the inner asteroid belt, allowing terrestrial planets to form from our initially compact disk and remain in stable orbits as far as 2.98 AU from the center of mass of the binary stars. The effects of the stellar perturbations on the inner edge of the planetesimal disk become evident in systems with larger $a_B$ (and $Q_B \\geq$ 0.3 AU) and in most of the simulations with $e_B \\textgreater$ 0. Terrestrial-mass planets can still form around binary stars with nonzero eccentricity, but the planetary systems tend to be sparcer and more diverse. Binary stars with $Q_B \\geq$ 0.3 AU perturb the accreting disk such that the formation of Earth-like planets near 1 AU is unlikely. Despite these constraints, at least one terrestrial planet (at least as massive as the planet Mercury) formed in each of our simulations." }, "0607/astro-ph0607544_arXiv.txt": { "abstract": "We investigate the dynamics of an injected outflow propagating in a progenitor in the context of the collapsar model for gamma-ray bursts (GRBs) through two dimensional axisymmetric relativistic hydrodynamic simulations. Initially, we locally inject an outflow near the center of a progenitor. We calculate 25 models, in total, by fixing its total input energy to be $10^{51}\\mbox{ergs\\ s}^{-1}$ and radius of the injected outflow to be $7\\times 10^7$ cm while varying its bulk Lorentz factor, $\\Gamma_{0} = 1.05\\sim 5$, and its specific internal energy, $\\epsilon_0/c^2 = 0.1\\sim 30$ (with $c$ being speed of light). The injected outflow propagates in the progenitor and drives a large-scale outflow or jet. We find a smooth but dramatic transition from a collimated jet to an expanding outflow among calculated models. The opening angle of the outflow ($\\theta_{\\rm sim}$) is sensitive to $\\Gamma_0$; we find $\\theta_{\\rm sim} < 2^\\circ$ for $\\Gamma_0 \\gtrsim 3$. The maximum Lorentz factor is, on the other hand, sensitive to both of $\\Gamma_0$ and $\\epsilon_0$; roughly $\\Gamma_{\\rm max} \\sim \\Gamma_0 (1+\\epsilon_0/c^2)$. In particular, a very high Lorentz factor of $\\Gamma_{\\rm max} \\gtrsim 100$ is achieved in one model. A variety of opening angles can arise by changing $\\epsilon_0$, even when the maximum Lorentz factor is fixed. The jet structure totally depends on $\\Gamma_0$. When $\\Gamma_0$ is high, a strong bow shock appears and generates a back flow. High pressure progenitor gas heated by the bow shock collimates the outflow to form a narrow, relativistic jet. A number of internal oblique shocks within the jet are generated by the presence of the back flow and/or shear instability. When $\\Gamma_0$ is low, on the contrary, the outflow expands soon after the injection, since the bow shock is weak and thus the pressure of the progenitor gas is not high enough to confine the flow. Our finding will explain a smooth transition between the GRBs, X-ray rich GRBs (XRRs) and X-ray Flashes (XRFs) by the same model but with different $\\epsilon_0$ values. ", "introduction": "The GRBs are to our best knowledge the most energetic phenomena in the Universe. So far intense efforts have been made both on the observational and theoretical grounds toward understanding of their natures, but their origins still remains to be investigated. One of the most important finding recently is that GRBs are involved with relativistic collimated flows. To account for the observations, an extremely high bulk Lorentz factor, typically more than 100, is required \\citep{Rees92}. GRBs are known to be composed of two classes: long-duration GRBs (with duration being longer than a few seconds) and short-duration GRBs (with duration being less than a second). At least, some of the long-duration GRBs are known to be associated with supernovae (SNe). Good examples of GRB-SN connection are GRB980425/SN1998bw \\citep{Galama98} and GRB030329/SN2003dh \\citep{Stanek03,Uemura03,Hjorth03,Price03}. These provide strong evidences that the central engines of (at least part of) the long-duration GRBs are SNe. Such association was theoretically predicted by \\citet{Woosley93} and \\citet{Paczynski98}. A signature of a supernova contribution is actually found in the afterglow spectra of these GRBs. These supernovae are categorized in the type Ic whose progenitor has lost its hydrogen and helium envelope when the core-collapse occurs. Another possible GRB associated with SN is GRB021211. A strong absorption feature is observed in the spectrum of the afterglow \\citep{DellaValle03}. They concluded that the absorption is due to CaII which is synthesized by the associated supernova explosion. Further, the discovery of the host galaxy of the long duration GRBs being star forming galaxies \\citep{Bloom02,LeFloch03,Tanvir04} also strengthens the idea of strong GRB-SN connection. It has been suggested that even supernovae, in which no associated GRB was found, might have a link with GRB. For example, the peculiar SN2002ap recorded a high velocity component of $0.23 c$ and huge kinetic energy of the jet of $5\\times 10^{50}$ ergs, at least. These values are similar to those of GRBs, indicating a similar explosion mechanism of SN2002ap to that of GRBs. \\citet{Totani03} concluded that SN2002ap is one example of the supernovae which failed to make a GRB. Recently, a number of X-ray-rich GRBs (XRRs) and XRFs, very similar phenomena to GRBs but with significantly lower peak energy, have been successively discovered thanks to the good performance of HETE-2 (see \\citet{Heise01} and \\citet{Sakamoto05}). Interestingly, the event rates of XRRs and XRFs are similar to that of the long duration GRBs \\citep{Sakamoto05}. The origin of these events is poorly understood, but similar burst properties of GRB, XRRs, and XRFs except for peak energy leads to an idea that they all might have the same origins but with different viewing angles (Nakamura 2000) or with variable opening angles (Lamb et al. 2005). In this paper, we elucidate the theory of XRFs, XRRs, and GRBs in the context of the so-called collapsar model. The collapsar is a death of a massive star in the last stage of the stellar evolution. The collapsar model for GRBs was proposed by Woosley (1993; see also MacFadyen \\& Woosley 1999), for a central engine of GRBs. In this model, strong outflows, or jets, emerge from deeply inside the collapsar and propagate into the interstellar medium (ISM), producing gamma-ray bursts. \\citet{MacFadyen99} performed two dimensional hydrodynamic simulations based on this model. Assuming annihilation of neutrino and anti-neutrino, they deposited thermal energy around the center of the core which had been collapsed and become a black hole. Their initial mass density profile is very flattened due to rotation of the progenitor. The gas around the center, where the high thermal energy is deposited, expands and forms very collimated outflow; i.e., a ``jet''. The outflow successfully became a bipolar outflow. Unfortunately, however, their calculations were not relativistic one, so the relativistic effects which are important to understand GRBs were not included. Since the mass density of the progenitor is quite high, it is not a trivial issue whether or not the formed outflow can always keep collimated structure and break out from the progenitor surface as a jet. It has been pointed out through the AGN jet simulations that the multi-dimensional effects are so important for the dynamics of the ``light jet'' into some dense gas (see, e.g., \\citet{Mizuta04}, and references therein). Here, light jet stands for the jet, the mass density of which is smaller than that of the ambient gas. At least two dimensional hydrodynamic calculations are indispensable to investigate the outflow propagation and its dynamics inside and outside the progenitor. Multi-dimensional, relativistic hydrodynamic simulations have been so far performed by several groups in the context of collapsar model \\citep{Aloy00,Zhang03,Zhang04,Umeda05}. \\citet{Aloy00}, for example, performed relativistic hydrodynamic simulations based on the model by \\citet{MacFadyen99}. They have found that a bipolar flow is created and it breaks out from the progenitor. Interestingly, the maximum Lorentz factor of about 40 has been achieved, when the jet breaks out of the progenitor. Another type of relativistic hydrodynamic simulations of the outflow have also been performed by \\citet{Zhang03}, \\citet{Zhang04}, and \\citet{Umeda05}. They modeled a mainly very hot jet, i.e. an initially thermal energy dominated jet. Injected jets from the computational boundary, which is assumed to be very close to the collapsed center of the progenitor, always successfully propagate in the progenitor keeping good collimation and break out the progenitor. It is still open question, however, whether a collimated outflow emerging from the center of the progenitor can always break out or not. Although several provenance studies have demonstrated successful propagation throughout the progenitor and breakout of the input outflow, they might have assumed unrealistically large energy input in the initial condition. We should be aware that the formation mechanism of the outflow from the center of the collapsed progenitor has been poorly understood. In other words, we still do not know the physical conditions (density, thermal energy, kinetic energy, magnetic energy, opening angle etc) for generating outflows. Further, it is not completely clear yet what discriminates between SNe associated with GRBs and those without GRB association. The connection between XRFs and GRBs is another important issue. A key factor may be attributed to the different dynamics of the outflow propagation. Motivated by these questions we perform series of relativistic hydrodynamic simulations of the outflow propagation in the progenitor and ISM. In these simulations, we fix the total input energy power of $10^{51} \\mbox{ergs s}^{-1}$ but vary the bulk Lorentz factor ($\\Gamma_0$) and the specific internal energy of the initial outflow ($\\epsilon_0$ which excludes the rest mass energy). We discuss what types of outflow can emerge from the central system for a wide range of parameters, $\\Gamma_0$ and $\\epsilon_0$. This paper is organized as follows. In Sec. \\ref{method} we introduce our model and then explain the numerical methods and the initial background and outflow conditions. The results of the numerical simulations are presented in Sec. \\ref{results}, where we will demonstrate the emergence of two distinct types of outflows: a collimated jet and an expanding outflow. We then discuss the dynamics and structure of the outflow, focusing on the distinctions between the two types of flows in Sec. \\ref{discussion}. The final section is devoted to conclusions. ", "conclusions": "\\label{conclude} We investigate the propagation and dynamics of the outflows in the progenitor in the context of collapsar model for a central engine of GRBs by means of hydrodynamical simulations. We assume a fixed power input of the initial outflow of $\\dot{E}_0=10^{51}\\mbox{ergs s}^{-1}$ and the radius of the injected outflow $R_0=7\\times 10^7$ cm, and follow the propagation of hot outflow for different values of $\\epsilon_0$ and $\\Gamma_0$ over the ranges of $1.05 \\le \\Gamma_{0} \\le 5 \\ (0.3 \\le v_0/c \\le 0.98)$ and $0.1 \\le \\epsilon_0/c^2 \\le 30$. The net energy for 10-second injection satisfies the explosive energy of so-called hypernovae $\\sim 10^{52}$ ergs. Our conclusions can be summarized as follows: \\begin{enumerate} \\item The propagation dynamics of the outflow dramatically changes from the collimated structure to the expanding one as $\\Gamma_0$ decreases. If the Lorentz factor is high enough, say $\\Gamma_0 > 3$, the outflow can propagate throughout the progenitor, keeping a very collimated structure. The half opening angle is $\\theta_{\\rm sim} < 2^\\circ$ for $\\Gamma_0 \\gtrsim 3$. But the opening angle has weak dependence on $\\epsilon_0$, as well; we get $\\theta_{\\rm sim} < 3^\\circ$ even for smaller $\\Gamma_0$ but with small $\\epsilon_0/c^2 \\lesssim 0.1$. The maximum Lorentz factor is, on the other hand, sensitive to both of $\\Gamma_0$ and $\\epsilon_0$; roughly $\\Gamma_{\\rm max} \\sim \\Gamma_0 (1+\\epsilon_0/c^2)$. \\item In the relativistic, collimated flow, a back flow, which is anti-parallel to the main jet, appears. During the propagation in the progenitor, we can see some internal structures caused by the instability grown by the shear flow in the jet or by the interaction between the jet and back flow in the collimated jets. Such oblique shocks can help the reconfinement of the jet. The maximum Lorentz factor of the jet follows a simple formula derived from energy conservation relation. After the breakout the outflow expands into the interstellar space, although there still remains a high velocity component along the $z$-axis. Its half opening angle if a few degrees. This could be observed as GRBs. Another flow component which surrounds the central high velocity component can also be seen. It originates from the back flow during the propagation in the progenitor and shocked progenitor gas. \\item When the Lorentz factor (or the initial velocity) of the outflow is not large, say, $\\Gamma_0 \\lesssim 1.4$ ($v_0/c \\lesssim 0.7$), the outflow no longer keeps the collimation and thus expands to the forward and lateral directions. Eventually the outflow breaks out like an aspherical supernova explosion. This is because with the small outflow velocity the bow shock is weak and cannot drive the progenitor gas to high enough pressure. As a result, the reconfinement shocks, which are necessary for the collimation, does not appear. Thus, the structure is relatively featureless in the outflow. As the cross section of the reverse shock increases with time, the mass is collected at the head of the outflow. This enhances lateral expansion. \\item High Lorentz factor ($>10$)is needed to explain energetic phenomena, such as GRBs and XRFs, but the different initial internal energy affects the opening angle of the outflow for injected outflows of smaller Lorentz factor, i.e. slower velocity, thereby producing a marked difference in its observable. Rather low internal energy, $\\epsilon_0/c^2 \\lesssim 0.1$, and relatively small Lorentz factor ($<$ 5) leads to collimated non-relativistic jets, which will be observed as a failed GRB. High internal energy, $\\epsilon_0/c^2 \\gtrsim 5$, leads to un-collimated relativistic jets, which could be observed as XRFs. We can thus phenomenologically explain different types of explosions, GRBs XRFs, and failed GRB along the same line but with different values of $\\epsilon_0$ for slower injected velocity. However, a cause of producing a variety of $\\epsilon_0$ and $\\Gamma_0$ is still unknown. It should be investigated in future work. \\end{enumerate} We thank W. Zhang for his comments on the definition of $f_0$ in their paper. We appreciate S. Woosley for his helpful comments to this manuscript. We gratefully acknowledge M. A. Aloy, E. M\\\"uller, A. Macfadyen, N. Ohnish, and M. Horikoshi for the discussion on the numerical method. One of the authors (A.M.) was supported by a Research Fellowship of the Japan Society for the Promotion of Science. This work was supported in part by the Grants-in-Aid of the Ministry of Education, Science, Culture, and Sport (14079205, A.M. and S.M), (16340057 S.M.) and (14102004, 14079202, and 16740134, S.N.), This work was supported by the Grant-in-Aid for the 21st Century COE \"Center for Diversity and Universality in Physics\" from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan. This work was carried out on NEC SX5, Cybermedia Center and Institute of Laser Engineering, Osaka University, SX8 at YITP in Kyoto University, and Fujitsu VPP5000 of National Observatory of Japan. We appreciate computational administrators for technical supports. \\begin{appendix}" }, "0607/astro-ph0607258_arXiv.txt": { "abstract": "{The TeV-emitting BL~Lac object Mkn~421 was observed with very long baseline interferometry (VLBI) at three closely-spaced epochs one-month apart in March--April 1998. The source was also monitored at very-high $\\gamma$-ray energies (TeV~measurements) during the same period in an attempt to search for correlations between TeV variability and the evolution of the radio morphology on parsec scales. While the VLBI maps show no temporal changes in the Mkn~421 VLBI jet, there is strong evidence of complex variability in both the total and polarized fluxes of the VLBI core of Mkn~421 and in its spectrum over the two-month span of our data. The high-energy measurements indicate that the overall TeV activity of the source was rising during this period, with a $\\gamma$-ray flare detected just three days prior to our second VLBI observing run. Although no firm correlation can be established, our data suggest that the two phenomena (TeV activity and VLBI core variability) are connected, with the VLBI core at 22~GHz being the self-absorbed radio counterpart of synchrotron self-Compton (SSC) emission at high energies. Based on the size of the VLBI core, we could derive an upper limit of 0.1~pc ($3\\times 10^{17}$~cm) for the projected size of the SSC zone. This determination is the first model-free estimate of the size of the $\\gamma$-ray emitting region in a blazar. ", "introduction": "Among active galactic nuclei (AGN), BL~Lac objects form a class of sources characterized by high radio and optical variability, dominance of continuum over line emission, and strong and variable polarization. Such extreme properties have already suggested long ago that substantial relativistic beaming most probably occurs in this type of source \\citep{br78,bk79}. This idea was confirmed by direct detection of apparent superluminal motion in the radio jets of many BL~Lac objects based on high-resolution imaging with the very long baseline interferometry (VLBI) technique \\citep{pm82,wsj88,gwr89,msb90,gpc00}. Building on these findings, AGN unification schemes have emerged, all basically describing BL~Lac objects as radio-loud AGN with relativistic jets pointing within a small angle towards the observer \\citep{up95,fgb95,g95,kpl96,gm98,gcf98,fmc98,scu00,bl01}. In the past decade, more evidence of relativistic beaming came from high-energy X-ray and $\\gamma$-ray data, which revealed intense fluxes and strong variability up to the TeV level, especially in the two nearby BL~Lac objects Mkn~421 and Mkn~501 \\citep{pac92, qab96}. In the case of Mkn~421, very high relativistic Doppler factors (between 10 and 15) are demanded in order to reproduce the dramatic rapid flares that have been observed on timescales as short as 30~min \\citep{gab96,c97,cfr98,c04,aaa05}. Similar or even higher Doppler factors (in the range 20--50) are also obtained when fitting basic synchrotron self-Compton models to the broadband spectrum of Mkn~421 \\citep{mft99,tkm00,ksk01,kmk03,ksk03}. The VLBI observations have sought evidence of superluminal motion in \\object{Mkn 421} ever since the early 1980's. A first series of VLBI maps at 5~GHz reported apparent motions of about~2c and an angle to the line of sight of $34^{\\circ}$ for the parsec-scale jet \\citep{bel81,b84,zb90}. On the other hand, the source was found to be unresolved at high frequency with a core size of 0.15~milliarcsecond (mas) at 22~GHz \\citep{zb91}. In the 1990's, the improved performances of the VLBI technique permitted the detection of a weak one-sided jet \\citep{pwx95,xrp95,emu98,kvz98,gfv99}. This jet shows wiggles starting at about 5~mas from the core, as well as strong distortions at a distance of 20~mas from the core. Based on these data, \\citet{gfv99} derived a viewing angle smaller than $30^{\\circ}$ and an apparent jet speed between $\\sim$0.8c and 1c. While faster speeds ($\\sim$2c) were reported, invoking possible earlier misidentification of several rapidly-evolving VLBI components \\citep{m96,m99}, another analysis based on a dense time coverage (15~observing epochs over 3~years) confirmed the existence of only subluminal apparent motion ($\\le 0.3$c), therefore implying a very small viewing angle to the line of sight of $0.4^{\\circ}$ for the VLBI jet \\citep{puw99}. This analysis has recently been refined with an extended data span (28~epochs over 8~years), leading to a revised apparent speed of only $0.1\\pm0.02$c for the fastest VLBI jet component \\citep{pe05}. While this value is smaller than the value found by \\citet{klh04}, who reported an average component speed of $0.4$c from 15~GHz monitoring over 6~years, the two results do rule out the existence of superluminal motion in the Mkn~421 jet. Interestingly, Mkn~421 is not a unique case and such low apparent speeds have been found in the other TeV blazars as well \\citep{pe04}. The apparent inconsistency between the low degree of relativistic beaming derived from the VLBI observations of TeV~blazars and the high value predicted by the theory may be explained if the Doppler factor decreases along the jet as a result of either jet curvature or jet deceleration \\citep{gk03,gtc05}. It is also possible that the measured VLBI jet speed does not correspond to the speed of the actual underlying jet but instead to the speed of a perturbed pattern along the jet \\citep{z97}. In this case, the radio core may still be efficiently boosted although there would be no evidence for relativistic beaming from apparent motions in the VLBI jet. Nevertheless, one would expect correlations between radio core properties and high-energy events if this hypothesis is correct. One indication of the existence of these correlations comes from the apparent connection between the epoch of emergence of new VLBI components and the occurrence of strong X-ray and $\\gamma$-ray flares in several $\\gamma$-ray emitting AGN \\citep{wuz93,prk95,uwl97,okk98,bwk98,b98,kko98,w99,mmm00}. Furthermore, it appears that the $\\gamma$-ray flares detected by the EGRET detector onboard the Compton Gamma-Ray Observatory seem to occur during the rising phase of high-frequency radio outbursts, again suggesting a possible connection between radio and high-energy properties \\citep{vtl96,vt96,lvt00}. This is especially true for Mkn~421, which occasionally shows multi-spectral flares (from radio to TeV energies), as reported by \\citet{ksk03}. Searching for additional clues along these lines with new multi-frequency data is important for investigating whether TeV blazars are indeed strongly-beamed sources, as presumed so far. This paper reports multi-frequency VLBI maps of Mkn~421 obtained at three closely-spaced epochs (one month apart) in~1998. These observations were arranged during regular monitoring of the source at very-high $\\gamma$-ray energy (TeV level) by the CAT Cherenkov imaging telescope. In the following sections, we present results of the VLBI and TeV observations, discuss the variability of the source, and investigate possible connections between VLBI properties and TeV activity. We also derive an upper limit for the projected size of the $\\gamma$-ray emission region based on the size of the compact radio core as measured from our high-resolution VLBI data. ", "conclusions": "Coordinated multi-frequency VLBI polarization and very-high-energy TeV observations of the $\\gamma$-ray blazar Mkn~421 have provided evidence of variability in the compact VLBI core on a timescale of a few weeks and at a time when the overall TeV activity of the source was rising. Although we cannot be entirely sure about the connection between the two phenomena, our measurements are consistent with a scenario in which the VLBI core at 22~GHz represents the self-absorbed part of the low-energy emission induced by a one-zone synchrotron self-Compton model consistent with simultaneous high-energy data. At the lower radio frequencies (15, 8, and 5~GHz), the VLBI core encompasses additional emission unrelated to the SSC phenomenon and most probably originating from regions outside the $\\gamma$-ray emitting zone. Based on these observations, we also derived an upper limit of 0.1~pc ($3\\times 10^{17}$~cm) for the projected size of the $\\gamma$-ray emitting region, in agreement with previous estimates but based solely on our radio-interferometric data and hence free of any theoretical model. Further similar coordinated VLBI and TeV observations on Mkn~421 should be primarily targeted at improving the VLBI time coverage and angular resolution. Observations a few days apart over several weeks would be important to confirm whether the VLBI core variability and TeV activity are indeed connected and whether there is any anticorrelation between such variations. Increasing the VLBI angular resolution, e.g. by observing at 43~GHz with a transatlantic network or at 86~GHz with the global millimeter VLBI array, would also be important to further constrain the location and size of the $\\gamma$-ray emission region in the framework of the proposed SSC scenario. By providing an angular resolution improved by a factor of up to~6, such high-frequency VLBI observations may perhaps even permit a direct measurement of the size of the SSC zone, which is on the order of $10^{16}$~cm according to present models. VLBI monitoring of Mkn~421 should also be pursued on the long-term to refine jet component proper motions and possibly obtain more clues as to why there are no superluminal motions in this presumably strongly-beamed source. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0,bb=10 27 535 385]{4078fig9.eps}} \\caption{Multi-frequency spectral energy distribution from radio to TeV energies and VLBI core flux for Mkn~421. The continuous line is a prediction based on the one-zone synchrotron self-Compton model derived by \\citet{tkm00} from the data of the multi-wavelength campaign conducted in late April~1998, while the dots represent our measured VLBI core fluxes at 5, 8, 15, and 22~GHz on April 26, 1998 (see Table~\\ref{tab:models}). The inset shows an enlargement of the radio region.} \\label{TakaSSC} \\end{figure}" }, "0607/astro-ph0607402_arXiv.txt": { "abstract": "We present an Ultraviolet (UV) selected sample of 268 objects in the two fields of the Great Observatories Origins Deep Survey (GOODS). We used the parallel observations taken with WFPC2 in the U--band (F300W) which covered 88\\% of the GOODS fields to identify sources and selected only objects with GOODS/ACS counterparts. Spectroscopic redshifts for 95 of these sources are available and we have used the multiwavelength GOODS data to estimate photometric redshifts for the others. Most of the objects are between $0.2$ 4.5 (i.e. starbursts) are marked in blue. The northern field objects are in circles and the southern field in triangles. Only galaxies with $z_{\\rm phot}$ $<2$ were included. Bluest objects are identified by numbers 1--9. \\label{plotubbvpz2}} \\end{figure} \\begin{figure} \\plotone{fig10.ps} \\caption{Rest-frame magnitude B versus B-V color. Galaxies with spectral types $>$ 4.5 (i.e. starbursts) are marked in blue. The northern field objects are in circles and the southern field in triangles. The three squares are values typical of E, Sa-Sb, Sc-Irr (clockwise); the cross on the top left corresponds to the dE and dSph; the box region corresponds to the strong star-forming galaxies (Bershady et al. 2000) which contains blue nucleated galaxies, compact narrow emission-line galaxies and small, blue galaxies at intermediate redshifts. Only galaxies with $z_{\\rm phot}$ $<2$ were included. Numbers 1--9 are the same objects as in Fig.~\\ref{plotubbvpz2}. \\label{plotbbvpz2}} \\end{figure} \\begin{figure*} \\plotone{fig11.eps} \\caption{Contours on top of z--band images of the bluest objects in the U--B versus B--V plot shown in Fig.~\\ref{plotubbvpz2}. All objects have spectral types of starbursts, except for object \\#2 which has spectral type of Im galaxies. Contours limits are mininum=0.002, maximum=0.01, 5 levels, except for object \\#3 and \\#5 which have maximum values 0.1 and 0.05 and 9 levels, respectively. Size of each image is 76 $\\times$ 53 pixels. \\label{bluetail}} \\end{figure*} \\begin{figure*} \\plotone{fig12.eps} \\caption{Contours on top of B--band image of the tadpole starbursts at $z\\sim 1$. Coordinates are given for each object. \\label{contours1}} \\end{figure*} \\begin{figure*} \\plotone{fig13.eps} \\caption{Contours on top of z--band image of the tadpole starbursts at $z\\sim 1$. Coordinates are given for each object. \\label{contours1z}} \\end{figure*}" }, "0607/astro-ph0607634_arXiv.txt": { "abstract": "We present the results of three-dimensional simulations of supersonic Euler turbulence with grid resolutions up to $1024^3$ points. Our numerical experiments describe nonmagnetized driven turbulent flows with an isothermal equation of state and an rms Mach number of 6. We demonstrate that the inertial range scaling properties of turbulence in this strongly compressible regime deviate substantially from a Kolmogorov-like behavior previously recovered for mildly compressible transonic flows. ", "introduction": "Understanding the nature of supersonic turbulence is of fundamental importance in astrophysics and in aeronautical engineering. In the interstellar medium, highly compressible turbulence is believed to control star formation in dense molecular clouds. A whole class of more {\\em terrestrial} applications deals with the drag and stability of projectiles traveling through the air at hypersonic speeds. Molecular clouds have a highly inhomogeneous structure and the intensity of their internal motions corresponds to an rms Mach number of the order of 20. \\citet{larson81} has demonstrated that within the range of scales from 0.1~pc to 100~pc, the gas density and the velocity dispersion tightly correlate with the cloud size.\\footnote{See \\citet{kaplan.53} for an earlier version of what is now known as Larson's relations.} Supported by other independent observational facts indicating scale invariance, these relationships are often interpreted in terms of supersonic turbulence with characteristic Reynolds numbers $Re\\!\\sim\\!10^8$. Within a wide range of densities around $10^3$~cm$^{-3}$ the gas temperature remains close to $\\!\\approx\\!10$~K since the thermal equilibration time at these densities is shorter than a typical hydrodynamic time scale. Thus, an isothermal equation of state can be used as a reasonable approximation. While self-gravity, magnetic fields, chemistry, cooling and heating, as well as radiative transfer should be ultimately accounted for in turbulent models of molecular clouds, we focus here specifically on hydrodynamic aspects of the problem. Numerical simulations of {\\em decaying} supersonic hydrodynamic turbulence with the piecewise parabolic method \\citep[PPM]{colella.84} in two dimensions were pioneered by \\citet{passot..88}\\footnote{See a review on compressible turbulence by \\citet{pouquet..91} for references to earlier works.} and then followed up with high resolution 2D and 3D simulations by \\citet{porter..92a,porter..92b,porter..94,porter..98}. \\citet{sytine....00} compared the results of PPM Euler computations with PPM Navier-Stokes results and showed that Euler simulations agree well with the high-$Re$ limit attained in the Navier-Stokes models. The convergence in a statistical sense as well as the direct comparison of structures in configuration space indicate the ability of PPM to accurately simulate turbulent flows over a wide range of scales. More recently, \\citet{porter..02} discussed measures of intermittency in simulated {\\em driven} transonic flows at Mach numbers of the order unity on grids up to $512^3$ points. \\citet{porter...99} review the results of these numerical studies focusing on the origin and evolution of turbulent structures in physical space as well as on scaling laws for two-point structure functions. One of the important results of this fundamental work is the demonstration of compatibility of a Kolmogorov-type \\citep[K41]{kolmogorov41a} spectrum with a {\\em mild} gas compressibility at transonic Mach numbers. Since most of the computations discussed above assume a perfect gas equation of state with the ratio of specific heats $\\gamma=7/5$ or $5/3$ and Mach numbers generally below 2, the question remains whether this result will still hold for near isothermal conditions and {\\em hypersonic} Mach numbers characteristic of dense parts of star forming molecular clouds where the gas compressibility is much higher. What kind of coherent structures should one expect to see within the inertial range of scales in highly supersonic isothermal turbulence? Do low-order statistics of turbulence follow the K41 predictions closely in this regime? How intermittent is the turbulence? The interpretation of astronomical data from new surveys of cold ISM and dust in the Milky Way by {\\em Spitzer} and {\\em Herschel} satellites requires more detailed knowledge of these basic properties of supersonic turbulence. In this paper we report first results from our large-scale numerical simulations of driven supersonic isothermal turbulence at Mach 6 with PPM and grid resolutions up to $1024^3$ points. We solve the Euler equations in a periodic box of linear size $L=1$ with initially uniform density distribution $\\rho\\equiv1$ and the sound speed $c\\equiv1$. We initialize the simulation on a grid of $512^3$ points with the random velocity field $\\pmb{u}_0$ that contains only large-scale power within the range of wavenumbers $k/k_{min}\\in[1,2]$, where $k_{min}=2\\pi$, and corresponds to the rms Mach number $M=6$. The same velocity field is then used, with an appropriate normalization, as a steady random force (acceleration) to keep the total kinetic energy within the box on an approximately constant level during the simulation. The random force is isotropic in terms of the total specific kinetic energy per dimension, $\\left=\\left=\\left$, but its solenoidal ($\\pmb{\\nabla}\\cdot\\pmb{u}_0^s\\equiv0$) and dilatational ($\\pmb{\\nabla}\\times\\pmb{u}_0^c\\equiv0$) components are anisotropic since one of the three directions is dominated by the large-scale compressional modes, while the other two are mostly solenoidal. The distribution of total specific kinetic energy $E\\equiv\\frac{1}{2}\\int u^2dV=E^s+E^c$ between the solenoidal $E^s$ and dilatational $E^c$ components is such that $\\chi_0\\equiv E_0^s/E_0\\approx0.6$. The driving field is helical, but the mean helicity is very low: $\\left\\ll\\sqrt{\\left}$, where the helicity $h$ is defined as $h\\equiv\\pmb{u}\\cdot\\pmb{\\nabla}\\times\\pmb{u}$. In compressible flows with an isothermal equation of state the mean helicity $\\left$ is conserved, as in the incompressible case, since the Ertel's potential vorticity is identically zero \\citep{gaffet85}. We first ran the simulation on a grid of $512^3$ points from the initial conditions up to five dynamical times to stir the gas within the box. The dynamical time is defined as $t_d\\equiv L/(2M)$. Then we doubled the resolution and evolved the simulation for another $5t_d$ on a grid of $1024^3$ points. We allow one dynamical time for relaxation at high resolution to reach a statistical steady state after regridding. The time-average statistics are computed using 170 snapshots evenly spaced in time over the final segment of $4t_d$. We use the full set of 170 snapshots to derive the density statistics, since the density field displays a very high degree of intermittency. This gives us a very large statistical sample, e.g. $\\sim2\\,10^{11}$ measurements are available to determine the probability density function (PDF) of the gas density. The time-average power spectra discussed below are also based on the full data set. The velocity structure functions are derived from a sample of 20\\% of the snapshots covering the same period of $4t_d$. The corresponding two-point PDFs are built on $2-4\\times10^9$ pairs per snapshot each, depending on the pair separation. ", "conclusions": "Using high-resolution numerical simulations of nonmagnetic highly compressible turbulence at rms Mach number of 6, we have demonstrated that scaling exponents of low-order statistics deviate substantially from Kolmogorov laws for incompressible turbulence. A much higher than $1024^3$ resolution is required to possibly trace a transition from a steeper supersonic inertial range scaling at lower $k$ to a flatter Kolmogorov-like transonic scaling at higher wavenumbers." }, "0607/astro-ph0607521_arXiv.txt": { "abstract": "We study the inverse Compton scattering of solar photons by Galactic cosmic-ray electrons. We show that the \\gray\\ emission from this process is substantial with the maximum flux in the direction of the Sun; the angular distribution of the emission is broad. This previously-neglected foreground should be taken into account in studies of the diffuse Galactic and extragalactic \\gray\\ emission. Furthermore, observations by GLAST can be used to monitor the heliosphere and determine the electron spectrum as a function of position from distances as large as Saturn's orbit to close proximity of the Sun, thus enabling unique studies of solar modulation. This paves the way for the determination of other Galactic cosmic-ray species, primarily protons, near the solar surface which will lead to accurate predictions of \\grays\\ from $pp$-interactions in the solar atmosphere. These albedo \\grays\\ will be observable by GLAST, allowing the study of deep atmospheric layers, magnetic field(s), and cosmic-ray cascade development. The latter is necessary to calculate the neutrino flux from $pp$-interactions at higher energies ($>$1 TeV). Although this flux is small, it is a ``guaranteed flux'' in contrast to other astrophysical sources of neutrinos, and may be detectable by km$^3$ neutrino telescopes of the near future, such as IceCube. Since the solar core is opaque for very high-energy neutrinos, directly studying the mass distribution of the solar core may thus be possible. ", "introduction": "Interactions of Galactic cosmic-ray (CR) nuclei with the solar atmosphere have been predicted to be a source of very high energy (VHE) neutrinos \\citep{Moskalenko1991,Seckel1991,Ingelman1996}, and \\grays\\ \\citep{Seckel1991}. They are the decay products of charged and neutral pions produced in interactions of CR nucleons with gas. The predictions for these albedo \\grays\\ give an integral flux $F_\\gamma(>100\\ {\\rm MeV})\\sim(0.2-0.7)\\times10^{-7}$ cm$^{-2}$ s$^{-1}$, while analysis of EGRET \\gray\\ telescope data has yielded only the upper limit $F_\\gamma(>100\\ {\\rm MeV})=2.0\\times10^{-7}$ cm$^{-2}$ s$^{-1}$ \\citep{Thompson1997}. At lower energies ($<$100 MeV) a contribution from CR electron bremsstrahlung in the solar atmosphere may exceed that from $\\pi^0$-decay \\citep[see Fig.~4 in][]{Strong2004}. Cosmic-ray electrons comprise $\\sim$1\\% of the total CR flux. However, they propagate over the heliospheric volume, which is large compared to the solar atmosphere where the albedo \\grays\\ are produced. The Sun emits photons that are targets for inverse Compton (IC) scattering by CR electrons. As a result the heliosphere is a diffuse source of \\grays\\ with a broad angular distribution. In this paper, we evaluate the importance of IC scattering within the heliosphere, and discuss the consequences of its measurement by such instruments as the upcoming Gamma Ray Large Area Space Telescope (GLAST) mission. In the following, we use units $\\hbar=c=m_e=1$. ", "conclusions": "In this paper we have studied the IC scattering of solar photons by CR electrons\\footnote{When this work had already been completed we learned about work by \\citet{orlando2006} on the same subject.}. We have shown that the emission is significant and broadly distributed with maximum brightness in the direction of the Sun. The whole sky is shining in \\grays\\ contributing to a foreground that would otherwise be ascribed to the Galactic and extragalactic diffuse emission. The IC emission from CR electrons depends on their spectrum in the heliosphere and varies with the modulation level. Observations in different directions can be used to determine the electron spectrum at different heliocentric distances. A sensitive \\gray\\ telescope on orbit could monitor the heliosphere, providing information on its dynamics. Such observations also could be used to study the electron spectrum in close proximity to the Sun, unreachable for direct measurements by spacecraft. The assumed isotropy of the electron distribution everywhere in the heliosphere is an approximation. Closer to the Sun, the magnetic field and non-isotropic solar wind speed affect the CR electron spectrum and angular distribution, which in turn will produce asymmetries in the IC emission. Observations of such asymmetries may provide us with information about the magnetic field and solar wind speed at different heliolatitudes including far away from the ecliptic. What are the immediate implications? Since solar modulation theory is well developed \\citep[e.g.,][]{Zank2003,Ferreira2004}, accurate measurement of the electron spectrum near the solar surface will open the way to derive the spectra of other Galactic CR species, primarily protons, near the Sun. The CR proton spectrum is the input to calculations of \\grays\\ from $pp$-interactions in the solar atmosphere. These predictions can be further tested using GLAST observations of pionic \\grays\\ from the Sun. In turn, this would provide information about the density profile of the solar atmosphere, magnetic field(s), and CR cascade development \\citep{Seckel1991}. The higher the energy of \\grays, the higher the energy of the ambient particles, and thus the depth of the layers tested is increased. In conjunction with other solar monitors this can bring understanding of the deep atmospheric layers, Sun spots, magnetic storms, and other solar activity. Furthermore, understanding the solar atmosphere is necessary to calculate the neutrino flux from $pp$-interactions at higher energies ($>$1 TeV), as the CR shower development depends on the density distribution and the underlying magnetic field structure (see \\citealt{Seckel1991} and \\citealt{Ingelman1996} for details). Although small, the VHE neutrino flux from the Sun is a ``guaranteed flux'' in contrast to other astrophysical sources of neutrinos. It may be detectable by km$^3$ neutrino telescopes of the near future, such as IceCube \\citep[e.g.,][]{Hettlage2000,Ahrens2004}, on a timescale of a year or several years. Therefore, only long term periodicities associated with the solar activity could be detectable. Since the solar core is opaque for VHE neutrinos \\citep{Seckel1991,Moskalenko1993,Ingelman1996}, observations of the neutrino flux may provide us with information about the solar mass distribution. To summarise, we have shown that the observation of the IC emission from CR electrons in the heliosphere and its distribution on the sky will open a new chapter in astrophysics. This makes a sensitive \\gray\\ telescope a useful tool to monitor the heliosphere and heliospheric propagation of CR, and to study CR interactions in the solar atmosphere." }, "0607/astro-ph0607467_arXiv.txt": { "abstract": "We present a two-dimensional grid-based hydrodynamic simulation of a thin, viscous, locally-isothermal corotating disk orbiting an equal-mass Newtonian binary point mass on a fixed circular orbit. We study the structure of the disk after multiple viscous times. The binary maintains a central hole in the viscously-relaxed disk with radius equal to about twice the binary semimajor axis. Disk surface density within the hole is reduced by orders of magnitude relative to the density in the disk bulk. The inner truncation of the disk resembles the clearing of a gap in a protoplanetary disk. An initially circular disk becomes elliptical and then eccentric. Disturbances in the disk contain a component that is stationary in the rotating frame in which the binary is at rest; this component is a two-armed spiral density wave. We measure the distribution of the binary torque in the disk and find that the strongest positive torque is exerted inside the central low-density hole. We make connection with the linear theory of disk forcing at outer Lindblad resonances (OLRs) and find that the measured torque density distribution is consistent with forcing at the 3:2 ($m=2$) OLR, well within the central hole. We also measure the time dependence of the rate at which gas accretes across the hole and find quasi-periodic structure. We discuss implications for variability and detection of active galactic nuclei containing a binary massive black hole. ", "introduction": "\\label{sec:intro} Circumbinary disks have been observed in young stellar binaries, such as the spatially-resolved pre--main-sequence binary star GG Tau (e.g., \\citealt{McCabe:02,Krist:05}). Circumbinary disks should also form around binary massive black holes when interstellar gas accretes into the binary's dynamical sphere of influence (e.g., \\citealt{Artymowicz:96,Ivanov:99,Armitage:02,Milosavljevic:05}). It is generally expected that the dynamics of a binary embedded in an accretion disk is similar to the dynamics of a planet embedded in a protoplanetary disk. The planet orbits the central star while exchanging gravitational torques with, and accreting from, a gaseous circumstellar disk. If the radius of the planet's Roche lobe exceeds the pressure scale height of the disk, the planet clears a gap in the disk (e.g.,~\\citealt{Goldreich:80,Takeuchi:96,Lubow:99}). When a similar-mass binary is surrounded by a thin corotating disk, this gap-opening condition is automatically satisfied. An understanding of circumbinary accretion flows is essential for future identification of binary massive black holes in astronomical surveys. Massive binary black holes have only tentatively been identified in astronomical observations (\\citealt{Komossa:03,Komossa:06,Rodriguez:06} report evidence for a binary with projected separation of $7.3\\textrm{ pc}$), but the theoretical case for their existence is very strong. A massive binary black hole forms during the conclusion of a galaxy merger and persists at parsec or subparsec separations over a cosmological time. Under propitious circumstances, gravitational radiation ultimately induces coalescence between the black holes \\citep[and references therein]{Begelman:80,Merritt:05}. At parsec separations, the black holes can be optically resolved only in nearby galaxies. Better resolution can be achieved at radio frequencies, but binary radio cores can be confused with structures in radio jets and chance superpositions. An unresolved binary can be detected by associated characteristic spectral and time-dependent signatures if a circumbinary accretion flow is present. For example, asymmetries in the accretion flow and time-dependent accretion onto the component black holes are sources of variability that can be detected in monitoring surveys such as with the Large Synoptic Survey Telescope (LSST)\\footnote{http://www.lsst.org}. To predict the spectral signatures of massive binary black holes, one must analyze the hydrodynamic structure of the circumbinary accretion flow and then synthesize the radiative spectra emitted by the flow. In this work we carry out the hydrodynamic analysis in the case of a thin corotating disk. We defer the spectral synthesis to subsequent work (see also \\citealt{Milosavljevic:05,Bogdanovic:07b}). The search for massive binary black hole spectral transients is most effective if one can identify robust transients that repeat continuously over an extended period. Transients that are restricted to a single moment in the life of a binary, such as its formation, will be rare and difficult to interpret. Quasi--steady-state circumbinary accretion disks are particularly attractive as they may persist long enough to render their astronomical detection probable. Previous numerical simulations of circumbinary disks based on the smoothed particle hydrodynamics method \\citep{Artymowicz:94,Escala:05} and grid-based hydrodynamics with explicit bulk viscosity \\citep{Guenther:02} suggest that the surface density of a circumbinary disk drops sharply inward of $r\\sim 1.5-2a$, where $a$ is the binary semimajor axis. The binary resides within a low-density hole, also called the circumbinary gap, at the center of the disk. Although the gas density within the hole is low, some mass transfer does take place from the disk onto the point masses \\citep{Guenther:04}. However the accretion rate is below the rate that would be expected neglecting the torque that the binary exerts on the disk. The reduced accretion was detected in simulations of massive planets embedded in protoplanetary disks (e.g, \\citealt{Lubow:99,Lubow:06}). The simulations of disks with planets also indicate that an initially circular disk becomes eccentric over time, even if the orbit of the planet is circular \\citep{Papaloizou:01,Kley:06}. The growth of disk eccentricity has been attributed to a dynamical instability \\citep{Lubow:91,Dangelo:06} driven by the planet's tidal potential. In turn, the disk eccentricity can excite eccentricity in the binary, since the two eccentricities are coupled \\citep{Papaloizou:01}. Even a weak residual eccentricity can be detected in the gravitational radiation emitted during black hole coalescence (\\citealt{Armitage:05}; see also \\citealt{Artymowicz:91}). Therefore an understanding of the circumbinary disk dynamics is key for electromagnetic and gravitational wave observations of coalescing massive black holes with the {\\it Laser Interferometer Space Antenna} ({\\it LISA}).\\footnote{http://lisa.nasa.gov} Simultaneous electromagnetic and gravitational detection can be used to measure the geometry of the universe and constrain properties of dark energy (e.g., \\citealt{Holz:05,Kocsis:06}). We present high resolution simulations of a corotating viscous circumbinary disk around a circular Newtonian equal-mass binary. The disk lies in the orbital plane of the binary. This is an astrophysically-motivated simplifying assumption; alignment of the orbital plane of the inner disk and that of the binary will take place because differential precession of a non-aligned disk caused by the binary's mass quadrupole moment forces the inner part of the disk to align with the orbital plane of the binary (e.g.,~\\citealt{Larwood:97}).\\footnote{The orbital plane of the binary will itself precess due to the relativistic spin-orbit coupling, if the binary components have spins misaligned with the binary's angular momentum.} In \\S~\\ref{sec:simulations} we describe the computational method and the binary-disk setup. In \\S~\\ref{sec:results} we present our results. We study the disk eccentricity, the spatial distribution of torque deposition in the disk, and the mass accretion across the central hole in the disk. In \\S~\\ref{sec:discussion} we compare the torque distribution measured in the simulation with predictions of the theory of linear response of tidally forced disks. We also present a tentative interpretation of the eccentricity excitation in terms of a mean motion resonance. In \\S~\\ref{sec:black_holes} we discuss implications for astronomical detection of binary massive black hole mergers in progress. Finally, in \\S~\\ref{sec:conclusions} we summarize the main conclusions. ", "conclusions": "\\label{sec:conclusions} We carried out a grid-based, two-dimensional hydrodynamic simulation of a thin, locally-isothermal, corotating, non--self-gravitating, viscous accretion disk around a Newtonian equal-mass binary on a fixed circular orbit. The disk is evolved over multiple viscous times of the inner disk. We study the quasi--steady-state structure of the circumbinary fluid flow and the mechanics of the binary-disk torque coupling. We here summarize the main conclusions of this work. 1. Fluid density within twice the binary's semimajor axis is significantly reduced. The point masses orbit within the central low-density hole in the disk. The disk first becomes elliptical and then eccentric. The eccentricity of the disk is strong and precesses slowly. Radial fluid motion within the central hole is supersonic and time-dependent. Strong shocks are driven into the low-density fluid by the tidal field. 2. The binary exerts torque on the disk by imparting gravitational kicks to the eccentric low-density fluid inside the central hole. Angular momentum is transferred to the disk as the kicked fluid impacts the edge of the hole. A two-armed spiral density wave that is stationary in the rotating frame in which the binary is at rest is driven into the disk. The structure of the two-armed density wave and that of the binary torque distribution in the disk suggest that the binary forces the disk at the $m=2$ outer Lindblad resonance, where the orbital period of the disk is $3/2$ times the binary's orbital period. 3. The accretion into the hole toward the two point masses is reduced relative to the rate expected for a disk with the same surface density and a torque-free inner edge. However, significant mass transfer from the disk onto the black holes does take place. The mass transfer rate is quasi-periodic and punctuated by outbursts. The accreting fluid carries angular momentum, but the amount of angular momentum accreted onto the binary is insufficient to offset the loss of angular momentum to the gravitational torque. If the binary remains circular, its separation decays on a viscous time of the disk reduced by the ratio of the mass at the inner edge of the disk to the binary mass. 4. The quasi-periodic accretion across the hole, the shocks in the low-density fluid inside, and the disk eccentricity are sources of variability that can be used to distinguish an AGN containing a binary massive black hole from a regular AGN containing a single black hole. Detailed modeling of the emission from these structures is required to identify the most promising indicators of binary massive black hole AGNs." }, "0607/astro-ph0607184_arXiv.txt": { "abstract": "The analysis of the density and brightness of big planets' satellites, main asteroid belt objects, Kuiper belt objects and centaurs has been carried out as well as the analysis of suspected unseen satellites of the stars. According to the date on the first of January, 2006 the catalogue of planetary objects has been compiled. ", "introduction": " ", "conclusions": "" }, "0607/astro-ph0607651_arXiv.txt": { "abstract": "We show that there are physically relevant situations where gravitational waves do not inherit the frequency spectrum of their source but its wavenumber spectrum. ", "introduction": "Let us consider a source of (weak) gravitational waves (GWs), an anisotropic stress $\\Pi_{ij}(\\bx,t)$. The induced GWs can be calculated via the linearized Einstein equations. In the case of an isolated, non-relativistic, far away source, one can derive the quadrupole formula (see e.g.~\\cite{straumann}), $r=|\\bx|, c=1$ \\be h_{ij}(\\bx,t) = \\frac{2G}{r}\\ddot Q_{ij}(t-r)~, \\ee where $Q_{ij}$ denotes the quadrupole of the source, and we are considering the perturbed metric $g_{\\mu\\nu}=\\eta_{\\mu\\nu}+h_{\\mu\\nu}$. The wave has the same time dependence as the source. If we have an isolated source which has a harmonic time dependence with frequency $\\omega_s$, $\\Pi_{ij}(\\bx,t) = \\Pi_{ij}(\\bx)e^{-i\\om_s t}$, the wave zone approximation gives, far away from the source, \\be h_{ij}(\\bx,t) = \\frac{4Ge^{-i\\om_s (t-r)}}{r} \\int d^3x'\\,\\Pi_{ij}(\\bx')e^{-i\\om_s \\hat{\\bx}\\cdot\\bx'}~. \\label{wavezone} \\ee Again, $h_{ij}$ has inherited the frequency of the source: it is a spherical wave whose amplitude in direction $\\hat\\bx$ is determined by $\\tilde\\Pi_{ij}(\\bk)$ with $\\bk=\\om_s\\hat\\bx$. As we will show in this brief report, this simple and well known fact from linearized general relativity and, equivalently, electrodynamics, has led to some errors when applied to GW sources of cosmological origin. As an example we consider a first order phase transition in the early universe. This can lead to a period of turbulent motion in the broken phase fluid, giving rise to a GW signal which is in principle observable by the planned space interferometers LISA or BBO \\cite{kamionkowski,kosowsky,dolgov,notari,dolgovgrasso,our,sargent}. One can describe this phase of turbulence in the fluid as a superposition of turbulent eddies: eddies of characteristic size $\\ell$ rotate with frequency $\\om_\\ell \\simeq v/\\ell\\neq 1/\\ell$, since $v<1$. The wavenumber spectrum of the eddies peaks at $k_L\\simeq 1/L$, where $L$ is the stirring scale, and has the usual Kolmogorov shape for $k\\gsim 1/L$ ($\\ell$ 10$^{25}$ W/Hz), characterised by a convex radio spectrum, peaking at frequencies between 100 MHz and a few GHz. Their radio morphologies appear to be the scaled-down version of powerful edge-brightened radio galaxies, with luminous mini-lobes ($\\sim$ 0.1 up to few kpc) and weaker jets and cores. Both kinematic (Polatidis \\& Conway \\cite{PC03}), and spectral (Murgia \\cite{Mu03}) studies strongly support the youth scenario, indicating ages of 10$^{3}$ -- 10$^{5}$ years. Given the relatively high detection rate of \\HI\\ absorption in these objects (Vermeulen et al. \\cite{rv03}; Pihlstr\\\"{o}m et al. \\cite{yp03}), it is of particular interest to investigate the characteristics of the nuclear interstellar medium (ISM) in even younger radio sources. In the youth scenario, the anti-correlation between the turnover frequency and the linear size (O'Dea \\cite{odea98}), which is indicative of the age, suggests that the youngest sources have the highest turnover frequency. Therefore, the ``High-Frequency Peaker'' (HFP) radio sources, characterised by the same properties of CSS/GPS, but with the spectral turnover occurring at frequencies higher than 5 GHz, are good candidates to be {\\it newly born} radio sources, with ages of about 10$^{2}$--10$^{3}$ years (Dallacasa \\cite{dd03}). This paper reports on the result of observations searching for \\HI\\ absorption in a sample of 6 HFP radio sources, selected from the Dallacasa et al. (\\cite{hfp0}) bright HFP sample, and suitable to be observed at the Westerbork Synthesis Radio Telescope (WSRT). \\begin{table*} \\begin{center} \\begin{tabular}{ccccccccccccc} \\hline \\hline &&&&&&&&&&&&\\\\ Source &Other &LS&$z_{\\rm opt}$&$ \\nu_{\\rm obs}$ & Resol.&r.m.s.&S$_{\\rm obs}$&S$_{\\rm HI}$&$\\tau_{\\rm peak}$&$\\Delta$v&Log(N$_{\\rm HI}$)&$z_{\\rm HI,peak}$\\\\ J2000 &name &pc& &MHz&km s$^{-1}$&mJy/b/ch &mJy&mJy& &km s$^{-1}$& & \\\\ (1) &(2) & (3) &(4)&(5)&(6)&(7)&(8)&(9)&(10)&(11)&(12)&(13)\\\\ \\hline & & & & & & & & & & & & \\\\ \\object{J0003+2129}& &22&0.452&977.96&4.3&6.1&50&$<$12.2&$<$0.15& &$<$21.43& \\\\ \\object{J0111+3906}&\\object{OC\\,314}&22$^{a}$&0.668&851.32&6.9&2.8&170&52.7&0.44&100&21.90&0.6687 \\\\ \\object{J0655+4100}& &$<$1&0.02156&1390.03&4.3&0.7&239&$<$1.4&$<$0.006& &$<$20.0& \\\\ \\object{J1407+2827}&\\object{OQ\\,208} &10&0.773&1318.60&4.5&1.1&826&5&0.005&1800&20.9&0.0769\\\\ \\object{J1511+0518}& &7&0.084&1309.96&9.5&1.1&80&$<$2.2&$<$0.02& &$<$20.6&\\\\ \\object{J1623+6624}& &$<$1&0.203&1180.38&4.5&0.7&129&$<$1.4&$<$0.01& &$<$20.3&\\\\ &&&&&&&&&&&&\\\\ \\hline \\end{tabular} \\vspace{0.5cm} \\end{center} \\caption{Physical and observational parameters of the 6 HFP galaxies observed with WSRT. Columns 1, 2: source names; Column 3: projected linear sizes (Orienti et al. 2006); Column 4: optical redshift; Column 5: central frequency; Column 6: channel resolution; Column 7: 1$\\sigma$ noise level in the line cube; Column 8: continuum flux density taken from our WSRT data at the observed frequency; Column 9: peak flux density of the absorption line, measured on the spectral image; Column 10: optical depth; Column 11: the width of the \\HI\\ absorption line: for \\object{J0111+3906} the FWHM is given, in the case of \\object{OQ\\,208} we give the FWZI, due to the complexity of the line profile. Column 12: \\HI\\ column density derived from $N_{\\rm \\HI} = 1.82 \\times 10^{18} T_{\\rm spin} \\tau_{\\rm peak} \\Delta V $ cm$^{-2}$, a T$_{\\rm spin}$ of 100 K has been assumed; Column 13: the redshift of the peak \\HI\\ absorption. The line flux density, the optical depth and the \\HI\\ column density upper limits have been computed assuming the 2$\\sigma$ noise level, a line width of 100 km s$^{-1}$ and T$_{\\rm spin}$ of 100 K, as in Vermeulen et al. (\\cite{rv03}). $a$: For the source \\object{J0111+3906}, the projected linear size is taken from Owsianik et al. (\\cite{ocp98}).} \\label{taboss} \\end{table*} ", "conclusions": "An \\HI\\ absorption search has been carried out with the WSRT for a sample of 6 HFP radio galaxies. We confirm the detection of \\HI\\ in absorption in 2 galaxies. One source, \\object{J0111+3906}, is characterised by a line width of $\\sim$ 100 \\kms\\ and a high optical depth of $\\tau = 0.44$. In the other source (\\object{OQ~208}), the line profile is very broad ($\\sim$ 1800 \\kms), blue-shifted and shallow, with a maximum optical depth of $\\tau = 0.005$. In the remaining 4 galaxies no evidence of \\HI\\ absorption has been detected. Although this result does not seem to follow the inverse correlation found (Pilhstr\\\"{o}m et al. \\cite{yp03}) between the column density and the linear sizes, it can be explained by orientation effects in a torus scenario, in which our line of sight intersects the torus in its inner region where the low optical depth is due to high spin and kinetic temperature. Since these 4 sources have faint flux densities, optical depths $\\leq$ 0.01 are not detectable due to sensitivity limitations. Therefore, the \\HI\\ absorption in an object with the same spectral characteristics of \\object{OQ\\,208} but with a fainter flux density, cannot be detected by our observations. Although HFP sources do not seem to follow the correlation between \\HI\\ column density and linear size found for CSS/GPS sources, this does not imply that we are looking at a different class of objects, instead of the youngest tail of a radio source population. Our results suggest that on linear scales smaller than few tens of parsecs, the \\HI\\ column density is much lower than one would have expected on the basis of the work of Pilhstr\\\"{o}m et al. (\\cite{yp03}). As a consequence, compact and rather faint (due to self-absorption) sources, such as these HFPs, are not the most suitable class of objects to investigate the \\HI\\ absorption on such a small scale." }, "0607/astro-ph0607301_arXiv.txt": { "abstract": "We use the Simon, Verde, \\& Jimenez (2005) determination of the redshift dependence of the Hubble parameter to constrain cosmological parameters in three dark energy cosmological models. We consider the standard $\\Lambda$CDM model, the XCDM parameterization of the dark energy equation of state, and a slowly rolling dark energy scalar field with an inverse power-law potential. The constraints are restrictive, consistent with those derived from Type Ia supernova redshift-magnitude data, and complement those from galaxy cluster gas mass fraction versus redshift data. ", "introduction": "Astrophysical and cosmological data gathered in the last decade strongly support a ``standard'' cosmological model dominated by dark energy. Supernova type Ia (SNIa) redshift-apparent magnitude data show that the Universe is now undergoing accelerated expansion \\citep[e.g.,][]{clocchiatti06, astier06, jassal06, conley06, calvo06, carneiro06}. Cosmic microwave background (CMB) data indicate that the Universe has negligible space curvature \\citep[e.g.,][]{podariu01b, durrer03, mukherjee03, page03, spergel06, baccigalupi06}. Many observations indicate that nonrelativistic matter contributes about $30\\enspace\\%$ of the critical density \\citep[and references therein]{chen03b}. These observational facts --- in the context of general relativity --- indicate that we live in a spatially-flat Universe with about $70\\enspace\\%$ of the total energy density of the Universe today being dark energy, a substance with negative effective pressure responsible for the current accelerated expansion. For reviews see \\citet{peebles03}, \\citet{carroll04}, \\citet{ perivolaropoulos06}, \\citet{padmanabhan06}, and \\citet{uzan06}, and for discussions of the validity of general relativity on cosmological scales see, e.g., \\citet{diazrivera06}, \\citet{stabenau06}, \\citet{sereno06}, and \\citet{caldwell06}. There are many different dark energy models.\\footnote{See \\citet{copeland06} for a recent review. For specific models see, e.g., \\citet{capozziello05}, \\citet{guo06}, \\citet{cannata06}, \\citet{grande06}, \\citet{szydlowski06}, \\citet{ nojiri06}, \\citet{brax06}, \\citet{ calgani06}, and \\citet{ guendelman06}.} Here we consider three simple, widely-used ones: standard $\\Lambda$CDM, the XCDM parameterization of dark energy's equation of state, and a slowly rolling dark energy scalar field with an inverse power-law potential ($\\phi$CDM). In all three cases we assume that the nonrelativistic matter density is dominated by cold dark matter (CDM). In the $\\Lambda$CDM model dark energy is Einstein's cosmological constant $\\Lambda$ and can be accounted for in the energy-momentum tensor as a homogeneous fluid with negative pressure $p_{\\Lambda}=-\\rho_{\\Lambda}$ where $\\rho_{\\Lambda}$ is the cosmological constant energy density \\citep{peebles84}. In the $\\phi$CDM scenario a scalar field $\\phi$ plays the role of dark energy. Here we consider a slowly rolling scalar field with potential energy density $V(\\phi)=\\kappa m_{\\rm p}^2\\phi^{-\\alpha}$ where $m_{\\rm p}$ is Planck's mass and $\\kappa$ and $\\alpha$ are non-negative constants \\citep{peebles88, ratra88}. In the XCDM parameterization dark energy is assumed to be a fluid with pressure $p_{\\rm x}=\\omega_{\\rm x} \\rho_{\\rm x}$ where $\\omega_{\\rm x}$ is time-independent and negative but not necessarily equal to $-1$ as in the $\\Lambda$CDM model. The XCDM parameterization can be used as an approximation of the $\\phi$CDM model in the radiation and matter dominated epochs, but at low redshifts, in the scalar field dominated epoch, a time-independent $\\omega_{\\rm x}$ is an inaccurate approximation \\citep[e.g.,][]{ratra91}. In the $\\phi$CDM and XCDM cases we consider a spatially-flat cosmological model while spatial curvature is allowed to be non-zero in the $\\Lambda$CDM case. We note that the $\\phi$CDM model at $\\alpha=0$ and the XCDM parameterization at $\\omega_{\\rm x}=-1$ are equivalent to a spatially-flat $\\Lambda$CDM model with the same matter density. Besides SNIa and CMB anisotropy, there are many other cosmological tests. Having many tests is important since this allows for consistency checks, and combined together they provide tighter constraints on cosmological parameters. Tests under current discussion include the redshift-angular size test \\citep[e.g.,][]{chen03a, podariu03, puetzfeld05, daly05, jackson06}, the galaxy cluster gas mass fraction versus redshift test \\citep{sasaki96, pen97, allen04, chen04, kravtsov05, laroque06}, the strong gravitational lensing test \\citep{fukugita90, turner90, ratra92, chae04, kochanek04, biesiada06}, the baryonic acoustic oscillation test \\citep[e.g.,][]{glazebrook05, angulo05, wang06, zhan06}, and the structure formation test \\citep[e.g.,][]{brax05, koivosta05, maor06, bertschinger06, mainini06}. For cosmological constraints from combinations of data sets see, e.g., \\citet{wilson06}, \\citet{wang06a}, \\citet{rahvar06}, \\citet{seljak06}, \\citet{xia06}, and \\citet{rapetti06}. Here we use a measurement of the Hubble parameter as a function of redshift to derive constraints on cosmological parameters \\citep{jimenez02}. \\citep[For related techniques see][and references therein.]{shafieloo05, daly05} In our analysis we use the Simon et al. (2005, hereafter SVJ) estimate for the redshift, $z$, dependence of the Hubble parameter, \\begin{equation} \\label{difage} H(z)=-\\frac{1}{1+z}\\frac{dz}{dt}, \\end{equation} \\noindent where $t$ is time. This estimate is based on differential ages, $dt/dz$, of passively evolving galaxies determined from the Gemini Deep Deep Survey \\citep{abraham04} and archival data \\citep{dunlop96, spinrad97, treu01, treu02, nolan03}. SVJ use the estimated $H(z)$ to constrain the dark energy potential and it's redshift dependence. This data has also been used to constrain parameters of holographic dark energy models \\citep{yi06}. Here we use the SVJ $H(z)$ data to derive constraints on cosmological parameters of the $\\Lambda$CDM, XCDM, and $\\phi$CDM models. In the next section we outline our computation, in $\\S\\:3$ we present and discuss our results, and we conclude in $\\S\\:4$. ", "conclusions": "We have used the SVJ Hubble parameter versus redshift data to constrain cosmological parameters of three dark energy models. The constraints are restrictive, and consistent with those determined by using Type Ia supernova redshift-magnitude data. The $H(z)$ data constraints complement those determined from galaxy cluster gas mass fraction versus redshift data. In combination with improved SNIa data \\citep[from, e.g., JDEM/SNAP, see, http://snap.lbl.gov/;][and references therein]{podariu01a, crotts05, albert05}, more and better $H(z)$ data will tightly constrain cosmological parameters. A large amount of $H(z)$ data is expected to become available in the next few years (R. Jimenez, private communication 2006). These include data from the AGN and Galaxy Survey (AGES) and the Atacama Cosmology Telescope (ACT), and by 2009 an order of magnitude increase in $H(z)$ data is anticipated. We acknowledge valuable discussions with Raul Jimenez, helpful comments from the referee, and support from DOE grant DE-FG03-99ER41093." }, "0607/astro-ph0607071_arXiv.txt": { "abstract": "We report spectroscopy of the newly discovered SU Ursae Majoris dwarf nova identified with the x-ray source RXS J053234.9+624755. Radial velocities of the H$\\alpha$ emission line in the quiescent state give an orbital period of 0.05620(4) d (80.93 min), which is among the shortest for SU UMa stars with determined periods. We also report UBVI magnitudes of the quiescent dwarf nova and surrounding stars. Using a previous measurement of the superhump period, we find the fractional superhump excess $\\epsilon$ to be 0.016(4), which is not atypical of dwarf novae in this period range. ", "introduction": "Cataclysmic Variables (CVs) are close binary systems that consist of an accreting white dwarf (the primary) and a secondary component that usually resembles a main-sequence star. \\citet{warn} comprehensively reviews CVs. Dwarf novae, or U Geminorum stars, are a subclass of CVs that can be further subclassified based on their outburst behavior. The SU Ursae Majoris stars (referred to as UGSU) form a subclass of dwarf novae that undergo occasional superoutbursts in addition to normal outbursts. Superoutbursts occur less frequently than normal outbursts, but are brighter and last longer. During a superoutburst, characteristic oscillations called ``superhumps'' develop in the light curves. The measured superhump period, $P_{\\rm sh}$, is usually a few percent longer than the measured orbital period, $P_{\\rm orb}$. Almost all known SU UMa stars have $P_{\\rm orb} < 2$ hr. The superhump period excess, $\\epsilon =$[($P_{\\rm sh}-P_{\\rm orb}$/$P_{\\rm orb}$], is an important quantity for these stars. \\citet{patt01} demonstrated that $\\epsilon$ correlates well with the mass ratio $q=M_2/M_1$, which is otherwise difficult to obtain. Another class of dwarf novae are the WZ Sagittae stars (UGWZ), which are extreme examples of the SU UMa-type stars. WZ Sge stars exhibit large-amplitude outbursts ($\\geq$ 6 mag) that occur less frequently than the those of the UGSU (\\citealt{osaki}). WZ Sge stars do exhibit superhumps in their light curves, but do not appear to have any `normal' outbursts. Here we present observations of the dwarf nova identified with the ROSAT x-ray source RXS J053234.9+624755 (hereafter RX0532+62). The discovery of this star, which lies in Camelopardalis, is described by \\citet{poy}. \\citet{bernhard} classified it as a U Gem-type dwarf novae with a recurrence time scale of 133 days. The relatively frequent outbursts show that this star is not a UGWZ. RX0532+62 underwent a well-observed superoutburst in 2005 March. During this outburst a superhump period, $P_{\\rm sh}$, of 0.0571(2) d (82.2 mins) was reported by Tonny Vanmunster (CBA Belgium) \\footnote{``Detection of Superhumps in the CV 1RXS J053234.9+624755'' can be found at http:$\\rm//users.skynet.be/fa079980/cv\\_2005/1RXSJ053234\\_2005\\_mar\\_18.htm$}; \\citet{poy} independently found a similar, but less accurate, value of $P_{\\rm sh}$. We undertook observations of this star to independently determine the orbital period, $P_{\\rm orb}$, and allow determination of the superhump period excess, $\\epsilon$. ", "conclusions": "Combining our $P_{\\rm orb}$ with the previously measured $P_{\\rm sh}$, we find $\\epsilon = 0.016(4)$. \\citet{patt03} plot log($\\epsilon$) against log($P_{\\rm orb}$) for a large number of systems with hydrogen-rich secondaries. On this plot, RX0532+62 lies near the short-period end, where systems appear to be evolving through the period minimum. These occupy a relatively wide range of $\\epsilon$ values. RX0532+62 is toward the top of this range, as if it has evolved into the turnaround region relatively recently. \\citet{patt01} fits the empirical relation between $\\epsilon$ and the mass ratio $q$ as $\\epsilon(q)=0.216q$. Using our $\\epsilon$ value, we determine $q=0.074(19)$, which is typical for dwarf novae with similar periods\\footnote{The uncertainty in $q$ is computed here using only the uncertainty in $\\epsilon$; imperfections in the empirical relation are ignored.}. We conclude that RX0532+62 is a typical SU UMa star lying near the minimum period for hydrogen-rich secondary systems. {\\it Acknowledgments}: We thank Holly Sheets for taking the 2006 January spectra. The National Science Foundation funded this research through award AST-0307413 and an REU supplement to that award. Travel for Ann Kapusta was made possible by a generous gift from Claudia and Jay Weed. \\clearpage" }, "0607/astro-ph0607592_arXiv.txt": { "abstract": "{} {We studied and compared the long-term average hard X-ray ($>$20\\,keV) spectra of a sample of twelve bright low-mass X-ray binaries hosting a neutron star (NS). Our sample comprises the six well studied Galactic Z sources and six Atoll sources, four of which are bright (\"GX\") bulge sources while two are weaker ones in the 2--10\\,keV range (H~1750--440 and H~1608--55). } { For all the sources of our sample, we analysed available public data and extracted average spectra from the IBIS/ISGRI detector on board \\textit{INTEGRAL}. } { We can describe all the spectral states in terms of the bulk motion Comptonisation scenario. We find evidence that bulk motion is always present, its strength is related to the accretion rate and it is suppressed only in the presence of high local luminosity. The two low-dim Atoll source spectra are dominated by photons up-scattered presumably due to dynamical and thermal Comptonisation in an optically thin, hot plasma. For the first time, we extend the detection of H~1750--440 up to 150\\,keV. The Z and bright \"GX\" Atoll source spectra are very similar and are dominated by Comptonised blackbody radiation of seed photons, presumably coming from the accretion disc and NS surface, in an optically thick cloud with plasma temperature in the range of 2.5--3\\,keV. Six sources show a hard tail in their \\emph{average} spectrum: Cyg~X-2 (Z), GX~340$+$0 (Z), GX~17$+$2 (Z), GX~5--1 (Z), Sco~X--1 (Z) and GX~13$+$1 (Atoll). This is the first detection of a hard tail in the X-ray spectrum of the peculiar GX~13$+$1. Using radio data from the literature we find, in all Z sources and bright \"GX\" Atolls, a \\emph{systematic} positive correlation between the X-ray hard tail (40--100\\,keV) and the radio luminosity. This suggests that hard tails and energetic electrons causing the radio emission may have the same origin, most likely the Compton cloud located inside the NS magnetosphere. } {} ", "introduction": "Low-Mass X-ray Binaries (LMXBs) are systems where a compact object, either a neutron star (NS) or a black hole candidate (BHC), accretes matter via Roche lobe overflow from a companion with a mass \\mbox{$\\textit{M}\\lesssim$ $1\\textit{M}_{\\odot}$}. NS LMXBs can be broadly classified according to their timing and spectral properties \\citep{hasinger89}. On the basis of this classification, NS LMXBs are divided in Z sources and Atoll sources from the shape of their track in the colour-colour diagram and from the different timing behaviour that correlates with the position on the tracks. The overall spectra of Z sources are very soft \\citep[][and references therein]{barret02} and can be described by the sum of a cool ($\\sim$1\\,keV) blackbody (BB) and a Comptonised emission from an electron plasma (\"corona\") of a few keV. Instead, Atoll sources perform quite dramatic spectral changes: when bright, they can have soft spectra (similar to Z sources) but they switch to low/hard spectra at low luminosities. \\cite{titarchuk05}, hereafter TS05, implemented a thorough analysis of spectral and temporal properties of the Atoll source 4U~1728--34. They show that the low/hard spectra at low luminosities can be described by the sum of up-scattered spectra related to the Comptonisation of the disc and NS surface soft photons. They found that the Compton cloud electron temperature is of the order of a few tens of keV. These spectra are very similar to the hard-state spectra of BHCs but they are softer (BHC photon index $\\Gamma \\sim 1.6\\pm 0.1$ vs NS $\\Gamma\\sim 2.1\\pm 0.1$) as expected from the theory \\citep[see, for example,][]{titarchuk04}. TS05 found that the high luminosity state spectrum of 4U~1728--34 consists of the sum of two pure blackbody-like spectra with colour temperatures of about 1\\,keV and 2.2\\,keV. The softer BB component is presumably related to disc emission as the harder one is related to the NS emission. It is worth noting that TS05 also found that when the source undergoes hard-soft transition, all power spectrum (PDS) frequencies (QPO and break frequencies) increase with the photon index, with no sign of saturation. Note that the index-QPO correlation observed in BHC shows the index saturation at high values of QPO frequency \\citep[see][]{shaposhnikov06}. So far, only Atoll sources (and more generally X-ray bursters) have been observed with low/hard spectra (i.e. Comptonising corona of few \\emph{tens} of keV). Z sources always have soft Comptonisation spectra (Comptonising corona of few keV) and can have an additional hard X-ray component dominating the spectrum above $\\sim$30\\,keV. This component is on top of the soft spectrum and is highly variable with most of the emission remaining soft \\citep[see][for a review on NS LMXB spectra]{barret01, disalvo02}. Hence, we would expect that the average high energy spectra of Atoll sources have a strong component above 30\\,keV and that in the soft Z sources this component is less prominent and smeared out in the time averaged spectrum. Z sources spend most of their time in the high/soft state, but they may show the transition to harder states at lower luminosity. The sensitivity of the past missions may have introduced an observational bias, similarly to the lack of a continuous coverage of the Galactic plane and Centre in the less explored hard X-ray range (above 20\\,keV). Moreover, the concentration of these sources towards the Galactic centre makes it difficult to observe them with non-imaging instruments. Consequently, data analysis and interpretation of such observations is extremely problematic. \\begin{figure*} \\centering \\includegraphics[width=1.0\\linewidth]{./5792f1.ps} \\caption{IBIS/ISGRI 22--40\\,keV mosaic of the Galactic bulge (about 5\\,Msec). Only the sources studied in this paper are labeled. The location of Cyg~X--2 and Sco~X--1 is not covered by this map. \\label{fig:ima1}} \\end{figure*} All these instrumental biases can be minimised with the use of the recently launched INTErnational Gamma-Ray Astrophysics Laboratory, \\textit{INTEGRAL} \\citep{winkler03}. The imager \\textit{INTEGRAL}/IBIS \\citep{ubertini03} has high sensitivity, about $\\sim$10 times better than \\textit{GRANAT}/SIGMA, coupled to imaging capability with 12$^{\\prime}$ angular resolution above 20\\,keV. In this paper we report the study of the average hard X-ray spectra of twelve NS LMXBs performed with the low energy (20--200\\,keV) IBIS detector, ISGRI \\citep{lebrun03}, using a coherent and large sample of data, free from systematic effects which play a role when combining data from different missions. The sample of the LMXBs chosen is given in Table~\\ref{tab:table1} and comprises six Galactic Z sources and six Atoll sources, four of which are bright (\"GX\") bulge sources while two are weaker ones in the 2-10\\,keV range. Our approach is two-fold: on one side, for comparison purposes, we study the average spectra in terms of phenomenological models as done in the literature, on the other, we study the sources in the frame of a physical model in the attempt to find a self-consistent scenario that describes all the spectral properties we observe. We discuss the similarities of such a scenario with the black hole LMXB case as well as the radio - X-ray correlation that is typical of LMXBs. \\begin{table} \\begin{center} \\caption{LMXBs studied in this paper. \\emph{D}: distance (in kpc) from references in \\cite{migliari06} except for (*) from \\cite{christian97} and (**) from \\cite{ford00}; \\emph{Rate}: average \\mbox{22--40\\,keV} counts/sec of the source as obtained from the mosaic image shown in Fig.~\\ref{fig:ima1}. Multiply by $\\sim$10 to obtain a flux estimate in units of mCrab. F(1mCrab)$_{22-40\\,keV}$ $\\sim$6.8$\\times$10$^{-12}$\\rm\\,erg\\,s$^{-1}$\\,cm$^{-2}$; \\emph{SNR}: signal to noise ratio in the 22--40\\,keV band; \\emph{MaxEn}: maximum energy channel (keV) with a signal to noise ratio higher than three in the average spectrum; \\emph{T$_{exp}$}: effective exposure time in ksec. }\\vspace{1em} \\renewcommand{\\arraystretch}{1.2} \\begin{tabular}[h]{cccccc} \\hline Source & D & Rate & SNR & MaxEn & T$_{exp}$\\\\ &(kpc) & (cps) & & (keV)\t& (ksec)\t\\\\ \\hline \\hline Z sources &\t\\\\ \\hline Sco~X--1& 2.8 & 58.3 & 1757 & 150 & 266\\\\ GX~340$+$0 & 11& 2.4 & 124 & 46 & 433\\\\ GX~349$+$2 & 5& 3.4& 189 & 39 & 265\\\\ GX~5--1 & 9.2 & 3.8 &330 & 80 & 1091\\\\ GX~17$+$2 &14 & 4.4& 203 & 80 & 248\\\\ Cyg~X--2 & 13.3& 2.2 & 57& 55& 149\\\\ \\hline Atoll sources &\\\\ \\hline H~1608--522 & 4** &0.8 &36 & 150 & 299\\\\ H~1705--440 &11** & 3 &144 & 150 & 307\\\\ GX~9$+$9& 5* &0.97 & 55 & 35 & 139\\\\ GX~3$+$1& 5.6*& 0.96 & 86 & 43 &2027\\\\ GX~9$+$1& 7* &1.2&91 & 37 & 491\\\\ GX~13$+$1 & 7 &0.96 & 57 & 80 & 290\\\\ \\hline \\end{tabular} \\label{tab:table1} \\end{center} \\end{table} ", "conclusions": "We have analysed IBIS/ISGRI available public data on a sample of twelve persistent LMXBs containing a neutron star. We focused our study on the average spectral behaviour of the sources and classified them in terms of spectral states. As shown by the light-curves we presented, the sources have some degree of variability, but the average spectrum is a representation of the sources above 20\\,keV and enables the study of the global properties. \\subsection{A scenario for spectral evolution in NS LMXBs} We observe three main spectral states: a \\emph{very soft} state with source spectra that can be well described by a single thermal Comptonisation component, an \\emph{intermediate} state (thermal Comptonisation plus power-law) and a \\emph{low/hard} state (single power-law). We have successfully studied these three spectral states in the frame of the Generic Comptonisation Model (\\BMC, TMK96). We present our scenario of the spectral evolution of NS LMXBs starting from the low/hard state. The {\\it low/hard state} is characterized by a low mass accretion rate in the disc. In this case the gravitational energy release in the disc is much smaller than the one in the optically thin outer boundary of the corona (Compton cloud). The coronal outer boundary is presumably related to the adjustment shock \\citep{titarchuk04}. The corona completely covers the seed photon area (high $\\log A$) and the emergent spectrum is a result of the up-scattering (Comptonisation) of the seed photons in the corona. At the very low level of accretion rate, the bulk motion is extremely weak and the emergent spectrum can be interpreted as pure thermal Comptonisation with no bulk signature (GX~354$-$0 in Fig.~\\ref{fig:comp}). The bulk motion effect increases with accretion rate and the power-law signature starts to be visible in the spectrum, the cut-off is moved at a higher energy than what expected from pure thermal Comptonisation (H~1608-522 in Fig.~\\ref{fig:comp}). In either case, one cannot see any trace of the seed photons in ISGRI (that give the BB bump). Open magnetic field lines are exposed to the observer but the outflow is weak because of the low mass accretion rate (there is not enough radiation in the disc to launch the wind) and the system is radio-quiet. In the {\\it intermediate state} the mass accretion rate increases with respect to the low/hard state. It leads to high efficiency of the Comptonisation, particularly bulk inflow Comptonisation that is seen as an extended hard tail in the spectrum. Thermal Comptonisation becomes less efficient because the coronal plasma is cooled down by the seed photons coming from the disc and NS surface. The corona consists of a quasi-spherical component (related to the closed field lines and bulk motion inflow) and of a cylindrical component (related to the open field lines and outflow, TS05). The vertical size of the cylindrical configuration is suppressed as the mass accretion increases. We start to see the seed blackbody bumps (only the higher energy one in the IBIS/ISGRI range), because the corona is cooled down and becomes more compact. The {\\it very soft} state is characterized by a high mass accretion rate that is very close to the critical (Eddington) values. The emergent spectrum is a sum of two blackbody-like spectra, one is related to the Comptonisation of the NS photons (visible in the IBIS/ISGRI range) and the other one is related to the Comptonisation of the disc photons. The seed photon and plasma temperatures differ by a factor of a few. The electron plasma and the photons in the Compton cloud are very close to equilibrium. In this state the corona is quasi-spherical, there is no bulk motion and no radio emission. In fact, the radiation pressure caused by the strong emission from the NS surface stops the bulk inflow and the high accretion rate changes the configuration of the field lines and the radio emission is quenched. It is important to note that what plays the key role to suppress the bulk motion is the \\emph{local} radiation pressure, hence the \\emph{local} accretion rate impinging the bulk inflow: Sco~X-1 has a total isotropically-estimated luminosity that is higher than e.g. GX~3$+$1, but we speculate that the inflow anisotropy is much higher in GX~3$+$1 than in Sco~X-1, resulting in a higher local radiation pressure, i.e. bulk suppression. \\begin{table}[t] \\caption{Comparison of the \\emph{average} radio (8.5\\,GHz) and X-ray hard tail (40--100\\,keV) flux of the Z and bright Atoll sources studied in this paper. The radio data are taken from \\cite{fender00} and \\cite{migliari06}. } \\label{tab:radio} \\begin{center} \\begin{tabular}{lccc} \\hline Source & Radio flux & Hard tail flux \\\\ & (mJy) & 10$^{-11}$\\rm\\,erg\\,s$^{-1}$\\,cm$^{-2}$ \\\\ \\hline \\hline Z sources & & & \\\\ \\hline Sco~X--1 & $10 \\pm 3$ & 18 \\\\ GX~17$+$2 & $1.0 \\pm 0.3$ &1.9 \\\\ GX~349$+$2 & $0.6 \\pm 0.3$ & $<$1.15 \\\\ Cyg~X--2 & $0.6 \\pm 0.2$ & 1.7 \\\\ GX~5--1 & $1.3 \\pm 0.3$ & 1.9\\\\ GX~340$+$0 & $0.6 \\pm 0.3$ & 2\\\\ \\hline Atolls & \\\\ \\hline GX~9$+$1 & $<$0.2 & $<$0.75 \\\\ GX~9$+$9 & $<$0.2 & $<$1.72 \\\\ GX~3$+$1 & $<$0.3 & $<$0.43 \\\\ GX~13$+$1 & $1.8 \\pm 0.3$ & 1.74 \\\\ \\hline \\end{tabular} \\end{center} \\end{table} \\begin{figure} \\centering \\includegraphics[width=1.0\\linewidth]{5792f27.ps} \\caption{Average radio luminosity plotted against the average hard \\mbox{X-ray} tail (40--100\\,keV) luminosity. The correlation is clearly visible also if we compare the source fluxes instead of the luminosities (see text). \\label{fig:corr}} \\end{figure} \\subsection{The radio emission - hard X-ray tail connection} In the scenario proposed in the previous Section there is a clear connection between the X-ray spectral states and the radio emission: low/hard states are associated to weak radio emission, the radio emission increases in the intermediate state and then is quenched in the very soft state. This trend is clearly met in the observations. We can imagine a continuous \"accretion line\" increasing from Atoll sources (island state to banana) moving to the Z sources (horizontal, normal and flaring branch). Only a few low-dim Atoll sources (typical low/hard state) have been detected in the radio band because of their low radio luminosity \\citep[][]{fender00, migliari06}. Z sources in their horizontal branch have a spectrum that corresponds to the intermediate state we defined here and indeed a clear radio emission has been reported in this state \\citep[][and references therein]{fender00, disalvo02}. The normal and flaring branch of Z sources corresponds to the very soft state defined here where no radio emission is expected and in fact Z sources in these branches have a severely reduced radio emission, if any. Z sources are much brighter radio emitters than Atolls and GX~13$+$1 has an important radio emission, thus behaves like a Z \\citep[][]{fender00, migliari06}. In average, unlike bright Atolls, Z sources have a hard tail besides the thermal Comptonisation by an optically thick plasma of $\\sim$3\\,keV \\citep[][and this work]{disalvo02}; GX~13$+$1 has a hard tail (this work), thus again behaves like a Z. We basically see that in GX~13$+$1, the link radio emission - hard X-ray tail (both \"against\" its originally declared nature to be an Atoll source) reveals itself in a solid way. Triggered by the GX~13$+$1 case, we compared the average radio emission of all the sources of our sample with the hard X-ray tail that appears on top of the $\\sim$3\\,keV Comptonisation spectrum. We note that the hard X-rays from the two soft X-ray dim Atoll sources, \\mbox{H~1608--522} and \\mbox{H~1705--440}, are mainly coming from Comptonisation of soft photons in a hot corona (i.e. the overall spectrum is hard, kT$_{e}$ $>$ 40\\,keV, see Table~\\ref{tab:fit}) and since we cannot disentangle the dominating thermal Comptonisation component from the dynamical (bulk) one, we do not include these two sources in the radio - hard tail study. Table~\\ref{tab:radio} reports the average radio flux as derived by \\cite{fender00} and \\cite{migliari06} and the 40--100\\,keV flux of the hard tails detected in this work. When a hard tail is detected, its 40--100\\,keV flux is computed after fixing the \\comptt~normalization to zero. For the sources where a hard tail was not detected, we computed 3$\\sigma$ upper limits (3$\\times$$error$ in the 40--100\\,keV range), assuming that all the flux in the 40--100\\,keV comes from the hard tail. Figure~\\ref{fig:corr} is visualising the relation between the average radio and hard tail \\mbox{(40--100\\,keV)} \\emph{luminosities} computed using the distances in Table~\\ref{tab:table1}. The correlation between the two is quite striking (Spearman rank order significance of 99.6\\%) and the upper limits are consistent with the trend. We note that we obtain a very good correlation (Spearman rank order significance 99.2\\%) also using the radio and hard tail \\emph{fluxes} instead of the luminosities. We chose to plot the luminosity correlation in order to show the intrinsic properties of the sources. In Fig.~\\ref{fig:corr} we have plotted the average radio and X-ray hard tail luminosities: the radio luminosities are the superposition of optically thick emission (compact jet) and optically thin flaring activity. The X-ray hard tail luminosities are the average of all hard tail states (from maximum to minimum strength). We are aware that these correlations are based on data that are not taken simultaneously, but they are all based on average fluxes that we can consider a good representation of the physics involved in these objects. For the sources where a hard tail has been observed, we compare the radio emission with what is \\emph{left} in the X-ray spectra of the sources once we \\emph{remove} the dominant Comptonisation component. In this respect, our detection of a hard tail in the spectrum of GX~13$+$1 is important: GX~13$+$1 has a similar X-ray flux (and luminosity) as the remaining bright Atoll sources (GX~3$+$1, GX~9$+$1, GX~9$+$9, see Table~\\ref{tab:table1}) but a much higher average radio and hard tail emission (Table~\\ref{tab:radio}), re-confirming the radio hard X-ray tail connection. In the presence of a cut-off, the hard tail can be explained via Comptonisation of soft seed photons in the jet and/or in the corona. For the case where no cut-off is detected, many models have been proposed: Comptonisation by a hybrid (thermal non-thermal) corona \\citep{coppi99}, synchrotron emission from the electrons of the jet \\citep{markoff05} or bulk motion inflow (dynamical) Comptonisation (TMK96, used in this work). The real test of any of these models can be done using the (variability) analysis of the power density spectrum (PDS) of the hard tail emission. The characteristic (break and QPO) frequencies of PDS do determine the geometric size of the configuration where the hard tail emission is formed. TS05 made this type of analysis for 4U~1728--34 and found evidence that the hard tails are formed in the compact Compton cloud with geometry changing from cylindrical-like in the low/hard state to the quasi-spherical one in the high/soft state. Note that the X-ray-radio correlation along with the QPO-radio correlation is well established in 4U~1728--34. TS05 presented an explanation of these correlations in the framework of an oscillation model using the observed correlations of QPO low frequencies and their ratio. The correlation we find between the radio and hard X-ray tail emission suggests that the hard tail formation area and the source of energetic electrons, ultimately causing the radio emission, are closely connected. The most probable site of this configuration is the NS magnetosphere. It can be suggested that the open magnetic field lines of the NS magnetosphere are the base of the jet seen in the radio emission. An increasing accretion rate leads to a more efficient radio emission (low/hard state to intermediate state) up to a point where the extremely high accretion rate (very soft state) changes the configuration of the field lines and the radio emission is quenched." }, "0607/astro-ph0607247_arXiv.txt": { "abstract": "s{Three different analysis techniques for Atmospheric Imaging System are presented. The classical {\\it Hillas parameters} based technique is shown to be robust and efficient, but more elaborate techniques can improve the sensitivity of the analysis. A comparison of the different analysis techniques shows that they use different information for gamma-hadron separation, and that it is possible to combine their qualities.} ", "introduction": "From the beginning of ground based gamma ray astronomy, data analysis techniques were mostly based on the ``Hillas parametrisation'' \\cite{Denaurois-Hillas-1985} of the shower images, relying on the fact that the gamma-ray images in the camera focal plane are, to a good approximation, elliptical in shape. More elaborate analysis techniques were pioneered by the work of the CAT collaboration on a model analysis technique, where the shower images are compared to a more realistic pre-calculated model of image. Other analysis techniques, such as the {\\it 3D Model analysis} were developed more recently with the start of the third-generation telescopes. The {\\it 3D Model analysis} is, for instance, based on the assumption of a 3 dimensional elliptical shape of the photosphere. These analysis techniques are complementary in many senses. We will show that they are sensitive to different properties of the shower, and can therefore be used to cross-check the analysis results or be combined together to improve the sensitivity. No analysis is currently really winning the race, and there is much space for further improvements. ", "conclusions": "We have presented three completely different analysis methods for Atmospheric Cerenkov Telescopes. These three methods show similar efficiencies, although they are sensitive to different properties of the shower. The intrinsic capabilities of each analysis (and it particular the hadronic rejection capabilities) can be combined together to improve the sensitivity of the analysis. Since these three analyses perform differently in different energy and impact parameter domain, more detailed studies should also allow to use the select on an event-per-event basis the optimal response and therefore improve the quality (angular resolution,...) of the analysis." }, "0607/astro-ph0607253_arXiv.txt": { "abstract": "In this work we study the physical and kinematical properties of the emission line region of Seyfert 1.5 galaxy Mrk 817 using three sets of observations, among which are high-resolution spectra obtained with the Isaac Newton Telescope on Canary Islands. We find that in Mrk 817 the Narrow (NEL) and Broad Emission Lines (BEL) are very complex, indicating that structure of both the Narrow (NLR) and Broad Line Region (BLR) is {\\bf complex} and consists of at least two sub-regions with different kinematical properties. We find that the BEL can be fitted with the two-component model, where the core of the line is coming from a spherical region with isotropic velocity distribution, and wings might be affected by a low inclined accretion disc (or disc-like emitting region). Also, we discuss the physical properties of the BLR. Moreover, we find that an outflow is present in the NLR, that may be driven by an approaching jet. ", "introduction": "The emission line region of Active Galactic Nuclei (AGN) is complex and is usually divided to the Narrow (NLR) and Broad Line Region (BLR), where the physics of the NLR is better understood then that of the BLR (see e.g. Sulentic et al. 2000). The most accepted scenario of the structure of AGN is the one in which AGN are powered by the accretion of matter from the host galaxy on to super-massive black hole. One of the way to study the inner emitting region of an AGN, one that is closest to the black hole, is by analyzing its broad emission lines. So far, in a small fraction of AGN double-peaked emission lines were detected (around 5 \\%, see Eracleous \\& Halpern~2003). Modeling these lines gave proof to the presence of an accretion disc in the AGN (Eracleous \\& Halpern~1994,~2003). Beside the disc, emission lines also imply presence of more kinematically different emission regions that contribute to formation of lines: complex broad and narrow line regions (Popovi\\'c et al.~2003). Even there are numerous papers devoted to the studies of the kinematical and physical properties of the NLR and BLR (Krolik ~1999, Kembhavi \\& Narlikar~1999, Sulentic et al.~2000), it is not clear yet what is the connection between these two kinematically different regions. One of the method to study the connection is to map a whole emission region of an AGN. For that one should have a high resolution spectra of the object that covers wide wavelength band. It is also needed that the AGN emits both narrow and broad lines. Accordingly, the spectroscopical investigations of Seyfert 1 galaxies that have strong narrow lines (as e.g. Seyfert 1.5) are important. Following these reasons, for our analysis we selected Mrk 817, that is a Seyfert 1.5 galaxy, with the redshift of 0.03145, {\\bf (Strauss \\& Huchra~1988) determined from the emission lines}\\footnote{\\bf A slightly different value of 0.031158, determined from both absorption and emission lines, can be found in Falco et al.~(1999), but since the redshift is not crucial for this study we adopt the value of 0.03145 for the systemic redshift of the galaxy.}, and which both broad and narrow emission lines are complex (Popovi\\'c \\& Mediavilla~1997; Peterson et al.~1998; Popovi\\'c et al.~2004). Also, for this galaxy the mass ($M_{\\rm BH}\\approx 4.9 \\times 10^7 M_{\\odot}$) and the BLR size ($R_{\\rm BLR}\\approx 15$ light days\\footnote{The radius of 15 light days corresponds to $3400 \\ R_{\\rm g}$, where $R_{\\rm g}=GM/c^2$ is the gravitational radius for the black hole mass of $4.9 \\times 10^7 M_{\\odot}$ (G is the gravitational constant, M is the mass of the black hole and c is the speed of light).}) have been estimated by reverberation mapping studies (Peterson et al.~1998,~2004; Kaspi et al.~2000). We observed the galaxy several times collecting the high resolution spectra in the H$\\alpha$ and H$\\beta$ wavelength band, as well as the low resolution spectra in the wide wavelength region. The aim of this work is to explore the properties of the whole emission region of the active galaxy Mrk 817. Using the high-resolution spectra (such as one obtained with the Isaac Newton Telescope) in analyzing the broad spectral line shapes and applying the two-component model of the BLR proposed by Popovi\\'c et al. ~(2004), we investigate the kinematical parameters of the BLR in Mrk 817. First we will present the Gaussian analysis of the H$\\alpha$ and H$\\beta$ lines and after that we will apply the two-component model for fitting the broad emission lines. Also, an estimate of the electron temperature in the BLR will be given. Finally we make a scheme of the emission line region of Mrk 817 and discuss the complex structure of that region. ", "conclusions": "In this contribution, we have reported the results of the emission line regions study of the Seyfert 1.5 galaxy Mrk 817. The Seyfert 1.5 galaxies are convenient for exploring the narrow and broad line emission regions since those galaxies have very strong both narrow and broad emission lines. In our work we have used several sets of different spectral observations with high spectral resolution. Our main finding is that both NLR and BLR are complex emission regions, where they are composed of at least two kinematically separated regions. We made a scheme of the emission line regions of Mrk 817 (Figure~\\ref{fig09}). We found that the BLR of Mrk 817 can be described with the two-component model, where the core of the line is coming from a spherical region with isotropic velocity distribution, and wings are probably affected by a low inclined accretion disc. {\\bf The ratio of the fluxes coming from the disc and from the spherical region did not depend on the parameters of the fit, it was almost the same in evert fit and equal to unity}. In the NLR, we found that an outflow is present. In principle we can conclude that emission line region of Mrk 817 is complex, heaving at least four kinematically and physically different regions." }, "0607/astro-ph0607586_arXiv.txt": { "abstract": "We propose an evolutional scenario of the universe which starts from quantum states with conformal invariance, passing through the inflationary era, and then makes transition to the conventional Einstein space-time. The space-time dynamics is derived from the renormalizable higher-derivative quantum gravity on the basis of a conformal gravity in four dimensions. Based on the linear perturbation theory in the inflationary background, we simulate evolutions of gravitational scalar, vector and tensor modes, and evaluate the spectra at the transition point located at the beginning of the big bang. The obtained spectra cover the range of the primordial spectra for explaining the anisotropies in the homogeneous CMB. \\vspace{5mm} \\noindent PACS: 98.80.Cq, 98.80.Qc, 04.60.-m, 98.70.Vc \\noindent Keywords: CMB anisotropies, inflation, quantum gravity ", "introduction": "\\setcounter{equation}{0} \\setcounter{figure}{0} \\noindent After passing many theoretical as well as experimental tests, the general theory of relativity has been established as the fundamental theory of gravity capable of describing the universe. On the other hand, tracing the history of the universe, the space-time in its early epoch would be totally fluctuating quantum mechanically so that geometry lose its classical meaning. This would imply that there is a transition from quantum space-time to classical space-time, and there should exist a dynamical scale separating these two phases. There is a possibility to observe the instance of the transition, because we can trace the past guided by known physical laws as far as the classical general relativity holds. Valuable information on physical processes taking place in the expanding universe has been recorded in the cosmic microwave background radiation (CMB) as tiny anisotropies. Angular power spectra of the anisotropies, recently observed by the Cosmic Background Explorer (COBE) \\cite{cobe} and the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{wmap,wmap3}, are, roughly speaking, projection of the history of universe for the period between its birth to the present. It gives us an amazing hope that if we believe the idea of inflationary universe \\cite{guth, starobinsky} meaning an extremely rapid expansion without global thermalization, the long-distance correlation in observed anisotropies can provide information about dynamics of the period before the universe grew to the Planck scale. We are now at the stage of revealing and verifying the quantum aspect of universe. The model of space-time transition proposed in this paper emerges from the renormalizable higher-derivative quantum theory of gravity developed on the basis of the conformal gravity in four dimensions. In this theory, at very high energies beyond the Planck scale, quantum fluctuations of the conformal mode in the metric field are dominated, and it is treated non-perturbatively. The space-time is described by a conformal field theory whose dynamics is governed by conformal invariant gravitational actions. The conformal field theory loses its validity at the dynamical scale indicated by the asymptotic freedom of the unique dimensionless coupling constant introduced for the traceless tensor mode. At about this point the universe is expected to make a transition from the quantum space-time to the classical Einstein space-time. There are three mass scales in the model, namely the Planck mass $M_\\P$, the dynamical scale $\\Lam_\\QG$, and the cosmological constant $\\Lam_{\\rm COS}$. We set their ordering as \\bb M_\\P \\gg \\Lam_\\QG \\gg \\Lam_{\\rm COS}^{1/4}. \\ee We shall obtain an evolutional scenario according to the order starting inflation driven by quantum effects of gravity without adding any artificial field by hands \\cite{hy}. The inflationary model induced by quantum effects of gravity was first proposed by Starobinsky \\cite{starobinsky}. We here develope along the idea and propose an evolutional scenario of the universe summarized as follows: the conformal symmetry begins to be broken about the Planck scale to form an inflationary universe with the expansion time constant of order of the Planck mass, and completely broken at the dynamical scale turning to the classical Einstein universe. Then, energies stored in extra degrees of freedom in higher-derivative gravitational fields shift to matter degrees of freedom, causing the big bang. The primordial power spectra obtained from the two-point correlation functions of gravitational fields show a significant character of the transition which is expected to be observed cosmologically. The aim of this paper is to clarify why we can observe the Planck scale phenomena today from the tiny CMB anisotropies. Evolution of the universe is described by the equations of motion taking effects of running coupling as a time-dependent function. We find an inflationary solution that starts from the Planck scale and end at the dynamical scale where the transition of space-time occurs. Since the inflationary homogeneous solution is stable, the fluctuations from this solution will be slowly diminishing during the inflationary era, and thus the linear perturbation about the inflationary solution becomes applicable. We evaluate gravitational fluctuations perturbed about the homogeneous solution: the scalar perturbation so-called Bardeen potential, the vector perturbation and the tensor perturbation. We obtain the spectra at the transition point, which should be used for the primordial spectra to analize the data observed by WMAP. The rest of the paper organized as follows: we summarize the model of quantum gravity in section 2 and construct an evolutional scienario of the universe in section 3. In section 4 we study the linear perturbation theory in the inflationary background. We then simulate evolutions of perturbations and obtain primordial spectra in section 5. We conclude in section 6. ", "conclusions": "\\setcounter{equation}{0} \\noindent We have constructed an evolutional model of the universe based on the conformal gravity with the symmetry breaking Einstein action. The universe starts from a quantum state at the Planck time and grows up exponentially. As a consequence of the asymptotically free dynamics, there appears a strong coupling phase where the structure of the space-time changes drastically. The transition takes place at the dynamical energy scale of gravity, where quantum space-time with conformal invariance makes transition to classical Einstein universe. Then the universe grows up to the present size, such that the size of the Planck length at the Planck time extends to the size of the order of $10$ Mpc distance today. We have suggested that at the transition, field fluctuations freeze to localized objects, and they eventually decay into the classical matter driving the universe into the big bang phase. It was shown that the linear perturbation is applicable on the inflationary background for the momentum range which covers the size of fluctuation observed as CMB anisotropies today by COBE and WMAP. The evolutions of gravitational scalar, vector and tensor fluctuations have been evaluated. Since the initial fluctuations are provided by conformal field theory, we expect that the amplitudes of tensor and vector fluctuations are relatively small in comparison to that of scalar fluctuations. The scalar fluctuations are getting small during the inflation, and the tensor fluctuation is preserved to be small, while the vector fluctuation is getting large near the transition point. However, since the vector fluctuation disappears in the Einstein era, the tensor fluctuation in addition to the scalar fluctuation may contribute to the primordial spectra constituting the observed CMB anisotropies. The condition $M_\\P \\gg \\Lam_\\QG$ is significant to make the inflationary scenario. It also implies that quantum effects turn on the size much larger than the Planck length so that not only the space-time singularity but also the horizon of an elementary excitation with the Planck mass disappear. Furthermore, the strong repulsive effect in quantum gravity which causes the inflation also erase the singularity inside a black hole by balancing the pressure of the collapsing matter. However, if the collapsing goes too far and exceedes the balance, the black hole may explode something like as mini-inflation. \\vspace{1cm} \\begin{center} {\\Large {\\bf Appendix}} \\end{center} \\appendix" }, "0607/astro-ph0607065_arXiv.txt": { "abstract": "We study the relation between the rms mass fluctuations on 8$h^{-1}$Mpc scales and $\\Omega_{\\rm m}$ using the recent clustering results of XMM-{\\it Newton} soft (0.5-2\\,keV) X-ray sources, which have a median redshift of $z\\sim 1.2$. The relation can be represented in the form $\\sigma_{8}=0.34 (\\pm 0.01) \\Omega_{\\rm m}^{-\\gamma}$ where $\\gamma\\equiv \\gamma(\\Omega_{\\rm m},w)$ and it is valid for all $w<-1/3$ models. By combining the X-ray clustering and SNIa data we find that the model which best reproduces the observational data is that with: $\\Omega_{\\rm m}\\simeq 0.26$, $w\\simeq -0.90$ and $\\sigma_{8}\\simeq 0.73$, which is in excellent agreement with the recent 3-year Wilkinson Microwave Anisotropy Probe results. ", "introduction": "The combination of the recently acquired, high quality, observational data on galaxy clustering, the SNIa Hubble relation and the CMB fluctuations, strongly support a universe with flat geometry and a currently accelerated expansion due to the combination of a low matter density and a dark energy component (eg. Riess, et al. 1998; Perlmutter et al. 1999; Percival et al. 2002; Efstathiou et al. 2002; Spergel et al. 2003; Tonry et al. 2003; Schuecker et al. 2003; Riess et al. 2004; Tegmark et al. 2004; Seljak et al. 2004; Allen et al. 2004; Basilakos \\& Plionis 2005; Blake et al. 2006; Spergel et al. 2006; Wilson, Chen \\& Ratra 2006, for a review see also Lahav \\& Liddle 2006). From the theoretical point of view various candidates of the exotic ``dark energy'' have been proposed, most of them described by an equation of state $p_{Q}= w\\rho_{Q}$ with $w<-1/3$ (see Peebles \\& Ratra 2003 and references therein). Note that a redshift dependence of $w$ is also possible but present measurements are not precise enough to allow meaningful constraints (eg. Dicus \\& Repko 2004; Wang \\& Mukherjee 2006). From the observational point of view and for a flat geometry, a variety of studies indicate that $w< -0.8$ (eg. Tonry et al. 2003; Riess et al. 2004; Sanchez et al. 2006; Spergel et al. 2006; Wang \\& Mukherjee 2006 and references therein) Another important cosmological parameter is the normalization of the CDM power spectrum in the form of the rms density fluctuations in spheres of radius 8$h^{-1}$Mpc, the so called $\\sigma_{8}$. A tight relation between $\\sigma_{8}$ and the $\\Omega_{\\rm m}$ has been derived mainly using the cluster abundance with $\\sigma_{8}\\simeq 0.52 \\Omega_{\\rm m}^{-0.52}$ for a $\\Lambda$ cosmology (Eke, Cole \\& Frenk 1996). Also, Wang \\& Steinhardt (1998) generalizing to take into account dark energy models (with $w\\ge -1$) found: $\\sigma_{8} \\simeq 0.5 \\Omega_{\\rm m}^{-0.21+0.22w-0.33\\Omega_{\\rm m}}$. In this letter we use the clustering of high-$z$ X-ray AGNs to estimate a new normalization of the CDM spectrum, valid for spatially flat cosmological models and also for $w\\le -1$ (the so called Phantom models). Finally, combining our results with SNIa data (Tonry et al. 2003), we put strong constraints on the value of the equation of state parameter. ", "conclusions": "We have combined the clustering properties of distant X-ray AGNs, identified as soft (0.5-2 keV) point sources in a shallow $\\sim$ 2.3 deg$^{2}$ XMM survey, with the SNIa data. From the X-ray AGN clustering likelihood analysis alone we have estimated the normalization of the CDM power spectrum and find that the rms density fluctuation in spheres of radius 8$h^{-1}$Mpc is fitted by: $$ \\sigma_{8}\\simeq 0.34(\\pm 0.01) \\Omega_{\\rm m}^{-0.22+0.40w+0.052\\Omega_{\\rm m}} $$ which is valid also for Phantom models ($w<-1$). Furthermore, a joined likelihood analysis between the X-ray and SNIa data provides % a best model fit with: $\\Omega_{\\rm m}\\simeq 0.26$ and $w\\simeq -0.90$, which corresponds to $\\sigma_{8} \\simeq 0.73$, in agreement with the recent 3-years WMAP results (Spergel et al. 2006)." }, "0607/astro-ph0607315_arXiv.txt": { "abstract": "Gravitational lensing distorts the cosmic microwave background (CMB) temperature and polarization fields and encodes valuable information on distances and growth rates at intermediate redshifts into the lensed power spectra. The non-Gaussian bandpower covariance induced by the lenses is negligible to $l=2000$ for all but the $B$ polarization field where it increases the net variance by up to a factor of 10 and favors an observing strategy with 3 times more area than if it were Gaussian. To quantify the cosmological information, we introduce two lensing observables, characterizing nearly all of the information, which simplify the study of non-Gaussian impact, parameter degeneracies, dark energy models, and complementarity with other cosmological probes. Information on the intermediate redshift parameters rapidly becomes limited by constraints on the cold dark matter density and initial amplitude of fluctuations as observations improve. Extraction of this information requires deep polarization measurements on only 5-10\\% of the sky, and can improve Planck lensing constraints by a factor of $\\sim 2-3$ on any {\\em one} of the parameters \\{$w_0,w_a,\\Omega_K,\\sum m_\\nu$\\} with the others fixed. Sensitivity to the curvature and neutrino mass are the highest due to the high redshift weight of CMB lensing but degeneracies between the parameters must be broken externally. ", "introduction": "Primary cosmic microwave background (CMB) anisotropy from recombination has proven itself to be a veritable gold mine of cosmological information. One of the most important secondary signals that should be detected by upcoming cosmic microwave background experiments is the distortion to the temperature and polarization fields due to gravitational lensing by the large-scale structure of the universe (see \\cite{ChaLew05} for a recent review). Lensing distortions add cosmological information on parameters such as curvature, neutrino masses and dark energy that change the expansion and growth rate at intermediate redshifts ($z \\lesssim 5$). This distortion in real space couples power in harmonic space and hence introduces non-Gaussianity into the CMB temperature and polarization fields. Beyond power spectra, this non-Gaussianity is a source of information in that it allows direct reconstruction of the convergence field \\cite{Ber98,ZalSel99,Hu01b,HuOka01,HirSel02}. On the other hand, for purposes of extracting cosmological information from lensed power spectra as considered here, this non-Gaussianity is largely an impediment as it makes power spectrum estimates covary across a wide range of multipoles. The purpose of this paper is twofold. First, we calculate the full non-Gaussian covariance between all combinations of temperature and polarization bandpowers in the lensed CMB. This extends previous work in which the temperature \\cite{Hu01,Coo02} and $B$-mode polarization covariance \\cite{SmiHuKap04} were calculated separately. Second, we present a general framework for studying the extra information on cosmological parameters that lensed CMB spectra supply, with particular attention to the impact of non-Gaussianity. Previous works have noted that the lensed CMB signal may be used to study the dark energy \\cite{Hu01c,Kap03,AcqBac05} and neutrino mass \\cite{KapKnoSon03}. These studies did not compute the non-Gaussian covariance but assumed either that the information is encoded in the unlensed primary CMB and a reconstruction of the lenses or by approximating the non-Gaussian covariances with a degradation factor from \\cite{SmiHuKap04}. Our results lend support to these analyses. We also study the sensitivity of lensing to curvature and find that future CMB measurements can provide interesting constraints on it. This paper is organized as follows. In \\S\\ref{sec:pscov}, we compute non-Gaussian contributions to the covariance between all lensed CMB temperature and polarization bandpowers. We then describe in \\S\\ref{sec:formalism} how this non-Gaussian covariance propagates into Fisher matrix parameter forecasts and present formal bounds on its impact. In \\S\\ref{sec:deobs}, we define two parameter independent observables which contain essentially all information from the lensed CMB and discuss their relationship to distance and growth as well as their degeneracy with parameters that control the matter power spectrum. Armed with this general framework, we show how constraints on these observables can be interpreted in the context of common parameterizations of the dark energy and dark matter in \\S\\ref{sec:exfisher}. Finally in \\S\\ref{sec:applications} we show how future CMB surveys can be optimized for sensitivity to the lensing observables. We conclude in \\S\\ref{sec:discussion} and briefly address the issues of goodness-of-fit in Appendix \\ref{sec:chisq} and scaling with the fiducial cosmology in Appendix \\ref{sec:fiducial}. ", "conclusions": "\\label{sec:discussion} We have provided a comprehensive study of the additional cosmological information supplied by lensed power spectra of the CMB temperature and polarization fields including the non-Gaussian covariance between bandpower estimates. This covariance originates from the sample variance of the degree scale lenses on the CMB fields at smaller scales. It is nearly irrelevant for the temperature and $E$-polarization fields out to $\\ellmax=2000$ due to the larger sample variance of the unlensed CMB. For the amplitude of the $B$-polarization field, it increases the variance by up to a factor of $\\sim 10$ and changes the optimal observing strategy to one that covers a factor of $\\sim 3$ times more sky area. The impact of non-Gaussianity on parameter estimation as well as the net information content of the lensed spectra is more subtle. These answers depend on the choice of parameters and the external priors associated with them. We have provided a framework of lensing observables that greatly simplifies these examinations. In this framework, lensed CMB power spectra provide information on only two observables, one which determines the lens power spectra at $l \\sim 100$ associated with the \\{$T,E$\\}-fields and one which determines it at $l \\sim 500$ associated with the $B$-field. The observables are constructed from the principal components of the lensing power spectrum $C_l^{\\phi\\phi}$. Non-Gaussianity is then automatically incorporated in the errors on the observables which will eventually approach, but never exceed, the sampling errors of the lenses as the measurements improve. This construction also illuminates the origin of parameter degeneracies which can rapidly become the limiting source of uncertainties for parameters of interest. Any combination of parameters that leaves the lensing observables and the CMB at recombination fixed within the errors cannot be determined. To illustrate these effects, we have isolated two parameters $\\wm$ and $\\ldz$ that determine the shape and amplitude of the matter power spectrum respectively, and marginalized their uncertainties assuming internal CMB determinations of each from the Planck satellite. These become the limiting uncertainties once the observables are determined to the several percent level and are only slowly improved as the lensing survey itself improves the nuisance errors. While $\\wm$ constraints can be improved externally to the CMB, those on $\\ldz$ are more difficult to improve and may be limited by our understanding of reionization. There are also degeneracies within the space of the parameters of interest that control the expansion rate and growth of structure at intermediate redshifts. When taken one at a time, uncertainties on the parameters \\{$w_0,w_a,\\Omega_K,\\sum m_\\nu$\\} can be improved by a factor of $\\sim 2-3$, relative to Planck alone, by a deep ground-based polarization survey on 5-10\\% of the sky. However \\{$w_0, w_a$\\} are nearly perfectly degenerate in the lensing observables as are \\{$\\Omega_K, \\sum m_\\nu$\\} separately. The degeneracy between two parameters in each pair is weakly broken by the two observables. For example, when errors on $\\sum m_\\nu$ are marginalized over $w_0$ they degrade by a factor of 2 for the reference survey (see Fig.~\\ref{fig:omnh2_w}). However sensitivity to the \\{$\\Omega_K, \\sum m_\\nu$\\} pair is much greater than to the dark energy parameters due to the high redshift weights of the lensing observables. When combining lensed CMB power spectra with other more incisive probes of the dark energy, lensing essentially fixes one well-defined combination of \\{$\\Omega_K,\\wn$\\} \\cite{HuHutSmi06}. Our conclusions have several caveats associated with them. The observables framework implicitly assumes that lensing is an independent and additive source of cosmological information that may be combined with the intrinsic CMB anisotropy. An important exception to this statement occurs for tensor modes, where lensing $B$-modes mask the intrinsic $B$-modes. Forecasts for tensor modes should be made employing lensed power spectra as a destructive contribution but here the Gaussian approximation suffices. The conversion between instrumental noise and errors on the observables depends only mildly on the fiducial model given current cosmological constraints but we give a crude scaling in Appendix \\ref{sec:fiducial}. Secondly, we have considered only the information contained in the lensed power spectrum. Beyond the power spectrum, non-Gaussianity from lensing allows a direct reconstruction of the lensing fields \\cite{HuOka01,HirSel02} which carries substantially more information that can break parameter degeneracies \\cite{Hu01c,KapKnoSon03}. It may also allow ``de-lensing'' techniques that recover the intrinsic $B$-modes from tensor modes \\cite{KnoSon02,KesCooKam02,SelHir03}. However techniques have yet to be developed that can remove systematics and contamination at the levels required. Thirdly, our parameter forecasts employ the Fisher matrix approximation. It is well known that Fisher matrix forecasts are not accurate along ill-constrained directions in the parameter space. Hence our results are only robust for quantities that lensed power spectra constrain well. Finally, we never consider CMB multipoles beyond $\\ellmax=2000$ in this paper. Well beyond this limit there is extra information on the high multipole structure of the lensing field but this is likely to prove difficult to extract in the presence of other secondaries and foregrounds." }, "0607/astro-ph0607123_arXiv.txt": { "abstract": "The supermassive black hole in the center of our Galaxy, Sgr A*, is unique because the angular size of the black hole is the largest in the sky thus providing detailed boundary conditions on, and much less freedom for, accretion flow models. In this paper I review advection-dominated accretion flow (ADAF; another name is radiatively inefficient accretion flow) models for Sgr A*. This includes the developments and dynamics of ADAFs, and how to explain observational results including the multi-waveband spectrum, radio polarization, IR and X-ray flares, and the size measurements at radio wavebands. ", "introduction": "The center of our Galaxy provides the best evidence to date for a massive black hole (e.g., Sch\\\"odel et al. 2002; Ghez et al. 2003), associated with the compact radio source, Sgr A* (see, e.g., Melia \\& Falcke 2001). Since the original discovery of Sgr A* in 1974, there have been intensive efforts in both observational and theoretical aspects, with dramatic progresses in the past few years. The reason why we are so interested in this object is because of its proximity, which allows us to determine observationally the dynamics of gas quite close to the BH, providing unique constraints on theoretical models of accretion flows. Before introducing the model, I first briefly review the main observational results of Sgr A*. As shown by the data points in Fig. 2, its radio spectrum consists of two components. The component below 86 GHz has a spectrum $F_{\\nu} \\propto \\nu^{0.2}$, while the high frequency component, the ``submm bump'', has a spectrum $F_{\\nu} \\propto \\nu^{0.8}$ up to $\\sim 10^3$ GHz (e.g., Falcke et al. 1998; Zhao et al. 2003). Variability at centimeter and millimeter wavelength are detected with a timescale ranging from hours to years with amplitude of less than 100\\% (Zhao et al. 2003, 2004; Herrnstein et al. 2004; Miyazaki et al. 2004). High level of variable linear polarization fraction ($\\sim 2\\%-10\\%$) at frequencies higher than $\\sim 150$ GHz puts a rotation measure upper limit of $7\\times 10^5 {\\rm rad~m^{-2}}$, which argues for a low density at the innermost region of ADAF (e.g., Aitken et al. 2000; Bower et al. 2003; Marrone et al. 2006; Macquart et al. 2006; Quataert \\& Gruzinov 2000). At IR wavelength, the source is highly variable. Genzel et al. (2003) detected Sgr A* at 1.6-3.8 $\\mu$m, with a factor of $\\sim 1-5$ variability on timescales of $\\sim 10-100$ min. Similarly, at 3.8 $\\mu$m, Ghez et al. (2004) found that the flux changes by a factor of 4 over a week, and a factor of 2 in just 40 min. If describing the IR spectrum with a power-law, Gillessen et al. (2006) found that the spectral index is correlated with the instantaneous flux but Hornstein et al. (2006) found it remains constant during the flare process. Sgr A* has been convincingly detected in the X-rays (Baganoff et al. 2001, 2003; Goldwurm et al. 2003). The X-ray emission has two distinct components. In ``quiescence,'' the emission is soft and relatively steady, with a large fraction of the X-ray flux coming from an extended region with a diameter $\\approx 1.4''$ (Baganoff et al. 2001, 2003). Several times a day, however, Sgr A* has X-ray ``flares'' in which the X-ray luminosity increases by a factor of a few -- 50 for roughly an hour. For the most flares, the spectrum is hard, with a photon index of $\\Gamma=1.3^{+0.5}_{-0.6}$. {\\em XMM}, however, detected a very bright and soft flare with $\\Gamma=2.5^{+0.3}_{-0.3}$ (Porquet et al. 2003). Recent several multiwavelength campaigns found that there is no time lag between the IR and X-ray flares (Eckart et al. 2004, 2005; Yusef-Zadeh et al. 2006). This strongly suggests a common physical origin. The short timescale argues that the emission arises quite close to the BH, within $\\sim 10 R_S$ (where $R_S$ is the Schwarzschild radius). Sgr A* is extremely dim overall, with a bolometric luminosity of only $L\\approx 10^{36}\\ergs \\approx 3\\times 10^{-9}L_{\\rm Edd}$. ", "conclusions": "The supermassive black hole in our Galactic center represents a unique opportunity to probe the physics of accretion, especially at extremely low accretion rates. In this review, I first briefly introduce the dynamics and evolution of the ADAF. I then have tried to argue that this model can provide a reasonable explanation to most of the current observations. {\\em Chandra} observations tell us the density and temperature of the hot gas at $\\sim 1^{\"}\\sim 0.04$ pc. This radius happens to be the Bondi radius where the gas is captured by the gravity of the center black hole and starts to be accreted. The Bondi accretion rate, $\\dot{M}_{\\rm Bondi}\\sim 10^{-5}\\mpy$ provides a good estimation to the real accretion rate. As a comparison, the numerical simulation gives $\\dot{M}\\sim 3\\times 10^{-6}\\mpy$. Since the bolometric luminosity of Sgr A* is only $10^{36}\\ergs$, the radiative efficiency should be very low, $\\sim 5\\times 10^{-6}$. The standard thin disk model (Shakura \\& Sunyaev 1976) is therefore ruled out immediately. ``Old'' ADAF models can naturally explain such a low efficiency (e.g., Narayan et al. 1995). However, these models fail to explain the high linear polarization at submm waveband because the density and further, the rotation measure, are too large. On the other hand, theoretical studies of ADAFs also indicates the presence of outflow (e.g., Stone, Pringle \\& Begelman 1999; Blandford \\& Begelman 1999). This feature is taken into account in the ``new'' ADAF models (or RIAF; YQN03). The YQN03 model can explain the quiescent state spectrum and the polarization, as shown by Figs. (2)\\&(3). The accretion rate close to the horizon is only $4\\times 10^{-8}\\mpy$. So the low radiative efficiency of the ADAF in Sgr A* is partly because of the outflow, which contributes a factor of $\\sim 10^{-2}$, and partly because of the energy advection, which contribute a factor of $\\sim 5\\times 10^{-4}$. The IR and X-ray flares are explained by the synchrotron and inverse-Compton emissions from the heated/accelerated electrons during the magnetic reconnection events in the innermost region of the ADAF (Figs. 4\\&5). Finally, the YQN03 model satisfactorily pasts the test of recent observations of the size of Sgr A* at 3.5 and 7 mm wavebands (Fig. 6). Moreover, this model predicts that the observation at 1.3 mm should be able to detect GR effects (Fig. 6). \\ack This work was supported in part by the One-Hundred-Talent Program and the National Natural Science Foundation of China (grants 10543003)" }, "0607/astro-ph0607409_arXiv.txt": { "abstract": "We present the first hydrodynamic N-body simulations of primordial gas clouds responsible for the reionisation process in dark energy cosmologies. We compare the cosmological constant scenario with a SUGRA quintessence model with marked dynamics in order to highlight effects due to the different acceleration histories imposed by the dark energy. We show that both the number density of gas clouds and their clumpiness keep a record of the expansion rate during evolution, similar to the non-linear dark matter profile at virialisation, as was recently demonstrated by Dolag et al. (2004). Varying the shape of the primordial power spectrum, we show how this effect is mitigated by a running spectral index decreasing the power at small scales. Our results demonstrate that, in order to constrain the dark energy from large scale structures, one must track its effects down to the distribution of luminous matter. ", "introduction": "\\label{sect:intro} Different observational data, like those from high-redshift supernovae \\citep{riess2004,astier2006}, the cosmic microwave background \\citep{spergel2003,spergel2006}, and large-scale structure \\citep{tegmark2004,cole2005}, are now giving a consistent picture of our universe, which can be described by the so-called `concordance' $\\Lambda$CDM model: a spatially flat universe whose expansion accelerates in the present epoch because of a dominant dark energy component. Historically, the first and simplest candidate for that has been the cosmological constant $\\Lambda$, corresponding to matter with an equation of state $p=w \\rho c^2$, with constant $w=-1$. However, observations suggest a value for $\\Lambda$ which is more than a hundred orders of magnitude smaller than the energy scales expected to be responsible in the very early universe. To avoid inelegant solutions based on parameter fine-tuning, the general idea of dark energy is extended to encompass the so-called quintessence, possibly corresponding to a suitably self-interacting scalar field, whose pressure and energy density evolve during cosmic history. The currently available observational data sets do not allow yet to place strong constraints on $w$. We know that it has to come close to $-1$ with a precision of about 10 per cent at the present epoch \\citep[see, e.g.,][]{riess2004,spergel2006}, but its redshift evolution is still poorly constrained and can hopefully be determined only with new-generation data; see \\citet{seljak2005} for one of the very first attempts to measure the high-redshift dark energy equation of state. One consequence of $w>-1$ earlier is that the formation of cosmic structures sets in earlier compared to the $\\Lambda$CDM model. At a given redshift, this corresponds to a higher abundance of more concentrated halos \\citep[see, e.g.,][]{dolag2004}. This may open alternative ways to study quintessence models based on galaxy-cluster counts \\citep{wang1998,haiman2001,battye2003,majumdar2003} and strong gravitational lensing \\citep{bartelmann2003,meneghetti2005}. The different history of structure formation must also affect the reionisation epoch. To produce the required ionising radiation, a sufficiently high number of early stellar sources is needed, such as super-massive Pop-III stars or more ``standard'', massive Pop-II stars in proto-galaxies. In particular, Pop-III stars are supposed to be composed mainly of hydrogen and helium with their primordial abundances, to have masses $\\gtrsim100\\,M_\\odot$ much larger than the standard Pop-I and -II stars, and to form in dark matter haloes with typical masses of $10^{5\\ldots6}\\,M_\\odot$ (the so-called mini-haloes). The recent analysis of the temperature-polarisation cross-correlation and the polarisation auto-correlation measured by the WMAP satellite now allows the estimation of the epoch when reionisation occurs. The Thomson optical depth of of $\\tau\\sim 0.17\\pm 0.04$ extracted from the first-year WMAP data had been interpreted as signalling the high redshift of $z=17\\pm5$ for the primordial epoch of global reionisation \\citep{kogut2003}. The recently released three-year WMAP data allowed an improved control in particular of the polarised foreground emission, yielding $\\tau=0.09\\pm0.03$. This corresponds to $z=10^{+2.7}_{-2.3}$ for the completion of the reionisation process, even if some level of parameter degeneracy still remains \\citep[see, e.g., Fig. 3 in][]{spergel2006}.\\footnote{We have also to notice that the three-year WMAP data are suggesting a lower value not only for $\\tau$, but also for the power spectrum normalization $\\sigma_8$: these two effects nearly cancel in terms of early structure formation, i.e. the inferred small $\\tau$ does not really allow slow reionization for a given $\\sigma_8$.} Different authors \\citep[see, e.g.,][]{wyithe2003,ciardi2003,sokasian2003,sokasian2004} suggested that these reionization data can place new complementary constraints on the cosmological parameters, and in particular on the nature of dark energy. So far, studies of the structure formation process in dark energy cosmologies were focused on modifications of dark matter structures due to the different expansion histories \\citep[see, e.g.,][]{klypin2003,dolag2004}. They revealed that virialised objects keep a record of the expansion rate at the time of their formation. The general picture is that in cosmologies with $w>-1$, the increase of the dark energy density with redshift enhances structure growth, and virialised haloes become more concentrated because of the denser environment they form in \\citep{dolag2004}. This effect is stronger in tracking quintessence models, which we describe in the next section, in which the increase of the dark energy density with redshift is enhanced. Complementary hydrodynamical simulations still need to be carried out. This work is a first step into this direction, beginning with the numerical study of the dependence of the primeval gas clouds responsible for the reionisation process on the cosmic expansion rate at their formation time. Due to the higher concentration of structures in dark energy scenarios, which has so far been verified only for the dark matter, it is mandatory for our study to control the shape of the primordial power spectrum. When combined with the Ly-$\\alpha$ forest data and with the analysis of the $2dF$ galaxy redshift survey, the first-year WMAP results supported a ``running'' spectral index which tilts the spectrum slightly towards small scales, starting at $k\\gtrsim1\\,\\mathrm{Mpc^{-1}}$. The three-years WMAP data are compatible with this \\citep{spergel2006}, albeit with less emphasis than the earlier results. A running index would cause a slower growth and evolution of small-scale cosmic structures compared to models with a constant index $n$. A reduction of power on small scales may alleviate several potential discrepancies in the $\\Lambda$CDM model, such as the abundance of substructures in galaxies and the high central concentration of galactic haloes. However, as shown by analytic and numerical work \\citep[see, e.g.][]{somerville2003,yoshida2003b}, models with a running spectral index (RSI) may have severe problems in producing enough objects to allow a global reionisation at high redshift. Since this might be balanced by the enhanced structure growth due to the dark energy, it is important to jointly study these two aspects. In the present paper, we study the high-redshift structure formation in quintessence models based on the results of high-resolution cosmological N-body simulations combined with hydrodynamics. Specifically, we consider four different cases combining two flat cosmological models, i.e.\\ the concordance $\\Lambda$CDM model and a quintessence model with a SUGRA potential, with two types of the primordial power spectrum, one with a constant spectral index $n=1$ and an RSI model assuming the best fit relation found by \\cite{spergel2003}. More detail will be given in the next section. The paper is organised as follows. In Sect.~\\ref{sect:sim}, we describe the general characteristics of the hydrodynamical simulations used below and introduce the quintessence models chosen. The techniques adopted for identifying dark matter haloes and the corresponding results are presented in Sect.~\\ref{sect:haloes}. The abundances of gas clouds are presented in Sect.~\\ref{sect:clouds} together with the clumping factors and recombination times as computed from the simulation outputs. The implications of the previous results in terms of reionisation are discussed in Sect.~\\ref{sect:reionization}. The final discussion and our main conclusions are drawn in Sect.~\\ref{sect:conclusions}. ", "conclusions": "\\label{sect:conclusions} We presented hydrodynamic N-body simulations of the formation of primordial gas clouds on scales of tens of kpc in a variety of cosmological models, characterised by different dynamics in the dark-energy component. Our main results are that the records of the modified expansion rate are well evident in the population and the clumpiness of such clouds. Cosmological models with the same power-spectrum normalisation at present show earlier cloud formation if the dynamics of the dark energy is enhanced, represented by an equation of state parameter $w>-1$ as in quintessence models. Within dark energy models compatible with the present data on cosmic microwave background and large scale structure, the difference in the integral population of clouds may vary by up to an order of magnitude, as a consequence of the different differential efficiency for structure formation. This is consistent with earlier results indicating a higher concentration in dark matter haloes under similar conditions \\citep{dolag2004}. Since abundance and clumpiness of structures are directly related to the amount of primordial power on the corresponding scales, we varied the shape of the primordial power spectrum by a running spectral index reducing power on small scales within the confidence level of the three-years WMAP data \\citep{spergel2006}. As expected, we find that the extra population and clumpiness of clouds produced by a higher dark energy abundance compared with its level today might be mitigated if the primordial spectral index is running, decreasing the power on small scales. Adopting a simple picture for the reionisation process, we derived consequences for the reionisation itself, leading to an earlier beginning of the reionisation process in models where cloud formation starts earlier. On the basis of these results, we are able to identify possible tension between the WMAP data on the reionisation optical depth and cosmological models whose dark energy is as dynamical as in SUGRA quintessence models. Our results demonstrate that the effects on cosmological structure formation from a modified expansion history through different dark energy models must be traced back to the formation of the first clouds. In turn, this means that constraints on the dark energy, and in particular its abundance at high redshifts, may be obtained by forthcoming experiments aiming at measuring the abundance and the clumpiness of primordial gas clouds. In particular, the Atacama Large Millimeter Array (ALMA\\footnote{www.eso.org/projects/alma/science/}) and the Mileura Widefield Array (MWA\\footnote{www.haystack.mit.edu/ast/arrays/mwa/site/index.html}) will probe the early stages of structure formation, say between $z=6$ and $10$, where the reionization in progress should keep a record of the population of reionizing primordial gas clouds." }, "0607/astro-ph0607645_arXiv.txt": { "abstract": "{} { We test the hypothesis that S0 galaxies are the descendants of fading spirals whose star formation has been shut down, by using the properties of their globular cluster systems. } { We estimate the amount by which the globular cluster specific frequency (number of globular clusters per unit $V$-band luminosity) is enhanced in S0s relative to spirals. If the transformation hypothesis is correct, and no clusters are created or destroyed in the process, then this difference provides a measure of the degree to which the S0's $V$-band luminosity has faded relative to that of its spiral progenitor, which we can compare with the independent values estimated from stellar population synthesis and the S0 Tully--Fisher relation. We also explore whether the degree to which the globular cluster specific frequency is enhanced in S0s correlates with the colour of the stellar population, as also predicted by this hypothesis in which galaxies become redder as they fade. } { We find that, on average, the globular cluster specific frequency is a factor $\\sim 3$ larger for S0s than for spirals, which can be interpreted as meaning that passively-evolving S0s have faded on average by about a factor of three from their spiral progenitors. This value fits remarkably well with the predictions of stellar population synthesis calculations, and the offset between the S0 and spiral Tully--Fisher relations, where the S0 $V$-band relation lies $\\sim 1.2$ magnitudes, or a factor of three, below the spiral relation. We also find that the global colours of S0 galaxies are strongly correlated with their globular cluster specific frequencies: the redder the stellar population of an S0, the larger its specific frequency, as we might expect if we are catching different S0s at different stages of passively fading and reddening. Comparison to the predictions of stellar population synthesis models show that this explanation works quantitatively as well as qualitatively. } { These tests strongly support the hypothesis that S0 galaxies were once normal spirals, whose star formation was cut off, presumably due to a change of environment. We are now in a position to start to make quantitative measurements of when this life-changing event occurred in different galaxies. } ", "introduction": "S0, or lenticular, galaxies live at the crossroads between elliptical and spiral galaxies in the traditional Hubble tuning-fork diagram, suggesting that they should play a key role in understanding the morphologies of galaxies. In fact, in some cluster environments, S0s are the single most common type of luminous galaxy \\citep{Dressler:1980}, so clearly understanding how they form and evolve is essential if we wish to have a complete picture of how galaxy morphology is related to galaxy formation and the environment. \\citet{Dressler:1980} also showed that as the fraction of S0s increases in the densest cluster environments, so the fraction of spirals decreases, naturally suggesting that S0s are simply ``dead'' spiral galaxies that have had their star formation shut off by their surroundings, and are now quietly fading away as their stellar populations age. Indeed, a number of mechanisms have been suggested that might cause such a transformation: ram-pressure stripping of disk gas could remove the raw material of the next stellar generation \\citep{Gunn_Gott:1972}, while the somewhat gentler removal of a larger-scale reservoir of halo gas could lead to a slower ``strangulation'' of star formation \\citep{Larson_etal:1980}. Although this scenario is quite plausible, we need some more direct evidence that it has actually occurred. Circumstantial evidence for such a transformation over cosmic timescales comes from observations which indicate that the proportion of S0 galaxies is substantially smaller in distant clusters than in nearby ones, while spirals show the opposite trend \\citep[see, for example,][]{Dressler_etal:1997}, but any inferences drawn from such studies always face the criticism that one may not be equating comparable systems at different redshifts. Further evidence that S0s are fading spirals comes from studies of the Tully--Fisher relation between luminosity and rotation speed for S0 galaxies, which is found to be rather broad and offset from the relation for spirals in the sense that the S0s are systematically fainter than the spirals. A recent analysis by \\citet{Bedregal_etal:2006} found that the S0 relation in the $B$-band is offset by $\\sim 1.4$ magnitudes from the \\citet{Sakai_etal:2000} relation for spirals, with a scatter of $\\sim 1$ magnitude. Such an offset can most straightforwardly be interpreted as arising from the fading of the S0s relative to their spiral progenitors, with the spread in the relation arising from the different epochs at which this fading commenced. Indeed, \\citet{Bedregal_etal:2006} were able to uncover some evidence that the magnitude of each S0's offset from the spiral Tully--Fisher relation depends on the time since star formation was cut off, as measured by the age of its stellar population, just as this picture would predict. However, the evidence is still uncomfortably circumstantial, as we have had to assume that the ancient progenitors of these S0s respect the same Tully--Fisher relation that we see in the nearby spirals, and that the difference between the relations for S0s and spirals arises from luminosity evolution: the offset could also arise because these systems have fundamentally different mass properties, shifting the relation in rotation speed. Even with these assumptions, we can only tie spirals and S0s together on a statistical basis, as we do not have a direct connection between any S0 and the properties of its individual progenitor. Ideally, we would like to find some historical record preserved in the current properties of an S0 that tells us about its luminosity during its early life as a spiral, and that will not have been defaced by the transformation process or any subsequent evolution of the galaxy. Fortunately, just such a record exists in the globular cluster (GC) population of the S0. As we will see below, GCs have a reasonably well defined specific frequency (number per unit galaxy $V$-band luminosity) for spiral galaxies, so the number of GCs provides a useful proxy for a spiral galaxy's luminosity. Further, the transformation from spiral to S0 is unlikely to alter the number of GCs significantly: the hydrodynamic processes that strip out the gas from the spiral will not have any impact on these collisionless stellar clusters, and the conversion is likely to be sufficiently benign that we are unlikely to lose clusters or produce any new ones as we find in more dramatic phenomena like mergers \\citep{Ashman_and_Zepf:1998}. In addition, the GCs are sufficiently old that their own passive fading will be very slow, so the observed number counts should not decrease significantly during the transformation. Thus, the number of GCs in an S0 offers a reasonably robust indicator of its spiral progenitor's luminosity, so the specific frequency of GCs in such a system provides a direct measure of the ratio of its progenitor's luminosity to its current luminosity, which we can compare to the less direct indicators of fading such as those provided by the Tully--Fisher relation. In this paper, we carry out such an analysis, comparing the specific frequencies of GCs in S0s to those in spirals in Section~\\ref{sec:GCSFcomp} to see how much these systems must have faded, and going on to look for evidence that different amounts of fading might be related to different epochs of transformation in Section~\\ref{sec:fadeage}. Conclusions are presented in Section~\\ref{sec:conc}. ", "conclusions": "\\label{sec:conc} The classic problem in trying to address questions related to galaxy evolution is that our snapshot view of the Universe means that we never get to view the process directly. Thus, although it seems quite plausible that spiral galaxies fade to become S0s when their interstellar media are stripped out and star formation ceases, there is little direct evidence to this effect. Here, we finesse this issue by making use of the historical record imprinted in a galaxy via its GC population, which should still be ``readable'' even after the transformation process. In particular, if the number of globular clusters is taken as a measure of the progenitor spiral galaxy's luminosity, and this number does not change as the galaxy passively evolves, then we have been able to show that S0s must have faded by an average factor of approximately three during this process. This result matches very neatly with the values predicted by stellar population synthesis models and those found by comparing the Tully--Fisher relations of spirals and S0s. Further, using the amount by which the stellar population has reddened as a measure of the time since the transformation into an S0 began, we have been able to show that individual galaxies are at different stages along this evolutionary track: the colours of these galaxies correlate with the specific frequency of GCs in exactly the way that one would expect if their star formation shut down at different times in the past. Thus, we are at the point of being able to follow the histories of individual S0 galaxies with the detail necessary to say which underwent their life-changing transformations first. The next logical step in writing the life histories of these galaxies is to obtain uniform samples of spectra of the quality necessary to derive stellar population ages for these systems \\citep[e.g.][]{Kuntschner_Davies:1998}. Not only would these data resolve any remaining ambiguity by lifting the degeneracy between age and metallicity effects in the interpretation of the broad-band colours used here, but they would also allow us to put much more accurate dates to the ``birthdays'' of S0 galaxies." }, "0607/astro-ph0607159_arXiv.txt": { "abstract": "We present results from Chandra and XMM-Newton observations of the bright group of galaxies HCG~62\\@. There are two cavities at about $30''$ northeast and $20''$ southwest of the central galaxy in the Chandra image. The energy spectrum shows no significant change in the cavity compared with that in the surrounding region. The radial X-ray profile is described by a sum of 3-$\\beta$ components with core radii about 2, 10, and 160 kpc, respectively. We studied radial distributions of temperature and metal abundance with joint spectral fit for the Chandra and XMM-Newton data, and two temperatures were required in the inner $r< 2'$ (35~kpc) region. The sharp drop of temperature at $r\\sim 5'$ implies the gravitational mass density even lower than the gas density, suggesting the gas may not be in hydrostatic equilibrium. Fe and Si abundances are 1--2 solar at the center and drop to about 0.1 solar at $r \\sim 10'$. O abundance is less than 0.5 solar and shows a flatter profile. Observed metal distribution supports the view that iron and silicon are produced by type Ia supernova in the central galaxy, while galactic winds by type II supernova have caused wide distribution of oxygen. The supporting mechanism of the cavity is discussed. Pressure for the sum of electrons and magnetic field is too low to displace the hot group gas, and the required pressure due to high energy protons are nearly 700 times higher than the electron pressure. This leaves the origin of the cavities a puzzle, and we discuss other possible origins of the cavities. ", "introduction": "Groups of galaxies hold significantly less amount of hot gas compared with rich clusters on the average, and the apparent deficiency of baryons in these low-mass systems is a problem in explaining the baryon budget in rich systems in terms of hierarchical merging scenario \\citep{Voit2005}. The observations of compact groups of galaxies are important to search for the hidden form of baryons and their release mechanism, which would be strongly connected with dynamics of gas and galaxies. In particular, some groups dominated by bright central galaxies have shown gas features which are strongly affected by the activity of the central galaxies. High sensitivity X-ray observations of groups of galaxies are powerful method to look into the role of central galaxies in terms of metal distribution and gas morphologies. The striking gas features most likely caused by the central galaxies are the X-ray cavities. The cavities are circular regions showing a significant depression of X-ray surface brightness. Nearly 20 cavities have been recognized in clusters and groups with high resolution images taken by ROSAT and Chandra (\\cite{Birzan2004}, hereafter B04; \\cite{Dunn2004}; \\cite{Dunn2005}, hereafter D05). They are located typically at 10--30 kpc from the central galaxies, and strong correlation with radio lobes are seen in about 10 systems. Remarkable cases are seen in the Perseus (e.g., \\cite{Boehringer1993,Fabian2000}) and Hydra A (e.g.\\ \\cite{McNamara2000}) clusters, both showing strong correlation with the 1.4~GHz radio lobes. Several giant elliptical galaxies with radio robes, e.g., M84 \\citep{Finoguenov2001}, NGC~4636 \\citep{Ohto2003}, are also known as nesting X-ray cavities. The remaining half of the cavities, on the other hand, are not associated with radio lobes, and they are designated as ghost cavities. The one in A~2597 \\citep{McNamara2001} or the outer depressions in Perseus \\citep{Fabian2000} are examples. Cavities are thought to be produced by jets or buoyant bubbles which are directly connected with the activity of central radio galaxies. A subsonic displacement of the gas would create a low density, rising bubble keeping the pressure balance with the surrounding ICM\\@. It appears to be in general supposed for cavities that non-thermal pressure originated in relativistic particles and/or magnetic fields in the radio lobe is probably large enough to balance with the surrounding ICM gas pressure \\citep{Fabian2002}. This pseudo-pressure balance is justified by the fact that there are no evidence for shock-heated gas around the radio lobes in almost all of the X-ray cavities observed so far, except for MKW~3s \\citep{Mazzotta2002}. This general scenario has been modeled theoretically, and has at least qualitatively reproduced the morphology of cavities (e.g., \\cite{Churazov2001}). Energy density of relativistic electrons inferred from the synchrotron radio emission is almost always smaller than that required to offset the hot gas by orders of magnitude, and it is discussed that energy density of protons are higher than those due to electrons by factors of 100--1000 (D05). However, there is no direct evidence indicating that such a high energy density is really carried by protons. This situation is the severest in the case of ghost cavities. In this view, it is important to examine ghost cavities in groups of galaxies where the gas is relatively cool and non-thermal effect can be recognized somewhat easily. Detailed studies on the metal distribution in clusters and groups have been carried out using ASCA, BeppoSAX, Chandra and XMM-Newton. Distribution of iron and silicon indicate strong central concentration in clusters and groups characterized by bright central galaxies, and the excess iron mass is found to correlate with the luminosity of the cD galaxy \\citep{DeGrandi2004}. This indicates that iron and silicon (main products from type Ia supernova; SN~Ia) trace the enhanced star-formation activity in bright galaxies. On the other hand, distribution of oxygen (i.e.\\ type II supernova product; SN~II) is not well understood. \\citet{Matsushita2003} showed that oxygen distribution around M~87 is flatter than those of iron and silicon, with the level about half as much as the others. Such low oxygen abundances are also derived in the centers of other clusters and groups (e.g.\\ \\cite{Buote2003b,Xue2004}). For the study of oxygen distribution, low temperature systems such as groups of galaxies are suitable targets. In this paper, we report the results from Chandra and XMM-Newton observations of HCG~62, which is one of the nearest Hickson compact galaxy groups \\citep{Hickson1989}. The whole group consists of 63 galaxies \\citep{Mulchaey2003} within a radius of 50$'$ (900~kpc), but the central region is dominated by 4 galaxies. HCG~62 is the brightest group of galaxies in the X-ray band, and the extended X-ray emission was first discovered by \\citet{Ponman_Bertram1993} from the ROSAT PSPC observation. Based on the ASCA observation, \\citet{Fukazawa2001} detected excess hard X-ray emission, and \\citet{Finoguenov_Ponman1999} report strong central concentration of iron. Using the high resolution image of Chandra, \\authorcite{Vrtilek2001}~(\\yearcite{Vrtilek2001},\\yearcite{Vrtilek2002}) detected two ghost cavities, which is the first report of cavities in groups of galaxies. This paper is organized as follows: In \\S\\,2 we describe the Chandra and XMM-Newton observations and the data reduction. In \\S\\,3 we give the image of HCG~62 of both Chandra and XMM-Newton, in \\S\\,4 we describe the X-ray cavity structure using Chandra image. In \\S\\,5 we present our results on the temperature profiles and the abundances profiles of Fe and $\\alpha$-elements (Si, Mg, and O). In \\S\\,6--8 we give discussions of the obtained results, and finally we summarize our conclusions in \\S\\,9. Throughout this paper we adopt $\\Omega_{\\rm \\Lambda}=1-\\Omega_{\\rm M}=0.73$ and $h_{70}\\equiv H_0/(70~{\\rm km~s^{-1}\\,Mpc^{-1}})=1$; $1\\arcmin$ corresponds to 17.8~kpc at $z=0.0145$. The quoted errors indicate the 90\\% confidence range, unless otherwise stated. We use the solar abundance ratio of \\citet{Anders_Grevesse_1989}. ", "conclusions": "\\label{sec:conc} \\begin{itemize} \\item We have carried out a detailed study on the hot-gas properties of the group of galaxies HCG~62, based on the data from Chandra and XMM-Newton. We confirmed the two cavities located almost symmetrically around the central galaxy. \\item The size of spherical hollow cavities are constrained from the surface brightness structure to be 12$''$--17$''$. The agreement with the observed angular size suggests that the gas density in the cavity is very low, less than 1/3 and consistent with zero. \\item The spectral fit indicated that the cavities were not caused by X-ray absorption. The observed temperature in the cavity region is consistent with that in the surrounding region. \\item The spectrum within $4'$ from the center requires two temperatures: 0.7~keV and 1.4~keV\\@. The cool component is centrally concentrated, narrower than the hot component, suggesting its association with the central galaxy HCG~62a. \\item The mass profiles were obtained for the gas and stars. The hot component is much more extended than the stars, and thought to trace the gravitational potential of the galaxy group. \\item The gravitational mass density drops steeply at about 5$'$ from the center. This is caused by the observed sharp drop of the temperature. There is a possibility that these regions are not in the hydrostatic equilibrium. \\item Abundance of O is $\\sim 0.3$~solar, $\\sim 3$ times less abundant than Fe and Si, and shows a flatter profile. The shallow potential of HCG~62 is unable to confine the SN~II products which should have been escaped in the form of galactic winds. The marginally higher Mg/O ratio of $3.3\\pm 2.2$ implies steeper IMF\\@. \\item Abundances of Fe and Si show concentration in the central region, and a high Ni/Fe ratio is suggested. These results are consistent with that they are synthesized by SN~Ia in the central galaxy. \\item The non-thermal energy density necessary to support the cavity implies $\\mathcal{K}/\\mathcal{F}=690$, namely almost 700 times larger energy than that of electrons needs to be contained the cavity. The lack of the central AGN or the trailing radio feature seems to suggests that the origin of the cavity in HCG~62 may not be directly related to AGN activities. \\item We looked into alternative scenarios for the cavity creation. A clump of very hot gas and fast motion of the central galaxy were considered, but more observational evidences are necessary to perform a quantitative evaluation. \\end{itemize} \\bigskip Thanks are given to an anonymous referee for useful comments which improved the original manuscript. Part of this work was financially supported by a Research Fellowship for Young Scientists from JSPS and Grant-in-Aid for Scientific Research (No.\\ 16340077) from the Japan Society for the Promotion of Science, and also by a Grant-in-Aid of the Ministry of Education, Culture, Sports, Science and Technology (14079103; 16340077). N.~O.\\ acknowledges support from the Special Postdoctoral Researchers Program of RIKEN." }, "0607/astro-ph0607190_arXiv.txt": { "abstract": "We present the V band variability analysis of the point sources in the Faint Sky Variability Survey on time scales from 24 minutes to tens of days. We find that about one percent of the point sources down to V = 24 are variables. We discuss the variability detection probabilities for each field depending on field sampling, amplitude and timescale of the variability. The combination of colour and variability information allows us to explore the fraction of variable sources for different spectral types. We find that about 50 percent of the variables show variability timescales shorter than 6 hours. The total number of variables is dominated by main sequence sources. The distribution of variables with spectral type is fairly constant along the main sequence, with 1 per cent of the sources being variable, except at the blue end of the main sequence, between spectral types F0--F5, where the fraction of variable sources increases to about 2 percent. For bluer sources, above the main sequence, this percentage increases to about 3.5. We find that the combination of the sampling and the number of observations allows us to determine the variability timescales and amplitudes for a maximum of 40 percent of the variables found. About a third of the total number of short timescale variables found in the survey were not detected in either B or/and I. These show a similar variability timescale distribution to that found for the variables detected in all three bands. ", "introduction": "There is a wide range of photometrically variable systems in the universe. The range of timescales on which these systems vary is as wide as the physical processes that produce their variability. For example we have intrinsically variable stars, where the variability is caused by changes in their internal structure or atmosphere that vary with timescales of minutes to years \\cite{bg94}. Other stars show variability because they rotate and their surface is inhomogeneous, e.g. because of star spots, \\cite{b05}, or because they form part of a binary or multiple system and their revolution around the centre of mass of the system results in changes on the detected flux due to the changing aspect of a non-isotropically emitting surface or eclipses. This is also the case for planets orbiting stars. The timescale of the variability in this case is dictated by the orbital parameters of the system and can range from seconds to years. Near Earth Objects (NEOs), such as asteroids, also show variability as they rotate and are non-spherical. We find photometric variability in extragalactic objects as well, such as quasars, where the variability is probably the result of material being accreted by the central engine, or ``one of'' systems such as gamma ray bursts (GRB) or supernovae (SNe) where the variability is produced by intrinsic changes in the structure of an astronomical object that take place only once. The study of variability provides important information about the physical nature of the variable objects, leads to the discovery of new classes of objects, helps to study the physical structure of stars, e.g. pulsating stars, allows us to obtain information on galactic structure through the use of variables such as RR~Lyrae as standard candles, and is the key to determining extra-galactic distances through the use of standard candles such as Cepheids and supernovae Type Ia. Most of our knowledge of variability is based on the study of apparently bright sources, which naturally selects members of {\\it intrinsically} bright populations. At present little is known about variability of intrinsically fainter populations because in bright samples they are lacking altogether or are only represented by a few members. The Faint Sky Variability Survey (FSVS; Groot et al. \\shortcite{groot03}) was designed to account for this deficit by studying two unexplored regions of the variability space: the short timescale variability region (down to tens of minutes) and the intrinsically faint variable sources (down to V = 24 mag) at mid and high Galactic latitudes. The FSVS also contains colour information for all targets, giving us the option of positioning objects in the colour-colour diagram as well as finding the variability timescales and amplitudes that characterise them. The main aims of the FSVS are thus to obtain a map of a region of the Galaxy ($\\sim$21\\,deg$^2$) in variability and colour space, to determine the population density of the different variable objects that reside in the Galaxy and to find the photometric signature of up-to-now unknown intrinsically faint variable populations. In this paper we explore these three goals. There are other surveys that study the variable optical sky, each emphasising one aspect or one particular region of this parameter space. The timescales sampled, depth and sky coverage of different variability surveys varies depending on the astronomical objects they are designed to study. For example, with a brightness limit similar to the FSVS, Street et al. \\shortcite{s05} study the variability around an open cluster with timescales longer than a few hours, and Ramsay \\&\\ Hakala \\shortcite{rh05} study the rapid variability (down to 2 minutes) of objects as faint as V$\\sim$22.5. Of great interest is the Deep Lens Survey (DLS) that, in a similar way to the FSVS, combines colour and variability information and explores similar variability timescales \\cite{dls}. Becker et al. \\shortcite{dls} also provide a comprehensive review of past and on-going variability surveys. The future of optical variability surveys looks quite promising with the advent of large aperture telescopes such as the Large-aperture Synoptic Survey 8.4\\,m Telescope \\cite{tyson02}, the 4\\,m telescope VISTA and the 2.5\\,m VLT Survey Telescope. ", "conclusions": "We have analysed the short timescale variability information contained in the FSVS and find that about 1 per cent of all point sources are variable. Of those variables, about 50 per cent show variability timescales shorter than 6 hours, 22 per cent show variabilities between 6 hours and 1 day, 20 per cent between 1 and 4 days and 8 per cent show periods longer than 4 days. The distribution of variables with spectral type is fairly constant along the main sequence, with 1 per cent of all the sources being variable, except at the blue end of the main sequence where the fraction of variable sources increases possibly due to contamination by non main sequence sources. Above the main sequence, beyond the blue cut-off at (B$-$V)$<$0.38, we find that the fraction of variables increases to 3.5 percent. The highest space density of variables found in the FSVS (i.e. 17 per deg$^2$) show periods below 12 hours. These include CVs, RR~Lyr stars, and other short period pulsators such as $\\delta$\\,Scuti stars. We find a density of 4 variables per deg$^2$ centred at a 1 day period which includes longer period CVs, RR~Lyr and other pulsators like $\\gamma$\\,Doradus stars and Pop II Cepheids. A space density of 2 variables per deg$^2$ at 3.75 days includes, some longer period CVs, $\\gamma$\\,Doradus stars, Pop II Cepheids and longer period pulsators such as subdwarf B stars. At 12.75 days we also find 2 variables per deg$^2$. These would be mainly binaries with those orbital periods and Pop II Cepheids. It is easier to compare these space densities with those expected for the mentioned populations when we combine the period information with the colours of the populations under study. The case of CVs and many pulsators is complicated as they appear mixed through several period and colour ranges and in many cases it is necessary to obtain spectra to confirm the nature of the variable source. The space densities of CVs and subdwarf B stars will be studied in detail in a future paper. In the case of RR Lyr stars, we find 3 certain members and 9 other candidates down to V = 21.6. Assuming we have detected all RR Lyr between V = 16--22, we determine a space density of $\\sim$10$^{-3}$kpc$^{-3}$ in agreement with the space density determined by Preston, Shectman \\&\\ Beers \\shortcite{psb91} at a distance of 100--150kpc from the Galactic Centre. By using the floating mean periodogram, we have determined the most likely periods and amplitudes of a fraction of the variables found in the FSVS. We find that we are complete down to V = 22 for CVs in the minimum period (80 min) as long as they show variability amplitudes of the order of 0.4 mag. We are complete down to V = 22 for periods between 80 min and 1 day in a 17.82\\,deg$^2$ area of the survey as long as the amplitude of the variability is at least 0.7 mag. This includes most RR Lyr stars. We will be able to detect RR Lyr also down to V = 23 when their variability amplitudes are at least 1.5 mag." }, "0607/astro-ph0607473_arXiv.txt": { "abstract": "Analysis of {\\it INTEGRAL}/IBIS survey observations has revealed that the rare intermediate polar and asynchronous polar cataclysmic variables are consistently found to emit in the 20--100~keV energy band, whereas synchronous polars and the common non-magnetic CVs rarely do so. From the correlation of a candidate {\\it INTEGRAL}/IBIS survey source list with a CV catalogue, 15 CV detections by IBIS have been established including a new {\\it INTEGRAL} source IGR~J06253+7334. The properties of these sources and 4 additional CV candidates are discussed in the context of their 20--100~keV emission characteristics and we conclude that the {\\it INTEGRAL} mission is an important tool in the detection of new magnetic CV systems. Furthermore, analysis of the time-averaged spectra of CVs detected by {\\it INTEGRAL} indicate that although there is little difference between the spectral slopes of the different sub-types, intermediate polars may be considerably more luminous than polars in the soft gamma-ray regime. We also present the detection of an unusual high-energy burst from V1223~Sgr discovered by inspection of the IBIS light-curve. Additionally, we have compared the IBIS and optical AAVSO light-curves of SS Cyg and extracted IBIS spectra during single periods of optical outburst and quiescence. We find that the 20--100~keV flux is an order of magnitude greater during optical quiescence. This is in agreement with previous studies which show that the hard X-ray component of SS Cyg is suppressed during high accretion states. ", "introduction": " ", "conclusions": "" }, "0607/astro-ph0607229_arXiv.txt": { "abstract": "{ We report the results of analyzing {\\it Swift}, {\\it RossiXTE}, and {\\it INTEGRAL} data of \\J1922, a likely transient X-ray source discovered by Swift/BAT. Both the fast variability measured by the {\\it RXTE}/PCA and the combined {\\it Swift}/XRT, {\\it RXTE}/PCA, and {\\it INTEGRAL}/ISGRI (0.5--100 keV) energy spectrum suggest that the system is a neutron-star low-mass X-ray binary or a black-hole candidate at low accretion levels. The non-simultaneous spectra are consistent with the same spectral shape and flux, suggesting little variability over the period July to October 2005, but the analysis of archival {\\it INTEGRAL} data shows that the source was not detected in 2003--2004, suggesting a transient or strongly variable behavior. ", "introduction": "\\label{sec:intro} \\J1922\\ was discovered during the {\\it Swift} Burst Alert Telescope (BAT) hard X-ray survey. The BAT survey covers the 15--200 keV band and the time range December 2004 - March 2005 where \\J1922\\ was reported as $>5.5\\sigma$ detection \\citep[][]{tueller06a,tueller06b}. The BAT detection was also confirmed in a follow-up observation with the {\\it Swift} X-ray telescope (XRT) in the 0.5--10 keV energy band. The source was found at the best-fit position $\\alpha_{\\rm J2000} = 19^{\\rm h}22^{\\rm m}37\\fs0$ and $\\delta_{\\rm J2000} = -17{\\degr}17\\arcmin02\\farcs6$ with an estimated uncertainty of $3\\farcs2$ (90\\% confidence). From the Ultra-Violet Optical Telescope (UVOT) onboard {\\it Swift}, the observed counterpart was not consistent with the Palomar survey \\citep{tueller06b}. Using the XRT data, the source spectrum was consistent with an absorbed power-law model with a photon index of $\\Gamma=2.05\\pm0.05$ and a relatively low equivalent hydrogen column density $N_H=(1\\pm0.5)\\times10^{21}$ cm$^{-2}$ \\citep[][]{tueller06a,tueller06b}. \\J1922\\ was also observed as a target of opportunity (ToO) performed on October 21, 2005 with the {\\it Rossi X-ray Timing Explorer} ({\\em RXTE}), and the data were made publicly available. The source was also detected serendipitously during the {\\it International Gamma-Ray Astrophysics Laboratory} ({\\em INTEGRAL}) ToO observation of HETE J1900.1-2455 performed from November 10--12, 2005. In this letter we report the result of the spectral and timing analysis of \\J1922\\ using the {\\it Swift}, {\\em RXTE}, and {\\em INTEGRAL} data. We analyze the broad band spectrum from 0.5--100 keV and perform a timing analysis to identify the nature of this new source. ", "conclusions": "\\label{sec:discussion} The broad-band spectrum (0.5--100 keV) allowed us to perform an improved spectral analysis for the new source \\J1922\\ using the {\\it Swift}/XRT, {\\it RXTE}/PCA, and {\\it INTEGRAL}/ISGRI data. The best fit to the data required a two-component model, a cutoff \\pl, or a thermal Comptonization model, together with a soft component, see Table \\ref{table:spec}. The hard spectral component contributes most of the observed flux (76\\%), even though a soft \\bb\\ component is needed by the data. Observationally, the best-fit parameters are similar to those observed in BHC in their low/hard state \\citep[see e.g.][]{wilms06} and in weakly magnetic neutron-star LMXB in their low-luminosity states \\citep[see e.g.][]{barret00}. In the low/hard or low-luminosity state hard X-ray components extending up to energies of a few hundred keV have been clearly detected in these systems. In NS systems, the hard spectrum is dominated by a \\pl-like component, with a typical slope of $\\Gamma \\sim1.5-2.5$, which is followed by an exponential cutoff at energies often $\\gtrsim$ 20 keV \\citep[see e.g.,][]{barret00}, with some contribution from an additional soft thermal component, $kT_{\\rm soft}$. In BH systems in their low/hard state, the slope is around $\\Gamma \\sim 1.5-1.6$, with a high-energy cutoff of a few dozen keV \\citep[see e.g.,][]{wilms06}, while contribution from a soft thermal component is observable only when the interstellar absorption is not too high \\citep[see e.g.,][]{frontera01}. The soft thermal emission, $kT_{\\rm soft}$, could be associated to the radiation from the accretion disc either for a NS of BHC source. The \\pl\\ is usually interpreted as Comptonization of seed photons in hot, optical-depth plasma. Using the \\comptt\\ model, the hard spectrum is described by unsaturated Comptonization of soft seed photons, $kT_{\\rm seed}\\sim 0.4$ keV, in the hot $kT_{\\rm e} \\sim 11$ keV optically-thick $\\tau \\sim 3$ plasma. An important parameter is the distance to the source. Its galactic coordinates are $l_{II}$=20.7, $b_{II}$=-14.5, therefore in the direction of the galactic bulge and substantially below the galactic plane. Were the source within the galactic plane, its distance would be rather small ($\\sim1.5$ kpc for a thin disc); the derived absorption of $(1-2)\\times 10^{21}$cm$^{-2}$ is compatible with the total galactic absorption in that direction as estimated from HI maps. Assuming a distance to the galactic center of 8 kpc and a bulge radius of 3 kpc, we conclude that the distance to \\J1922\\ is between 5 and 11 kpc. With these distance estimates, the unabsorbed 0.1--100 keV flux of (3.3--3.9)$\\times 10^{-10}$erg cm$^{-2}$ s$^{-1}$ (considering the different models) translates to a luminosity of (1--5)$\\times 10^{36}$ erg s$^{-1}$, compatible with both NS and BH sources at low accretion-rate levels. We assume the source is probably located in the galactic bulge. At the minimal distance of 5 kpc, the \\bb\\ fits would imply an emission radius of only 3--4 km. This hints at a larger source distance. The radii obtained with the \\dbb\\ fits are also very small and depend crucially on the inclination of the source. Using a mean inclination angle of 60$^{\\circ}$ and a source distance of 10 kpc, an acceptable $R_{\\rm in}\\sim 13-18$ km can be obtained. We note that the distance estimates assume that the observed emission originates from the entire \\bb\\ surface facing the observer. The presence of an obscuring structure, such as an accretion disc or stream, will affect those estimates. We conclude that we observed \\J1922\\ in its hard state, during which it emits hard X-rays up to 100 keV. The soft excess, $kT_{\\rm soft}\\sim 0.4$ keV detected at low energies is most likely originating in the accretion disc. As for the hard component, it most likely originates in the Comptonization of soft seed photons, $kT_{\\rm seed}\\sim 0.53$ keV in a hot plasma. Our results are compatible with the soft emission being the source of the seed photons, i.e. $kT_{\\rm soft}=kT_{\\rm seed}$. From the best-fit spectral parameters and the timing properties, we cannot tell whether the objects harbors a NS or a BH. Interestingly, we could obtain a satisfactory fit with the same model as all three non-simultaneous spectra ({\\it Swift}/XRT, {\\it RXTE}/PCA, and {\\it INTEGRAL}/ISGRI), indicating that its spectrum most likely remained constant if the source varied between the observations. Once again, this is a typical observational fact both in NS and BH systems at low luminosity. The measured \\nh\\ value is consistent with the expected Galactic value in the source direction \\J1922. This consistency suggests the absence of intrinsic absorption and, together with the position in the sky, indicates that it is most likely located in the Galactic bulge. The lack of iron-line emission is consistent with the absence of a reflected component in the broad band spectrum of the source. The fact that \\J1922\\ was not detected in the ISGRI data from March 2003 to October 2004, for a total 20--60 keV exposure time of 420 ks, indicates that the system is either very variable on long time scales or transient. Across the period spanned by our observations (2005 July to October), both flux level and spectrum were compatible with being constant." }, "0607/astro-ph0607535_arXiv.txt": { "abstract": "Beijing-Arizona-Taiwan-Connecticut (BATC) multi-band photometric data in the field of open cluster M48 are used to determine its membership. By comparing observed spectral energy distributions (SEDs) of stars with theoretical ones, membership probabilities of 750 stars with limiting magnitude of 15.0 in BATC $c$ band ($\\lambda_{eff}=4194$\\,\\AA) are determined. 323 stars with membership probabilities higher than 30\\% are considered as candidate members of M48. Comparing membership probabilities of 229 common stars obtained by the present method and the proper-motion based methods, a 80\\% agreement among these methods is obtained. ", "introduction": "Open clusters (OCs) have long been recognized as important tools in the study of the Galactic disk. They have been used to determine spiral arm structure, to map the rotation curve of the Galaxy, to investigate the mechanisms of star formation, to constrain the initial luminosity and mass functions, and to define disk abundance gradients and age-metallicity relationship \\citep{fr95,fr02,ch03,sal04,von05,bo06}. The first step to determine the physical parameters of an open cluster is to select probable members in the vicinity of this cluster. Omitting the parameter of position which provides us with the first clue as to the existence or nonexistence of a cluster, there are two types of independent methods by which to establish cluster membership: photometric and kinematic \\citep{ca90}. When kinematic data are available, it is accepted that the membership probabilities obtained from the analysis of proper motions or radial velocities are more reliable. Unfortunately, few clusters have a large, homogeneous radial velocity data to permit a detailed analysis of the entire cluster. It is well known that the proper motion analysis is at present the most valuable criterion to establish membership probabilities in OCs \\citep{sl77,ca85}. The first attempt to statistically determine the membership of an open cluster based on proper motion data was made by \\citet{va58}. They point out that the cluster and field probability density functions can be modeled as bivariate Gaussian distributions: a circular normal distribution for the cluster population, and an elliptic normal distribution for the field population. A maximum likelihood principle was developed to obtain the distribution parameters of clusters and membership probabilities of individual stars \\citep{san71}. The parametric Vasilevskis-Sanders method has been frequently used to derive the membership of star clusters \\citep{wu02a, wu02b}. However, the hypotheses for cluster and field stars distribution in the parametric Vasilevskis-Sanders method are not always true. Even if the hypotheses are realistic for some clusters, the parametric method will fails when the cluster member-to-field star ratio is small \\citep{ca90}. This method does not work in the case of significant internal motion in a cluster or its rotation \\citep{sl77,ja06}. In order to overcome some difficulties that arise from the parametric Vasilevskis-Sanders method, \\citet{ca90} developed a non-parametric approach to the membership problem. The key of their method is to perform an empirical determination of the probability density functions, without relying on any previous assumption about their profiles \\citep{ga98}. The non-parametric method has been used in recent years to determine membership of several OCs \\citep{ga98, ba04, ba05}. More recently, \\citet{ja06} presented a new method that enlarges the statistical distance between the cluster members and field stars by revealing the group of stars with the least relative velocities without any assumptions about the distribution of field stars. Although proper-motion based methods have been considered to be the most reliable for cluster -- field segregation, however, these methods need high precision proper motion data. Several decades or even more than one hundred years are needed to obtain these high precision proper motion data. So, membership determination based on proper motions is very time-consuming. On the other hand, most of proper motions of stars are derived from plate data, high precision results can obtained only for bright stars. Color-magnitude diagrams (CMDs) or color-color diagrams are used to derive the fundamental parameters of clusters, and in general, are also used to determine the membership of cluster. If a field of cluster and, a comparison field which only including field stars, are both obtained for study, the CMD or color-color diagram of comparison-field stars are subtracted from that of field including the cluster, the resulting difference diagram will show most of members of this cluster \\citep{mei00,von02}. However, the method obviously requires that the foreground and background of both the cluster and the comparison field, are essentially identical, but above conditions are not satisfied due to the non-uniformity of the background \\citep{ba83}. On the other hand, if no comparison field is obtained, theoretical isochrones are used to match the observed star sequence, and stars residing in the neighborhood of the best-fitting isochrone are considered as members of this cluster. At the mean time, the fundamental parameters of this cluster such as age, metallicity, distance, and reddening can be derived \\citep{jef01}. In general, it can't be ensured reasonably to derive the four fundamental physical parameters just using observational results in two or three bands without knowing any one of those parameters in advance. So, it is much more difficult to determine the members of cluster just based only on CMDs or color-color diagrams without knowing any physical parameters of cluster. More recently, using $UBVRI$ and $uvby\\textrm{H}\\beta$ photometric data, \\citet{sar99} and \\citet{twa00} present new method to distinguish cluster and field stars based only on observed star sequences in CMDs or color-color diagrams. \\citet[and references therein]{kal04} present the synthetic CMD method to derived the fundamental parameters of cluster. The BATC photometric system including 15 bands provides a sort of \\textsl{low-resolution spectroscopy} that defines the spectral energy distributions (SEDs) of each star. If the membership are known with the proper-motion based method, fitting the observed SEDs of cluster member stars with theoretical models has an advantage that it has more observational data than the number of free parameters to be solved. Fitting the observed SEDs with theoretical ones can also be used to derive both the membership and fundamental parameters of a cluster at the same time. In this paper, we develop a method based on fitting SEDs of stars in a field of star cluster with theoretical ones to determine the membership and fundamental parameters of this cluster at the same time. Comparing with traditional photometric membership determination methods based on CMDs or color-color diagrams, the advantage of the present method is that no comparison field is needed and the fundamental parameters of cluster can also be derived more reliable. Moreover, this method can present the membership probabilities of stars. Comparing proper-motion based methods, the observational data used by the present method can be obtained within one month, no any assumptions about the distribution profiles of cluster and field stars are needed, and membership probabilities of stars with more fainter magnitude can be obtained. We apply this method to open cluster M48 with 13 bands of BATC photometric data. Section \\ref{data} describes the proper motion and photometric data as well as theoretical model used for applying the present method. The details of our present method are presented in Section \\ref{method}. We apply our present method to open cluster M48 and compare our results with those derived by proper-motion based methods in Section \\ref{disc}. Finally, a summary is presented in Section \\ref{conc}. ", "conclusions": "} In this paper, using BATC 13 bands photometric data, we develop a SED method to determine membership and fundamental parameters of a star cluster simultaneously. Membership probabilities of 750 stars with limiting magnitude 15.0 in BATC $c$ band are derived for open cluster M48, 323 stars with membership probabilities higher than 30\\% are considered as member stars of M48. Comparing with the membership determinations of 229 common stars taken from the parametric \\citep{wu02b} or non-parametric \\citep{ba05} proper-motion based methods, we get 80\\% agreement with them. The present SED method can investigate membership for stars with fainter magnitude than that of proper-motion based methods (about 2 magnitude deeper in the case of M48) and membership probabilities of 521 stars in the field of M48 are derived at the first time. At the same time, the fundamental parameters of M48 are also derived and consist with the previous determinations \\citep{ri04, wu05, ba05}." }, "0607/astro-ph0607044_arXiv.txt": { "abstract": "We investigate the physical mechanism of the GZ-effect that could explain the production of multiple primaries from an event initiated outside the Earth's atmosphere. In this case, there would correspondingly be multiple extensive air showers in temporal coincidence at ground, even for detectors separated by many kilometers, and also showers initiated by primaries of different energies could consequently have a common source. We analyse the perspectives and limits of some models and discuss the experimental counterparts. ", "introduction": "The history of a Cosmic ray from the production point and through the acceleration sites undergoes many changes in velocity and eventually in chemical structure. The nature of such a primary, either a heavy nucleus or proton or gamma-ray or whatever, influences strongly the byproduct of its interaction with the medium: interstellar matter, Cosmic Microwave Background photons (CMB) or local interstellar/intergalactic/galactic magne\\-tic fields \\cite{GZK66,PSB76,MS98,SS99,BILS02}. A great interest is therefore devoted to understanding the abundances of protons and relative chemical composition of the CR's flux, in terms of other elements or ions. \\\\ Among the various kinds of projectiles hitting the Earth's atmosphere, we will consider the fragments deriving from the photo disintegration of heavy nuclei (for example Fe) when interacting with the solar magnetic field \\cite{GZ60}: the crucial aspect of this fragmentation relies in the possibility of detecting on Earth two (or eventually more) of them. In fact, the influence of the solar field permeates the space surrounding the Sun up to distances limited to 3 or 4 AU. This relatively small volume allows the fragments to arrive on Earth almost simultaneously and spaced ranging some up to few thousands Km. Thus, the Extensive Air Showers (EAS) generated by the two projectiles when hitting the atmosphere would be temporally as well spatially correlated and detectable when having many detectors placed at different distances, some closely and some widely spaced. \\\\ The arrival rate computed in the present paper heavily depends on the energy of the incoming particles and for this reason we include in our estimate small variations (some units) providing important differences in the expectation values. \\\\ What is strongly encouraging in the experimental search for this peculiar phenomena is the possibility of being detected by the new experiment ``Extreme Energy Events'' (EEE) which is starting in Italy \\cite{eee}. In fact, the disposition of the particle detectors is planned inside numerous High Schools over all the Italian territory (about 300.000 Km$^2$), more densely inside the cities, from south to north. ", "conclusions": "The event rates are strongly dependent on the primary energy, as can be seen in the case of a very energetic one, since the flux varies accordingly by orders of magnitude. The fragments can generate showers over a wide range distances between them, from few to thousands Km, thus not requiring a peculiarly spaced array of detectors, but moreover a large area. If we consider primaries with a given $E$ and flux, our calculations split essentially in two classes: the first relies on the events characterized by a large separation of fragments on Earth $\\delta >200 $ Km, i.e. events originated closer to the solar influence, at distances of order $>2 ~AU$ from the Earth; the second is given by a smaller separation of the fragments $\\delta < 15 $ Km, corresponding to primaries photo disintegrated at distances $< 2 ~AU$.\\\\ Considering a relatively modest variation for the incoming particle energy (about few units) and for the values computed for $\\eta_{GZ}$, we find that the rate of GZ events expected can be effectively high (see Tables), leaving very promising work for the observation of the GZ effect." }, "0607/astro-ph0607272_arXiv.txt": { "abstract": "\\hspace{3mm} Using data from the DEEP2 Galaxy Redshift Survey and \\textit{HST}/ACS imaging in the Extended Groth Strip, we select nearly 100 interacting galaxy systems including kinematic close pairs and morphologically identified merging galaxies. $Spitzer$ MIPS \\24m fluxes of these systems reflect the current dusty star formation activity, and at a fixed stellar mass (\\sm) the median infrared luminosity (\\lir) among merging galaxies and close pairs of blue galaxies is twice (1.9$\\pm$0.4) that of control pairs drawn from isolated blue galaxies. Enhancement declines with galaxy separation, being strongest in close pairs and mergers and weaker in wide pairs compared to the control sample. At $\\overline{z} \\sim 0.9$, $7.1\\%\\pm4.3\\%$ of massive interacting galaxies (\\sm $>$ $2\\times10^{10} M_{\\odot}$) are found to be ULIRGs, compared to $2.6\\%\\pm0.7\\%$ in the control sample. The large spread of \\speir among interacting galaxies suggests that this enhancement may depend on the merger stage as well as other as yet unidentified factors (e.g., galaxy structure, mass ratio, orbital characteristics, presence of AGN or bar). The contribution of interacting systems to the total IR luminosity density is moderate ($\\la 36\\%$). ", "introduction": "Galaxy-galaxy interaction has long been regarded as a key process in galaxy evolution, especially as a mechanism for enhancing star formation during mergers \\citep{lar78, bar00}. Hydrodynamic $N$-body simulations show that active star formation can be triggered by gaseous inflows resulting from mergers of gas-rich galaxies \\citep{mih96, bar04,cox04}. Interaction-triggered star formation is also thought to be responsible for luminous infrared sources. In the local universe, luminous infrared galaxies (LIRGs) and ultra luminous infrared galaxies (ULIRGs) are primarily merging systems \\citep{san88,bor99}. Nevertheless, by studying various star formation indicators of interacting galaxies and normal galaxies, \\citet{ber03} concluded that galaxy interactions in general are inefficient triggers of starbursts; interactions are a necessary but not sufficient condition to trigger violent starbursts. The importance of galaxy interactions in the volume-averaged galaxy star formation rate (SFR) remains an open question. Recent studies of mid-IR (MIR) sources at a median redshift of $z\\sim$0.7 suggest that the IR density at that epoch is dominated by morphologically normal galaxies instead of strongly interacting galaxies \\citep{bel05,mel05}. Two main factors may contribute to this result. First, only a small fraction of the galaxy population may be undergoing a major merger at any given time \\citep{car00,bun04,lin04}. Second, the overall SFR in normal galaxies at $z\\sim$0.7 may be enhanced relative to the local population \\citep{bel05}, perhaps as the result of internal processes such as a higher gas fraction leading to a higher SFR, such that galaxy interactions may have less dramatic effects on triggering star formation at that epoch and/or may be harder to identify. It is the aim of this Letter to examine this second hypothesis. This Letter presents an analysis of the IR properties of close kinematic galaxy pairs, morphologically selected merging galaxies, and a control sample of randomly selected pairs of isolated galaxies, in the range $0.1$1~Gyr-old field dwarfs. Adopting a radius from current models of sub-stellar evolution, we hence obtain that the effective temperature of HD~203030B is only $1206^{+74}_{-116}$~K, markedly lower than the $\\approx$1440~K effective temperatures of field L/T transition dwarfs. The temperature discrepancy can be resolved if either: (1) the ages of field brown dwarfs have been over-estimated by a factor of $\\approx$1.5, leading to under-estimated radii, or (2) the lower effective temperature of HD~203030B is related to its young age, implying that the effective temperature at the L/T transition is gravity-dependent. ", "introduction": "} After more than a decade of study, the fundamental parameters of brown dwarfs remain poorly constrained. Unlike hydrogen-burning stars, whose effective temperatures and luminosities are approximately age-independent (at a fixed mass and composition) on the main sequence, brown dwarfs cool and dim continuously as they age. Only one of the fundamental sub-stellar parameters, luminosity, can be estimated independently of the others, through the use of trigonometric parallax and empirically determined bolometric corrections. Any attempts to resolve the remaining degeneracies among temperatures, ages, masses, and radii rely on the fortuitous discovery of brown dwarfs in association with stars or clusters of known age. In such cases, brown-dwarf ages can be inferred by demonstrating the physical association of the brown dwarfs, usually as common proper motion secondaries or as members of stellar associations, with stars that have known ages. Brown-dwarf masses can be measured from orbital astrometry and radial-velocity monitoring of close ($<$5--10~AU) binary systems, in which at least one component is sub-stellar. Finally, brown-dwarf radii, and hence, effective temperatures and surface gravities, can be measured in sub-stellar eclipsing binaries. The union of these three fortuitous cases, eclipsing brown-dwarf binaries associated with stars, offers the best opportunity for empirical determination of sub-stellar parameters. However, such unusual systems are extremely rare---the first one, 2MASSJ~05352184--0546085, has only recently been reported \\citep*{stassun_etal06}. In all other cases involving brown dwarfs, the estimation of their properties relies to various degrees on the use of theoretical models of sub-stellar evolution \\citep[e.g.,][]{burrows_etal97, burrows_etal01, chabrier_etal00, baraffe_etal03}. Consequently, sub-stellar evolution models remain, for the most part, empirically unconfirmed. Partial tests have been carried out for systems other than the ideal eclipsing double-line systems. \\citet{close_etal05} and \\citet*{luhman_etal05} used astrometric measurements of the multiple system AB~Dor A/B/C to test the mass---age---luminosity relation at the stellar/sub-stellar boundary. The two analyses offer examples of how sub-stellar evolution models can be tested without considering all degenerate sub-stellar parameters---in this case, excluding effective temperature. We note, however, that the conclusions of the two teams differ \\citep[see][for further discussion]{nielsen_etal05, luhman_potter06}, in part because the stellar parameters themselves (in particular, the stellar ages), although used as a reference in either study, may often not be known to the desired or believed level of accuracy. In a separate example, \\citet{mohanty_etal04b} and \\citet*{mohanty_etal04} tackle the full set of degenerate sub-stellar parameters using high-resolution spectroscopic observations of brown dwarfs in the Upper Scorpius and Taurus associations. Their studies rely on model spectra of brown-dwarf photospheres, constructed independently of sub-stellar evolutionary models, to resolve the mass---effective temperature---luminosity degeneracy. Even though model-dependent, they allow a self-consistent comparison between sub-stellar cooling and photospheric models in the context of empirical data. In the present paper we take a similar approach and discuss the degeneracies among sub-stellar age, luminosity, and effective temperature (i.e., excluding mass) in the context of models of sub-stellar evolution. We address these fundamental properties with the help of a new brown dwarf that we discovered as a common proper motion companion to the 130--400~Myr-old main-sequence G8~V star HD~203030 (HIP~105232; \\S\\ref{sec_observations}). The companion, HD~203030B, has a spectral type of L7.5, and thus lies very near the transition between L- and T-type sub-stellar photospheres, characterized by the settling of dust and the appearance of methane absorption in the near-IR spectra of brown dwarfs. Because dust-settling occurs over a very narrow range of effective temperatures, 1500--1300~K (spectral types L6--T4), a fact inferred both theoretically \\citep{ackerman_marley01, tsuji02} and semi-empirically \\citep{golimowski_etal04, vrba_etal04}, the effective temperature of HD~203030B is expected to be constrained very well. Given the known age and luminosity of HD~203030B, we can find its radius and thus offer a constraint on the theory (\\S\\ref{sec_properties}). Our approach is not fully empirical, because it relies on the semi-empirical results of \\citeauthor{golimowski_etal04} and \\citeauthor{vrba_etal04}; both studies rely on evolution models for the sub-stellar age---mass---radius relations. However, by adopting the same theoretical evolution models \\citep{burrows_etal97, burrows_etal01} as in these two studies, our analysis provides a test of the self-consistency of the theoretical models (\\S\\ref{sec_lt_teff}). ", "conclusions": "We have discovered a proper motion sub-stellar companion to the 130--400~Myr-old Sun-like star HD~203030. The companion, HD~203030B, is 11$\\farcs$9 away from the primary, at a projected orbital separation of 487~AU. Assuming coevality with the primary star, sub-stellar evolution models place the companion mass at $0.023^{+0.008}_{-0.011}\\Msun$. From $K$-band AO spectroscopy, we determine a spectral type of L7.5$\\pm$0.5 for HD~203030B. The brown dwarf is thus near the spectroscopic L/T transition, characterized by diminishing amounts of dust in sub-stellar photospheres. Because of its association with a main-sequence star of a known age and heliocentric distance, HD~203030B offers a rare chance to probe the L/T transition in a setting in which most of the degeneracies characteristic of sub-stellar models are resolved. We find that the bolometric luminosity of HD~203030B is comparable to that of $>$1~Gyr-old field dwarfs of similar spectral types, despite the $\\approx$0.5~dex lower gravity anticipated at the young age of HD~203030B. As a result, theoretical models of sub-stellar evolution predict a $\\approx$240~K cooler effective temperature for HD~203030B compared to the effective temperatures of its older counterparts in the field. We consider three hypotheses for the discrepancy: (1) that the bolometric luminosities and effective temperatures of late-L dwarfs in the field have been over-estimated because of unresolved binarity; (2) that the effective temperatures of late-L dwarfs in the field may have been over-estimated (equivalently, their radii under-estimated) because theoretical models may have over-predicted their ages; and (3) that the spectral type---effective temperature for late-L dwarfs is not age- and gravity-independent. Based on multiplicity studies of field brown dwarfs from high-angular resolution observations with {\\sl HST} and AO, we do not find a significant discrepancy between the effective temperatures (or luminosities) of brown dwarfs with known and unknown multiplicities, and rule out the first possibility. We test the second hypothesis by comparing the effective temperatures of late-L secondaries of main-sequence stars (with known ages and distances) to late-L dwarfs in the field. We find evidence that the late-L companions are $\\geq$100~K cooler and $\\sim$1.5 times younger than their isolated counterparts. However, because ages of individual stars cannot be claimed with such accuracy, the significance of this result is only marginal. The last hypothesis is a compromise between the discrepant effective temperatures of HD~203030B and late-L dwarfs in the field, and draws a parallel with the dependence of spectral type and effective temperature on surface gravity in stars \\citep{gray92a}. Both remaining hypotheses await testing from larger samples of brown dwarfs with known distances and ages, either as members of known stellar associations or as companions to stars." }, "0607/astro-ph0607208_arXiv.txt": { "abstract": "We review current observational constraints on the polarization of the Cosmic Microwave Background (CMB), with a particular emphasis on detecting the signature of primordial gravitational waves. We present an analytic solution to the Polanarev approximation for CMB polarization produced by primordial gravitational waves. This simplifies the calculation of the curl, or B-mode power spectrum associated with gravitational waves during the epoch of cosmological inflation. We compare our analytic method to existing numerical methods and also make predictions for the sensitivity of upcoming CMB polarization observations to the inflationary gravitational wave background. We show that upcoming experiments should be able either detect the relic gravitational wave background or completely rule out whole classes of inflationary models. ", "introduction": "The cosmic microwave background (CMB) is one of the most powerful and precise cosmological probes. Because the CMB photons we observe today probe the physics of the early universe during the epoch of linear gravity, the CMB is often referred to as a ``snapshot'' of the primordial universe. The CMB has the promise to address the most fundamental cosmological questions: the geometry and age of the universe, the matter-energy content of the universe, the ionization history and the spectrum of primordial perturbations. This review addresses the theoretical foundations of CMB polarization generated by cosmological gravitational waves with particular emphasis given to analytic and numerical results which encode its behavior most cogently. Gravitational waves, in contrast to adiabatic perturbations (which are the dominant source of CMB temperature and polarization anisotropy) imprint a unique divergence-free pattern of polarization on the sky\\cite{polnarev85,seljak1997}. This pattern is called ``B-mode''\\cite{Zalandseljak97} or ``curl-mode''\\cite{kamkossteb} polarization. Although there may be a substantial GWB contribution to CMB temperature anisotropy, its effect on the CMB temperature is nearly completely degenerate with other cosmological parameters\\cite{kamionkowski97,selzal97}. This is unsurprising since the CMB temperature is a scalar quantity. However, the tensorial nature of CMB polarization permits separation of scalar fluctuations from tensor, GWB-generated, fluctuations. CMB polarization maps can be decomposed into two terms\\cite{kamionkowski97,selzal97}. One term is the gradient of a scalar potential and is invariant under parity transforms (often called ``E-mode\" in analogy to the electric field/scalar potential). The second component is the curl of a vector potential (``B-mode\"). Scalar perturbations have no handedness so the primary CMB curl-mode component exists only if there is a GWB. Besides the obvious importance of a new method of (indirect) detection of gravitational waves, detection of the B-mode signal provides the cleanest, and perhaps only window into unique predictions of the inflationary cosmological paradigm. Numerical methods for simulating CMB temperature and polarization power spectra have revolutionized cosmology \\cite{SeljakZal1996,Lewis:2002ah}. Without such codes parameter estimation from CMB data sets would be difficult, if not impossible. However, the complexity of these codes has hidden the underlying physics. For this reason there is considerable interest in analytical approaches for calculating CMB polarization caused by cosmological gravitational waves (see for example Ref.~\\refcite{polnarev85,basko,ColeFrewinPolnarev95,keating1998,pritchardkamionkowski2005,zhaozhang2005}). This review elucidates an analytical approach to the problem using the fundamental physical characteristics of gravitational waves to significantly simplify calculations. Our method simplifies the comparison between theoretical predictions and future observational results. We compare our results to existing numerical methods and summarize current observational results. We conclude with by making predictions for ground, balloon, and space-based observations in the upcoming decade. ", "conclusions": "We have developed an analytic method to generate the predictions of the imprint of gravitational waves on the CMB. Using phenomenological models of inflation we predict both the CMB polarization spectra and the derived inflationary parameters: the tensor-to-scalar ratio and spectral index of the scalar perturbations. The combination of the later two observables allows for reconstruction of the dynamics of inflation. With these predictions in hand we have shown that upcoming CMB polarization observations will be able to detect or constrain the cosmological GWB and hence, inflation itself. These new technological advances now position observational cosmology at the threshold of an exhilarating era -- one in which CMB polarization data will winnow down inflation's vast model-space and test models of the early universe at energy scales approaching the GUT-scale; nearly one trillion times higher energy than accessible from particle accelerators." }, "0607/astro-ph0607500_arXiv.txt": { "abstract": "New {\\it Spitzer} imaging observations have revealed the structure around the Mira variable star R Hya to be a one-sided parabolic arc 100 arcsec to the West stretching from North to South. We successfully model R Hya and its surroundings in terms of an interaction of the stellar wind from an asymptotic giant branch (AGB) star with the interstellar medium (ISM) the star moves through. Our three-dimensional hydrodynamic simulation reproduces the structure as a bow shock into the oncoming ISM. We propose this as another explanation of detached shells around such stars which should be considered alongside current theories of internal origin. The simulation predicts the existence of a tail of ram-pressure-stripped AGB material stretching downstream. Indications for such a tail behind R Hya are seen in {\\it IRAS} maps. ", "introduction": "Large detached shells have been observed around several asymptotic giant branch (AGB) stars. They have been seen in IRAS images of dust emission \\citep{waters94,izumiura97}, CO line emission \\citep{olofsson96,schoier05}, and in a few cases [Na\\,{\\sc i}] and [K\\,{\\sc i}] emission as well as in the optical continuum \\citep{gonzalez01,gonzalez03}. \\cite{olofsson90} suggested that such shells are the result of mass-loss variations and in particular, a thermal pulse or He-flash. During a He-flash, an intense short-lived mass ejection is driven by the star reaching a critical luminosity. Thermal pulses are separated by phases of quiescent hydrogen burning lasting $10^4$--$10^5$\\,yr. The stellar evolution tracks calculated by \\cite{vassiliadis93} confirmed that mass-loss fluctuations during the thermal pulse cycle can lead to detached circumstellar shells. Hydrodynamic simulations by \\cite{steffen00} showed that a brief period of high mass loss can translate into a geometrically thin shell expanding around the star. This has become the standard explanation of detached dust shells around stars and observations have been interpreted as such \\citep{zijlstra92,speck00}. A separate explanation for large detached shells is the interaction of the AGB wind with the interstellar medium (ISM) \\citep{young93}. \\cite{zijlstra02a} used this for a giant ($\\sim4$ pc at D $\\sim700$ pc) detached shell surrounding an M3 III AGB star. They proposed that its AGB wind has been stopped by the surrounding ISM and the swept-up `wall' is now expanding at the local sound speed. Simulations by \\cite{villaver03} and \\cite{wareing06} confirmed the viability of this mechanism. It is likely that both mechanisms occur, but whether the external mechanism of an ISM wall, or the internal mechanism of long-term mass-loss variations is the dominant cause of detached shells is not known. Sch{\\\"o}ier et al. (2005) found that the derived masses of the shells increase and the expansion velocities decrease with increasing radial distance from the star. They suggest that the shell is sweeping up surrounding material from an earlier mass loss phase. However, ISM sweep-up could yield similar effects. A difference between the two mechanisms is that the internal one will normally give a spherical shell, while external mechanism will give a shape which depends on the motion of the star through the ISM. This is a testable prediction, if the proper motion of the star is known. Here we show that the detached shell around the Mira variable R Hya is due to an ISM bow shock. ", "conclusions": "\\subsection{The bow shock} The two-wind model can reproduce the appearance of the circumstellar structure around R Hya, with physical dimensions matching that of the bow shock. The head of the bow shock is in good agreement with the direction of motion. Small regions of emission downstream can be explained in terms of regions of higher density ISM encountered by the stellar wind. Such high density regions in the ISM have been shown to survive ablation by stellar winds \\citep{pittard05}. In our simulation, we find a high temperature of 35\\,000 K at the head of the bow shock, with material cooling rapidly as it moves down the tail. The high temperatures are consistent with the observed H$\\alpha$ emission \\citep{gaustad01,ueta06}. The high temperatures will affect the dust by collisional heating. However, it is impossible for us to say whether this heat from the local gas is more important than the stellar radiation or the interstellar radiation field (Speck et al. 2000) in controlling the dust temperature. The {\\it Spitzer} and {\\it IRAS} images indicate that the bow shock region is detected through dust emission; the fact that the emission is strongest at 60$\\mu$m indicates a dust temperature of order 60 K. But other emission processes may play a role. Shocked regions can show [O\\,{\\sc{i}}] 63$\\mu$m and 146$\\mu$m lines. The {\\it IRAS} 100$\\mu$m detection \\citep{hashimoto98}, in a band without strong lines, shows that dust contributes to the far-IR emission. \\subsection{Other cases: mass-loss variations versus ISM interaction} Cases of detached shells around AGB stars include U Ant \\citep{izumiura97}, U Hya \\citep{waters94}, Y CVn \\citep{izumiura96}, R Hya \\citep{hashimoto98} and IRAS02091+6333 \\citep{zijlstra02a} observed in the far-infrared; detached molecular (CO) gas shells around TT Cyg, S Sct, R Scl, U Ant, U Cam, V644 Sco and DR Ser (Olofsson et al. 1996, Sch{\\\"o}ier et al. 2005). All are carbon stars, apart from R Hya and IRAS02091+6333. It has been suggested that the thermal pulse scenario may only lead to mass-loss spikes for carbon stars (Sch{\\\"o}ier et al. 2005) but this is very much dependent on the assumed mass loss prescription which is not well known. {\\it IRAS} colours indicate that detached shells do also exist around oxygen-rich stars \\citep{zijlstra92} and many more stars are likely to show large shells (Young et al. 1993). An interesting comparison can be made between R Hya and U Hya. U Hya has a thin circular shell with a radius of 120 arcsec \\citep{waters94}. It has a space velocity of 50 \\kms\\ and the parallax gives a distance of 162 pc \\citep{perryman97}. The circular nature of the shell suggests an internal origin such as the mass-loss variations proposed by, for example, \\cite{zijlstra92} and Sch{\\\"o}ier et al. (2005). The age of the shell in this particular case is 6000 years. The ISM interaction may be radially further away from the star. The star TT Cyg is surrounded by a circular detached shell at a radius of approximately 40 arcsec \\citep{olofsson96}; the physical size of the shell is very similar to U Hya, in view of the larger distance of TT Cyg. The thin symmetrical appearance supports an internal origin. Interestingly, the space motion of TT Cyg (50 \\kms) is almost all in the radial direction away from us and any bow shock formed by an interaction with the ISM would appear circular on the sky. But the fact that the slight offset of the star from the centre of the shell is at right angles to the direction of proper motion, favours an internal origin in mass-loss variations. For other stars there is insufficient data to decide on the cause of the detached shells. It may be that the well-studied carbon stars with detached shells are mostly due to internal mechanisms, i.e. thermal pulses, while the fainter, less studied shells are dominated by ISM interactions. CO observations of shells around carbon stars reveal relatively high velocities (about 20 \\kms) which favours thermal pulse origins (Olofsson et al. 1996, Sch{\\\"o}ier et al. 2005). Multiple detached shells do exist in the AGB phase of evolution, e.g. S Sct and U Ant \\citep{gonzalez03} and also R Hya: thus, both mechanisms may occur simultaneously. Young et al. (1993) in their analysis of 76 AGB stars resolved by {\\it IRAS} found no evidence for distortions by interaction with the ISM. Since R Hya is included in their sample, this lack of evidence can be attributed to the poor spatial resolution and image quality of the {\\it IRAS} instrument. R Hya can be considered a typical case for mass-loss rate and ISM density although its space velocity is perhaps high. Lower velocity objects will have a bow shock and with higher mass-loss rates this will be located further from the star. In a zero velocity case, the bow shock transforms into a spherical swept-up shell of ISM material (Young et al. 1993, Speck et al. 2000, Zijlstra \\& Weinberger 2002). We predict that all AGB stars will show some degree of an AGB-ISM interaction, although bow shocks will be rarer. The interpretation of a detached shell in terms of a mass-loss variation must be considered with this ISM interaction in mind. \\subsection{Mass loss history} This model has shown that information usually gleaned from circumstellar dust shells around Mira variables can no longer be inferred in this situation. The bow shock has destroyed any mass loss history older than about $10^4$ years in the case of R Hya. Higher mass-loss rate and/or lower lower ISM density could increase this timescale as found by Young et al. (1993) and Zijlstra \\& Weinberger (2002). \\begin{figure} \\begin{center} \\includegraphics[width=7cm]{figure3.eps} \\caption{A figure showing an {\\it IRAS} $60 \\mu$m observation of the area around R Hya. North is up and East to the left. Evidence for a tail of material is indicated by the ellipse.} \\label{iras} \\end{center} \\end{figure} The simulations predict the occurence of a tail, consisting of swept-back ISM and stellar wind gas. The mass loss history can in principle still be traced down the length of the tail. {\\it IRAS} maps provide some indication for a tenuous detection of material downstream of R Hya as shown in Fig. \\ref{iras}. This material stretches up to 30 arcmin away or 1.5 pc at a distance of 165 pc. At 49.5 \\kms\\ this implies a minimum tail age of 30\\,000 years. Adding the 25\\,000 years it takes to form the stable bow shock, to represent the travel time from the star to the bow shock and down the tails, we predict R Hya has been losing mass for at least 55\\,000 years. If we consider it has taken 25\\,000 years to form the bow shock and after this the bow shock is in a steady state, an appreciable amount of mass is in the bow shock. The estimate from the simulation is $2.6\\times10^{-3}$ \\msun\\ in the region of the bow shock defined as a hemispherical upwind shell centred on the star with inner radius of 0.08 pc and a thickness of 0.1 pc. Our analytical estimate of the mass in the bow shock is consistent with this estimate, suggesting there is four times more stellar mass in the bow shock than ISM mass. At a dust temperature of 60 K and assuming a dust to gas ratio of 100, we predict a 100 $\\mu$m flux of 15 Jy, in broad agreement with the fluxes observed by \\cite{ueta06}. We use a constant stellar mass-loss rate and wind velocity. However, the dust shell models of various people (\\protect{Zijlstra} \\& Weinberger 1992, Sch{\\\"o}ier et al. 2005) indicate that these quantities are variable, particularly during the thermal pulses. Models indicate that in the course of thermal pulses the stellar radius changes temporarily by a factor of 2, and the mass-loss rate by an order of magnitude, causing a density and velocity spike moving radially away from the star - considered as the origin of detached dust shells. Such events could severely disturb the pressure balance that dictates the well defined location of the bow shock when one of these `wind'-shells runs into the bow shock. There is likely to be a complex time dependence of the wind-ISM interaction on the timescales of the thermal pulses, as has been considered in 2D simulations by Villaver et al. (2003)." }, "0607/astro-ph0607670_arXiv.txt": { "abstract": "{% {} {We report on a high-spatial-resolution survey for binary stars in the periphery of the Orion Nebula Cluster, at 5--15~arcmin (0.65 -- 2\\,pc) from the cluster center.} {We observed 228 stars with adaptive optics systems, in order to find companions at separations of $0\\farcs13$ -- $1\\farcs12$ (60 -- 500\\,AU), and detected 13 new binaries. Combined with the results of Petr\\ (1998), we have a sample of 275 objects, about half of which have masses from the literature and high probabilities to be cluster members. We used an improved method to derive the completeness limits of the observations, which takes into account the elongated point spread function of stars at relatively large distances from the adaptive optics guide star.} {The multiplicity of stars with masses $>2\\,M_{\\sun}$ is found to be significantly larger than that of low-mass stars. The companion star frequency of low-mass stars is comparable to that of main-sequence M-dwarfs, less than half that of solar-type main-sequence stars, and 3.5 to 5 times lower than in the Taurus-Auriga and Scorpius-Centaurus star-forming regions. We find the binary frequency of low-mass stars in the periphery of the cluster to be the same or only slightly higher than for stars in the cluster core ($<$3~arcmin from $\\theta^1$C~Ori).} {This is in contrast to the prediction of the theory that the low binary frequency in the cluster is caused by the disruption of binaries due to dynamical interactions. There are two ways out of this dilemma: Either the initial binary frequency in the Orion Nebula Cluster was lower than in Taurus-Auriga, or the Orion Nebula Cluster was originally much denser and dynamically more active.} }% ", "introduction": "Over the past decade it has become clear that stellar multiplicity can be very high among young low-mass stars, with companion star frequencies close to 100\\,\\% for young stars in well-known nearby star-forming T~associations (Leinert et al.\\ \\cite{Leinert93}, Ghez et al.\\ \\cite{Ghez93}, Ghez et al.\\ \\cite{Ghez97}, Duch\\^ene \\cite{Duchene99a}). Thus, our current understanding is that star formation resulting in binary or multiple systems is very common, if not the rule. However, the multiplicity of low-mass main-sequence field stars is significantly lower% , only $\\sim 55\\,\\%$ for solar-type stars (Duquennoy \\& Mayor \\cite{DM91}), and $\\sim35$ to $42\\,\\%$ for M-dwarfs (Reid \\& Gizis \\cite{RG97}, Fischer \\& Marcy \\cite{FM92}). On the other hand, high binary frequencies are {\\em not\\/} observed among low-mass stars in stellar clusters. Binary surveys in the center of the young Trapezium Cluster (e.g. Prosser et al.\\ \\cite{Prosser94}, Padgett et al.\\ \\cite{Padgett97}, Petr et al.\\ \\cite{Petr98}, Petr \\cite{PetrPhD}, Simon et al.\\ \\cite{Simon99}, Scally et al.\\ \\cite{Scally99}, McCaughrean \\cite{MJM2001}), which is the core of the Orion Nebula Cluster (ONC); and in the young clusters IC\\,348 and NGC\\,2024 (Duch\\^ene et al.\\ \\cite{Duchene99b}, Beck et al.\\ \\cite{BeckSC03}, Liu et al.\\ \\cite{Liu2003}, Luhmann et al.\\ \\cite{Luhman2005}), as well as those in older ZAMS clusters (Bouvier et al.\\ \\cite{Bouvier97}, Patience et al.\\ \\cite{Patience98}) show binary frequencies that are comparable to that of main-sequence fields stars, i.e.\\ lower by factors of 2 -- 3 than those found in loose T associations. The reason for this discrepancy is still unclear. Theoretical explanations include: \\begin{itemize} \\item Disruption of cluster binaries through stellar encounters. Kroupa (\\cite{Kroupa95}) and Kroupa et al.\\ (\\cite{Kroupa99}) suggested that dynamical disruption of wide binaries (separations ${}>100\\rm\\,AU$) through close stellar encounters decreases the binary fraction in clusters. If the primordial binary output from the star-formation process is the same in dense clusters and in loose T~associations, then the number of ``surviving'' binaries depends on the number of interactions of a binary system with other cluster members that occurred since the formation of the cluster. This number is derived from the age of the cluster divided by the typical time between stellar interactions. The typical time between interactions is inversely proportional to the stellar volume density of the cluster (Scally et al.\\ \\cite{Scally99}), thus binaries at the cluster center get destroyed more quickly than those at larger radii. Observing various subregions of a single star-forming cluster representing different stellar number densities will therefore reveal different binary fractions if this mechanism is dominant in the evolution of binary systems. \\item Environmental influence on the initial binary fraction. Durisen \\& Sterzik (\\cite{DurSterz94}) and Sterzik et al.\\ (\\cite{Sterzik2003}) suggested an influence of the molecular cores' temperature on the efficiency of the fragmentation mechanism that leads to the formation of binaries. Lower binary fractions are predicted in warmer cores. Assuming the ONC stars formed from warmer cores than the members of the Taurus-Auriga association, less initial binaries are produced in this scenario. Observations of different subregions of the ONC should therefore reveal the same (low) binary frequency (if the molecular cores in these regions had the same temperature -- a reasonable, though unverifiable assumption). \\end{itemize} These theoretical concepts make different predictions that can be tested observationally. Indeed, measuring the binary fraction as a function of distance to the cluster center will provide important observational support for one or the other proposed theoretical explanation. \\begin{table}[htp] \\caption[]{Fields observed for this work. Name is the designation in Jones \\& Walker (\\cite{JW88}) or Parenago (\\cite{Parenago}) of the central star that was used to guide the adaptive optics system, $r$ is the distance to $\\theta^1$C~Ori. In the last columns, we list the number of stars actually observed in this field and the date(s) of observation.} \\label{ObsTab} \\setlength{\\tabcolsep}{3.5pt} \\begin{tabular}{lccrcc} \\noalign{\\vskip1pt\\hrule\\vskip1pt} Name & $\\alpha_{2000}$ & $\\delta_{2000}$ & $r$ [\\arcmin] & Targets & Obs.~Date\\\\ \\hline \\object{JW0005} & 5:34:29.243\t & -5:24:00.37 \t & 11.8 & 6 & 09.12.2001 \\\\ \\object{JW0014}\t& 5:34:30.371\t & -5:27:30.46\t & 12.2 & 2 & 11.12.2001 \\\\ \\object{JW0027}\t& 5:34:34.012\t & -5:28:27.72\t & 11.8 & 4 & 11.12.2001 \\\\ \\object{JW0045}\t& 5:34:39.774\t & -5:24:28.27\t & 9.2 & 4 & 10.12.2001 \\\\ \\object{JW0046}\t& 5:34:39.917\t & -5:26:44.70\t & 9.7 & 5 & 10.12.2001 \\\\ \\object{JW0050}\t& 5:34:40.831\t & -5:22:45.07\t & 8.9 & 1 & 11.12.2001 \\\\ \\object{JW0060}\t& 5:34:42.187\t & -5:12:21.55\t & 14.0 & 2 & 10.12.2001 \\\\ \\object{JW0064}\t& 5:34:43.496\t & -5:18:30.01\t & 9.6 & 3 & 17.02.2002 \\\\ \\object{JW0075}\t& 5:34:45.188\t & -5:25:06.33 \t & 8.0 & 7 & 09.12.2001 \\\\ \\object{JW0108}\t& 5:34:49.867\t & -5:18:46.79\t & 8.1 & 3 & 10.12.2001 \\\\ \\object{JW0116}\t& 5:34:50.691\t & -5:24:03.25 \t & 6.5 & 5 & 11.12.2001 \\\\ \\object{JW0129}\t& 5:34:52.347\t & -5:33:10.38\t & 11.5 & 4 & 10.12.2001 \\\\ \\object{JW0153}\t& 5:34:55.390\t & -5:30:23.42\t & 8.8 & 2 & 17.02.2002 \\\\ \\object{JW0157}\t& 5:34:55.936\t & -5:23:14.58\t & 5.1 & 4 & 17.02.2002 \\\\ \\object{JW0163}\t& 5:34:56.560\t & -5:11:34.84\t & 12.8 & 3 & 11.12.2001 \\\\ \\object{JW0165}\t& 5:34:56.601\t & -5:31:37.48\t & 9.6 & 2 & 11.12.2001 \\\\ \\object{JW0221}\t& 5:35:02.202\t & -5:15:49.28\t & 8.4 & 4 & 11.12.2001 \\\\ \\object{JW0232}\t& 5:35:03.090\t & -5:30:02.27 \t & 7.5 & 5 & 09.12.2001 \\\\ \\object{JW0260}\t& 5:35:04.999\t & -5:14:51.45\t & 9.0 & 3 & 11.12.2001 \\\\ \\object{JW0364}\t& 5:35:11.455\t & -5:16:58.33\t & 6.5 & 13 & 09.12.2001 \\\\ \\object{JW0421}\t& 5:35:13.706\t & -5:30:57.76\t & 7.6 & 5 & 11.12.2001 \\\\ &\t\t &\t\t &\t & & 17.02.2002 \\\\ \\object{JW0585}\t& 5:35:18.275\t & -5:16:37.83\t & 6.8 & 17 & 09.12.2001 \\\\ \\object{JW0666}\t& 5:35:20.936\t & -5:09:15.92\t & 14.2 & 8 & 09.12.2001 \\\\ \\object{JW0670}\t& 5:35:21.043\t & -5:12:12.52\t & 11.2 & 2 & 17.02.2002 \\\\ \\object{JW0747}\t& 5:35:23.929\t & -5:30:46.82\t & 7.6 & 5 & 10.12.2001 \\\\ \\object{JW0779}\t& 5:35:25.417\t & -5:09:48.98\t & 13.8 & 11 & 09.12.2001 \\\\ \\object{JW0790}\t& 5:35:25.953\t & -5:08:39.42\t & 14.9 & 5 & 10.12.2001 \\\\ \\object{JW0794}\t& 5:35:26.121\t & -5:15:11.27\t & 8.5 & 6 & 09.12.2001 \\\\ \\object{JW0803}\t& 5:35:26.565\t & -5:11:06.84 \t & 12.5 & 4 & 10.12.2001 \\\\ \\object{JW0804}\t& 5:35:26.666\t & -5:13:13.97\t & 10.5 & 1 & 17.02.2002 \\\\ \\object{JW0818}\t& 5:35:27.716\t & -5:35:19.01\t & 12.3 & 2 & 11.12.2001 \\\\ \\object{JW0866}\t& 5:35:31.077\t & -5:15:32.23\t & 8.7 & 4 & 09.12.2001 \\\\ \\object{JW0867}\t& 5:35:31.116\t & -5:18:55.12\t & 5.8 & 3 & 09.12.2001 \\\\ \\object{JW0873}\t& 5:35:31.521\t & -5:33:07.91 \t & 10.5 & 6 & 11.12.2001 \\\\ \\object{JW0876}\t& 5:35:31.627\t & -5:09:26.88\t & 14.4 & 3 & 11.12.2001 \\\\ \\object{JW0887}\t& 5:35:32.487\t & -5:31:10.05\t & 8.8 & 1 & 17.02.2002 \\\\ \\object{JW0915}\t& 5:35:35.509\t & -5:12:19.38\t & 12.0 & 2 & 11.12.2001 \\\\ \\object{JW0928}\t& 5:35:37.385\t & -5:26:38.20 \t & 6.2 & 4 & 10.12.2001 \\\\ \\object{JW0950}\t& 5:35:40.519\t & -5:27:00.48 \t & 7.0 & 4 & 11.12.2001 \\\\ \\object{JW0959}\t& 5:35:42.019\t & -5:28:10.95\t & 8.0 & 5 & 10.12.2001 \\\\ \\object{JW0963}\t& 5:35:42.528\t & -5:20:11.73\t & 7.3 & 5 & 09.12.2001 \\\\ \\object{JW0967}\t& 5:35:42.803\t & -5:13:43.69\t & 11.7 & 1 & 17.02.2002 \\\\ \\object{JW0971}\t& 5:35:43.476\t & -5:36:26.22\t & 14.7 & 4 & 10.12.2001 \\\\ \\object{JW0975}\t& 5:35:44.558\t & -5:32:11.40\t & 11.3 & 3 & 11.12.2001 \\\\ \\object{JW0992}\t& 5:35:46.845\t & -5:17:54.69\t & 9.4 & 6 & 09.12.2001 \\\\ \\object{JW0997}\t& 5:35:47.326\t & -5:16:56.01\t & 10.1 & 4 & 11.12.2001 \\\\ \\object{JW1015}\t& 5:35:50.509\t & -5:28:32.56\t & 10.0 & 2 & 10.12.2001 \\\\ \\object{JW1041}\t& 5:35:57.831\t & -5:12:52.17\t & 14.8 & 1 & 17.02.2002 \\\\ \\object{Par1605} & 5:34:47.201\t & -5:34:16.76\t & 13.1 & 3 & 10.12.2001 \\\\ % \\object{Par1744} & 5:35:04.822\t & -5:12:16.61\t & 11.5 & 4 & 11.12.2001 \\\\ % \\object{Par2074} & 5:35:31.223\t & -5:16:01.54 \t & 8.2 & 13 & 09.12.2001 \\\\ % &\t\t &\t\t &\t & & 17.02.2002 \\\\ \\object{Par2284} & 5:35:57.539\t & -5:22:28.21\t & 10.3 & 2 & 17.02.2002 \\\\ % \\noalign{\\vskip1pt\\hrule} \\end{tabular} \\end{table} \\begin{figure*}[t] \\centerline{\\includegraphics[angle=270,width=0.9\\hsize]{4561f01.ps}} \\caption{Fields observed in this work (large boxes) and Petr (\\cite{PetrPhD}, small boxes in the central region). Dots mark the positions of all stars in the list of Jones \\& Walker (\\cite{JW88}), these are {\\em not} the stars observed by us. The stars used to guide the adaptive optics are marked by star symbols and their number in Jones \\& Walker or Parenago (\\cite{Parenago}). The large boxes show the areas where we selected our target stars. These fields were not fully covered by our observations, we only observed stars visible in the 2MASS quicklook images (see Sect.~\\ref{ObsSect}).} \\label{ObsFieldsFig} \\end{figure*} The ONC is the best target for this study. Its stellar population is very well studied, more than 2000 members are known from extensive near-infrared and optical imaging (Hillenbrand \\cite{H97}, Hillenbrand \\& Carpenter 2000). To date, binary surveys of the cluster have focused on the central 0.25\\,pc core, where the stellar density reaches as high as $2$ -- $5\\times 10^4\\rm\\,pc^{-3}$ (McCaughrean \\& Stauffer \\cite{MJMStau94}, Hillenbrand \\& Hartmann \\cite{Hillenbr98}). The typical time between interactions for a binary with a separation of 250\\,AU ($\\sim0\\farcs5$ at the distance of the ONC) in the core is $\\sim 1\\rm\\,Myr$, the age of the cluster. Therefore, most 250\\,AU binaries are likely to have experienced at least one close encounter. However, for the observed stellar density distribution of the cluster, which can be described by an isothermal sphere with $n \\propto r^{-2}$ outside 0.06\\,pc (Bate et al.\\ \\cite{Bate98}), the volume density at $\\sim$\\,1\\,pc distance from the center ($\\sim8\\arcmin$ at the distance of the ONC) is roughly $200\\rm\\,pc^{-3}$ and the interaction timescale for our 250\\,AU binary would be $>250$\\,Myr, hundreds of times the age of the cluster. We also know that dynamical mass segregation in the cluster has not yet occurred (Bonnell \\& Davies \\cite{BonnellDavies98}) and that the ejection of single stars from the inner parts of the cluster (where many binary disruptions have already occurred) has not been efficient to populate the outer regions if the whole cluster is roughly virialized (Kroupa et al.\\ \\cite{Kroupa99}). For these reasons, the binary fraction in the outer parts of the ONC is unlikely to have been modified by the dynamical evolution of the cluster, and should be the intrinsic value resulting from the fragmentation process, while the binary frequency of stars in the cluster core has already been lowered by dynamical interactions. We have measured the frequency of close binaries among stars in a number of fields in the outer part of the ONC using adaptive optics imaging in the K-band. The results we obtain in this survey for the outskirts of the Orion Nebula Cluster will be compared with a similar study of the ONC core, carried out by Petr et al.\\ (\\cite{Petr98}) and Petr (\\cite{PetrPhD}). Since the same instrument and observing strategy was used, we incorporate their results, and constrain the radial distribution of the binary frequency from the ONC core to the cluster's periphery. ", "conclusions": " \\begin{itemize} \\item The binary frequency of low-mass stars in the periphery of the cluster is slightly higher than in the core, albeit with a low statistical significance of less than $1\\,\\sigma$. \\item The binary frequency of low-mass stars in the periphery of the ONC is lower than that of young stars in Taurus-Auriga, with a statistical significance on the $2\\,\\sigma$ level. \\item The binary frequency of stars with masses $>2\\,M_{\\sun}$ in the periphery is lower than in the center, but the difference is not statistically significant due to the small number of objects. \\end{itemize} These results do not support the hypothesis that the initial binary proportion in the ONC was as high as in Taurus-Auriga and was only later reduced to the value observed today. In that case, we would expect a much higher number of binaries in the periphery than observed. There are models that can explain the observations with a high initial binary frequency that was reduced by dynamical interactions, e.g.\\ a cluster that was much denser in the past, or a hierarchical formation with many dense subclusters. However, the simplest explanation with our current knowledge is that the initial binary frequency in the ONC was lower than in Taurus-Auriga. This suggests that the binary formation rate is influenced by environmental conditions, e.g.\\ the temperature of the parental molecular cloud." }, "0607/astro-ph0607393_arXiv.txt": { "abstract": "We present chemical abundances in K and M red-giant members of the Galactic bulge derived from high-resolution infrared spectra obtained with the Phoenix spectrograph on Gemini-South. The elements studied are carbon, nitrogen, oxygen, sodium, titanium, and iron. The evolution of C and N abundances in the studied red-giants show that their oxygen abundances represent the original values with which the stars were born. Oxygen is a superior element for probing the timescale of bulge chemical enrichment via [O/Fe] versus [Fe/H]. The [O/Fe]-[Fe/H] relation in the bulge does not follow the disk relation, with [O/Fe] values falling above those of the disk. Titanium also behaves similarly to oxygen with respect to iron. Based on these elevated values of [O/Fe] and [Ti/Fe] extending to large Fe abundances, it is suggested that the bulge underwent a more rapid chemical enrichment than the halo. In addition, there are declines in both [O/Fe] and [Ti/Fe] in those bulge targets with the largest Fe abundances, signifying another source affecting chemical evolution: perhaps Supernovae of Type Ia. Sodium abundances increase dramatically in the bulge with increasing metallicity, possibly reflecting the metallicity dependant yields from supernovae of Type II, although Na contamination from H-burning in intermediate mass stars cannot be ruled out. ", "introduction": "The idea that galaxies consist of distinct stellar populations goes back to the classic observations of M31, M32, and NGC 205 by Baade (1944). Within Baade's concept of populations, the Milky Way can be divided somewhat crudely into the halo, thick disk, thin disk, and bulge. Understanding the ages, kinematics, and chemical enrichment histories of the various populations of a galaxy provides insight into and constraints on models of galaxy formation and evolution. The elemental abundance distributions, in particular the abundance ratios of certain critical elements in stars are sensitive to global variables such as star formation histories and chemical enrichment timescales within the various galactic components. Within the Milky Way, there are numerous abundance studies of the disk (e.g., Edvardsson et al. 1993; Reddy et al. 2003), thick disk (e.g. Prochaska et al. 2000; Bensby et al. 2003), or halo (e.g., Fulbright 2002; Fulbright \\& Johnson 2003), all of which provide an increasingly detailed view of chemical evolution and evolutionary timescales in the Galaxy. Noticeably absent from the many lists of detailed abundance studies are corresponding abundance analyses of the bulge population. This is due to its distance ($\\sim$8 kpc) and interstellar absorption, both of which combine to render bulge stars, even the bright K and M giants, relatively faint for high-resolution spectroscopic analyses. The earliest determinations of abundances in bulge stars were derived from low-resolution spectra and consisted of overall metallicities as characterized by [Fe/H]. These initial abundance studies were by Rich \\& Whitford (1983), Rich (1988; 1990), Terndrup, Sadler \\& Rich (1995), or Sadler, Rich, \\& Terndrup (1996). The early heroic study by McWilliam \\& Rich (1994) was the pioneering effort to probe elemental abundance distributions and chemical evolution in the bulge population, but suffered from somewhat low spectral-resolution and signal-to-noise (S/N). The dearth of bulge high-spectral resolution abundance studies is now being remedied by prototype studies that have much higher-quality spectra at their disposal. The most recent works are the detailed study of stellar parameters and iron abundances in 27 bulge K-giants by Fulbright, McWilliam \\& Rich (2006), or the analysis of 14 M-giants by Rich \\& Origlia (2005). Fulbright et al. (2006), in their first paper of a planned series, focus on carefully deriving the fundamental stellar parameters of effective temperature, surface gravity, iron abundance, and microtubulent velocity. Their discussion then centers on comparing the iron abundances in their 27-star sample with the same stars in the larger, low-spectral resolution studies of Rich (1988) and Sadler et al. (1996). By transforming the Fe-abundances from the older (but with larger stellar samples) studies onto their Fe-abundance scale as set by their sample from high-quality, high-resolution spectra, Fulbright et al. derive an updated metallicity (i.e. [Fe/H]) distribution for the bulge. Rich \\& Origlia (2005) concentrate on bulge M-giants that have a rather restricted range in metallicity ([Fe/H] $\\sim$ -0.3 to 0.0) but the sample size in this range allows them to evaluate the scatter in abundance ratios. In this study, we present abundances in 7 bulge red-giant members; 5 K-giants and 2 M-giants. Our spectral types thus overlap with Fulbright et al. (2006) and Rich \\& Origlia (2005), as well as spanning a wide metallicity range. We are thus in a position to probe the chemical evolution of a small, but key set of elements in the bulge. ", "conclusions": "$\\bf Iron$: Our bulge sample, although admittedly small, spans a significant range across the metallicity distribution as found from low resolution studies of the bulge. The sampled metallicity range thus provides the opportunity to infer characteristics of bulge chemical evolution as defined by the abundances summarized below. {\\bf Carbon, Nitrogen and the CN cycle}: The $^{12}$C-depletion and $^{14}$N-enhancements are indicative of CN-cycled material. The conclusion of this investigation into the $^{12}$C and $^{14}$N abundances is that the $^{16}$O abundances in the studied bulge red giants are not measurably altered from their primordial values, as the mixing is not nearly extensive enough to have altered the initial $^{16}$O measurably. Therefore, the derived oxygen abundances represent good monitors of chemical evolution within the bulge population. $\\bf Oxygen$ : At low metallicity, the [O/Fe] results agree with halo-like [O/Fe] abundances. The [O/Fe] values then decline as [Fe/H] increases. The decline in [O/Fe] in the bulge, however, is not as large as for the thin and thick disk at the highest metallicities. The most straightforward explanation of this trend is that the bulge underwent more rapid metal enrichment than the halo, but that star formation did continue over timescales that may include the onset of SN Ia. $\\bf Sodium$: Near solar metallicity ([O/H] $\\sim$ 0.0), the bulge stars overlap the field-star distributions but the lowest metallicity bulge star in our sample (IV-003) falls well below the trend. Since the Na/O yields from SN II are metallicity sensitive, the low value of Na/O for IV-003 may indicate that the initial enrichment of the bulge to [O/H] $\\sim$ -1 was quite rapid and dominated by very metal-poor massive stars (with [m/H] less than -2). This would be an example in which a metallicity-dependant ratio (Na/O) retained the metallicity signature of its parent SN II. At the highest metallicities, the sodium abundances are very high, defining a sharp upward trend with the oxygen abundance. It is interesting to note that this pattern of high sodium was also found for the metal rich open cluster NGC 6791 by Peterson \\& Green (1998). $\\bf Titanium$: The Ti abundances obtained define a somewhat different trend than any of the other populations in the Milky Way. There is a tendancy of the bulge red giants to retain elevated [Ti/Fe] values past the downturn shown by the halo-disk field stars. In addition, the bulge stars are significantly more elevated relative to [Ti/Fe] versus [Fe/H] in dwarf spheroidals. The behavior of [Ti/Fe] as a function of [Fe/H] in the bulge is reminiscent of that of [O/Fe]. The suggestion would be that the bulge underwent a more rapid metallicity enrichment from SN II than the halo, but in the most metal-rich bulge stars there is a decline in [$\\alpha$/Fe] (found in our sample for both O and Ti)." }, "0607/astro-ph0607446_arXiv.txt": { "abstract": "We use numerical simulations to study the kinematic structure of remnants formed from mergers of equal-mass disk galaxies. In particular, we show that remnants of dissipational mergers, which include the radiative cooling of gas, star formation, feedback from supernovae, and the growth of supermassive black holes, are smaller, rounder, have, on average, a larger central velocity dispersion, and show significant rotation compared to remnants of dissipationless mergers. The increased rotation speed of dissipational remnants owes its origin to star formation that occurs in the central regions during the galaxy merger. We have further quantified the anisotropy, three-dimensional shape, minor axis rotation, and isophotal shape of each merger remnant, finding that dissipational remnants are more isotropic, closer to oblate, have the majority of their rotation along their major axis, and are more disky than dissipationless remnants. Individual remnants display a wide variety of kinematic properties. A large fraction of the dissipational remnants are oblate isotropic rotators. Many dissipational, and all of the dissipationless, are slowly rotating and anisotropic. The remnants of gas-rich major mergers can well-reproduce the observed distribution of projected ellipticities, rotation parameter ($V/\\sigma$)$^*$, kinematic misalignments, $\\Psi$, and isophotal shapes. The dissipationless remnants are a poor match to this data. We also investigate the properties of merger remnants as a function of initial disk gas fraction, orbital angular momentum, and the mass of the progenitor galaxies. Our results support the merger hypothesis for the origin of low-luminosity elliptical galaxies provided that the progenitor disks are sufficiently gas-rich, however our remnants are a poor match to the bright ellipticals that are slowly rotating and uniformly boxy. ", "introduction": "\\label{sec:intro} The observed absorption-line spectra and red colors of elliptical galaxies suggest that their stars were formed at high redshift ($z\\geq1$) and that very little star formation has occurred in them since then. According to the ``merger hypothesis'' \\citep{TT72,T77}, these red elliptical galaxies are produced by the collision and merger of spiral galaxies, and hence the progenitors of present day ellipticals may be high-redshift spirals. While relatively little is known about disk galaxies at high redshift, it is likely that these disks were more concentrated and gas-rich than their low-redshift counterparts. Indeed, preliminary observational evidence \\citep{Erb06m} indicates that galaxies at redshift $z\\approx 2$ do have large gas fractions $f_{\\rm gas}\\sim0.5$, with some approaching $f_{\\rm gas}\\sim0.8 - 0.9$. Thus, any attempt to understand the formation, properties and scaling relations of the present day population of elliptical galaxies, within the context of the ``merger hypothesis'', must consider gas-rich mergers. Requiring that the disk galaxy progenitors contain a significant fraction of gas is nothing new to the ``merger hypothesis''. One of the main objections to this mechanism of producing ellipticals, argued by, for instance, \\citet{Ost80}, is that the centers of ellipticals are more concentrated than local spirals. Cast in terms of phase-space density, this objection states that the high central phase-space density of ellipticals cannot be produced by the merger of low phase-space spirals because, according to Liouville's Theorem, phase-space density is conserved during a collisionless process \\citep{Car86}. However, this argument breaks down when the merger constituents contain gas, which can radiate energy, and hence increase the phase space density \\citep{Lak89}. An estimate of how much gas is required to match the central densities of ellipticals was provided by \\citet{HRemIII}, who used N-body simulations and analytic arguments to suggest that mergers of spiral galaxies containing $\\geq30$\\% gas would be sufficient to account for the high phase space densities of ellipticals. Within the context of the hierarchical theory of structure formation, gas-rich major mergers may play a much larger role than just resolving the central phase-space densities of elliptical galaxies. Previously, we have described a ``cosmic cycle'' of galaxy formation and evolution in which gas-rich mergers drive the evolution of quasars \\citep{Hop06big}, induce the growth of supermassive black holes \\citep{dMSH05}, and produce red elliptical galaxies \\citep{SdMH05red,Hop06red} that obey many of the observed scaling relations \\citep{Rob06fp,Rob06b}. A schematic view of this picture is presented in Fig. 1 of \\citet{Hop06big}. However, the success of this scenario also must be gauged by its ability to produce remnants that have kinematic and morphological properties characteristic of observed ellipticals. It is this question that we address in the current paper. Observations indicate that galaxy spheroids can be classified into two groups \\citep[][and references therein]{Dav83,Ben89,Ben88, Fab97,KB96}. Large, luminous spheroids have hot gaseous halos, box-shaped isophotes, surface-brightness profiles with flat ``cores,'' show very little rotation and are almost uniformly classified as ellipticals. Less luminous spheroids tend to have little, if any, hot gas, disk-shaped isophotes, power-law surface-brightness profiles, and exhibit rotation that is along the photometric major axis. The latter group encompasses many low-luminosity ellipticals, bulges and S0s. The last of these properties, the alignment and rotation of spheroidal galaxies, may be the result of a more fundamental dichotomy among elliptical galaxies; that is, the isotropy of their velocity distributions. Because spheroids can be flattened for (at least) two reasons, rotation and velocity anisotropy, the lack of rotation in large ellipticals suggests that these systems have significant velocity anisotropy, while many low-luminosity ellipticals and bulges are consistent with being isotropic systems flattened by their observed rotation. One viewpoint is that these two classes of elliptical galaxies exhibit different properties because they are formed via different mechanisms. Along this line of reasoning, \\citet{NB03} argue that large ellipticals are formed by the dissipationless merger of two comparable-mass disk galaxies, i.e., major mergers, while low-luminosity ellipticals are produced by the dissipationless merger of unequal mass disks, i.e., minor mergers. To demonstrate this possibility, \\citet{NB03} used numerical simulations of dissipationless disk galaxy mergers to show that remnants of major mergers rotate very little and are, in general, boxy, similar to luminous ellipticals. On the other hand, the simulated minor mergers rotate significantly and have disky isophotes, similar to low-luminosity ellipticals. However, these simulations did not include gas physics, and hence would not induce starbursts, quasar activity, nor satisfy the scaling relations of elliptical galaxies \\citep{Rob06fp}. Moreover, if most ellipticals are relatively old, i.e. have mainly old stellar populations, and were formed long ago by mergers, it is likely that the progenitor galaxies would have been gas-rich and had a higher gas fraction than local large, star-forming galaxies. In this paper we use a large suite of numerical simulations to explore the kinematic properties of remnants produced by the merger of comparable-mass disk galaxies. We specifically address the differences between remnants formed via dissipationless mergers versus those produced in gas-rich mergers that include the cooling of gas, star formation and feedback. We find that gas-rich mergers can successfully reproduce many of the kinematic properties of observed elliptical galaxies, while dissipationless remnants provide a poor match the data. The organization of the rest of this paper is as follows. In \\S~\\ref{sec:meth} we describe the numerical simulations, including the disk galaxy models and the galaxy collisions (\\ref{ssec:sims}), followed by the techniques employed to analyze individual merger remnants (\\ref{ssec:anal}). In \\S~\\ref{sec:results} we present the results of our analysis for the entire series of merger remnants. To begin, we report the aggregate of measured properties for our entire series of simulated merger remnants (\\ref{ssec:basics}). Following this, we show the remnant rotational support (\\ref{ssec:majrot}), including where this rotation originates (\\ref{ssec:origin}), and what individual remnants are like (\\ref{ssec:ind}). To better characterize the remnants we also analyze their anisotropy (\\ref{ssec:anisot}), shape (\\ref{ssec:shapes}), minor-axis kinematics (\\ref{ssec:minrot}), and how these quantities depend on our various input assumptions (\\ref{ssec:dep}). Finally, in \\S~\\ref{sec:disc} we discuss the implications of our results for the formation of elliptical galaxies and we conclude in \\S~\\ref{sec:conclusions}. ", "conclusions": "\\label{sec:conclusions} In this paper we have analyzed a series of numerical simulations to investigate the kinematic structure of galaxies formed from the collision of equal-mass disk galaxies. In particular, we determine that remnants produced by the collision of gas-rich disk galaxies are smaller, rounder, have higher central velocity dispersions (on average) and larger rotation speeds than the remnants of dissipationless disk galaxy mergers. The larger rotation present in gas-rich merger remnants owes its origin to dissipation. Stars formed during the merger rotate faster than stars present before the merger began. Dissipation and star formation also produce remnants that are closer to oblate and are uniformly more isotropic than dissipationless remnants. Slightly more than one-quarter of the dissipational merger remnants are consistent with being oblate isotropic rotators. Compared to observed ellipticals, the remnants formed from dissipational mergers are a significantly better match to the distribution of ellipticities, ($V/\\sigma$)$^*$, and kinematic misalignment angles $\\Psi$ of elliptical galaxies than dissipationless merger remnants. In particular, dissipationless remnants demonstrate significant minor axis rotation and appear to be flattened much more than observed ellipticals. We also calculate the isophotal shapes of the simulated merger remnants. Dissipational remnants tend to be disky while dissipationless remnants tend to be boxy. Observed ellipticals are evenly distributed between disky and boxy. While both remnants appear to be a sufficient match to the observed distribution, when comparing the correlation between disky/boxy and rotation, the dissipationless merger produce a significant number of slowly rotating disky remnants where there are no observed analogs. Both dissipational and dissipationless mergers produce remnants that are both disky and boxy, and thus these mechanisms have difficulty reproducing the luminous, slowly rotating ellipticals that are observed to be uniformly boxy. In general, our results suggest that dissipationless disk galaxy mergers {\\it cannot} be the dominant mechanism to form elliptical galaxies. Dissipational mergers, on the other hand, appear to be a viable mechanism to produce elliptical galaxies, specifically oblate isotropic rotators (i.e., low-luminosity ellipticals), and thus our results lend support to the ``merger hypothesis'' provided that the progenitor disk galaxies are gas-rich. As mentioned in \\S~\\ref{sec:intro}, additional evidence for the dissipative origin of ellipticals comes from their high phase space density compared to spiral galaxies. Gas dissipation provides a natural mechanism to increase the phase space density during the merger and also appears to be necessary for reproducing the scaling relation of elliptical galaxies \\citep{Rob06fp}. Our modeling suggests several avenues for further testing this hypothesis. Mergers between gas-rich spirals will imprint subtle features into the remnants. The central starburst will modify the inner profiles of the remnants \\citep{MH94dsc}, perhaps accounting central light excesses seen in merging systems \\citep{RJ04,RJ06}. In principle, this can be tested by comparing predictions for metallicity and color gradients with observations \\citep{MH94popg}. Dissipational merger may also provide a natural mechanism to produce kinematic subsystems in elliptical galaxies \\citep{HB91, BB00}. The shells, ripples, loops and other fine structures seen around many relaxed ellipticals \\citep{Schweizer98} are a natural consequence of mergers involving disk galaxies \\citep{HS92}, but that do not form in major mergers between hot stellar systems. Determining the ubiquity of fine structure in red galaxies would further constrain the importance of disk mergers to the formation of ellipticals. Placed within the ``cosmic cycle'' of galaxy formation, we can now argue that gas-rich major mergers trigger quasars and starbursts, fuel the growth of supermassive black holes, and produce remnant galaxies which have the colors, scaling relations, and kinematics akin to present day ellipticals." }, "0607/nucl-th0607034_arXiv.txt": { "abstract": "The $^1S_0$ pairing in neutron matter has been investigated in presence of realistic two-- and three--nucleon interactions. We have adopted the Argonne $v_{8^\\prime}$ NN and the Urbana IX 3N potentials. Quantum Monte Carlo theory, specifically the Auxiliary Field Diffusion Monte Carlo method, and Correlated Basis Function theory are employed in order to get quantitative and reliable estimates of the gap. They both fully take into account the medium modifications due to the interaction induced correlations. The two methods are in good agreement up to the maximum gap density and both point to a slight reduction with respect to the standard BCS value. In fact, the maximum gap is about $2.5~\\text{MeV}$ at $k_F \\sim\\, 0.8~\\text{fm}^{-1}$ in BCS and $2.3$--$2.4~\\text{MeV}$ at $k_F \\sim\\, 0.6~\\text{fm}^{-1}$ in correlated matter. At higher densities the Quantum Monte Carlo gap becomes close to BCS. In general, the computed medium polarization effects are much smaller than those previously estimated within \\emph{all theories}. Truncations of Argonne $v_{8^\\prime}$ to simpler forms give the same gaps in BCS, provided the truncated potentials have been refitted to the same NN data set. Differences among the models appear in the correlated theories, most of the reduction being attributable to the tensor force. The three--nucleon interaction provides an additional increase of the gap of about 0.35 MeV. ", "introduction": " ", "conclusions": "" }, "0607/astro-ph0607320_arXiv.txt": { "abstract": "We present timing measurements, astrometry, and high-resolution spectra of a number of nearby, thermally emitting, isolated neutron stars. We use these to infer magnetic field strengths and distances, but also encounter a number of puzzles. We discuss three specific ones in detail: (i) For RX J0720.4$-$3125 and RX J1308.6+2127, the characteristic ages are in excess of 1~Myr, while their temperatures and kinematic ages indicate that they are much younger; (ii) For RX J1856.5$-$3754, the brightness temperature for the optical emission is in excess of that measured at X-ray wavelengths for reasonable neutron-star radii; (iii) For RX J0720.4$-$3125, the spectrum changed from an initially featureless state to one with an absorption feature, yet there was only a relatively small change in $T_{\\rm eff}$. Furthermore, we attempt to see whether the spectra of all seven sourced, in six of which absorption features have now been found, can be understood in the context of strongly magnetised hydrogen atmospheres. We find that the energies of the absorption features can be reproduced, but that the featureless spectra of some sources, especially the Wien-like high-energy tails, remain puzzling. ", "introduction": "\\label{intro} One of the great benefits of the {\\em ROSAT} All-Sky Survey \\citep{rbs} is that is has provides an unbiased sample of all classes of nearby neutron stars (limited only by their age and distribution of the local interstellar medium). Particularly interesting is the discovery of the group of seven nearby, thermally emitting, isolated neutron stars (INS; for a review, see Haberl, these proceedings). The INS form the majority among the nearby neutron stars (typical distances are less than $\\sim\\!500$~pc; \\citealt{kvka02}; see also Posselt, Popov, these proceedings), yet are atypical of the neutron-star population represented by radio surveys: while pulsars detected by their thermal emission all have normal periods of less than a second, five out of the seven INS have periods about ten times longer (the remaining two appear to have no pulsations despite intensive searches; \\citealt{rgs02,vkkd+04}). A number of models --- accretors \\citep{w97}, middle-aged magnetars \\citep{hk98,hh99}, long-period pulsars \\citep{kkvkm02,zhc+02} --- have been suggested to explain these objects. A prime reason for studying the INS is the hope of constraining fundamental physics at very high densities: neutron stars are natural laboratories for quantum chromodynamics \\citep{rho00}. The overall goal is to determine the masses and radii of a number of neutron stars and hence constrain the equation of state (EOS) of ultra-dense matter (\\citealt{lp00}; Lattimer, these proceedings). For the majority of known neutron stars (i.e., radio pulsars), this is complicated by the non-thermal emission that dominates the spectrum, but for the INS this is not the case: the X-ray spectra show thermal emission only. Hence, much effort has been spent trying to derive constraints from the INS \\citep{bzn+01,bhn+03,dmd+02,pwl+02}. The constraints have not been very meaningful, however, because the data could not be interpreted properly: they just do not fit any current realistic models \\citep{mzh03,ztd04}. To make progress in understanding the thermal emission, we need first to know the basic ingredients: the elemental abundances, the temperature distribution, and the magnetic field strength. Furthermore, to use the thermal emission to infer radii, we need information about the distance. Fortunately, observational clues are now becoming available: broad absorption features at energies of 0.3--0.7~keV have been discovered in the spectra of six of the seven INS (\\citealt{hsh+03,hztb04,hmz+04,vkkd+04,zct+05}; see Haberl, these proceedings), and, as described below, magnetic field strengths have been inferred from timing solutions and new or improved parallaxes have been measured. The outline of this contribution is as follows. First, in \\S\\ref{sec:timing}, we present timing solutions for RX J0720.4$-$3125 and RX J1308.6+2127, and discuss the resulting estimates of the magnetic field strengths and characteristic ages. Next, in \\S\\ref{sec:d}, we describe new parallax distance measurements for RX J1856.5$-$3754 and RX J0720.4$-$3125. For the former, these resolve previous conflicting results, but also raise a puzzle: a rather large radius or high brightness temperature inferred for the optical emission. In \\S\\ref{sec:spectra}, we turn to high-resolution X-ray spectra, comparing spectra of RX J0720.4$-$3125, before and after its spectral change, with those of RX J1308.6+2127. In \\S\\ref{sec:atmosphere}, we attempt to interpret the observations assuming the sources have gaseous atmospheres, focussing on hydrogen, but also briefly discussing the possibility of helium. We summarise and discuss future work in \\S\\ref{sec:discussion}. From here on, we will refer to the INS in the text using abbreviated names: J0420 for RX J0420.0$-$5022, J0720 for RX J0720.4$-$3125, J0806 for RX J0806.4$-$4123, J1308 for RX J1308.6+2127 = RBS 1223, J1605 for RX J1605.3+3249, J1856 for RX J1856.5$-$3754, J2143 for RX J2143.0+0654 = RBS 1774 = RXS J214303.7+065419. \\begin{figure} \\begin{center} \\includegraphics[width=\\hsize]{ppdot.eps} \\caption{$P$-$\\dot P$ diagram, showing radio pulsars (points) and magnetars (diamonds); selected objects are labeled. Also shown are the five INS with periodicities: RX J0720.4$-$3125 and RX J1308.6+2127 are shown by the stars, while RX J0420.0$-$5022, RX J0806.4$-$4123, and RX J2143.0+0654 are the arrows at the top (since $\\dot P$ is unknown). The diagonal lines show loci of constant dipole magnetic field and spin-down age, as labeled.} \\label{fig:ppdot} \\end{center} \\end{figure} ", "conclusions": "\\label{sec:discussion} Of the four main parameters mentioned in the introduction that determine the properties of the thermal emission from INS, we now appear to have reasonable handles on three: the shapes of the X-ray spectra indicate temperatures around $10^6~$K, period derivatives imply magnetic field strengths of a few $10^{13}~$G, and parallax measurements show that a fair fraction of the surface is emitting X-ray radiation. The main unknown appears to be the composition. We found that the energies of the observed absorption features can be matched fairly easily for hydrogen atmospheres. However, reproducing the smooth, featureless spectra of some INS, and the Wien-like high-energy side of the X-ray spectra in general, appears problematic, nor is it clear how the spectrum of J0720 could change from featureless to one that has an absorption line. Fortunately, it should soon become clear whether these issues are real problems or not, since great progress is being made in constructing more reliable strongly magnetised hydrogen model atmospheres (Lai, Potekhin, these proceedings). From \\S\\ref{sec:atmosphere}, it seems particularly important to include in full detail transitions to the auto-ionising levels, verify that all sources of opacity, including from (traces of) molecules are included, and check the influence, in particular on the temperature profile, of high-density effects and vacuum resonance mode conversion. At the same time, it would seem worthwhile to consider atmospheres of other elements; for the INS, He might be most relevant, but it would be good to check heavier elements as well, since these may cause the absorption features seen in 1E 1207.4$-$5209 (\\citealt{hm02}). From the observational side, the easiest route to further progress would appear to be timing. With further estimates of the magnetic fields, one can test the predictions based on hydrogen atmospheres, that J0420 has a field about as strong as that of J0720 and J1308, J0806 a stronger one, approaching $10^{14}~$G, and J2143 the strongest, in excess of $10^{14}~$G. For the X-ray spectra, further monitoring is useful, but perhaps the largest advance will come from the unified analysis of all sources, which allows one to exclude instrumental effects. This is already well underway for the EPIC-PN data (Haberl, these proceedings), and similar studies of the LETG and RGS data should prove fruitful. As present, first steps are being taken in detailed modelling of the phase-resolved spectra (Haberl, Zane, these proceedings), and this should help obtain stronger constraints on the thermal distribution over the surface. Finally, in the optical-ultraviolet regime, it would be good to complete the census of the sources, and obtain at least rough spectral energy distributions, to determine whether the emission is thermal, or whether there are non-thermal components. For sources that are sufficiently bright, proper motion measurements can help determine true ages and parallax measurements can help determine distances." }, "0607/astro-ph0607116_arXiv.txt": { "abstract": "In this letter we present ground-based subarcsecond mid-infrared imaging and spectroscopy of young super star clusters in the overlap region of the merging galaxies NGC4038/4039 (the Antennae) obtained with the VLT Imager and Spectrometer for mid-Infrared (VISIR). With its unprecedented spatial resolution VISIR begins to resolve the H\\,{\\sc ii}/PDR complexes around the star-forming regions for the first time. In the N-band spectra of two young star clusters unexpectedly low polycyclic aromatic hydrocarbon (PAH) emission is observed, compared to what is seen with the Infrared Space Observatory (ISO) and with the Spitzer Space Telescope. We conclude that a large fraction of the PAH emission cannot directly be associated with the super star clusters, but originate from an extended region of at least 215 pc radius around the clusters. In the distribution of [Ne\\,{\\sc ii}] 12.81 \\um ~emission a highly obscured cluster is discovered that does not have an optical or near-infrared counterpart. ", "introduction": "Starburst galaxies experience a phase of rapid evolution. The rate at which their gas reservoir is turned into stars cannot be maintained for long, which makes the starburst phase by definition a transient one. The resulting stellar clusters make starbursts unique laboratories for the study of star formation, stellar populations and the evolution of galaxies as a whole. Since the earliest stages of massive star formation are generally heavily enshrouded by dust, infrared observations are essential to detect the youngest stellar populations and provide various diagnostic features to study the properties of the interstellar matter (ISM) and the underlying stellar population. Recently, a new generation of ground- and space-based instruments working at mid-infrared wavelengths has become operational, giving this field of research a large impulse. The Antennae (NGC4038/4039, Arp244) is the nearest major merger of two large spirals \\citep{Toomre:1972}. Since the beginning of the interaction the system went through several episodes of violent star formation of which the last one is probably still ongoing \\citep{Vigroux:1996}. The resulting (super) star clusters have been studied extensively throughout the electromagnetic spectrum \\citep{Whitmore:2005, Wang:2004, Gilbert:2000}. Radio and mid-infrared observations showed that the region between the two spirals (the overlap region) hosts spectacular obscured star formation. The brightest mid-infrared component produces 15\\% of the total 15 \\um ~luminosity of the entire system \\citep{Mirabel:1998, Hummel:1986}. This region is covered by a prominent dustlane, and can be associated with a faint, red source in the HST images \\citep[number 80 of][hereafter WS95]{Whitmore:1995}, illustrating how optical data alone are insufficient to identify and study the youngest star-forming regions. An ISOCAM spectrum of this region shows strong fine-structure emission lines ([Ne\\,{\\sc ii}] and [Ne\\,{\\sc iii}]) and pronounced emission from polycyclic aromatic hydrocarbons (PAHs) \\citep{Mirabel:1998}. We use the recently commissioned VLT Imager and Spectrometer for mid-Infrared (VISIR) \\citep{Lagage:2004} at the Very Large Telescope (VLT) of the European Southern Observatory (ESO) to study the most luminous super star clusters and the surrounding matter in the Antennae overlap region in detail. ", "conclusions": "" }, "0607/astro-ph0607266_arXiv.txt": { "abstract": "We describe briefly the properties of the recently completed Southern African Large Telescope (SALT), along with its first light imager SALTICAM. Using this instrument, we present 4.3 hr of high speed unfiltered photometric observations of the eclipsing polar SDSSJ015543.40+002807.2 with time resolution as short as 112 ms, the highest quality observations of this kind of any polar to date. The system was observed during its high luminosity state. Two accreting poles are clearly seen in the eclipse light curve. The binary system parameters have been constrained: the white dwarf mass is at the low end of the range expected for cataclysmic variables. Correlations between the positions of the accretion regions on or near the surface of the white dwarf and the binary system parameters were established. The sizes of the accretion regions and their relative movement from eclipse to eclipse were estimated: they are typically $4^{\\rm o}$ -7$^{\\rm o}$ depending on the mass of the white dwarf. The potential of these observations will only fully be realised when low state data of the same kind are obtained and the contact phases of the eclipse of the white dwarf are measured. ", "introduction": "\\label{salt} The Southern African Large Telescope, colloquially known by its acronym SALT, has recently been completed. It is now in the final stages of commissioning and first stages of science operations. This paper describes the properties of the telescope and its imaging camera SALTICAM. It then presents the first science observations of an eclipsing magnetic cataclysmic variable, SDSS J015543.40+002807.2, analyses the data and shows how our knowledge of this star has been advanced. Although SALT has been described in the proceedings of SPIE conferences (e.g. Stobie et al. 2000; Meiring et al. 2003; Meiring \\& Buckley 2004), these publications are not available in all astronomical libraries and not accessible on the Internet. Thus, a brief description of SALT is appropriate. Its basic principle of operation is similar to that of the Arecibo radio telescope, employing a spherical primary mirror and a payload carrying the instrumentation which tracks the spherically-shaped focal surface. This surface is concentric with the primary mirror and located at a distance of 13.08 m (half the radius of curvature) from the primary mirror. The overall optical design of SALT is described in Swat et al. (2003). The primary mirror is tilted so that its axis of symmetry is fixed at an inclination of 37$^{\\rm o}$ to the zenith. The telescope can move in azimuth, but not in elevation so it can observe any celestial object which reaches an altitude of 53$^{\\rm o}$. The tracking capability described below extends this range of altitudes from 47$^{\\rm o}$ to 59$^{\\rm o}$. The basic concept for this telescope was proposed by Ramsey \\& Weedman (1984) and its first implementation was the construction of the Hobby-Eberly Telescope at McDonald Observatory (HET). SALT is the second such telescope and started off its development process using the design of the HET. However, many aspects of this design were improved upon. \\subsection{Primary mirror} The primary mirror is comprised of 91 hexagonal spherical segments 1 m across, 50 mm thick, and made of Astro-Sitall (Swiegers \\& Gajjar 2004). The arrangement of the primary segments is also hexagonal such that a circle of diameter 11 m would pass through the vertices of the hexagon. The individual segments are shaped to a global radius of curvature of the primary of $26.165\\pm0.0005$ m, and figured to better than 1/15th wave rms. Each segment is supported on a mirror mount with nine points of support; three actuators move the segment in tip/tilt and piston. The mirror mounts are ``plugged\" in to a steel space-frame truss which itself is kinematically supported on the telescope structure. Alignment of the mirror segments with respect to each other to an accuracy of 0.06\" rms takes place in the evening twilight and is achieved using a Shack-Hartmann wavefront camera, located at the centre of curvature of the primary mirror in a tower alongside the telescope building. The alignment is maintained using a system of capacitive edge sensors which are bonded to the six sides of each mirror segment, generating tiny but detectable changes in capacitance as the segments move with respect to each other during the night. These signals are fed to a closed-loop control system which adjusts the segment positioning every 20 s to maintain the original alignment of the mirror. \\subsection{Tracker and Payload} The payload at the prime focus is positioned using a ``Tracker\" which is capable of movement in x, y, z, tip, tilt and rotation with accuracies of $\\sim5$ microns and 1 arcsec (despite carrying a weight of more than a metric ton!). Tracking celestial objects requires adjustment of x and y, along with concomitant changes in tip and tilt. The focus of the telescope is controlled by movement in z. The payload is also rotated about the telescope optical axis during observation to remove field rotation at the focal plane. The payload includes a 4-mirror spherical aberration corrector (SAC) (O'Donoghue 2000) to correct the huge spherical aberration of the primary (the circle of least confusion of the uncorrected prime focus is about 20 arcmin in diameter) and feed the corrected prime focus with an f/4.2 beam. The SAC yields a circular, flat, science field of view of 8 arcmin in diameter, with a 1 arcmin annulus around it for guide stars used in closed loop guiding. The entrance pupil of the SAC is 11 m in diameter and can thus accept light from celestial targets from the entire primary mirror. However, as the Tracker tracks celestial objects, the alignment of the SAC entrance pupil must necessarily be offset with respect to the primary mirror, so that parts of the primary mirror outside the entrance pupil do not contribute to light reaching the focal plane. Fortunately, the exit pupil of the optical system is readily accessible and a moving hexagonal baffle at this position is kept aligned with the primary mirror, thus preventing stray light from the periphery of the primary mirror reaching the focal plane. The telescope is equipped with an atmospheric dispersion compensator (O'Donoghue 2002) to enable access to wavelengths as short as 320 nm without image quality degradation arising from atmospheric dispersion. As mentioned, the axis of symmetry of the primary mirror is tilted with respect to the zenith by 37$^{\\rm o}$, so that the accessible declination range is +10.5$^{\\rm o}$ to -75.3$^{\\rm o}$. Celestial targets are accessible to the telescope when they enter an annular region which is centred on the zenith and has an angular distance (with respect to the zenith) of 37$^{\\rm o}$. The range of the Tracker is $\\pm6^{\\rm o}$ so that a celestial target is accessible when its zenith distance lies in the range 31$^{\\rm o}$ to 43$^{\\rm o}$. These constraints permit objects to be tracked for a duration of 1 hr at the northern end of the accessible declination range, and more than 3 hr at the southern end. The mode of operation of the telescope is straightforward: the telescope is slewed to the appropriate azimuth for a given celestial target. The primary mirror is stationary at this azimuth during the observation. When the desired celestial target enters the annular region of accessibility, it is tracked using the Tracker. Most celestial targets are accessible twice in a given night: once in the eastern sky and once in the western sky. At the extreme north and south ends of the declination range, targets can be re-acquired by re-positioning of the telescope in azimuth, thus permitting more than two pointings. ", "conclusions": "SALT has been completed and first science observations of the polar SDSS J015543+002807 have been obtained. These are high speed unfiltered photometric observations with time resolution as short as 112 ms and are the highest quality observations of this kind of any polar to date. The target was observed in its high luminosity state and shows rapid eclipses of two compact regions comprising unequivocal evidence for two accretion regions on or near the surface of the white dwarf. These account for all the optical light emitted by the system except for 1.5 per cent at mid eclipse and $\\sim3-4$ per cent from the photosphere of the white dwarf (as distinct from the spots). The binary system parameters, $q$ and $i$ have been constrained. Further limits, if M$_2 \\sim 0.07$ M$_\\odot$, are $0.5 < M_{wd} < 0.7\\ M_\\odot$, $q < 0.14$; or, if M$_2 \\sim 0.10$ M$_\\odot$, $0.5 < M_{wd} < 0.8\\ M_\\odot$, $q < 0.20$. The correlation between $i$, $R_{wd}$ and the co-latitudes of the accretion regions has been explored, and the relative sizes and movements from eclipse to eclipse of the spots have been determined. In the absence of the detection of the white dwarf, there are few secure constraints on the co-latitudes of the spots. The symmetry of the eclipses of each spot suggest that they are located within a small range in longitude. The dominance of the accretion regions in the luminosity budget of SDSS J015543+002807, combined with its eclipses, offers unusual promise in exploring the details of the accretion regions in a polar. Further observations involving time-resolved spectrophotometry, polarimetry and spectropolarimetry, all of which are available on SALT, are likely to uncover even more information. This potential can be fully exploited only when high time resolution observations during the low luminosity state of the system are obtained. It is likely that the contact points of the white dwarf will then be visible and this information will allow the position of the accretion regions with respect to the white dwarf to be fixed." }, "0607/astro-ph0607099_arXiv.txt": { "abstract": "An \\xmm\\ observation of the bright QSO \\pg\\ was previously reported to show evidence for a massive, energetic outflow, with an outflow velocity of v$\\sim$0.1c based on the identification of blue-shifted absorption lines detected in both EPIC and RGS spectra. Subsequently, an order-of-magnitude lower velocity has been claimed from an ion-by-ion model fit to the RGS data. We show here, in a re-analysis of the higher signal-to-noise EPIC data, that the high velocity is confirmed, with the resolution of additional absorption lines yielding a revised outflow velocity in the range $\\sim$0.13-0.15c. Confirmation of a massive and energetic outflow in a non-BAL AGN has important implications for metal enrichment of the IGM and for the feedback mechanism implied by the correlation of black hole and galactic bulge masses. We note the near-Eddington luminosity of \\pg\\ may be the critical factor in driving such an energetic outflow, a situation likely to be common in AGN at higher redshift. ", "introduction": "An analysis of the EPIC and RGS spectra from an \\xmm\\ observation of the bright narrow emission line QSO \\pg\\ in 2001 provided evidence for a highly ionised outflow with a velocity of $\\sim$0.1c (Pounds \\et\\ 2003; hereafter P03), though a lower velocity has recently been claimed from a separate analysis, principally based on the low signal-to-noise RGS data (Kaspi and Behar 2006). Confirmation of the high velocity outflow is important since the mechanical energy in the flow, if not highly collimated, is a significant fraction of the bolometric luminosity of \\pg\\ and could be typical of AGN accreting near the Eddington rate (King and Pounds 2003), while also providing an example of the feedback required by the linked growth of SMBH in AGN with their host galaxy (King 2005). Subsequently, the same \\xmm\\ observation of \\pg\\ has been used by Gierlinski and Done (2004) to suggest how strong absorption of the intrinsic X-ray continuum in a `velocity-smeared', high column, of moderately ionised gas can provide a physically preferred explanation (to Comptonisation) for the strong soft excess widely seen in type 1 AGN (Wilkes and Elvis 1987, Turner and Pounds 1989). A similar study by Chevallier \\et\\ (2006), which also considered an ionised reflection origin of the soft excess, concluded that absorption was the more likely cause of a strong soft excess (as in \\pg). In the present paper we re-examine the \\xmm\\ EPIC data of \\pg\\ which formed the strongest evidence for the high velocity claimed in P03. We use the latest calibration files and - in particular - take advantage of the higher spectral energy resolution of the MOS cameras demonstrated in recent studies of the Type 2 Seyferts Mkn3 (Pounds and Page 2005) and NGC1068 (Pounds and Vaughan 2006). We assume a redshift for \\pg\\ of $z=0.0809$ (Marziani \\et\\ 1996). ", "conclusions": "(1) A previous analysis of the 2001 \\xmm\\ observation of the bright quasar \\pg\\ reported evidence of a high velocity ionised outflow, with a mass and kinetic energy comparable to the accretion mass and bolometric luminosity, respectively (P03). (2) This finding is now confirmed, with the previous uncertainty in the derived velocity removed by securing the identification of the main observed absorption lines. (4) We suggest that fast, energetic outflows may be a typical signature of type 1 AGN accreting at or close to the Eddington limit." }, "0607/astro-ph0607050_arXiv.txt": { "abstract": "{} {We quantify the effect of gravitational redshift on emission lines to explore the transition region from the Newtonian to the Einsteinian regime. With the emitting region closer to the Kerr black hole, lines are successively subjected to a stronger gravitationally induced shift and distortion. Simulated lines are compared to broad, optical emission lines observed in Mrk~110.} {We simulate relativistic emission line profiles by using Kerr ray tracing techniques. Emitting regions are assumed to be thin equatorial rings in stationary Keplerian rotation. The emission lines are characterised by a generalized Doppler factor or redshift associated with the line core.} {With decreasing distance from the black hole, the gravitational redshift starts to smoothly deviate from the Newtonian Doppler factor: Shifts of the line cores reveal an effect at levels of 0.0015 to 60~\\% at gravitational radii ranging from $10^{5}$ to 2. This corresponds to fully relativistic Doppler factors of 0.999985 to 0.4048. The intrinsic line shape distortion by strong gravity i.e.\\ very asymmetric lines occur at radii smaller than roughly ten gravitational radii.} {Due to the asymptotical flatness of black hole space--time, GR effects are ubiquitous and their onset can be tested observationally with sufficient spectral resolution. With a resolving power of $\\sim100000$, yielding a resolution of $\\approx0.1$ {\\AA} for optical and near--infrared broad emission lines like H$\\beta$, HeII and Pa$\\alpha$, the gravitational redshift can be probed out to approximately 75000 gravitational radii. In general, gravitational redshift is an important indicator of black hole mass and disk inclination as recently demonstrated by observations of optical lines in Mrk~110. Comparing our simulated lines with this observations, we independently confirm an inclination angle of 30 degrees for the accretion disk. Redshift deviations induced by black hole spin can be probed only very close to the black hole e.g. with X--ray iron lines. } ", "introduction": "Active galactic nuclei (AGN) such as Seyfert galaxies and quasars are powered by accreting supermassive black holes (SMBHs) following the standard model that has been developed over four decades \\citep{Lynden-Bell1969, Lynden-Bell1971}. The masses of SMBHs lie in the range $10^{6}$ to $10^{10}$ M$_{\\odot}$, see e.g \\cite{Netzer2003}. In the standard model, clouds moving in the gravitational potential of the black hole are photoionized by the central AGN continuum, thereby producing Doppler broadened emission lines with widths of typically 10$^{3}$--10$^4$ km\\,s$^{-1}$ \\citep[~as pioneering studies]{Woltjer1959}. The region where the broad lines originate is usually referred to as the broad--line region (BLR). The scale of the BLR is believed to be $10^{15}$ to $10^{17}$ cm, corresponding to $\\sim 10^{3}$ to $\\sim 10^{5} \\ r_\\mathrm{g}$ or 0.6 to 60 light days for a 10$^{7}$ M$_{\\odot}$ black hole. Here the gravitational radius is defined as $r_{\\rm g}= GM/c^{2}$ with Newton's constant ${\\rm G}$, vacuum speed of light ${\\rm c}$ and black hole mass $M$. \\cite{Robinson1990} have presented complex models involving spherical or disk geometries for the BLR as well as rotational and radial cloud kinematics. They studied the influence of continuum variability on line profiles and found diverse line shapes exhibiting spikes, bumps and shoulders though in a non--relativistic regime. Additional velocity components in the BLR caused by disk winds \\citep{Konigl1994} and a radial component have been suggested from accretion theory and radial velocity maps of the narrow--line region \\citep{Ruiz2001}. However, the detailed structure and velocity field of the BLR remain unclear \\citep{Collin2006}. \\\\ The broad--line clouds respond to variations in the central photoionizing continuum as suggested by strong correlations between H$\\beta$ response times and non--stellar optical continuum fluxes, $\\tau_\\mathrm{cent}\\propto\\sqrt{F_\\mathrm{UV}}$ \\citep{Peterson2002}. This phenomenon is exploited in reverberation mapping techniques to determine both the scale of the BLR and the black hole mass \\citep{Blandford1982, Peterson1993, Kaspi2000}. \\\\ The idea that gravitational redshift may influence optical lines causing line asymmetries was raised by \\citet{Netzer1977}. In the Seyfert--1 galaxy Akn~120, a slight redward displacement of the H$\\beta$ line was reported, amounting to $\\Delta z\\sim 0.0013$, interpreted as the result of gravitational redshift \\citep{Peterson1985}. However, such effects may also arise from attenuation of the BLR or light--travel time effects, as discussed by \\cite{Peterson1985}. Similar studies that assume that observed effects are a result of gravitational redshift have been done for a quasar sample where the SMBH mass of QSO~0026+129 could be roughly estimated to be $2\\times 10^9\\sol$ \\citep{Zheng1990}; this is still the current value within a factor of 2 \\citep{Czerny2004}. Recently, several BLR optical emission lines in the narrow--line Seyfert--1 galaxy Mrk~110 were investigated \\citep[see][ ~K03 hereafter]{Kollatschny2003}. In that work, H$\\alpha$, H$\\beta$, HeI$\\lambda$5876 and HeII$\\lambda$4686 emission lines were found to possess a systematic shift to the red, with higher ionization lines showing larger shifts as expected in a BLR with stratified ionization structure. \\\\ In this paper, we study the gravitational redshift over a large range of distances from the central black hole. We quantify the relativistic gravitational redshift on emission lines until GR fades beyond the current observable limit. The investigation is carried out in a very general form by discussing the observed line profile as a function of the generalized GR Doppler factor ($g$--factor) for Kerr black holes and an arbitrary velocity field of emitters, see e.g.\\ \\citet[ ~M04 hereafter]{Mueller2004}. Pioneering work on relativistic spectra was performed by \\citet{Cunningham1975} using transfer functions. However, the considerations of the $g$--factors in this work were restricted to minimum and maximum values of $g$ on infinitesimally narrow and thin stationary rings. Furthermore, the distance range of interest for BLRs, $10^{3}$ to $10^{5} \\ r_\\mathrm{g}$, has not been investigated in detail. \\\\ \\cite{Corbin1997} studied relativistic effects on emission lines from the BLR by assuming Keplerian orbits for the emitting clouds in a Schwarzschild geometry. It was found that line profiles decrease in both, width and redward centroid shift when the line emitting region moves away from the black hole. \\\\ Our goal is to accurately quantify the effects of gravitational redshift in the vicinity of a Kerr black hole. After a very general consideration that holds for any classical black hole of arbitrary mass, a more specific treatment involving optical emission lines from BLRs is addressed. For the case study of Mrk~110, it is even demonstrated how the mass of the SMBH and the inclination of the inner disk can be determined. ", "conclusions": "\\label{sec:conc} Line cores at distances from 2 to 100000 $r_\\mathrm{g}$ from a rotating black hole have been analysed using relativistic ray tracing simulations in the Kerr geometry. The line cores are gravitationally redshifted by $z_\\mathrm{core}\\simeq \\ 10^{-5},\\,10^{-4},\\,10^{-3},\\,10^{-2},\\,10^{-1}\\,10^{0}$ at distances of $100000,\\,10000,\\,1000,\\,100,\\,10,\\,2 \\ \\mathrm{r_g}$ from the black hole, respectively. This $z\\approx\\frac{1}{\\mathrm{distance[r_g]}}$ behaviour at large radii is a straightforward consequence of the Schwarzschild factor. Lines characterised by a core energy $g_\\mathrm{core}$ confirm this scaling behaviour. Gravitational redshift occurs in two modes. One regime starts at larger distances from the black hole and shifts only the line as a total feature while conserving its intrinsic shape; the amount of redshift can be looked up in Tab. \\ref{tab:01}. The other regime, which is the strong gravity regime, dominates at $r\\lesssim 10 \\ \\mathrm{r_g}$ or at $r \\lesssim 10^{13}$ cm for a ten million solar mass black hole. Relativistic emission lines originating in this region are strongly deformed and suppressed and differ substantially from the corresponding line profiles in the emitter's rest frame. The present work has demonstrated that the onset of GR becomes important at distances smaller than $75000 \\ \\mathrm{r_g}$ if assuming an optical spectral resolution of 0.1 \\AA. It is stressed that this critical radius depends on the astronomical resolving power and lies farther out if the resolution is higher. However, even with high--resolution spectroscopy it remains a challenge to probe gravitationally redshifted spectral lines due to the fact that competing effects and more complex physics are likely to be involved: a flat Keplerian BLR model may be too simple to model the complex BLR velocity field and both spherical BLR structure and a wind component may play a crucial role. In addition, narrow-line components and e.g.\\ contamination of the H$\\beta$ line by FeII may also complicate both the analysis and the interpretation. Nevertheless, the theory of General Relativity predicts that the gravitational redshift is present, and a valuable ansatz for probing it is to search for systematic shifts with varying distance, as done by K03. It is suggested in this work to supplement this by multi--wavelength observations that should all point towards the same central mass and inner inclination. We confirm the analysis by K03 of the NLS--1 galaxy Mrk~110 here, using a more general treatment with stationary $g$--factors that are in concordance with observational data. Ray tracing simulations in the Kerr geometry support an inclination of $i=30^{\\circ}$ for the inner disk of Mrk~110. Reversely, if the inclination is known, this can be exploited to determine the black hole mass from the fitting procedure outlined here. Whether fitting $i$ or $M$ (or both) -- such techniques may help explore AGN unification schemes: multi--wavelength studies allow for studying the inclination deep into the AGN and to probe orientation and luminosity--dependence of AGN types. Furthermore, we show that broad optical lines can not serve as a probe of black hole spin because frame--dragging effects only occur very close to the black hole. In this region, only hot emission lines (such as the Fe K$\\alpha$ line in X--rays) or other relativistic spectral features indicate black hole spin. The analysis presented here is not only valid for supermassive black holes but also for stellar--mass black holes in X--ray binaries or intermediate--mass black holes that may be found in ultra--luminous X--ray sources or globular clusters." }, "0607/hep-ph0607086_arXiv.txt": { "abstract": "In this article, we explore the ability of direct and indirect dark matter experiments to not only detect neutralino dark matter, but to constrain and measure the parameters of supersymmetry. In particular, we explore the relationship between the phenomenological quantities relevant to dark matter experiments, such as the neutralino annihilation and elastic scattering cross sections, and the underlying characteristics of the supersymmetric model, such as the values of $\\mu$ (and the composition of the lightest neutralino), $m_A$ and $\\tan \\beta$. We explore a broad range of supersymmetric models and then focus on a smaller set of benchmark models. We find that by combining astrophysical observations with collider measurements, $\\mu$ can often be constrained far more tightly than it can be from LHC data alone. In models in the $A$-funnel region of parameter space, we find that dark matter experiments can potentially determine $m_A$ to roughly $\\pm 100$ GeV, even when heavy neutral MSSM Higgs bosons ($A$, $H_1$) cannot be observed at the LHC. The information provided by astrophysical experiments is often highly complementary to the information most easily ascertained at colliders. ", "introduction": "A great deal of effort has been directed to developing methods of detecting particle dark matter. Over the past years and decades, numerous studies have been conducted to assess the prospects for these various techniques \\cite{review}. If dark matter consists of neutralinos, or another weakly interacting particle with a TeV-scale mass, it is likely that one or more of these techniques will make the first detection of dark matter particles within the next several years. Of all of the candidates for dark matter that have been proposed, none has received as much attention as the lightest neutralino in models of supersymmetry. Supersymmetry is theoretically attractive for a variety of reasons. Among the most compelling is its ability to provide a natural solution to the hierarchy problem~\\cite{susyreview}, and a common scale for the unification of the forces of the Standard Model~\\cite{gut}. From the standpoint of providing a dark matter candidate, the lightest neutralino is naturally stable by virtue of R-parity conservation \\cite{neutralinodm}, and in many models is produced in the early universe in a quantity similar to the measured density of cold dark matter~\\cite{wmap}. Astrophysical techniques for detecting neutralinos include direct and indirect detection experiments. Direct detection experiments attempt to observe neutralinos scattering elastically off of target nuclei. Indirect detection experiments, in contrast, attempt to detect the annihilation products of neutralinos, including gamma-rays, neutrinos, positrons, anti-protons and anti-deuterons. Direct and indirect detection measurements are of critical importance in determining the identity of dark matter. Even if collider experiments were to observe a long-lived, weakly interacting, massive particle that appears to be a suitable dark matter candidate, such experiments will never be able to determine whether a particle is stable over cosmological timescales. To determine whether the dark matter of our universe is made up of such a particle (either entirely, or in part), direct and indirect detection experiments will be needed. But looking beyond the mere detection of dark matter, what will these astrophysical observations reveal to us about the nature of particle dark matter? In particular, is it possible to determine the properties of the lightest neutralino and the corresponding supersymmetry model by direct or indirect dark matter detection experiments? In this article, we attempt to address these questions. While extracting such information from direct and indirect dark matter detection experiments is challenging, it may be possible in many scenarios. We have found that direct detection measurements of the neutralino's spin-independent elastic scattering cross section, rates in neutrino telescopes, and the brightness of gamma-ray lines (from $\\chi^0_1 \\chi^0_1 \\rightarrow \\gamma \\gamma$ and $\\chi^0_1 \\chi^0_1 \\rightarrow \\gamma Z$) are among the most useful astrophysical probes for determining the properties of supersymmetry. In most cases, the value of $\\mu$ (or alternatively, the composition of the lightest neutralino) can be constrained far more tightly if astrophysical data is included than it can be by LHC data alone. Furthermore, in some models (those in the $A$-funnel region of supersymmetric parameter space) the mass of the CP-odd Higgs boson can be determined by astrophysical experiments to within roughly $\\pm 100$ GeV, even if it cannot be observed at the LHC. ", "conclusions": "If low energy supersymmetry exists in nature, it will likely be discovered in the next few years at the LHC (or possibly at the Tevatron). Over roughly the same period of time, the prospects for direct and indirect searches for neutralino dark matter are also very encouraging. Each of these windows into the characteristics of supersymmetry can provide us with useful and complementary information. In this paper, we have studied the ability to constrain the parameters of supersymmetry, beyond what can be done at the LHC, using direct and indirect dark matter experiments. Although the presence of a neutralino LSP can in most cases be confirmed at the LHC, the composition and corresponding couplings of such a state will likely go unconstrained by such an experiment. Direct and indirect dark matter experiments, on the other hand, can probe the lightest neutralino's interactions with nucleons and with themselves (annihilations), potentially allowing for a determination of its composition (and the value of $|\\mu|$). We also find that direct and indirect dark matter measurements, combined with relic abundance considerations, can in some cases (the $A$-funnel region) determine the mass of the CP-odd Higgs boson ($A$), even when beyond the reach of the LHC." }, "0607/astro-ph0607376_arXiv.txt": { "abstract": "{} {We characterize the importance of metallicity on the presence of molecular hydrogen in damped Lyman-$\\alpha$ (DLA) systems.} {We construct a representative sample of 18 DLA/sub-DLA systems with log~$N$(H~{\\sc i})~$>$~19.5 at high redshift ($z_{\\rm abs}$~$>$~1.8) with metallicities relative to solar [X/H]~$>$~$-$1.3 (with [X/H] = log~$N$(X)/$N$(H)$-$log(X/H)$_{\\odot}$ and X either Zn, S or Si). We gather data covering the expected wavelength range of redshifted H$_2$ absorption lines on all systems in the sample from either the literature (10 DLAs), the UVES-archive or new VLT-UVES observations for four of them. The sample is large enough to discuss for the first time the importance of metallicity as a criterion for the presence of molecular hydrogen in the neutral phase at high-$z$. } {From the new observations, we report two new detections of molecular hydrogen in the systems at $z_{\\rm abs} = 2.431$ toward Q\\,2343$+$125 and $z_{\\rm abs}= 2.426$ toward Q\\,2348$-$011. We compare the H$_2$ detection fraction in the high-metallicity sample with the detection fraction in the overall sample from Ledoux et al. (2003). We show that the fraction of DLA systems with log~$f$~=~log~2$N$(H$_2$)/(2$N$(H$_2$)~+~$N$(H~{\\sc i}))~$>$~$-$4 is as large as 50\\% for [X/H]~$>$~$-$0.7 when it is only $\\sim$5\\% for [X/H]~$<$~$-$1.3 and $\\sim$15\\% in the overall sample (with $-$2.5~$<$~[X/H]~$<$~$-$0.3). This demonstrates that the presence of molecular hydrogen at high redshift is strongly correlated with metallicity. } {} ", "introduction": "Early searches for molecular hydrogen in DLAs, though not systematic, have lead to either small values of or upper limits on the molecular fraction of the gas. For a long time, only the DLA at $z_{\\rm abs}=2.811$ toward Q\\,0528$-$250 was known to contain H$_2$ molecules (Levshakov \\& Varshalovich 1985, Foltz et al. 1988). Ge \\& Bechtold (1999) searched for H$_2$ in a sample of eight DLAs using the MMT moderate-resolution spectrograph (${\\rm FWHM}=1$ \\AA ). Apart from the detection of molecular hydrogen at $z_{\\rm abs}=1.973$ and 2.338 toward, respectively, Q\\,0013$-$004 and Q\\,1232$+$082, they measured in the other systems upper limits on $f$ in the range $10^{-6}-10^{-4}$. A major step forward in understanding the nature of DLAs through their molecular hydrogen content has recently been made possible by the unique high-resolution and blue-sensitivity capabilities of UVES at the VLT. In the course of the first large and systematic survey for H$_2$ at high redshift, we have searched for H$_2$ in DLAs down to a detection limit of typically $N($H$_2)=2\\times 10^{14}$ cm$^{-2}$ (Petitjean et al. 2000, Ledoux et al. 2003). Out of the 33 surveyed systems, eight had firm detections of associated H$_2$ absorption lines. Considering that three detections were already known from past searches, H$_2$ was detected in $\\sim 15$\\% of the surveyed systems. The existence of a correlation between metallicity and depletion factor, measured as [X/Fe] (with X~=~Zn, S or Si) was demonstrated (see also Ledoux et al. 2002a) and the DLA and sub-DLA systems where H$_2$ was detected were usually among those having the highest metallicities. However, the high metallicity end of the Ledoux's sample was biased by the presence of already known detections. Therefore, to investigate further this possible dependence with metallicity, and derive what is the actual molecular content of the high-redshift gas with highest metallicity, we have searched a representative sample of high metallicity DLAs for H$_2$. The number of H$_2$ measurements in systems with [X/H]~$>$~$-$1.3 is twice larger in our sample compared to previous surveys. We describe the sample and the observations in Section~2, present two new detections of H$_2$ in Section~3 and the results of the survey and our conclusions in Section~4. ", "conclusions": "" }, "0607/astro-ph0607189_arXiv.txt": { "abstract": "We investigate spherical accretion to a rotating magnetized star in the \"propeller\" regime using axisymmetric resistive magnetohydrodynamic simulations. The regime is predicted to occur if the magnetospheric radius is larger than the corotation radius and smaller than the light cylinder radius. The simulations show that accreting matter is expelled from the equatorial region of the magnetosphere and that it moves away from the star in a supersonic, disk-shaped outflow. At larger radial distances the outflow slows down and becomes subsonic. The equatorial matter outflow is initially driven by the centrifugal force, but at larger distances the pressure gradient force becomes significant. We find the fraction of the Bondi accretion rate which accretes to the surface of the star. ", "introduction": "Rotating magnetized neutron stars pass through different stages in their evolution (Shapiro \\& Teukolsky 1983, Lipunov 1992). Initially, a rapidly rotating ($P \\le 1{\\rm s}$) magnetized neutron star is expected to be active as a radio-pulsar. The star spins down owing to the wind of magnetic field and relativistic particles from the region of the light cylinder $r_L$ (Goldreich \\& Julian 1969). However, after the neutron star spins-down sufficiently, the light cylinder radius becomes larger than magnetospheric radius $r_m$ where the ram pressure of external matter equals the magnetic pressure in the neutron star's dipole field. The relativistic wind is then suppressed by the inflowing matter (Shvartsman 1970). The external matter may come from the wind from a binary companion or from the interstellar medium for an isolated neutron star. The centrifugal force in the equatorial region at $r_m$ is much larger than gravitational force if $r_m$ is much larger than the corotation radius $r_{cor}$. In this case the incoming matter tends to be flung away from the neutron star by its rotating magnetic field. This is the so called ``propeller\" stage of evolution (Davidson \\& Ostriker 1973, Illarionov \\& Sunyaev 1975). The ``propeller\" stage of evolution, though important, is still not well-understood theoretically. We discuss results of axisymmetric, two-dimensional, resistive MHD simulations of accretion to a a rotating magnetized star in the ``propeller\" regime. We treat the case when matter accretes spherically with the Bondi accretion rate (Bondi 1952). Bondi accretion to a non-rotating and a slowly rotating star was investigated by Toropin et al. (Toropin at al. 1999) and Toropina et al. (Toropina at al. 2003). Investigation of the accretion to a rotating star in the ``propeller\" regime was started by Romanova et al. (Romanova et al. 2003). \\begin{figure} \\plotone{toropina1.eps} \\caption{Matter flow in the ``propeller\" regime for a star rotating at $\\Omega_*=0.5\\Omega_{K*}$ after $6.9$ rotation periods of the star. The axes are measured in units of the star's radius. The background represents that the density and the length of the arrows is proportional to the poloidal velocity. The thin solid lines are magnetic field lines.} \\end{figure} ", "conclusions": "Axisymmetric magnetohydrodynamic simulations of Bondi accretion to a rotating magnetized star in the propeller regime of accretion have shown that: (1) A new regime of matter flow forms around a rotating star. Matter falls down along the axis, but only a small fraction of the incoming matter accretes to the surface of the star. Most of the matter is expelled radially in the equatorial plane by the rotating magnetosphere of the star. A low-density torus forms in the equatorial region which rotates with velocity significantly larger than the radial velocity. Large scale vortices form above and below the equatorial plane. (2) The accretion rate to the star is much less than the Bondi accretion rate and decreases as the star rotation rate increases ($\\propto \\Omega_*^{-1.0}$), (b) as the star's magnetic moment increases ($ \\propto \\mu^{-2.1}$), and as the magnetic diffusivity decreases [$\\propto (\\eta_m)^{0.7}$]. (3) Because the accretion rate to the star is less than the Bondi rate, a shock wave forms in our simulations and propagates outward. It has the shape of an ellipsoid flattened along the rotation axis of the star." }, "0607/astro-ph0607140_arXiv.txt": { "abstract": "EF Eri is a magnetic cataclysmic variable that has been in a low accretion state for the past nine years. Low state optical spectra reveal the underlying Zeeman-split white dwarf absorption lines. These features are used to determine a value of 13-14 MG as the white dwarf field strength. Recently, 5-7 years into the low state, Balmer and other emission lines have appeared in the optical. An analysis of the H$\\alpha$ emission line yields the first radial velocity solution for EF Eri, leading to a spectroscopic ephemeris for the binary and, using the best available white dwarf mass of 0.6M${\\odot}$, a mass estimate for the secondary of 0.055M${\\odot}$. For a white dwarf mass of 0.95M${\\odot}$, the average for magnetic white dwarfs, the secondary mass increases to 0.087M${\\odot}$. At EF Eri's orbital period of 81 minutes, this higher mass secondary could not be a normal star and still fit within the Roche lobe. The source of the Balmer and other emission lines is confirmed to be from the sub-stellar secondary and we argue that it is due to stellar activity. We compare EF Eri's emission line spectrum and activity behavior to that recently observed in AM Her and VV Pup and attributed to stellar activity. We explore observations and models originally developed for V471 Tau, for the RS CVn binaries, and for extra-solar planets. We conclude that irradiation of the secondary in EF Eri and similar systems is unlikely and, in polars, the magnetic field interaction between the two stars (with a possible tidal component) is a probable mechanism which would concentrate chromospheric activity on the secondary near the sub-stellar point of the white dwarf. ", "introduction": "EF Eridani has become a binary star of renewed interest in recent years. This is primarily due to the fact that the mass accretion from the low-mass secondary to the highly magnetic white dwarf has been essentially stopped for the past 9 years. During this period of very low accretion, observers have flocked to telescopes around the world in hope of detecting and understanding the component stars that make up EF Eri. As we will see, it has taken a small 1.5-m telescope, part of the SMARTS consortium, to provide the first new insights into the true nature of this fascinating (and confounding) binary. EF Eri was discovered over 30 years ago as a hard X-ray source (2A 0311 -227; Cooke et al. 1978). It was quickly identified optically as a 14th magnitude blue source (Griffiths et al. 1979) with properties similar to AM Herculis (Charles and Mason 1979). AM Her, the not-so-typical prototype of the AM Her class of interacting binaries, was the first member of a new class of cataclysmic variable (Tapia 1977). Today, these magnetic cataclysmic variable (CV) binary systems are called polars (due to their polarized light), systems in which matter is accreted onto the more massive primary from the lower mass secondary. The primary star is a highly magnetic white dwarf (B$\\sim$10 MG to 250 MG) and the mass donor ranges from normal M stars (as in AM Her) to sub-stellar brown dwarf-like objects which are believed to reside in EF Eri and other ultra-short period CVs (see Warner 1995 for a review). Polars show large, random changes in their brightnesses on weeks to months, and even year timescales. In the high state, when mass accretion is in full swing, their luminosities are dominated by an accretion-produced continuum from the X-ray to the IR, superposed with strong emission lines from hydrogen and helium. When the mass transfer slows or stops, the underlying stars are revealed, and these dominate the optical and infrared spectral energy distribution. The origin for the starting and stopping of the mass transfer from the low mass companion is not understood, but stellar activity cycles in the synchronously locked, rapidly spinning late type secondary has been a long standing explanation (see Bianchini, 1992; King and Cannizzo 1998; Hessman et al. 2000) which has recently gained some observational support (Kafka et al. 2005, 2006; Mason et al. 2006). A number of recent papers have been dedicated to EF Eri (e.g., Beuermann et al. 2000; Harrison et al. 2003, 2004) and they track the evolution of this system into its low state, where the system evolved to the point where the optical spectrum was emission-line free, and dominated by a cool white dwarf that showed Zeeman-split Balmer absorption lines (e.g., see Fig. 2 in Harrison et al. 2004). The secondary in EF Eri has not yet been conclusively identified, but both Beuermann et al. (2000), and Harrison et al. (2004) suggest that it must be brown dwarf-like object, and recent $Spitzer$ observations detect strong mid-infrared emission that suggest this object is an L or T dwarf (Howell et al., 2006). We have been monitoring EF Eri with regularity over the past 9 years, using its continued low state as an invitation from the binary to attempt to learn about its stellar components. During the Fall of 2004, we noticed that EF Eri, while becoming no brighter, began to show weak, narrow H$\\alpha$ emission. The remainder of the Balmer series, and other spectral lines, are also now present. Detailed analysis of the H$\\alpha$ emission line yields the first dynamical orbital velocity solution for EF Eri. We use this information and the remaining spectral features to determine the source of the emission and show that it is consistent with chromospheric activity on the sub-stellar secondary. ", "conclusions": "\\subsection{Low States} Today, we know that most polars, including EF Eri, AM Her, and VV Pup have extended low states lasting years and that, in general, most polars show long term low states and/or spend much of their time in low states (Gerke et al. 2006 and refs therein). It is believed that all CVs (polars, dwarf novae, etc.) should have low states but that we only notice them in polars due to their lack of an accretion disk. In CVs with disks, if the mass transfer stops, the optical light may not show a significant dimming, as the optical light is dominated by or has a large contribution from the accretion disk. The stoppage of mass transfer in a non-magnetic CV may not be easily noted and if mass transfer restarts within a short time period ($\\sim$2-4 weeks), the entire event might escape detection. King and Cannizzo (1998) provide a theoretical framework for this idea and explore possible observational ramifications that could result. They conclude that current observations do not not agree with their predictions of how a disk system would react to a stoppage of mass transfer. Given their work and the fact that polars often show extended low states it is hard to reconcile the idea that CVs with accretion disks (i.e., the non-magnetic white dwarf CVs) have low states similar to polars. The VY Scl stars may be the exception, as we discuss below. It has been argued that the low states are a result of stellar activity cycles on the secondary star, in particular starspots migrating to the L1 region and stopping mass transfer for some period of time. The idea that ``solar cycles\" and starspots cause low states has been around for decades. Van Buren and Young (1985) suggested that changes in radius of the magnetically-active secondary in RS CVn systems drove the observed orbital period variations. The basic idea is that, during a stellar activity cycle, the magnetic pressure due to the enhanced subsurface magnetic fields displace gas, resulting in a fractionally larger stellar radius. The change in the moment of inertia drives the system out of synchronous rotation, and tidal torques then quickly bring the system back into synchrony. In the case of the RS~CVn secondaries, the fractional change in the orbital periods ($\\sim$10$^{-6}$) requires a similar fractional change in the stellar radius. Applegate \\& Patterson (1987) applied this type of model to the cataclysmic variables, and predicted that there should be periodic $O-C$ variations on the magnetic activity period of the secondary. These magnetic activity periods are observed to be of order a decade long, within a factor of 2 of the length of the Solar cycle. In the case of the polars, it may be this fractional change in the stellar radius that comes into play. Since the secondary is filling its Roche lobe when in the high state, a small reduction in the stellar radius as the magnetic activity wanes may lead to a cessation of the accretion. The implication is that as the magnetic activity picks up, the star should then expand and eventually accretion should resume. This seems consistent with what we observed in the EF Eri secondary: the magnetic activity was at a minimum when the accretion ceased, and then increased prior to the start of the current high state. The few direct observational results we do have that reveal stellar activity are all from polars in low states. We have seen that EF Eri entered a low state and its secondary was in-active, it turned on, and EF Eri stayed in a low state for another 1.5 years. On the other hand, AM Her and VV Pup with normal active M star secondaries, went into low states (recently, twice each) and from start to end during the low state the secondary was active, at a constant level, throughout. Single active M stars display spectroscopic activity indicators almost all the time but when they go into flaring or super-active episodes, their chromospheres can expand locally and eject material. This idea of active chromospheric expansion driving L1 mass loss was proposed by Howell et al. (2000) as the possible cause of high/low state. The model removed the need for a starspot at L1 but still kept the idea that stellar activity/flaring somehow triggers state changes. If some magnetic cycle connection is invoked to drive mass flows, then it may be that even in a high state, the photosphere need not fill its Roche lobe The evidence for starspots at L1 and/or stellar activity {\\it directly} starting and stopping high/low states seems not to exist. Assuming, for the moment, that non-magnetic CVs do not undergo low states, can we explain such a phenomena? Polar secondaries have been shown to be normal main sequence stars (Harrison et al. 2005) while dwarf nova secondary stars, to date at least for systems above the period gap, have been shown to have spectra that reveal odd abundance patterns and are consistent with CNO processed material. Evidence has been presented that argues for a divergence in the evolutionary history of polars and non-magnetic CVs (see Schmidt et al. 2005; Harrison et al. 2005). If the secondaries in non-magnetic CVs are in fact remnant He cores from a past time when they were more massive, as has been suggested by Beuermann et al. (1998), Howell (2005), and Harrison et al. (2005), possibly evolving from Algols, then maybe normal stellar activity is not possible. If high/low states are somehow related to activity cycles on the mass donor, this idea may provide a natural explanation as to why dwarf novae and other non-magnetic CVs do not have polar-like low states. Having said this, we must consider the VY Scl type of CV. These systems are believed to have non-magnetic white dwarfs and accretion disks, yet they show low state behavior that may be similar to that seen in polars. VY Scl stars have been used as ``proof\" of starspots and stellar activity but to date, this is more speculation than anything else. This sub-class of CV lies in the 3-4 hour period range, directly above the period gap. Howell et al. (2001) suggests these stars behave as they do because their secondaries are far from thermal equilibrium and mass transfer can be highly modulated by thermal timescale changes caused by mass loss. If this is true or not, is yet to be seen, but observations of the secondary in VY Scl stars are likely to offer important clues as to the cause of high/low states. \\subsection {Related Binaries} What does the secondary star do in terms of its activity cycles before, during, and after high/low state transitions? We have seen that stellar activity (as indicated by the usual optical emission lines) may be a constant feature in the normal M star secondaries of polars (i.e., VV Pup and AM Her) and that it can ``turn on\" during a low state (or at other times?) in at least one sub-stellar mass object, the secondary in EF Eri. Let us examine a few other binary systems in which magnetic fields and stellar activity play important roles. These objects may provide valuable clues toward our understanding of polars. \\subsubsection {RS~CVn Binaries} Low mass stars in close binaries exhibit enhanced activity, which is generally attributed to the amplification of the magnetic dynamo by tidally-enforced rapid rotation. The activity is generally saturated, with activity-related emission at the level of 10$^{-2}$ to 10$^{-3}$ of the bolometric flux. The secondaries of polars have similarly-rapid rotation and, if magnetic dynamo processes are present they may also be expected to exhibit activity at the saturated levels. However, the analogy is not exact, as the polar secondaries may have lost a significant fraction of their convective zones, with an effect on their ability to maintain or amplify a magnetic dynamo. Stellar magnetic field emergence is known to be concentrated at certain active longitudes, both in the Sun and more active stars. The active RS~CVn binaries are characterized by a ``photometric wave\", interpreted as a spotted (and magnetically active) hemisphere that migrates around one of the stars with respect to the binary phase with a period of order a decade. The stars are tidally locked in such close binary systems; the migration is interpreted to represent differential rotation of the star, as the starspots migrate equator-ward as they do during a solar cycle. While there is evidence for such a photometric wave with a 6-9 month period in the pre-cataclysmic variable V471 Tau system (\\.{I}banogl\\u{u} et al.\\ 1994), there is also evidence that the chromospheric H$\\alpha$ is enhanced near the substellar longitudes (Rottler et al.\\ 2002). Because the strength of the emission decreased with time, Rottler et al.\\ dismissed irradiation as the cause, and suggested that there is a permanent active longitude at the substellar point, caused by tidal distortion of the convective dynamo. It is not clear whether this is in conflict with the observed wave migration (see \\S4.2.2). There is spectroscopic evidence for enhanced flaring activity at the substellar longitudes in the RS~CVn system UX~Arietis (Simon, Linsky, \\& Schiffer 1980), and photometric evidence for enhanced flaring activity at the substellar longitudes in the close binary dMe system YY~Gem (Doyle \\& Mathioudakis 1990). This is presumably attributable to recombination between the extended magnetic loops of the two stars, as in the models by Uchida and Sakurai (1985) (see Figure 11). There is no evidence addressing the long term variability of the flaring rates. If the recombination involves large-scale loops with sizes of order the binary separation, then flaring rates may remain more-or-less constant, but the rates may be enhanced when the active longitudes coincide with the substellar point. Eclipse-mapping observations of RS CVn systems, of Algol, and of YY~Gem have been brought to bear on the question of the presence of magnetic connections between stars, as might be expected if the substellar magnetic activity enhancement is indeed permanent. Pre\\'s, Siarkowski, \\& Sylwester (1995) and Siarkowski et al.\\ (1996) modeled X-ray eclipse observations of the RS CVn systems TY~Pyx and AR~Lac, respectively. They concluded that in both cases a significant fraction of the emission arises between the stars, in magnetic loops connecting the stars. However, their unconstrained maximum entropy modeling solutions are not unique, and other equally good solutions without inter-stellar emission exist (see the review by G\\\"udel [2004]). \\subsubsection {V471 Tau} For V471 Tau, often used as the prototype pre-CV, the hot white dwarf (T=35,000K) was originally thought to irradiate the K secondary star. The evidence for this was H$\\alpha$ emission from the secondary star which was observed to be concentrated toward the white dwarf with an equivalent width that peaked near orbital phase 0.5 and the emission line disappeared completely when looking at the back side of the secondary. However, multi-year observations of V471 Tau by Rottler et al. (1998, 2002) showed that in 1987, the H$\\alpha$ emission was consistent with an irradiation interpretation while in 1990 the emission was much weaker and heavily concentrated toward the white dwarf and in 1992 the emission was completely absent. The (non-magnetic) white dwarf in V471 Tau was observed to be constant throughout this entire 5 year period showing that the secondary star emission was not due to irradiation by the hot white dwarf. Rottler et al. believe their observations are consistent with a ``solar cycle\"-like change that occurred in the K star. Given their result with V471 Tau, Rottler et al. (2002) (re)examined a number of similar WD + RD systems all of which were supposed to show irradiation induced emission from the non-interacting secondary star. Of the ten hot white dwarf (T=30,000 to 60,000K) plus red dwarf stars, only one (NN Ser, a WD+RD pair with T$_{WD}$=55,000K) is probably a true case of irradiation. The others are shown to have not enough UV flux to cause irradiation and the variable H$\\alpha$ emission was consistent with activity induced emission lines. Using a model based on the number of $<$912\\AA\\ photons available for irradiation of the secondary, they show that an incident UV flux at the surface of the secondary star of $\\la$1$\\times$10$^{10}$ ergs sec$^{-1}$ cm$^{-2}$ ster$^{-1}$ is not sufficient to produce irradiation induced emission lines. This value exceeds that present in EF Eri from its cooler but closer white dwarf (F$_{UV}$ $\\sim$1$\\times$10$^{9}$ ergs sec$^{-1}$ cm$^{-2}$ ster$^{-1}$) and, in fact, it exceeds nearly every CV when not in outburst. \\subsubsection {SDSS J121209.31+013627.7} It is interesting to note here the discovery of an object that is very similar to EF Eri in the low state, SDSS J121209.31+013627.7 (Schmidt et al. 2005). These authors discuss this 90 minute, cool (T$\\sim$10,000K), orbitally synchronized magnetic (7-13 MG) white dwarf binary and present optical spectroscopy closely approximating those of EF Eri in a low state. They conclude, however, that SDSS J1212's H$\\alpha$ emission is most likely due to irradiation of the probable brown dwarf secondary based on a single epoch set of phase-resolved optical spectroscopy. The H$\\alpha$ emission completely disappears when viewing the back end of the secondary star, a similar result to that observed once in V471 Tau. We therefore can ask, given the nearly similar nature of J1212 and EF Eri, if the secondary in J1212 is irradiated, why isn't the secondary in EF Eri? With stellar activity seemingly being concentrated on the secondary at the sub-stellar point of the white dwarf, we will need to be careful in our observational interpretation of any detected emission from the secondary. Radiatively heated atmospheres (irradiation) and active stellar chromospheres (starspots etc.) may be hard to distinguish. If a star is irradiated, the apparent spectral type (photospheric temperature) should vary as the star rotates. Active chromospheres do fill in lines, which mimics an earlier spectral type in the K stars, but the veiling is wavelength-dependent. In M stars, the veiling might give a spectral type earlier than T$_{eff}$, since molecular band strengths are increasing with decreasing T$_{eff}$. Realistic, 3-D, magneto-hydrodynamic models of the secondary star and its magnetic, tidal, and radiative interaction with the primary are needed to fully understand the observations. J1212 and EF Eri are good starting points to use as proxies for our understanding of irradiation, stellar activity on brown dwarf-like secondary stars, and exoplanet physics. With no mass transfer currently underway in J1212, this system is an ideal candidate for multi-epoch observations to monitor and detail the nature of the secondary star emission lines. \\subsubsection{Extra-Solar Planets} The analogy between polars and the magnetically-active binary systems is imperfect. In none of these cases does the active star fill its Roche lobe, in none of these cases does the strength of the photospheric magnetic field exceed a few kG, and in none of these cases is the secondary star bathed in a strong external magnetic field. The tidal forces are also much stronger in EF Eri than even in V471 Tau, where the period is 8 times longer and the mass of the cool star is over an order of magnitude larger. A better analogy may be between a star and a planet. Perhaps the answer to starting and stopping mass transfer lies in the area of magnetic (re)connection between the white dwarf and the secondary star. It was observed (Shkolnik et al. 2003) that for ``hot Jupiter\" type extra-solar planets, the host star Ca II H\\&K emission lines are sometimes modulated on the orbital period of the planet. A model proposed by Ip et al. (2004) explains this phenomenon as an interaction of the exoplanet magnetosphere with that of the parent star, in which a magnetic flux loop reconnects using the nearby planet as a conductor. In a polar, which has a much stronger field, it seems obvious that magnetic reconnection and closed field loops would have to pass through the secondary. This idea may also explain why the onset of stellar activity alone is not the direct cause of high/low states and why polars in low states seem to have a residual amount of mass transfer, probably magnetically connected secondary star wind accretion. A possible observation of magnetic reconnection and closed field loops in the region near L1 but between the two stars (as illustrated in Fig. 11) was noted by Kafka et al. (2005, 2006) in their low state observations of stellar activity in AM Her. The H$\\alpha$ emission line profile is triple peaked and leads to a model in which the regions on the secondary star where stellar activity occurs seem to be preferentially on the side facing the white dwarf and contain loop type structures surrounding the secondary. A similar white dwarf facing activity concentration and a similar multi-peaked H$\\alpha$ line has been observed in V471 Tau (Young et al. 1991). The H$\\alpha$ emission line satellites have been stable for over 2 years in AM Her and VV Pup showed a similar H$\\alpha$ emission line structure during one of its recent low states (Mason et al. 2006). \\subsection {Additional Observational Study} Stellar activity, as evidenced by emission lines such as H and Ca, is usually an optical bandpass specific proxy. Even very active stars, (see Fig. 12), show little obvious evidence for chromospheric activity in their near-IR and IR spectra. The reason that activity emission is so weak in the near-IR and IR bands is both a contrast effect and one of line formation. The activity induced emission lines in an active binary are often ``filled in\" by the emission from the bright photospheric continuum. Additionally, the typical spectral lines present in the IR region that one might expect to be in emission and associated with stellar activity (e.g., $J$- to $K$-band H I Paschen series and Ca lines) form too low in the stellar atmosphere to be greatly modulated or affected by an active chromosphere. In CVs, the contrast effect just mentioned can be provided by the high state accretion flux or accretion disk light, hiding not only the secondary but any possible activity-induced lines. Additionally, CV emission lines are generally very broad due to the high velocities in the disk or stream, again able to hide weaker, narrow lines caused by activity. Polars in low states offer the lack of additional (accretion) flux contribution allowing secondary star activity indicators to be observed in the optical. However, formation of the higher energy IR lines deeper in the stellar atmosphere still renders the 1-3 micron region a poor choice in the search for chromospheric activity. Indeed, near-IR spectroscopy of EF Eri obtained after the optical emission lines appeared (Johnson et al. 2005), show no sign of hydrogen or any other emission lines. ST LMi as well, was observed at 2.2 microns during a low state and no emission lines were seen, although Balmer emission was present (Howell et al. 2000). Thus, searches for stellar activity in CV secondaries are probably limited to low state observations in the optical, and as such, are mainly restricted to polars. Low state observations, such as those presented herein, are generally difficult to gather as they require optical spectroscopy, often as target of opportunity observations, with fairly large telescopes, as low state polars tend to be faint. Phase resolved spectroscopy is required as the origin and cause of the spectral emission (or absorption, see Mason et al. 2006) features observed must be firmly determined. A single spectrum or even a single epoch of observations can easily confuse activity with irradiation. Polars are probably the easiest targets for this purpose as they show their high/low states directly while the VY Scl stars might be considered prime targets to go after as we know little about their low states or their secondary stars. We have provided new direct spectroscopic evidence into the mix of understanding high/low states of polars and if and how these mass transfer changes are related to stellar activity. The simple idea of the turn on of stellar activity {\\it directly} starting and stopping mass transfer (i.e., high and low states) seems ruled out as activity turned on in EF Eri, yet it remained low for another 1.5 years, while stellar activity seems ever present in the secondary during low states of the polars AM Her and VV Pup. Just as we finish this paper, we can report that EF Eri has reentered a high state, reaching V=15.6 on 2006 March 04 (Stubbings, priv. comm.)\\footnote{ Note that there seems to be nothing special about the location of the start of the high state (JD 2453798.) in Fig. 10.}. Data obtained up to the time of this rebrightening (see Figs. 8, 9, 11) reveal no obvious change in the radial velocity, emission distribution on the secondary, or the EW of H$\\alpha$ directly before the high state started. Our monitoring programs show EF Eri to be currently providing all its usual high state properties observed in the past, even after its nine year hiatus." }, "0607/astro-ph0607230_arXiv.txt": { "abstract": "{} { In this paper, we present a new method to perform numerical simulations of astrophysical MHD flows using the Adaptive Mesh Refinement framework and Constrained Transport. } { The algorithm is based on a previous work in which the MUSCL--Hancock scheme was used to evolve the induction equation. In this paper, we detail the extension of this scheme to the full MHD equations and discuss its properties. } { Through a series of test problems, we illustrate the performances of this new code using two different MHD Riemann solvers (Lax--Friedrich and Roe) and the need of the Adaptive Mesh Refinement capabilities in some cases. Finally, we show its versatility by applying it to two completely different astrophysical situations well studied in the past years: the growth of the magnetorotational instability in the shearing box and the collapse of magnetized cloud cores.} { We have implemented a new Godunov scheme to solve the ideal MHD equations in the AMR code RAMSES. We have shown that it results in a powerful tool that can be applied to a great variety of astrophysical problems, ranging from galaxies formation in the early universe to high resolution studies of molecular cloud collapse in our galaxy.} ", "introduction": "Developing efficient numerical algorithms for the equations of magnetohydrodynamics (MHD) is of great astrophysical interest. Magnetic fields are ubiquitous in a great variety of environments. They are important components of the dynamics in such places as the early universe, the interstellar and intergalactic medium, the environment and interior of stars and the accretion flow around young stellar objects. In the last few decades, finite differences methods have been widely used in investigations of a number of astrophysical situations in which the magnetic field is important with such codes as ZEUS \\citep{stone&norman92a, stone&norman92b}, NIRVANA \\citep{ziegler&yorke97} or the Pencil Code \\citep{brandenburg&dobler02} for example. Even though, as expected, the numerical method breaks down in some circumstances \\citep{falle02}, a considerable amount of progress have been made in our understanding of MHD in astrophysics. A few attempts have also been made to try to extend the Smoothed Particle Hydrodynamics (SPH) method to MHD \\citep{phillips&monaghan85, price&monaghan04a, price&monaghan04b}. At the moment, it is not clear, however, how efficient the resulting codes will prove to be in the future. In the last few years, several attempts have been made to try to extend the standard Godunov approach \\citep{toro97}, initially designed to solve the Euler equations, to MHD. In addition to the accurate description of new waves that are peculiar to MHD (Alfv\\'en waves, the slow and fast modes), one of the most dramatic challenge in the development of such schemes comes from the solenoidality constraint, which states that the divergence of the magnetic field has to vanish everywhere at all times. The first algorithms that attempted to solve this problem kept the cell centering strategy of the standard Godunov approach. They used either a ``divergence cleaning'' step (see for example \\citeauthor{Brackbill80} \\citeyear{Brackbill80} or \\citeauthor{ryu98} \\citeyear{ryu98}), or various reformulations of the MHD equations including additional divergence-waves \\citep{powell99} or divergence-damping terms \\citep{Dedner02} to enforce the solenoidality constraint. A novel cell-centered MHD scheme has been recently developed by \\cite{Crockett05} that combines most of these ideas into one single algorithm. Alternative approach used the ``staggered'' discretisation of the grid commonly used in ``ZEUS--like'' codes along with the more geometrical Constrained Transport (CT) algorithm \\citep{evans&hawley88}. This is for example the case of \\citet{balsara&spicer99}, \\citet{toth00} and \\citeauthor{londrillo&delzanna00} (\\citeyear{londrillo&delzanna00}, \\citeyear{londrillo&delzanna04}). \\citet{gardiner&stone05a} also explored the possibility of combining the CT algorithm with the PPM scheme in the new code ATHENA. Recently, we proposed to extend the well--known MUSCL--Hancock algorithm originally designed for the Euler equation to the induction equation \\citep{teyssieretal06}. We showed that three variants of our scheme have good performances. Two are compatible with the Adaptive Mesh Refinement (AMR) algorithm implemented in RAMSES \\citep{teyssier02}. This first part was limited to the induction equation, and could only be applied to situations where the magnetic field does not affect the flow. This is enough, however, to capture the physics of fast dynamos, especially with the help of the AMR. Here we extend our approach to the full set of MHD equations and implement it in RAMSES. The plan of the paper is as follows: in section~\\ref{num method}, we present the details of the numerical algorithm. The discussion is based on our earlier work \\citep{teyssieretal06}, where the technical details of the scheme are presented. In section~\\ref{test section 1D} and \\ref{test section 2D}, we illustrate the properties of the code on standard 1D and 2D test problems. In section~\\ref{astro appli}, it is used to study a few 3D flows of astrophysical significance: the growth of the magnetorotational instability in accretion disks and the collapse of magnetized cloud cores. Finally, we summarise the properties of the code and highlight future possible developments in section~\\ref{conclusion}. ", "conclusions": "\\label{conclusion} In this paper, we have presented an extension of RAMSES to MHD. The algorithm is based on the MUSCL-Hancock approach already used in the hydrodynamic version of RAMSES \\citep{teyssier02}. The induction equation is evolved in time using the standard CT scheme \\citep{evans&hawley88}. To do so, time averaged EMFs are computed on cell edges by solving a 2D Riemann problem, as described in \\citet{londrillo&delzanna00}. Several tests are presented that illustrate the properties and robustness of the code. In particular, we show that the AMR scheme implemented in RAMSES can be crucial to describe accurately the propagation of some unusual waves peculiar to MHD like the compound waves. We also demonstrate the versatility of RAMSES by studying two problems of astrophysical significance: the development of MHD turbulence in accretion disk and the collapse of dense core in the interstellar medium. In both cases, we report results that are consistent with previous studies published in the literature. These two applications show that RAMSES is well suited to study a wide variety of problems involving MHD in astrophysics. In future studies, several improvements will now be investigated. It will be particularly useful, for example, to develop a proper 2D Riemann solver to calculate the time averaged EMFs, instead of making linear combination of 1D solvers as it is done now. Nonlinear Riemann solvers could also be implemented, like HLLC \\citep{miyoshi&kusano05} for example. Obviously, an extension to curvilinear coordinates would also be very interesting, particularly for applications involving accretion disks or galaxies. Finally, it will be necessary in some cases to go beyond the ideal MHD framework and to implement new physics like ohmic dissipation or ambipolar diffusion." }, "0607/astro-ph0607006_arXiv.txt": { "abstract": "The IC~1396N cometary globule (CG) within the large nearby HII region IC~1396 has been observed with the ACIS detector on board the $Chandra$ X-ray Observatory. We detect 117 X-ray sources, of which $\\sim 50-60$ are likely members of the young open cluster Trumpler~37 dispersed throughout the HII region, and 25 are associated with young stars formed within the globule. Infrared photometry (2MASS and $Spitzer$) shows the X-ray population is very young: 3 older Class III stars, 16 classical T Tauri stars, 6 protostars including a Class 0/I system. We infer a total T Tauri population of $\\sim 30$ stars in the globule, including the undetected population, with a star formation efficiency of $1-4\\%$. An elongated source spatial distribution with an age gradient oriented towards the exciting star is discovered in the X-ray population of IC~1396N, supporting similar findings in other cometary globules. The geometric and age distribution is consistent with the radiation driven implosion (RDI) model for triggered star formation in CGs by HII region shocks. The inferred velocity of the shock front propagating into the globule is $\\sim 0.6$~km/s. The large number of X-ray-luminous protostars in the globule suggests either an unusually high ratio of Class I/0 vs. Class II/III stars, or a non-standard IMF favoring higher mass stars by the triggering process. The $Chandra$ source associated with the luminous Class 0/I protostar IRAS~21391+5802 is one of the youngest stars ever detected in the X-ray band. We also establish for the first time that the X-ray absorption in protostars arises from the local infalling envelopes rather than ambient molecular cloud material. ", "introduction": "} It has long been recognized that star formation in molecular clouds can be triggered by ionization or shock fronts produced by nearby massive stars \\citep{Elmegreen77,Habing79,Elmegreen02}. This leads to triggered star formation at the interface between HII regions and molecular clouds and, on large scales, to the sequential formation of star clusters in molecular cloud complexes. Two mechanisms for triggered star formation have been discussed: the radiation driven implosion (RDI) model and the `collect-and-collapse' model \\citep[see review by][]{Elmegreen98}. In the RDI model, photoevaporation of the cloud outer layers induces a shock that compresses the cloud interior leading to gravitational collapse \\citep{Reipurth83,Sugitani89,Lefloch94,Gorti02}. In the `collect-and-collapse' model, the HII region compresses the cloud, triggering the formation of self-gravitating cores on timescales of $\\sim 10^5$ yr \\citep[see review by][]{Henney06}. Modern models indicate that the time-dependent relationships between dissociation, ionization, and shock fronts can be complex: gravitational collapse can form stars in a thin dense layer between the shock and ionization fronts propagating through the cloud \\citep{Hosokawa05,Hosokawa06,Zavagno06}. Triggered star formation has been reported on the edges of large HII regions in the Carina Nebula \\citep{Smith00}, Orion clouds \\citep{Stanke02,Lee05}, M~16 \\citep{Fukuda02}, M~17 \\citep{Jiang02}, 30~Doradus \\citep{Walborn02}, NGC~3603 \\citep{Moffat02}, RCW~49 \\citep{Whitney04}, W~5 \\citep{Karr05}, the Gum Nebula \\citep{Kim05}, and samples of more distant HII regions \\citep{Deharveng05}. Most of these are large star-forming complexes where the geometry and conditions are difficult to ascertain. Bright-rimmed cometary globules (CGs) are simpler structures where triggering processes may be active and thus offer an opportunity for understanding better the mechanisms at work. These are small, isolated clouds with dense cores surrounded by ionized rims facing the exciting star and tails extending in the opposite direction \\citep[e.g.][]{Loren78,Sugitani91}. CGs likely originate as dense clumps in the parental molecular clouds that have emerged after the dispersion of lower density gas by ultraviolet radiation from OB stars. The large nearby HII region IC~1396 has a rich population of bright-rimmed and cometary globules seen in silhouette against the emission nebula (Figure \\ref{spat_distrib1_fig}$a$). Over a dozen contain IRAS sources and are likely sites of star formation \\citep{Schwartz91}. Roughly 3$^\\circ$ in diameter, IC~1396 is excited by the O6.5f star HD~206267 in the Trumpler~37 cluster, which lies at the center of the Cepheus OB2 Association \\citep{Simonson68,Walborn84}. We study here globule IC~1396N \\citep[labeled E in Figure 2 of ][]{Weikard96} located $\\sim 11$~pc projected distance north of HD~206267 with the bright rim tracing the ionization front facing HD~206267 (Figure \\ref{spat_distrib1_fig}$b$). Signs of ongoing star formation in the globule include the luminous far-infrared source IRAS~21391+5802, H$_2$O maser sources, molecular outflows, HH flows, and clusters of near-infrared (NIR) embedded sources and radio-mm protostars \\citep{Sugitani89,Schwartz91,Tofani95,Slysh99, Codella01,Nisini01,Beltran02,Reipurth03}. Mass estimates for the globule range from 300 to 500~M$_{\\odot}$ and the absorption through its core is $A_V \\sim 9-10$~mag \\citep{Wilking93,Serabyn93,Nisini01,Froebrich05}. We adopt a distance of 750~pc \\citep{Matthews79} for compatibility with most recent studies, though we note that the $Hipparcos$ parallactic distance measurement is $\\sim 615$~pc \\citep{deZeeuw99}. Our study is unusual in that we use an X-ray telescope to find the young stellar population around IC~1396N. X-ray surveys are complementary to optical and infrared surveys because they trace magnetic activity (mainly plasma heated in violent magnetic reconnection flares) rather than photospheric or circumstellar disk blackbody emission of low-mass pre-main sequence (PMS) stars \\citep[see reviews by][and references therein]{Feigelson99, Feigelson06}. In regions like IC~1396, NIR surveys are overwhelmed by unrelated field stars; many old background stars may penetrate through the cloud and mimic young embedded objects (\\S \\ref{nisini_cluster_subsection}). Foreground and background Galactic stars have much less impact on X-ray studies, as magnetic activity in PMS stars is elevated $10^1-10^4$ above main-sequence levels \\citep{Preibisch05a}. The main X-ray contaminants are extragalactic objects, which are uncommon and can be identified with some reliability \\citep{Getman05b,Getman06}. X-ray surveys penetrate deeply into obscuring material ($A_V > 100$~mag) and thus are effective in detecting embedded objects \\citep{Getman05b,Grosso05}. X-ray observations also are not hampered by two problems that affect NIR and optical studies of young stellar populations in HII regions: they are not biased towards stars with protoplanetary disks, and are not subject to confusion from bright diffuse emission by heated gas and dust. The {\\em Chandra X-ray Observatory}, with its excellent high-resolution mirrors, is often effective in resolving crowded fields down to $\\simeq 0.7$\\arcsec\\/ scales \\citep{Getman05b}. The central region of IC~1396 was previously studied with the lower resolution $ROSAT$ observatory \\citep{Schulz97} but the northern part of the nebula with the IC~1396N globule was off the field. The $Chandra$ observation of IC~1396N and source list are described in \\S \\ref{observation_section}. Over 100 X-ray sources are detected, from which we identify a cluster of 25 young objects associated with the globule (\\S \\ref{globule_cluster_section}). Based on location, absorption, and infrared photometry, we distinguish embedded protostellar and more evolved T~Tauri sub-populations (\\S \\ref{ir_prop_subsection}). We compare the X-ray cluster to the NIR observations of \\citet{Nisini01} in \\S \\ref{nisini_cluster_subsection} and study X-ray properties of the cluster in \\S \\ref{x_ray_prop_section}. A particularly important intermediate-mass Class 0/I protostar is presented in \\S \\ref{srcsixtysix_section}. We end with discussion of how the rich cluster of protostars in IC~1396N fits within the larger picture of X-ray detections from protostars in other star forming regions, and with discussion of the implications for understanding triggered star formation in CGs (\\S \\ref{trigger_section}). ", "conclusions": "} We report results from a 30~ks $Chandra$ study of the stellar population in the cometary globule IC~1396N, previously recognized as a site of star formation triggered by an HII region expanding into an inhomogeneous molecular cloud. Of the 117 X-ray sources in the $17\\arcmin \\times 17\\arcmin$\\/ field, 25 appear associated with young stars formed in the globule. Although relatively bright ($10 10^{18}$ eV, we will probe for the first time the kinematical regime where $Q_{\\mathrm{sat}} > \\mu_c$. Therefore, at these energies one can expect a large modification of the charm quark total cross sections and, consequently, on the flux of prompt leptons at the Earth. Here we restrict ourselves to estimate the heavy quark production considering the present understanding of the high energy regime of the theory of strong interactions and to predict the magnitude of the saturation effects in the main quantities which are used as input in the atmospheric particle shower routines. As the theoretical predictions of the prompt leptons depend strongly on the behavior of the total cross section for heavy quark production at high energies, we believe that this partial calculation allow us to obtain a reasonable estimate of the magnitude of the saturation effects on prompt lepton flux at the Earth. Of course, more detailed studies are necessary in future in order to get precise predictions. This paper is organized as follows. In the next section we present a brief review of the heavy quark production in the color dipole picture, demonstrating the direct relation between the total heavy quark cross sections and the dipole-target cross section $\\sigma_{dip}$. The QCD dynamics is discussed in Section \\ref{section3} and the phenomenological model for $\\sigma_{dip}$ used in our calculations is presented. In the Section \\ref{section4} we present our results and our main conclusions are summarized in the Section \\ref{section5}. ", "conclusions": "\\label{section5} The determination of the prompt lepton flux is fundamental in order to establish, for instance, the background for ultra high energy neutrinos from cosmological sources. At high energies the lepton fluxes are quite sensitive to the heavy quark cross section, which implies that the choice of an appropriate theoretical framework to estimate the cross section in this energy range is fundamental. Previous calculations have considered perturbative QCD at next-to-leading order, assuming the validity of the collinear factorization and of the DGLAP dynamics in the kinematical range of very high energies. However, current accelerator data already have indicated the presence of new dynamical effects, associated to saturation physics. As the contribution of these effects increases with the energy, we can expect that it cannot be disregarded in the description of the interaction of ultra high energy cosmic rays with the atmosphere. Here we have estimated the heavy quark production in the interaction of cosmic rays in the atmosphere taking the primary cosmic ray as a proton or a photon. At ultra high energies and charm production the saturation scale stays above the semihard scale $\\mu_c^2$ and the process contains sizable contribution from the saturation regime. In particular, geometric scaling for charm production on the scaling variable $\\tau_c$ is demonstrated using small-$x$ DESY-HERA data. Within the color dipole approach and CGC formalism for dipole-target interaction, the parton saturation corrections are huge, suppressing the cross section by one order of magnitude for ultra-high energy primaries at $E_{\\gamma}\\approx 10^6$ GeV. We also predict a factor ten of suppression of the prompt lepton fluxes for a pure proton component for the primary cosmic ray, increasing if the primary component changes as appropriate for a Top-Down model. This suppression is also present in the $x_F$-distribution, which one the main inputs to the calculate the prompt lepton fluxes. The resulting predictions for the total cross section at very high energies are similar to those obtained previously by Thunman, Ingelman and Gondolo, which have been considered a lower bound. As there is a strict relation between the charm production and the prompt lepton fluxes, we believe that the resulting lepton fluxes obtained using our predictions for charm production as input of the atmospheric particle showers routines should be similar. In other words, we expect a suppression of the prompt lepton fluxes associated to the saturation physics when compared with those resulting from NLO calculations using the collinear factorization. Of course, a more detailed analyzes is necessary in order to quantify precisely this suppression. It is important to emphasize that these predictions should be considered a lower bound in the suppression, since in the CGC formalism the breaking of the factorization for heavy quark production is predicted for $Q_{\\mathrm{sat}} \\gg \\mu_c$ \\cite{gelis}, which implies a larger suppression. We also estimate the bottom production considering the dipole picture and saturation physics. In this case, we have that the saturation effects can be disregarded and the color transparency limit determines the behavior of the total cross section in the energy range of interest in this paper. However, as the $B$-hadron decays should contribute significantly for the flux of high energy $\\nu_\\tau$ neutrinos, the quantification of this cross section using the color dipole picture, which is expressed in terms of the eigenstates of interaction in QCD, is useful. We present a parameterization for the $x_F$-distribution for charm and bottom production resulting of our calculations. It can be useful for future calculations of the prompt lepton fluxes. The main conclusion of our phenomenological analyzes is that the saturation effects implies a suppression of prompt lepton fluxes. Of course, in a full calculation we should include the fragmentation of the heavy quark pairs into hadrons and their subsequent semileptonic decays, as well as, the propagation of the high energy particles through the atmosphere. We postponed these analyzes for a future publication." }, "0607/astro-ph0607283_arXiv.txt": { "abstract": "The large majority of extragalactic star cluster studies performed to date essentially use multi-colour photometry, combined with theoretical stellar synthesis models, to derive ages, masses, extinction estimates, and metallicities. M31 offers a unique laboratory for studies of globular cluster (GC) systems. In this paper, we obtain new age estimates for 91 M31 globular clusters, based on improved photometric data, updated theoretical stellar synthesis models and sophisticated new fitting methods. In particular, we used photometric measurements from the Two Micron All Sky Survey (2MASS), which, in combination with optical photometry, can partially break the well-known age-metallicity degeneracy operating at ages in excess of a few Gyr. We show robustly that previous age determinations based on photometric data were affected significantly by this age-metallicity degeneracy. Except for one cluster, the ages of our other sample GCs are all older than 1 Gyr. Their age distribution shows populations of young and intermediate-age GCs, peaking at $\\sim 3$ and 8 Gyr respectively, as well as the ``usual'' complement of well-known old GCs, i.e., GCs of similar age as the majority of the Galactic GCs. Our results also show that although there is significant scatter in metallicity at any age, there is a noticeable lack of young metal-poor and old metal-rich GCs, which might be indicative of an underlying age-metallicity relationship among the M31 GC population. ", "introduction": "\\label{Introduction.sec} M31 is the nearest large spiral galaxy, at a distance of about 780 pc \\citep{sg98, mac01}. It contains more than 337 confirmed globular clusters (GCs) and about 688 GC candidates \\citep{gall04}, i.e., significantly more than in our own Galaxy. Thus, M31 provides an excellent opportunity to study the properties of a large sample of GCs. From the observational evidence collected thus far \\citep[see, e.g.,][]{rich05}, the M31 GCs and their Galactic counterparts reveal some striking similarities \\citep{ffp94,dj97,bhh02}. For example, both GC systems seem to have similar mass-to-light ratios, structural parameters, and velocity dispersion -- luminosity relations \\citep[see also][]{degrijs05}. Studies of GCs in M31 can not only throw light on lots of questions about the formation, evolution and properties of M31 itself, including its mass, dynamics and chemical composition, they can also improve our understanding of the formation and structure of galaxies in general \\citep{batt80}. In addition, GCs can provide us with good samples of Population II stars characterised by homogeneous abundances and histories, and with unique stellar dynamical conditions for our study \\citep{barmby01}. Therefore, GCs are so important that they are considered as the fossils of the earliest stages of galaxy formation \\citep{bh00}. However, at the distance of M31, construction of colour-magnitude diagrams (CMDs) below the main sequence turn-off, the most reliable method for age determinations of stellar populations, is extremely challenging for current state-of-the-art instrumentation. Here, one suffers from the dual effects of crowding and the intrinsic faintness of the cluster stars \\citep{rich96,rich05,stephens01,bb04}, although we note that \\citet{brown04} presented a CMD down to the turn-off for a $\\sim 10$ Gyr-old M31 GC. Since the pioneering work of \\citet{Tinsley68,Tinsley72} and \\citet{ssb73}, evolutionary population synthesis modeling has become a powerful tool to interpret integrated spectrophotometric observations of galaxies as well as their components \\citep[see, e.g.,][]{Anders04}. Comprehensive compilations of relevant current model sets, such as e.g. developed by \\citet[][henceforth BC93, BC96]{bc93,bc96}, \\citet{Leitherer95}, and \\citet{frv97}, were provided by \\citet{lei96} and \\citet{ken98}. The evolution of star clusters is usually modeled based on the ``simple stellar population'' (SSP) approximation\\footnote{An SSP is defined as a single generation of coeval stars formed from the same progenitor molecular cloud (thus implying a single metallicity), and governed by a given initial mass function (IMF).}, which is a highly robust approximation for old GCs in particular \\cite[e.g.][]{bh01}. \\citet{Ma01,Ma02a,Ma02b,Ma02c} and \\citet{jiang03} estimated the ages of, respectively, 180 star clusters in M33 and 172 GC candidates in M31 by comparing the SSP synthesis models of BC96 with the clusters' integrated photometric measurements in the Beijing-Arizona-Taiwan-Connecticut (BATC) photometric system. \\citet{Ma06a} also determined the ages and metallicities of 33 M31 GCs and candidates using the updated method and updated SSP synthesis models of \\citet[henceforth BC03]{bc03}, and \\citet{Ma06b} estimated the age and reddening value of the M31 GC 037-B327 based on photometric measurements in a large number of broad-band passbands from the optical to the near-infrared. From an observational point of view, the study of M31 GCs is complicated, since in most cases we only have access to their integrated spectra and photometry, and cannot study the resolved stellar population. Therefore, we can only obtain the key physical parameters, such as the age and metallicity, by analysis of the integrated spectra or photometry. However, a large body of evidence suggests that there is a strong age-metallicity degeneracy if only optical photometry is used \\citep{ar96,wor94,kaviraj06}. A very useful method to break this degeneracy is through the application of particular spectral diagnostics based on the occurrence of individual stellar absorption-line features \\citep[e.g.,][]{fab73,rose84,rose85,dtt89,worth94,jw95,va99,bc03}. At the same time, observational GC spectral energy distributions (SEDs) are affected by reddening, an effect that is also difficult to separate from the combined effects of age and metallicity \\citep{calz97,vazde97,orig99}. However, if the metallicity and reddening are derived accurately (and, ideally, independently), these degeneracies are largely (if not entirely) reduced, and ages can then also be estimated accurately based on a comparison of multi-colour photometry spanning a significant wavelength range \\citep{degrijs03b,Anders04} with theoretical stellar population synthesis models. In this paper, we present new age estimates for 91 GCs of the \\citet{jiang03} sample, based on their 13 intermediate-band photometry in the BATC system, combined with additional broad-band optical and 2MASS near-infrared photometry, and on the updated SSP synthesis models of BC03. Section \\ref{Photometry.sec} describes the intermediate-band, broad-band and 2MASS photometry of our GC sample. In Section \\ref{Metal&Red.sec}, we describe the metallicity and reddening data of the sample GCs; Section \\ref{age.sec} includes a description of the SSP models used, and of our method to estimate the ages of the sample GCs. The main results and a discussion are also presented in this section. We summarise and conclude the paper in Section \\ref{Conclusions.sec}. ", "conclusions": "\\label{Conclusions.sec} In this paper, we accurately re-determined the ages of 91 M31 GCs, based on improved data, updated theoretical stellar synthesis models and sophisticated fitting methods. In particular, we used photometric measurements of the 2MASS, which can partially break the age-metallicity degeneracy, in combination with optical photometry. We showed robustly that previous age determinations based on photometric data were affected significantly by this age-metallicity degeneracy. Except for one cluster, the ages of our other sample GCs are all older than 1 Gyr. Their age distribution shows populations of young and intermediate-age GCs, peaking at $\\sim 3$ Gyr and $\\sim 8$ Gyr respectively, as well as the ``usual'' complement of well-known old GCs, of similar age as the majority of the Galactic GCs. The young-age peak at $\\sim 3$ Gyr we detect in this paper has not been discussed before; in view of the small age uncertainties at young ages, it is unlikely that this population represents the $\\la 2$ Gyr-old population of BLCCs of \\citet{fp05}. Instead, we argue that there may have been an additional violent star-forming event that triggered the formation of a GC subpopulation in the disc of M31 some 3 Gyr ago. The distributions of the ages of both the GCs and the field stars, combined with the existence of a metal-poor GC population in M31 obeying thin-disc kinematics, poses serious problems for our understanding of the formation and evolution of the galaxy's disc. While the tightly constrained kinematics argue for a relatively quiescent thin-disc evolution since early times, the formation of massive star clusters requires violent conditions (such as galaxy mergers) rather than {\\it in situ} formation. Our results also show that although there is significant scatter in metallicity at any age, there is a noticeable lack of young metal-poor and old metal-rich GCs, which might be indicative of an underlying age-metallicity relationship among the M31 GC population." }, "0607/astro-ph0607626_arXiv.txt": { "abstract": "\\noindent We discuss the physics of backreaction-driven accelerated expansion. Using the exact equations for the behaviour of averages in dust universes, we explain how large-scale smoothness does not imply that the effect of inhomogeneity and anisotropy on the expansion rate is small. We demonstrate with an analytical toy model how gravitational collapse can lead to acceleration. We find that the conjecture of the accelerated expansion being due to structure formation is in agreement with the general observational picture of structures in the universe, and more quantitative work is needed to make a detailed comparison. ", "introduction": "\\label{sec:intro} \\paragraph{Evidence for acceleration.} There is a large body of observational evidence supporting the claim that the expansion of the universe has accelerated in the recent past, and may be accelerating today. This conclusion has been bolstered by the verification of the prediction of the location of the baryon acoustic peak in the matter power spectrum \\cite{BAO}, in a convincing demonstration of concordance. In addition, the worrisome feature that nearby and distant populations of type Ia supernovae used to have different absolute magnitudes and were both individually consistent with deceleration has disappeared with new and better data \\cite{Padmanabhan} (though see \\cite{Shapiro:2005, Elgaroy:2006}). The \\LCDM model, where the acceleration is driven by vacuum energy (or the cosmological constant, which is the equivalent modification of gravity) agrees well with most observations, with the notable exception of the low CMB multipoles \\cite{CMBmaps, Hansen:2004, Bielewicz:2005, dipolesyst, Copi:2006} (it has also been argued that cluster observations support a non-accelerating universe \\cite{Blanchard}). However, given the lack of theoretical understanding about the parameters of the \\LCDM model (notably the vacuum energy density), it is a phenomenological fit rather than a well-founded theory, and its success does not rule out the possibility that quite a different model can also be a good fit to the data. (The values obtained for the parameters of a cosmological model by fitting to observations should not be mistaken for measurements, as model selection studies show; see e.g. \\cite{Mukherjee:2005, Shapiro:2005, Elgaroy:2006}.) In particular, while the observation that there is accelerating expansion seems robust, the nature of the acceleration is not well constrained. In the \\LCDM model, the transition to acceleration is gradual, but a rapid transition is not ruled out \\cite{trans, Ichikawa:2006}. In fact, from the SNIa data it is difficult to say anything beyond that the universe has accelerated in the recent past, even whether the expansion is still accelerating \\cite{Shapiro:2005, Elgaroy:2006, Gong:2006}. Keeping to the assumption that the universe is completely homogeneous and isotropic, any explanation of the acceleration has to involve either a medium with negative pressure or modified gravity. Such models in general, and the \\LCDM model in particular, suffer from {\\it the coincidence problem}: why does the acceleration happen around a redshift of unity, at around 10 billion years? In other words, why are we seeing a very particular phase in the evolution of the universe, when the inferred energy density of the source driving the acceleration has recently become equal to the energy density of matter? The clearest qualitative change in the late-time universe is the formation of non-linear structures. It therefore seems a natural possibility that the observed deviation from the prediction of homogeneous and isotropic cosmological models with normal matter and gravity could be related to the known breakdown of the assumption that the universe is homogeneous and isotropic (rather than to a speculated failure in the description of the matter content or the theory of gravity). \\paragraph{The inhomogeneous universe.} One possible avenue is trying to explain the observations without having any accelerated expansion. Cosmological information is borne to us by light along null geodesics (apart from information carried by neutrinos and cosmic rays). The standard analysis of light propagation assumes that the universe is perturbatively near a homogeneous and isotropic Friedmann--Robertson--Walker (FRW) model, which is manifestly not true on scales smaller than 70-100 \\mpc \\cite{Hogg:2004, Pietronero, morphology}. It is therefore possible that the propagation of light would be affected by the inhomogeneities and/or anisotropies in a way that looks like acceleration when interpreted in the context of an FRW model. Studies of the Lema\\^{\\i}tre--Tolman--Bondi (LTB) model \\cite{LTB} (see \\cite{Krasinski:1997} for a review), the spherically symmetric dust solution of the Einstein equation, have demonstrated that the effect of inhomogeneity on the luminosity distance can mimic acceleration \\cite{LTBgeo, Biswas:2006} (see \\cite{Biswas:2006} for more references). Even though spherical symmetry is a questionable assumption for the entire universe, it could be a good first approximation for the local region. In any case, one would expect a qualitatively similar effect to be present also in more realistic and less symmetric spacetimes \\cite{Bolejko} -- arguably, the effect of clumpiness could even be stronger when there is less symmetry. The effect of inhomogeneity and anisotropy on the luminosity distance has also been studied in perturbed FRW models \\cite{Barausse:2005, Bonvin}. An explanation of the apparent acceleration in terms of inhomogeneity and/or anisotropy could solve the coincidence problem, since inhomogeneity and anisotropy become important only in the late-time universe. However, inhomogeneity and anisotropy affect different observations in different ways, and it would require an odd coincidence for all the various indicators of expansion rate (SNIa luminosity distances, the cosmological microwave background (CMB) anisotropies, large scale structure (LSS), and so on) to be affected in a way that would be consistently interpreted as acceleration when fitting to a FRW model. One proposed possibility is that we live in an underdense region, a 'Hubble bubble' \\cite{hubbub} (for more references, see \\cite{Biswas:2006}). In this proposal, the local matter density today is $\\Omn\\approx$ 0.15-0.35, as indicated by local observations \\cite{Peebles:2004}, while the global value is $\\Om=1$. The SNIa luminosity distances as well as the difference between the local and global values of the expansion rate could be explained in terms of inhomogeneity, while a global model with no acceleration can fit most other observations, including the CMB and LSS \\cite{Sarkar} (though it is not supported by studies of the local and global expansion rate \\cite{Sandage:2006}). However, it would be difficult to explain the baryon acoustic peak \\cite{Blanchard:2005}: one would have to appeal to inhomogeneities (or features in the primordial power spectrum) to supply a pattern that by coincidence happens to fit the expectations of an accelerating FRW model. One way to phrase the issue is that cosmological observations involve a larger number of a priori independent parameters than the \\LCDM model. Therefore the \\LCDM model implies relations between observationally independent parameters. (For an early discussion of cosmological observations in an inhomogeneous and anisotropic spacetime which makes the issue transparent, see \\cite{Kristian:1966}.) It is not surprising that models with more degrees of freedom, such as the LTB model which involves two arbitrary functions, could fit the data as well as FRW models. However, it would be an unlikely coincidence for them to also produce the same relations between observables as FRW models. This is generally true for any models where the explanation of the luminosity distances is decoupled from the explanation of the low matter density, baryon acoustic peak and so on (such as mixing of photons with axions \\cite{Csaki:2001} or with the gauge bosons of a new $U(1)$ gauge group \\cite{Evslin:2005}). \\paragraph{The fitting problem.} While the FRW scale factor has been very successful in fitting observations, it is difficult to understand the matter content implied by the FRW equations which relate the scale factor to the energy--momentum tensor. Note that the fact that the mean properties of the universe are well described by an overall scale factor does not imply the stronger statement that the scale factor evolves according to the FRW equations, since the universe is not completely homogeneous and isotropic. The idea that the average behaviour of inhomogeneous and/or anisotropic spacetimes is in general different from the behaviour of homogeneous and isotropic spacetimes goes back to at least 1963 \\cite{Shirokov:1963}. The first comprehensive discussion was given in 1983 by George Ellis \\cite{fitting}, who called the task of finding the smooth metric which best fits the real clumpy universe {\\it the fitting problem}. The influence of inhomogeneity and/or anisotropy on the average behaviour is also known as backreaction \\cite{Kasai, Buchert:1995, Woodard, Ehlers:1996, Unruh:1998, Takada:1999, Buchert:1999, Sicka, Taruya:1999, Buchert:2000, Kerscher:2000, Carfora, Buchert:2001, Tatekawa:2001, Geshnizjani:2002, Schwarz, Buchert:2002, Brandenberger:2002, Geshnizjani:2003, Rasanen:2003, Buchert:2003, Hosoya:2004, Rasanen:2004, Kolb:2004a, Kolb:2004b, Barausse:2005, Kolb:2005a, Flanagan:2005, Geshnizjani:2005, Hirata:2005, Notari:2005, Rasanen:2005, Siegel:2005, Martineau:2005a, Tsamis:2005, Kolb:2005b, global, Ishibashi:2005, Losic, Martineau:2005b, Kolb:2005c, Kolb:2005d, Chuang:2005, Kasai:2006, Parry:2006, Kai:2006, Paranjape:2006, Rasanen:2006, Buchert:2006}; see \\cite{Krasinski:1997, Rasanen:2003} for further references, in particular early ones, and \\cite{Ellis:2005} for an overview. The idea that perturbations with wavelengths smaller than the Hubble radius could lead to acceleration for the scale factor (as opposed to merely mimicking the appearance of acceleration via changing null geodesics) was studied in the context of linear perturbation theory in \\cite{Rasanen:2003} (the possibility had been earlier touched upon in \\cite{Buchert:2000, Tatekawa:2001}; see also \\cite{Schwarz}). The calculation was then extended to second order \\cite{Barausse:2005, Kolb:2004a}, and it was suggested that linear perturbations with wavelengths much larger than the Hubble radius could lead to acceleration \\cite{Barausse:2005, Kolb:2004a, Kolb:2004b, Kolb:2005a}. It is now agreed that this is not possible \\cite{Flanagan:2005, Geshnizjani:2005, Hirata:2005, Rasanen:2005, Kolb:2005b}\\footnote{Even though super-Hubble perturbations do not contribute to acceleration in a dust universe, they could lead to deceleration during inflation driven by a scalar field or a cosmological constant, and super-Hubble scalar field perturbations left over from inflation could still be important today \\cite{Woodard, Unruh:1998, Geshnizjani:2002, Brandenberger:2002, Geshnizjani:2003, Martineau:2005a, Tsamis:2005, Losic, Martineau:2005b, Kolb:2005d}.}. As for perturbations smaller than the Hubble radius, there is no acceleration to at least second order in perturbation theory \\cite{Rasanen:2003, Kolb:2004a, Kolb:2005b, Kasai:2006, Parry:2006}. If inhomogeneity and anisotropy are to explain the observed acceleration, the only possibility is via non-linear sub-Hubble perturbations, that is, the process of structure formation, as proposed in \\cite{Schwarz, Rasanen:2003}. It has been analytically shown in the LTB toy model how backreaction of non-linear perturbations can modify the Hubble law \\cite{Rasanen:2004}, and acceleration has also been numerically demonstrated in the LTB model \\cite{Chuang:2005, Kai:2006, Paranjape:2006}, but the physical meaning of inhomogeneity- and anisotropy-driven acceleration and the connection to structure formation has been unclear. We will discuss the relation between homogeneity and isotropy, the overall scale factor, the FRW metric and the equations which describe the average expansion of the universe. We will then look at an exact toy model of structure formation, explain the physics of acceleration driven by inhomogeneity and anisotropy, and note that structure formation involves a preferred time near the observed acceleration era. In particular, we will clarify two apparent paradoxes of backreaction-driven acceleration: how the average expansion of a manifold can accelerate even though the local expansion rate decelerates everywhere, and how collapse implies acceleration. These issues were earlier discussed in the brief essay \\cite{Rasanen:2006}. In \\sec{sec:smoo} we discuss the assumptions underlying FRW models, and go through the derivation of the Buchert equations which describe the average behaviour of an inhomogeneous and/or anisotropic dust spacetime. In \\sec{sec:acc} we consider an exact toy model where gravitational collapse produces acceleration, discuss how this mechanism may operate in the real universe, and compare this picture with some observations and simulations of structures in the universe. In \\sec{sec:con} we summarise the situation with regard to the conjecture that the observed acceleration is due to backreaction. ", "conclusions": "\\label{sec:con} \\paragraph{Acceleration and inhomogeneity/anisotropy.} The observational evidence for the acceleration of the universe is usually interpreted in the framework of linearly perturbed Friedmann--Robertson--Walker (FRW) models, which describe a universe that is everywhere almost homogeneous and isotropic. In the context of such models, a medium with negative pressure or modified gravity is needed to explain the observations. This leads to the coincidence problem: why has the exotic matter or strange gravity become important only recently? The most significant qualitative change in the universe around the era where acceleration has been observed is the formation of non-linear structures, so it seems a natural possibility that the observed deviation from the general relativistic prediction of the homogeneous and isotropic cosmological models with normal matter could be related to the breakdown of homogeneity and isotropy. The issue of cosmological homogeneity and isotropy has been extensively discussed over the years by George Ellis and collaborators, notably in the context of the observational program of cosmology \\cite{Ellis:1975, obscos, hom, Matravers:1995}. One of the issues they have highlighted is that averaging and applying the field equations do not commute: in other words, the average properties of an inhomogeneous and/or anisotropic spacetime do not satisfy the Einstein equation. The task of finding the model that best describes the average behaviour of the inhomogeneous universe has been termed the fitting problem. The relativistic equations which describe the behaviour of average quantities in an inhomogeneous and/or anisotropic, but irrotational, ideal fluid universe have been derived by Thomas Buchert \\cite{Buchert:1999, Buchert:2001}. The Buchert equations show that it is possible for inhomogeneities and/or anisotropies to lead to accelerating average expansion in a dust universe, even though the local acceleration decelerates everywhere. They also show that the fraction of space occupied by non-linear regions is the determining quantity, not the size of the individual regions. Even when the average properties of space can be described in terms of an overall scale factor, the evolution of the scale factor does not necessarily follow the FRW equations. The possibility that inhomogeneities and/or anisotropies could lead to acceleration was studied in the context of linear perturbation theory in \\cite{Rasanen:2003}, and it was suggested that acceleration could be due to perturbations which have entered the non-linear regime but haven't yet stabilised. The possibility of acceleration via backreaction has been numerically verified \\cite{Chuang:2005, Kai:2006, Paranjape:2006}. However, the physics of how structure formation leads to acceleration and the question of why acceleration begins much later than structure formation have been unclear. We have now clarified these issues, which turn out to be intimately associated with the process of gravitational collapse. With a simple toy model, we have explicitly shown how overdense regions can first slow down the expansion, which then accelerates as these regions shrink and their contribution to the expansion rate decreases rapidly as they collapse. We have also noted that the matter-radiation equality scale imprinted on the dark matter power spectrum leads to a preferred time for structure formation that is near the observed acceleration era. The typical size of collapsing structures relative to the visual horizon grows rapidly at the start of structure formation, but then slows down, saturating around 10-100 billion years. A naive look at observations and simulations of structure in the universe shows that the degree of inhomogeneity required for backreaction to yield acceleration is plausible. \\paragraph{The backreaction conjecture.} The backreaction conjecture for the acceleration is simple. According to the Buchert equations, large variance of the expansion rate leads to acceleration. The physical interpretation is simply that the relative volume taken up by the regions of space which are expanding faster will come to dominate over the slower expanding regions, so the average expansion rate will rise. Collapsing regions, i.e. regions with a negative expansion rate, give a large contribution to the variance, since they contribute positively to the mean square but negatively to the square of the mean. Structure formation involves gravitational collapse, and the size of the collapsing regions is largest at late times when acceleration has been observed. Such an explanation keeps the phenomenological successes of the FRW scale factor in fitting the observations, while avoiding the failure of the FRW equations, which has required the introduction of a medium with negative pressure or modified gravity. This is in contrast to models which propose explaining the observations by the effect of inhomogeneities on the propagation of light without having accelerated expansion \\cite{LTBgeo, Biswas:2006}, where the success of the ansatz that one needs only to look at an overall scale factor is accidental. In the context of FRW models, there have been attempts to connect the late-time acceleration to inflation (via making the same scalar field responsible for both), the era of matter-radiation equality (via a tracker field which reacts to the change in the background equation of state) and dark matter (via unified dark matter and dark energy). Backreaction involves a subtle link to all these issues. Inflation determines the initial amplitude of the density perturbations, matter-radiation equality starts the clock for structure formation, and the nature of dark matter determines the processed form of the power spectrum and the time of formation of the first generation of structures. With many previously unclear conceptual and qualitative issues settled, the task is now to build a realistic model and make quantitative estimates that can be compared with observations. The relevant aspects of observations and simulations should also be understood better. On the basis of general considerations we can already state that we should have $H t<1$, and that there should be observable amounts of spatial curvature (assuming that vorticity is negligible and that matter can be treated as dust) \\cite{Rasanen:2005}. There may be a slowdown period preceding the acceleration, and the expansion may oscillate between deceleration and acceleration, but these issues have to be worked out in the context of a detailed model. Note that there are no new fundamental parameters to adjust, and any unknowns are due to existing uncertainties about the power spectrum, the modelling of structure formation and so on: the backreaction conjecture is eminently falsifiable. Backreaction analysis simply entails doing the usually implicit averaging in cosmology in a way that is both mathematically consistent and takes into account the structures that are known to be present in the universe, as has been advocated over the years in the context of the program of observational cosmology and related work. Backreaction offers an elegant possible explanation for late-time acceleration. Whether or not this possibility turns out to be realised, the effect of structure formation on the expansion rate should be carefully evaluated to solve the fitting problem and complete the program of determining the right equations for describing the overall behaviour of the universe. \\ack I thank Thomas Buchert for helpful correspondence, support as well as comments on the manuscript, several colleagues including John Dubinski and Constantinos Skordis for clarifying discussions, and numerous other people for stimulating criticism. I am also grateful to Michael Rauch for providing and explaining the simulation data of \\cite{Rauch:2005}. This paper is dedicated to the victims of operation ``Summer Rain'' and operation ``Just Reward''.\\\\ \\appendix \\setcounter{section}{1}" }, "0607/astro-ph0607410_arXiv.txt": { "abstract": "{We present a multicomponent model to explain the features of the pulsed emission and spectrum of the Crab Pulsar, on the basis of X and $\\gamma$-ray observations obtained with BeppoSAX, INTEGRAL and CGRO. This model explains the evolution of the pulse shape and of the phase-resolved spectra, ranging from the optical/UV to the GeV energy band, on the assumption that the observed emission is due to more components. The first component, $C_O$, is assumed to have the pulsed double-peaked profile observed at the optical frequencies, while the second component, $C_X$, is dominant in the interpeak and second peak phase regions. The spectra of these components are modelled with log-parabolic laws and their spectral energy distributions have peak energies at 12.2 and 178 keV, respectively. To explain the properties of the pulsed emission in the MeV-GeV band, we introduce two more components, $C_{O\\gamma}$ and $C_{X\\gamma}$, with phase distributions similar to those of $C_O$ and $C_X$ and log-parabolic spectra with the same curvature but peak energies at about 300 MeV and 2 GeV. This multicomponent model is able to reproduce both the broadband phase-resolved spectral behaviour and the changes of the pulse shape with energy. We also propose some possible physical interpretations in which $C_O$ and $C_X$ are emitted by secondary pairs via a synchrotron mechanism while $C_{O\\gamma}$ and $C_{X\\gamma}$ can originate either from Compton scattered or primary curvature photons.} ", "introduction": "The origin of the high energy emission of rotation-powered pulsars is still an unsolved problem. One of the main difficulties is related to the description of the phase and energy distributions of the pulsed signal, which depends on both a physical and geometrical modelling of the magnetosphere. It is still unclear, for example, whether electrons (or positrons) are accelerated and radiate streaming out the polar cap regions (Ruderman \\& Sutherland 1975, Salvati \\& Massaro 1978, Sturner \\& Dermer 1994, Daugherty \\& Harding 1994, 1996; Muslimov \\& Harding 2003) or in the outer gaps (Cheng, Ho \\& Ruderman 1986a,b; Chiang \\& Romani 1994, Romani \\& Yadigaroglu 1995, Cheng, Ruderman \\& Zhang 2000; Zhang \\& Cheng 2002, Hirotani, Harding \\& Shibata 2003). Another important problem concerns how quantum processes, like magnetic pair production (\\cite{erber66}) and photon splitting (\\cite{adler71}), modify the high energy $\\gamma$-ray spectrum. The Crab pulsar (PSR B0531+21) has been the best studied object of this class since its discovery (Staelin \\& Reifenstein 1968) and the amount of data collected is rich enough to search for a detailed physical picture of its emission properties. It is well known that the pulse shape of Crab has a characteristic double peak structure, with a phase separation of 0.4, detected from the radio band to $\\gamma$ rays and changing with energy. A very remarkable feature is that the so-called second peak, hereafter P2, in the X and soft $\\gamma$-ray ranges becomes progressively higher than the first peak (P1). A similar increase is also evident in the emission between the peaks, usually named the Interpeak region (Ip) or bridge (see Fig. 1). Above about 10 MeV, P1 is again the dominant feature. A satisfactory explanation for these changes has not been found so far. On the basis of high quality BeppoSAX data, covering a wide energy range from 0.1 to about 300 keV, we proposed a two component model (Massaro et al. 2000, hereafter MCLM) to interpret this behaviour. In the same paper we studied in detail the energy spectrum of the core of P1 (corresponding to a phase interval having a width of only 0.027 around the maximum) which shows a continuous steepening at high energies. We found that in the energy range 0.1--300 keV this spectral distribution is well represented by a parabolic law in a double logarithmic plot (hereafter log-parabola) with a rather mild curvature. The extrapolation of this model in the $\\gamma$-ray range, however, fails to reproduce the data and a more complex modelling is required. \\begin{figure*} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[angle=90]{MC_f1.ps}} \\caption{The pulse profile and the phase dependent photon index of the Crab pulsar observed with the four Narrow Field Instruments of BeppoSAX and with ISGRI-INTEGRAL experiment. The respective energy ranges are indicated in the upper panels. All the profiles are normalized to unity at the maximum of P1. Note the change of relative intensity of P2 and Ip with respect to P1 and the increase of the photon indices. } \\label{fig1} \\end{figure*} In a subsequent paper Kuiper et al. (2001) introduced three components to describe the spectrum up to the $\\gamma$-ray data obtained with COMPTEL and EGRET on board ComptonGRO. These authors based the analysis mainly on finding best fits of the spectral distributions in rather narrow phase intervals and found that the spectral variation of the pulsed emission with phase can be modelled by two log-parabolic components with a relative normalization changing with phase. A further third power law component, having a photon index equal to 2.07, was necessary for the emission of P1 and P2 to to match the EGRET data at energies higher than about 10 MeV. \\\\ In this paper we develop a model able to describe the phase and spectral distributions of the emission over a frequency interval from the optical frequencies to GeV range. This model is an extension of that presented in MCLM and it based on the results of a new detailed data analysis of many BeppoSAX observations which includes PDS data from March 1999 to April 2001 not considered by MCLM. The timing accuracy has been verified using detailed pulse profiles obtained from RXTE archive data. Moreover, to extend the energy range, we analysed several more recent observations performed with the IBIS-ISGRI experiment on board the INTEGRAL satellite and considered the results of Kuiper et al. (2001) on the COMPTEL and EGRET observations up to a few GeV. The main goal of our work is the definiton of a scenario that can be used to develop more detailed physical models of the Crab pulsar high-energy emission. ", "conclusions": "The development of a detailed physical model for the spectral and phase distributions of the broad-band emission from the Crab pulsar is a difficult problem. It requires a precise geometrical definition of the regions inside the magnetosphere where the observed radiation originates, and the knowledge of parameters like the orientation angles between the magnetic axis and the line of sight to the spin direction. Several models have appeared in the literature based on either polar cap or outer gap geometries. Usually, these models are focused on reproducing either the total spectrum or the phase profile and generally they are not fully satisfactory in explaining the complex observational picture. Moreover, the possibility that the observed features of the pulsed signal can arise from the superposition of two or more distinct components is not taken into account. We followed another approach and searched for a possible interpretation of the Crab signal based on the superposition of two or more components that provides a consistent description of the spectral and phase distributions. MCLM showed that a possible explanation of the energy dependence of the pulse shape of Crab in the soft to hard X-rays is that we are observing two emission components with different phase and energy distributions. In that paper we introduced an empiric model, based on a collection of BeppoSAX observations, covering the energy range from 0.1 to about 300 keV, from which it was possible to estimate some properties of the main components. Moreover, we showed that the X-ray spectrum of P1 presents a mild curvature well fitted by a log-parabola. At energies higher than 30 MeV, however, this two component model fails to represent both pulse profiles and spectra, as observed by EGRET-CGRO (Fierro et al. 1998). To take into account $\\gamma$-ray data, Kuiper et al. (2001) proposed the existence of three components, two of them having log-parabolic spectra, while the third one with a power law spectrum and a phase modulated intensity which reaches the highest level in correspondence of the two peaks and is almost absent in the Ip region. According to our point of view, however, this hypothesis is not fully consistent with the data, because the flux increases of Ip and P2 are very similar, suggesting that their X-ray emission is dominated by the same physical mechanism. Moreover, the parameters describing the curvatures of their two log-parabolic distributions are very different, while the analysis of X-ray spectra reported in Sect. 6.2 indicates that the curvature is rather stable with phase at an intermediate value between them. Finally, the power-law component does not match well the data in some phase intervals at $\\gamma$-ray energies where its contribution is dominant. To achieve a more consistent scenario, in the present paper we propose a model of the pulsed emission from Crab based on four components. It is properly a two double-component model because each component pair has similar phase distributions and spectra shifted in energy. At present, our model gives an empiric description of the broad-band properties and its validity is based on the limited number of assumptions and on the resulting capability to obtain a consistent description of the observations outside the energy ranges used to evaluate the parameters. For instance, the model extrapolates well the pulse profiles in the MeV range (see Sect.6.4) and the change with energy of the P2/P1 and Ip/P1 ratios. We stress that Crab is not the only pulsar that shows a behaviour that can be interpreted by a multicomponent model. Harding et al. (2002) proposed that the X-ray pulsed emission from Vela originates from two non-thermal components, one coincindent in phase with the $\\gamma$-ray pulse profile and the other one with the optical. The physical processes at the origin of these components, and the location in the magnetosphere where they occur, must be further investigated. The development of a detailed physical model of the high energy emission in the Crab magnetosphere is beyond the aim of the present paper, however some general indications on it can be derived from our conclusions. The same phase distribution assumed for each pair of components, such as $C_O$ and $C_{O\\gamma}$, suggests that their angular pattern and emission sites must be coincident or very close, otherwise the aberration effects would modify the pulse shapes. However, this is not necessarily true because on the trailing last open field line, aberration and propagation time effects would cancel to form a caustic, as shown by Dyks \\& Rudak (2003). \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{MC_f15.ps}} \\caption{The ratio between the fluxes of P2 and P1 phase regions (P1: -0.06--0.04; P2: 0.32--0.43), compared to the predictions of the model. The data points come from various experiments (Kuiper et al. 2001). The various extrapolations above 1 GeV correspond to different values of the cut-off energy of the $C_{O\\gamma}$ spectrum: 15 GeV (red), 13 GeV (green), 11 GeV (blue) and 9 GeV (violet) - from bottom to top.} \\label{rappP2P1} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{MC_f16.ps}} \\caption{The ratio between the fluxes of Ip and P1 phase regions (P1: -0.06--0.04; Ip: 0.14--0.25), compared to the predictions of the model. The data points come from various experiments (Kuiper et al. 2001).} \\label{rappIPP1} \\end{figure} According to a widely accepted scenario, primary electrons accelerated in a magnetospheric gap emit high energy photons which produce $e^{\\pm}$ pairs against the magnetic field. These secondary particles emit synchrotron radiation in the optical to MeV energy range. A first possibility to be considered is that $C_O$ and $C_X$ are synchrotron radiation from two different places whereas the corresponding high energy components are inverse Compton upscattered photons. A very early inverse Compton model for Crab was proposed by Zheleznyakov \\& Shaposhnikov (1972), when the $\\gamma$-ray emission properties were known very poorly, and another model was applied to the Vela pulsar by Morini (1983). A more recent development is that of Cheng and Wei (1995). Their model is based on the outer gap geometry and assumes that optical to hard X-ray photons are emitted by secondary $e^{\\pm}$ pairs created outside the accelerating region by high-energy primary curvature photons. $\\gamma$-rays are then emitted via a synchrotron self-Compton (SSC) process. In this case we expect that Compton scattering of hard X-ray photons occurs mainly in the Klein-Nishina regime, and therefore the observed high energy cut-off gives an estimate of the maximum energy of electrons. The fact that it is in the GeV range could be consistent with their origin from magnetic pair production. There is, however, another possibility for the origin of the high energy components. In a polar cap scenario these photons could be those emitted by primary electrons via curvature radiation and not absorbed by magnetic pair production throughout the magnetosphere. The attenuation length for this process can be approximated by (Erber 1966): \\begin{equation} x(B_{\\perp}, E_{\\gamma}) = 2.3 \\times 10^{-8} \\frac{B_{cr}}{B_{\\perp}}~ \\exp\\left[\\frac{4 B_{cr}} {3 B_{\\perp}} \\frac{2 m c^2}{E_{\\gamma}} \\right] \\mbox{\\ \\ \\ \\ cm\\,,} \\end{equation} where $B_{\\perp}$ is the transverse magnetic field seen by a photon of energy $E_{\\gamma}$, and $B_{cr}$=4.414$\\times$10$^{13}$ is the quantum critical field. It is easy to verify that a mean free path of the order of 10$^4$ cm for 5 GeV photons is obtained in a transverse field $B_{\\perp}\\simeq$ 7.5$\\times$10$^8$ G and that it increases very rapidly even for decreasing $E_{\\gamma}$. This cut-off in the curvature $\\gamma$-ray spectra at GeV energies has been verified since early numerical calculations, including an accurate evaluation of the absorption coefficient, for a polar cap acceleration (Salvati \\& Massaro 1978, Massaro \\& Salvati 1979). Although our multi-component model gives a consistent picture of the pulsed emission from Crab it is not yet completely determined by observational data. An important test will be the study of the spectra and pulse profiles at energies higher than a few GeV: in particular, we expect that above about 5 GeV, P2 would be the dominant feature, while a good measure of the flux in the Ip region will be very useful to draw the phase structure of $C_{X\\gamma}$. High quality data in this energy range will be obtained with the LAT telescope on board the GLAST mission to be operative next year. Another interesting test of the model could be obtained from phase resolved polarization measures in the X-ray range. We know that the optical linear polarisation of the Ip differs from those of P1 and P2 both in strength and direction (Smith et al. 1988; Kanbach et al. 2003). According to our model the polarization of the P2 X-ray emission must be more similar to that of Ip, because of the higher contribution of $C_X$. New generation high sensitivity polarimetry for X-ray astronomy, such as that proposed by Costa et al. (2001), could be very useful. Finally, the observation of other young spin powered pulsars, if their high energy emission is similar to Crab, could also be very useful because we expect to observe them from different directions and therefore to see other pulse shapes, depending on the various combinations of the two components." }, "0607/astro-ph0607404_arXiv.txt": { "abstract": "We constrain the form of the primordial power spectrum using Wilkinson Microwave Anisotropy Probe (WMAP) 3-year cosmic microwave background (CMB) data (+ other high resolution CMB experiments) in addition to complementary large-scale structure (LSS) data: 2dF, SDSS, Ly-$\\alpha$ forest and luminous red galaxy (LRG) data from the SDSS catalogue. We extend the work of the WMAP team to that of a fully Bayesian approach whereby we compute the comparative Bayesian evidence in addition to parameter estimates for a collection of seven models: (i) a scale invariant Harrison-Zel'dovich (H-Z) spectrum; (ii) a power-law; (iii) a running spectral index; (iv) a broken spectrum; (v) a power-law with an abrupt cutoff on large-scales; (vi) a reconstruction of the spectrum in eight bins in wavenumber; and (vii) a spectrum resulting from a cosmological model proposed by \\citet{Doran} (L-D). Using a basic dataset of WMAP3 + other CMB + 2dF + SDSS our analysis confirms that a scale-invariant spectrum is disfavoured by between 0.7 and 1.7 units of log evidence (depending on priors chosen) when compared with a power-law tilt. Moreover a running spectrum is now significantly preferred, but only when using the most constraining set of priors. The addition of Ly-$\\alpha$ and LRG data independently both suggest much lower values of the running index than with basic dataset alone and interestingly the inclusion of Ly-$\\alpha$ significantly disfavours a running parameterisation by more than a unit in log evidence. Overall the highest evidences, over all datasets, were obtained with a power law spectrum containing a cutoff with a significant log evidence difference of roughly 2 units. The natural tilt and exponential cutoff present in the L-D spectrum is found to be favoured decisively by a log evidence difference of over 5 units, but only for a limited study within the best-fit concordance cosmology. ", "introduction": "The recent release of 3-year Wilkinson Microwave Anisotropy Probe (WMAP3; \\citealt{WMAP3}) data have provided precise measurements of temperature fluctuations in the cosmic microwave background (CMB). The accepted inflationary paradigm suggests that a primordial spectrum of almost scale-invariant density fluctuations produced during inflation went on to produce the observed structure in the CMB and that seen on large-scales in the current distribution of matter. Now, for the first time a purely scale invariant primordial spectrum is ruled out at $1\\sigma$ (\\citealt{SpergelII}; \\citealt{Parkinson}) in favour of a `tilted' spectrum with $n < 1$. The WMAP team have already attempted limited constraints on the form of the spectrum and \\citet{Parkinson} have conducted a model selection study to ascertain the necessity of a tilt in the spectrum with the new data. In this paper we extend both studies to a suite of models covering a wide variety of possibilities based on both physical and observational grounds. We use a fully Bayesian approach to determine the model parameters and comparative evidence to ascertain which model the data actually prefers. Our previous paper \\citep{Bridges} [Bridges06] used WMAP 1-year data (WMAP1; \\citealt{WMAP1}) to constrain the same set of models. These generalisations were motivated principally by observations of a decrement in power on large-scales from WMAP1 and a tilting spectrum on small-scales from high resolution experiments such as the Arcminute Cosmology Bolometer Array (ACBAR; \\citealt{ACBAR}), the Very Small Array (VSA ; \\citealt{VSA}) and the Cosmic Background Imager (CBI; \\citealt{CBI}). With two more years observing time and improved treatment of systematic errors the decrement in power on large scales is now somewhat reduced, yet still evident in WMAP3 and is now constrained almost to the cosmic variance limit while the tilting spectrum on small-scales is now seen even without the aid of high-resolution small scale experiments, due to tighter constraints on the second acoustic peak. On physical grounds we test a broken spectrum caused perhaps by double field inflation \\citep{barriga} and a spectrum predicted by \\citet{Doran} (L-D) naturally incorporating an exponential cutoff in power on large scales by considering the evolution of closed universes out of a big bang singularity, with a novel boundary condition that restricts the total conformal time available in the universe. We also aim to reconstruct the spectrum in a number of bins in wavenumber $k$. ", "conclusions": "A scale-invariant spectrum is now largely disfavoured by the dataset I with a spectral index $n_s=0.95 \\pm 0.02$ deviating by at least 2$\\sigma$ from $n_s=0$. Moreover a running spectrum ($n_{run} = -0.038 \\pm 0.030$) is now significantly preferred but only using the most constraining prior. The addition of Ly-$\\alpha$ forest data improves all constraints but does not alter the preferred spectral tilt greatly. It does however, along with LRG data, suggest a significantly smaller running index ($n_{run} = -0.015 \\pm 0.015$, $n_{run} = 0.01 \\pm 0.05$). This tension has previously been analysed by \\citet{Seljak} who conclude that such discrepancies, even though at the 2$\\sigma$ level are consistent with normal statistical fluctuations between datasets. A power law spectrum with a cutoff provides the best evidence fit in our full parameter space study with a significant evidence ratio of roughly 2 units across all three datasets. The similarity of this cutoff model with the L-D spectrum suggests the latter should also provide a very good fit. This is indeed borne out with decisively large evidence ratios within our limited primordial-only analysis." }, "0607/astro-ph0607318_arXiv.txt": { "abstract": "Diffuse emission in the mid-infrared shows a wealth of structure, that lends itself to high-resolution structure analysis of the interstellar gas. A large part of the emission comes from polycyclic aromatic hydrocarbons, excited by nearby ultra-violet sources. Can the observed diffuse emission structure be interpreted as column density structure? We discuss this question with the help of a set of model molecular clouds bathed in the radiation field of a nearby O-star. The correlation strength between column density and ``observed'' flux density strongly depends on the absolute volume density range in the region. Shadowing and irradiation effects may completely alter the appearance of an object. Irradiation introduces additional small-scale structure and it can generate structures resembling shells around HII-regions in objects that do not possess any shell-like structures whatsoever. Nevertheless, structural information about the underlying interstellar medium can be retrieved. In the more diffuse regime ($n(\\mbox{HI})\\lesssim 100$cm$^{-3}$), flux density maps may be used to trace the 3D density structure of the cloud via density gradients. Thus, while caution definitely is in order, mid-infrared surveys such as GLIMPSE will provide quantitative insight into the turbulent structure of the interstellar medium. ", "introduction": "} Diffuse emission in the infrared seems like a perfect laboratory to study the dynamics of the interstellar medium. Recent large-scale surveys by the Spitzer Space Telescope, specifically the GLIMPSE project \\citep{BEA2003} have provided us with unprecedented high resolution data of the diffuse emission in the mid-infrared (MIR). At first glance the wealth of structure exhibited in the flux density maps seems a striking argument by itself for structure analysis. However, the conspicuous structures themselves -- namely shells, bubbles, filaments and dark clouds (see e.g. \\citealp{CEA2004,CPA2006,HWI2006,JEA2006,MEA2006}) raise the question of how much of the observed structure actually corresponds to physical structure. Flux density maps contain information about volume density, column density and excitation, but to extract one of them is only possible under assumptions. For the structure analysis, ideally, we are interested in volume density, which is accessible only indirectly, leaving us with column density as a second best at most. In fact, over a broad range of wavelengths from ultraviolet to infrared, the opportunities seem to be rather rare where we can interpret observed intensity maps of diffuse emission as information about the underlying column density structure. More often than not the medium is optically thick for the emitted radiation, or denser components of the ISM act as absorbers. Optical depth effects become weaker with increasing wavelength, which is why the mid-IR (~$\\sim5\\mu$m) takes a somewhat special position (for a study of correlation between emission and column density in the far-IR see e.g. \\citealp{BZH2004} and \\citealp{SBG2006}). At longer wavelengths from the far-IR through millimeter, current observatories have rather poor spatial resolution, precluding study of small-scale interstellar structure. With a dust/PAH extinction cross section of $C_{ext}\\approx 10^{-23}$cm$^2$ per H-atom \\citep{LDA2001,LDB2001,DRA2003}, the optical depth for MIR emission in molecular clouds should range below or around $1$, which would encourage a direct interpretation of flux density as column density. For the near-infrared (NIR) this possibility has been discussed and supported by \\citet{PJP2006} to interpret so-called ``cloudshine'' observations by \\citet{FOG2006}, although there, the column densities would have to be substantially smaller than in the MIR. Although applicable to a wider range of surveys, this paper focuses on the MIR diffuse emission as seen by the IRAC camera of the Spitzer Space Telescope. To a large extent, the emission in the [5.8] and [8.0] bands comes from polycyclic aromatic hydrocarbons (PAHs) (to a lesser extent they also contribute to the [3.6] band, \\citealp{DRA2003}), mostly excited in the environment of nearby UV sources. Two (not unrelated) issues are raised in the MIR: (1) The medium is generally optically thick for the soft UV-radiation longward of $912${\\AA} that excites the PAH emission. Thus, only the outer layers of any structures -- unless really diffuse -- are excited and thereby traced out in the MIR, giving the object a filamentary appearance. (2) Observationally, there seems to be a strong morphological bias to shells in the PAH-emission. These could be dynamical shells (see e.g. \\citealp{CPA2006}), like e.g. windblown bubbles or HII regions. However, the destruction of PAHs around the UV source could also lead to a shell-like structure. And finally, these shells could be irradiation effects as noted under (1). And, of course, a combination of mechanisms is possible also. Moreover, the usual projection problem introduces a bias to interpreting objects as being two-dimensional, while they could be very ``3D'': More diffuse material gets ionized or blown away first, leaving the ubiquitous elephant trunk remnant structures (for an impressive example of this problem see the GLIMPSE study of RCW49, \\citealp{CEA2004}). We investigate the appearance of a model molecular cloud in the radiation field of a nearby UV-source, in order to quantify the correspondence of various measures and tracers between the original column density maps and the derived flux density maps. Rescaling the density range in the models allows us to mimic different physical environments, from diffuse clouds to dense molecular clouds and cores. The appearance of the original cloud can be completely altered by irradiation. Only for densities of up to $n(\\mbox{HI})\\approx 100$cm$^{-3}$ can the flux density actually be interpreted as column density. Above that, MIR self-extinction and strong shadowing effects in the UV will let the maps diverge. However, even for the highest density range (up to $n(\\mbox{HI})=10^5$cm$^{-3}$), where the flux density maps bear no resemblance to the column density, the structural properties of the original column density distribution still can be retrieved from the flux density. Our results demonstrate that the diffuse emission data as made available by GLIMPSE can serve as a powerful means to analyze the (dynamical) structure of the interstellar medium. This study aims at pointing out possible pitfalls and at giving a rule-of-thumb estimate where flux density structure could be trusted to represent column density. In the next section (\\S\\ref{s:glimpse}), we will give some observational motivation in the form of diffuse emission maps from the GLIMPSE data. We are deferring the full structure analysis of the GLIMPSE data to a future paper. The models and the details of the radiative transfer treatment are described in \\S\\ref{s:irradmodels}. The results (\\S\\ref{s:results}) are summarized in \\S\\ref{s:summary}. ", "conclusions": "} The abundance of structure in MIR diffuse emission as observed in e.g. the GLIMPSE data seems to offer a perfect laboratory to study the dynamics of the dense ISM. However, a large part of the observed structure could be irradiation effects due to PAH emission. PAHs are excited by UV photons from nearby stellar sources or the interstellar radiation field, and re-emit in the MIR. Because the respective cross sections differ by a factor of approximately $100$, gas, which is optically thick in the UV, can still be optically thin in the (emitted) MIR. Thus, PAH emission is often seen in filamentary structures, probably the ``rims'' of denser clouds. These irradiation effects might spoil the opportunity to study the ISM gas structure, since this requires interpretation of the observed flux density in terms of volume or column density. Motivated by a few examples taken from GLIMPSE, we identified possible limitations of this interpretation. We quantified the reliability of flux-density maps of diffuse emission in the MIR to reproduce the underlying (column) density information. We used two model sets, one corresponding to a (more or less) isolated ``molecular'' cloud, and the other imitating a region deep inside a molecular cloud, both irradiated by an O5 star. PAHs absorb the UV and re-emit the energy in the MIR, which then is integrated along the line-of-sight, including MIR extinction. In the following, we will summarize how reliable flux density maps reproduce column density in our models, and how this affects the structure analysis, with possible applications to MIR observations. \\subsection{Morphology: column density and flux density} If the medium is optically thin for the irradiating UV, then the MIR emission maps could be used for a high-resolution study of the column density structure of the medium. \\citet{PJP2006} supported this possibility for the NIR, based on the observations of ``cloudshine'' by \\citet{FOG2006}. Since they were interested in the appearance of the diffuse emission, they used an isotropic radiation field mimicking a UV background, in contrast to our models that employ a point source for irradiating the surrounding medium. For higher-density environments such as molecular cloud in the vicinity of strong UV sources, interpreting MIR emission as column density requires some caution: The transition from $\\tau<1$ to $\\tau>1$ can lead to strong signals in the MIR flux density, but not necessarily paired with a corresponding signal in the column density: The maps bear no resemblance to the column density (Figs~\\ref{f:irradcloudA0}--\\ref{f:irradcloudA3} and \\ref{f:irradcloudB}). In the extreme case, PAH-emission is only excited at the rims of clouds (e.g. \\citealp{CPA2006}), causing the impression of a highly filigree structure in the gas: irradiation introduces more small-scale structure than observable in the underlying column density maps. As soon as shadowing is obvious, the structure seen in emission will generally not represent column density. Some of the observed shell-like structures could be just irradiation effects, and by themselves indicate that shadowing (or ``rimming'') has set in (Figs.~\\ref{f:g30.7-0.0}, \\ref{f:irradcloudA1} and \\ref{f:pahcuts}). Since our models do not include PAH destruction around the star, we expect them to exhibit fewer shell-like structures than observed. PAH cavities would lead to ``shells'' even if the cavities were not associated with e.g. wind-blown bubbles. \\subsection{Structural Properties} Power spectra are only partially useful for an analysis of diffuse emission structure. Their well-known main drawback is that they tend to confuse the information about extent of a region and the separation of regions. Furthermore, masking is always an issue in power spectra, since it tends to introduce a signal by itself. Power spectra containing the central source are completely dominated by that (Fig.~\\ref{f:specA}, left column), while spectra of residual maps are slightly flatter than the underlying column density distribution. This could be a projection effect and/or the result of additional small-scale structure traced out by irradiation. The spectral slope is pretty much insensitive to the density contrast within the error bars, which are significant. For collapsed regions (Fig.~\\ref{f:specA}, right column), the column density spectrum flattens considerably because of the strong point source contribution. Compared to that, the flux density spectra steepen because the point sources are not fully irradiated and thus do not show up (except in extinction). Structure functions seem to be a more viable tool to investigate localized diffuse emission. Despite the fact that the irradiation may modify the underlying density information beyond the point of recognizability, the resulting structure functions still retrieve the salient scale information -- given that the field investigated is small enough not to be contaminated by global irradiation effects. There is little hope to retrieve the {\\em large-scale} information accurately by applying a global structure measure such as power spectra. Structure in extinction can be used as a continuation of structure seen in emission, although this raises the issue of an appropriate choice of background for the extincted region (Figs.~\\ref{f:g26.9-0.3} and \\ref{f:irradcloudA1}). The deviations (Fig.~\\ref{f:structchecksum}) are within the errors on the mean of the structure function. Since the overall structure in the ISM tends to be anisotropic, applying two-point correlators seems at least questionable. Averaging in $k$ or $l$ space leads to substantial errors on the mean, which themselves indicate that the underlying structure is anisotropic to a large extent. The application of these results to actual observational data (GLIMPSE) and the discussion of anisotropy we defer to a future paper. \\subsection{Flux density as gradient indicator} Since the conversion of UV to MIR will occur predominantly at regions of large positive radial density gradients (as long as there are photons left), the flux density maps might offer the opportunity to gather information about the (3D) density structure of the cloud. To test this, we compared the radial volume and column density gradients to the flux density maps. For lower density contrasts, the flux density maps tend to trace out the 3D structure of the cloud, and in fact they can be used as 3D gradient indicators. For higher density contrasts, they revert to an indicator of the column density gradients: the structures seemingly become two-dimensional. Thus, more diffuse regions are intrinsically ``more 3D'', while higher-density environments tend to be 2D. \\subsection{Limitations} Beside the irradiation effects discussed here, there are other limitations to an interpretation of flux density maps as column density. (1) If the volume density is low enough that the exciting UV can irradiate the whole cloud, one might question how long the line-of-sight actually is, and whether angular effects leading to scale-mixing in a structure analysis would play a role. (2) On the other hand, the volume density might be large so that the irradiating UV will be absorbed more or less directly ``at the rim'' (if such a thing exists) of the cloud. Then, depending on the geometry of observer, irradiated medium and irradiation source, the observer might see predominantly 1D structures or 2D structures, implying projection effects in the spectral information. \\subsection{Conclusions} Depending on the diagnostics, MIR flux density maps of diffuse emission from PAHs excited by a nearby UV source can be used to extract information about the density structure of the underlying (molecular) cloud, though this statement needs some qualification. (1) Flux density maps need not correspond ``by eye'' to column density maps: due to irradiation effects they tend to show more small-scale structure. (2) Irradiation by a point source can produce shell-like structures, mimicking physical shells, even in objects which do not have any shell-like properties. (3) As long as structure studies are restricted to areas small enough not to be contaminated by any large-scale effects, flux density and column density show similar structural properties. However, the application of unmodified two-point correlation functions introduces substantial errors on the mean due to the underlying anisotropy in the ISM structure. (4) MIR flux density maps tend to trace out {\\em gradients} in the three-dimensional density distribution. Employing MIR diffuse emission to extract structure information about the underlying interstellar medium requires close attention to the environment. This study attempts to provide some guidelines to chose appropriate locations. Bearing the limitations in mind, analyzing the ISM structure with the help of the GLIMPSE data will be a promising task." }, "0607/astro-ph0607542_arXiv.txt": { "abstract": "The technique of weak-lensing aperture mass densitometry, so called the $\\zeta$-statistic, has recently been popular in actual observations for measurement of individual cluster mass. It has however been anticipated that the line-of-sight projection by foreground and background matter can adversely affect the cluster mass determination with not only substantial error dispersion but also a sizable positive systematic bias. Additionally, the finite number of background galaxies even at a reasonable observing depth can also introduce Poisson noise to the mass estimate. In this paper, we quantitatively investigate the degree of errors separately contributed by the two sources to the mass determination of those galaxy clusters with $M_{200}>10^{14}M_{\\odot}$. We find that the aperture mass of $\\zeta$-statistic turns out to be a mass estimator of much reduced systematic bias, due to the cancellation by the positively biased local background mass sheet. However, the error dispersion of $M_{200}$ arising from both projection effect and Poisson noise is found to be still sizable ($40\\%-90\\%$), even for the shear-selected, clean sample where multiple clusters located within a suitable projected aperture are removed. We also investigate how to remedy this large-error problem in weak lensing measurements, and propose a plausible alternative mass estimator, $M(<\\theta_{1000})$, an aperture mass measured within about half the virial radius. The aperture mass $M(<\\theta_{1000})$ is free of bias and has a substantially reduced error dispersion, $39\\%$ for the worst case of high-$z$, low-mass clusters, that can be smaller than the error dispersion of $M_{200}$ as much as a factor 3. ", "introduction": "The mass function of galaxy clusters has long been recognized as the most convenient and important indicator for probing the evolution of structure formation, thereby helping determine the cosmological parameters. In addition, the cluster halo mass, when combined with the cluster gas mass via the Sunyaev-Zel'dovich effect, provides the opportunity for probing the cluster baryon fraction (Umetsu, et al. 2005). Conventional techniques, such as measuring the velocity dispersion of gravitationally bound galaxies and the X-ray emission profile, have long been employed to measure the cluster mass, assuming cluster galaxies and X-ray emitting plasmas to be dynamically relaxed within the cluster gravitational potential. On the other hand, the new technique of mass measurement through weak gravitational lensing has been gaining popularity in recent years, with the advantage of not having to assume the dynamical equilibrium in the cluster (e.g., Umetsu, Tada, \\& Futamase 1999; Bartelmann \\& Schneider 2001; Schneider 2005). This methodology was first pioneered by Tyson et al. (1990). Various refined techniques for the weak-lensing mass determination were later proposed by several groups (Fahlman et al. 1994; Kaiser 1995; Bartelmann 1995; Seitz \\& Schneider 1996; Squires \\& Kaiser 1996; Broadhurst, Takada, Umetsu et al. 2005). Among these techniques, the $\\zeta-$statistic was particularly devised to measure the lens mass directly from the tangential component of local gravitational image distortions without involving a non-local mass reconstruction (Fahlman et al. 1994; Kaiser 1995). Schneider (1996) extended the $\\zeta$-statistic by generalizing its kernel, which allows one to define an optimal measure for the detection of mass concentrations, and this aperture mass technique has been applied to deep optical imaging data to search for clusters (e.g., Erben et al. 2000; Umetsu \\& Futamase 2000; Wittman et al. 2001, 2003; Miyazaki et al. 2002; Dahle et al. 2003; Hetterscheidt et al. 2005; Schirmer et al. 2006). King et al. (2001) investigated the cluster mass measurement influenced by interior substructures, and found the measured mass as accurate as within $10\\%$. Clowe et al. (2004) studied the effect of asphericity on the cluster mass determination, and concluded, under the assumption of an NFW profile, that the asphericality effect generally changes the mass estimate by $5\\%$ to $10\\%$. In addition, several authors have compared the weak-lensing mass with the mass determined by the galaxy kinematics and X-ray observations. Reblinsky \\& Bartelmann (1999) concluded that the mass estimates using $\\zeta$-statistic are significantly more accurate than those obtained from the galaxy kinematics. Ettori \\& Lombardi (2003) studied the mass distribution of the rich cluster MS 1008.1-1224 at $z=0.302$ based on {\\it Chandra} X-ray and FORS1-VLT multicolor-imaging data, and they found that the two mass profiles obtained from X-ray and weak-lensing analyses up to $550 h^{-1}$ kpc are consistent with each other within $1\\sigma$ uncertainty. Irgens et al. (2002), assuming a singular isothermal sphere model for the cluster mass profile, compared spectroscopic velocity dispersions, $\\sigma_p$, of 13 X-ray luminous clusters around $z \\sim 0.3$ with $\\sigma_{\\rm WL}$ of these clusters determined by weak-lensing tangential shear measurements out to the cluster virial radius. It was found that among all, two clusters are in strong discrepancy, with $\\sigma_{\\rm WL} > 2\\sigma_p$ and $\\sigma_{\\rm WL}\\approx 2\\sigma_p$, whereas the rest are in fair agreement, with $\\langle \\sigma_{p}/\\sigma_{\\rm WL}\\rangle \\approx 1$. Though these exceptional clusters may be in dynamical non-equilibrium, another possibility may arise from the projection effects of other mass concentrations (Cen 1997; Reblinsky \\& Bartelmann 1999; White, van Waerbeke, \\& mackey 2002; Padmanabhan, Seljak, \\& Pen 2003; Hamana, Takada, \\& Yoshida 2004; Henawi \\& Spergel 2005) and/or local filamentary structures (Metzler et al. 1999; Metzler, White, \\& Loken 2001) along the line-of-sight. Using $N$-body simulations White et al. (2002) studied the completeness and efficiency of weak-lensing cluster surveys on the basis of their {\\it mass-selected} mock cluster sample and found that the line-of-sight projection effects can be quite serious due to the broad lensing kernel. In the cluster mass estimate based on the convergence map, they found a positive bias of $\\sim 20-30\\%$ with a substantially larger error dispersion that can even occasionally yield negative lens masses. Metzler et al. (1999; 2001) studied the projection effects on weak-lensing mass estimates for massive clusters caused by the local large-scale filamentary structures. Including the projection effects from local matter within a sphere of $128 h^{-1} {\\rm Mpc}$ radius, they found the lensing convergence maps to yield an positive mass bias of $\\sim 30\\%$ and the mass error dispersion of $\\sim 0.3$ for massive clusters at a redshift of $z=0.5$. In fact, these problems of weak-lensing mass determination, i.e., positive mass bias and large mass error dispersion, have been alluded in earlier works (Cen 1997; Reblinsky \\& Bartelmann 1999). Further, the cluster halo triaxiality itself can cause a bias in the lensing-based mass estimation (Clowe et al. 2004; Hamana et al. 2004; Oguri et al. 2005), while it is likely to have less effect on the X-ray cluster mass estimate (Gavazzi 2005). Thus, despite that weak lensing offers a unique tool for the measurement of cluster masses without any assumption of their equilibrium state, it can however suffer from the projection effects. Such problems are less significant in X-ray or spectroscopic velocity-dispersion measurements. The present study aims to investigate the errors in weak lensing cluster mass measurements as well. However, this work differs from the aforementioned previous works, in that we attempt to simulate the actual wide-field weak lensing measurements, with numerically simulated shear data as closely resembling the observing data as possible. In particular, we shall focus on the bias errors and random errors pertinent to local weak-lensing measurement of $\\zeta$-statistic. Moreover, at a given observing depth, the finite number of background galaxies can introduce non-negligible Poisson noise convolved with the projection error in the measured data. We shall quantify, in this work, the regime for which the projection effect dominates, and the other regime where the Poisson noise dominates. This paper is organized as follows. In Sect.~2 we describe our cosmological $N$-body simulations, weak lensing simulations, and the construction of a mock cluster catalog. Details of our shear-based mass estimator and mock observations are presented in Sect.~3. In Sect.~4 we apply the shear-based mass estimator to our simulated weak lensing observations, and examine the statistical properties of the errors in weak lensing cluster mass estimates. The radial mass error profiles are discussed in Sect.~5. Based on the radial error profiles, we propose a plausible alternative cluster mass indicator that has much reduced error dispersion. We finally present the discussions and a summary in Sect.~6. ", "conclusions": "This paper reports a systematic study on the accuracy of cluster mass measurements through weak lensing observations. Specifically this work takes into account the mass errors introduced by the projection effect and by the Poisson noise of finite number of randomly oriented background galaxies. Among these sources of mass errors, the projection effect has been reported to yield a non-negligible systematic positive bias ($\\sim 20\\%$) in the cluster mass estimate based on weak lensing. Adopting the local shear-field measurement using $\\zeta$-statistic, which is often used in actual observations, we nevertheless found that such a positive bias can be largely canceled by the positively-biased local background mass sheet. That is, the $\\zeta$-statistic measurement can provide a bias-free cluster mass estimate. In this paper, we also report that the error in $M_{200}$ determination is expected to exceed $50\\%$ even for a moderately deep observation ($R\\simeq 25.5$ mag), regardless of lens mass and redshift. Even after a clean procedure that removes clusters with detectable companions in the projected map, the mass error is still substantial, exceeding $40\\%$ for a $R\\simeq 25.5$ observation. Mass errors can scatter observed data of one mass bin into another bin when constructing the cluster mass function, and smear out the mass function. In the mass range where the mass function has a large gradient, i.e., $M > M_{\\star}$, such mass errors can greatly distort the mass function. The measurement error will eventually propagate into the determination of cosmological parameters, such as the matter density parameter $\\Omega_m$ and the matter fluctuation amplitude $\\sigma_8$, which rely critically on an accurate mass function. The so-called self-calibration was devised to correct for the systematic errors of this kind (Hu 2003), but the random errors are un-removable even with the self-calibration. To significantly reduce the mass error in weak lensing measurements, we also suggest the possibility of an alternative lens mass, which has a considerably smaller error. It is an interior mass well inside $\\theta_{200}$. For example, Fig.~\\ref{fig:massratiodiag} shows that the error in $M(<\\theta_{500})$ (mass interior to approximately $0.7 \\theta_{200}$) is already noticeably smaller than $M_{200}$, the error dispersion of $M(<\\theta_{1000})$ is at most $39\\%$ for all detectable clusters and the error dispersion of $M(<\\theta_{1500})$ at most $32\\%$. Comparing Table \\ref{table:cleansigma500} with Table \\ref{table:cleansigma}, we find $M_{1000}$ to have more than a factor 2 in error reduction for all detectable clusters; for low-$z$ lenses, the error reduction can be as large as a factor 3. Given these results, we suggest that $M(<\\theta_{1000})$ be a better mass variable than $M_{200}$ for constructing the mass function. Concerning the mass function of $M(<\\theta_{1000})$, it should be reminded that the mass function of $M_{200}$ cannot be analytically derived, and needs to be determined empirically from $N$-body simulations. From this spirit, the mass function of $M(<\\theta_{1000})$ can also be obtained from $N$-body simulations, in a similar manner as the mass function of $M_{200}$. To the best of our knowledge, there have not been investigations on the mass function for mass different from $M_{200}$ in literature. Such a new class of mass function, containing less mass scatter than the conventional mass function, may be more useful for constraining the cosmological parameters. However, whether or not the more accurately measured $M(<\\theta_{1000})$ function can actually be more useful than the less accurately measured $M_{200}$ function really depends on the detailed form of the $M(<\\theta_{1000})$ function that contains the cosmology-parameter-sensitive feature, similar to a break at $M_{\\star}$ for the $M_{200}$ function. Therefore our suggestion at this point in favor of an alternative mass estimator for a new mass function should be regarded as plausible but still preliminary. Finally, though the depth of observation has been fixed to $R\\simeq 25.5$ mag in this work for which $n_g \\simeq 30$ arcmin$^{-2}$, for evaluation of mass errors, we may relax this constraint to assess other observing depths straightforwardly. At a medium depth $R\\simeq 24.5$ mag, suitable for wide-field surveys, the background galaxy density $n_g \\simeq 20$ arcmin$^{-2}$ (Fontana et al. 2000). One can quickly estimate from Fig.~\\ref{fig:massratiodiag} with the Poisson statistics that the galaxy ellipticity noise is still much less than the errors at $\\theta_{500}$, $\\theta_{1000}$, and $\\theta_{1500}$ introduced by the projection effect for clusters of $M>2\\times10^{14}M_{\\odot}$ and $0.24\\times10^{14}M_{\\odot}$ and $0.2$30\\% in the infrared and b) the infrared is considerably less affected by dust absorption than the optical. The all sky survey by IRAS provided the first opportunity to classify a large number of galaxies based on the shape of their global spectral energy distribution (SED). \\citet{deGrijp87} showed that a ratio of S60$/$S25$>$0.26\\footnote{Where S25 and S60 are the 25 and to 60 $\\mu$m IRAS flux densities.} could result in a 70\\% success rate in identifying previously unknown Seyferts. More recent developments on broadband mid-IR diagnostics are presented in another review on this volume. \\vspace*{-0.3cm} ", "conclusions": "\\vspace*{-0.3cm} It has become evident from the numerous contributions presented during this meeting that the high quality of Spitzer/IRS spectroscopy is opening new horizons in the use of the infrared as a tracer of the properties of nearby and high-redshift galaxies. With more than three years of mission to go, we have just glimpsed on the possibilities that lie ahead." }, "0607/astro-ph0607538_arXiv.txt": { "abstract": "In this paper we study the evolution of core and corona of nine open clusters using the projected radial density profiles derived from homogeneous CCD photometric data obtained through the 105-cm Kiso Schmidt telescope. The age and galactocentric distance of the target clusters varies from 16 Myr to 2000 Myr and 9 kpc to 10.8 kpc respectively. Barring Be 62, which is young open cluster, other clusters show a uniform reddening across the cluster region. The reddening in Be 62 varies from $E(B-V)_{min}$= 0.70 mag to $E(B-V)_{max}$= 1.00 mag. The corona of six of the clusters in the present sample is found to be elongated, however on the basis of the present sample it is not possible to establish any correlation between the age and shape of the core. The elongated core in the case of young cluster Be 62 may reflect the initial conditions in the parental molecular cloud. The other results of the present study are (i) Core radius `$r_c$' and corona size $`r_{cn}$'/cluster radius $`r_{cl}$' are linearly correlated. (ii) The $r_c/r_{cn}/r_{cl}$ are linearly correlated with the number of stars in that region. (iii) In the age range 10-1000 Myr, the core and corona shrink with age. (iv) We find that in the galactocentric distance range 9 - 10 kpc, the core and corona/cluster extent of the clusters increase with the galactocentric distance. ", "introduction": "The study of galactic open clusters is of great interest in several astrophysical aspects. Young open clusters provide information about current star formation processes and are key objects for clarifying questions of galactic structure, while observations of old and intermediate age open cluster play an important role in studying the theories of stellar and galactic evolution. The nucleus and the corona (extended region of the star cluster) are two main regions in open clusters (Kholopov 1969). The nucleus of a cluster contains relatively bright and massive ($\\ge$ 3 $M_\\odot$) stars whereas corona, which contains a large number of faint and low mass ($\\le$ 1 $M_\\odot$) stars, has important bearing on studies related to the mass function (MF), the structure and the evolution of open clusters. A detailed analysis of the structure of corona of open clusters is needed to understand the effects of external environments like the galactic tidal field and impulsive encounters with interstellar clouds etc. on dynamical evolution of open clusters (Pandey et. al. 1990). Extensive studies of the coronal regions of clusters have not been carried out so far mainly because of non-availability of photometry in a large field around open star clusters. The $2K\\times2K$ CCD mounted on Schmidt telescopes (Kiso, Japan), covering $\\sim 50'\\times 50'$ field can be used to get photometry in a large field around open star clusters. The ability to obtain improved photometry of thousands of stars means that large-scale studies of open clusters can be conducted to study the spatial structure and stability of galactic open clusters. With the addition of photometry of nearby field region, it is possible to construct luminosity function (LF)/mass function (MF), which are useful to understand the cluster-formation processes and the theory of star formation in open clusters (Miller \\& Scalo 1979). Considering the importance of low mass stars in the corona of star clusters, we have generated wide field photometric database around 9 open star clusters with the aim to re-investigate the cluster's parameters e.g. reddening, distance, age, their size and LF/MF using a homogeneous data base. The basic parameters of the cluster taken from WEBDA\\footnote{http://obswww.unige.ch/webda} (Mermilliod 1995) are given in Table 1. ", "conclusions": "" }, "0607/astro-ph0607012_arXiv.txt": { "abstract": "{}{We reexamine the theoretical instability domain of pulsating PG1159 stars (GW Vir variables).} { We performed an extensive $g$-mode stability analysis on PG1159 evolutionary models with stellar masses ranging from $0.530$ to $0.741 M_{\\odot}$, for which the complete evolutionary stages of their progenitors from the ZAMS, through the thermally pulsing AGB and born-again phases to the domain of the PG1159 stars have been considered.} {We found that pulsations in PG1159 stars are excited by the $\\kappa$-mechanism due to partial ionization of carbon and oxygen, and that no composition gradients are needed between the surface layers and the driving region, much in agreement with previous studies. We show, for the first time, the existence of a red edge of the instability strip at high luminosities. We found that all of the GW Vir stars lay within our theoretical instability strip. Our results suggest a qualitative good agreement between the observed and the predicted ranges of unstable periods of individual stars. Finally, we found that generally the seismic masses (derived from the period spacing) of GW Vir stars are somewhat different from the masses suggested by evolutionary tracks coupled with spectroscopy. Improvements in the evolution during the thermally pulsing AGB phase and/or during the core helium burning stage and early AGB could help to alleviate the persisting discrepancies.}{} ", "introduction": " ", "conclusions": "\\label{conclusions} In this paper we re-examined the pulsational stability properties of GW Vir stars. We performed extensive nonadiabatic computations on PG1159 evolutionary models with stellar masses ranging from $0.530$ to $0.741 M_{\\odot}$. For each sequence of models, we computed the complete evolutionary stages of PG1159 progenitors starting from the Zero Age Main Sequence. Evolution was pursued through the thermally pulsing AGB and born-again (VLTP) phases to the domain of the PG1159 stars. The employment of such full evolutionary PG1159 models constitutes a substantial improvement over previous studies on GW Vir stars regarding the stellar modelling. Numerous detailed investigations about pulsating PG1159 stars have been performed on the basis of artificial stellar models. In spite of the fact that the significant pulsation damping and driving occur in PG1159 envelope stars, the employment of such simplified stellar configurations appear not well justified in the case of these stars. This is in contrast to the situation of their more evolved counterparts, the white dwarf stars, for which their thermo-mechanical structure has relaxed to the correct one by the time the pulsational instability domains are reached. The main goal of the present work has been to assess to what degree the conclusions arrived at in previous studies on PG1159 stars change when realistic stellar configurations are adopted. Our study confirms the following results, already known from previous studies: \\begin{itemize} \\item $g$-modes in PG1159 models are excited by the $\\kappa$-mechanism due to partial ionization of carbon and oxygen. No abundance gradients between the driving region and the stellar surface are necessary to drive $g$-mode pulsations at the correct effective temperatures and period ranges. \\item There exists a well-defined instability domain with a blue edge which is strongly dependent on the stellar mass. \\item Different surface helium abundances lead to sizeable differences in the precise location of the theoretical blue edge of the instability domain. \\item The instability domain splits into two separated regions, one of them at high luminosities characterized by long periods, and the other at low luminosities, corresponding to shorter periods, as recently demonstrated by GAS05. \\item All pulsating PG1159 stars lay into the predicted instability domain in the $\\log (T_{\\rm eff})-\\log g$ plane. \\item There is a very good agreement between the full period spectrum observed in GW Vir stars and the theoretical ranges of unstable periods. \\item The pulsation periods of excited modes decrease with decreasing luminosity (increasing surface gravity), in line with the observational trend. \\end{itemize} As for our new findings, we mention: \\begin{itemize} \\item There exists a red edge of the instability domain at the high-luminosity (low-gravity) regime. This red edge is mass-dependent. \\item The border of the instability domains in the $\\log T_{\\rm eff}-\\log \\Pi$ plane at the high-luminosity, long-period regime is well delineated. \\item The pulsating PG1159 Longmore 4 is located at the very red edge of the instability strip at high luminosities, a fact that could be reflecting the surprising behaviour observed in the spectral type of this star (Werner et al. 1992). \\item Some non-variables occupying the instability strip have standard helium abundances and the presence of them between pulsators can not be explained through the argument of Quirion et al. (2004). \\item The pulsating nature and also the range of observed periods of PG 1153-035 --- the prototype of the GW Vir class --- are naturally accounted for by pulsationally unstable PG1159 models with a stellar mass of $\\sim 0.53-0.54 M_{\\odot}$. \\end{itemize} Finally, we found that generally the seismic masses (as inferred from the period spacings) are somewhat different from the spectroscopic masses, although the disagreement for the PG 1159-035 star is somewhat alleviated according to our calculations. The persisting discrepancies could be attributed to a number of factors. On the observational side, possible systematics errors in the spectroscopic determination of $g$ and $T_{\\rm eff}$, and/or errors in the measurement of the period spacings of pulsating PG1159 stars. On the other hand, differences in the microphysics or the previous evolution may alter the location of the post-AGB tracks (Blo\\\"ecker 1995). In fact, it has been argued by Werner \\& Herwig (2006) that the evolution during the TP-AGB (concerning third dredge up efficiency and TP-AGB lifetimes) may be key in determining the location of post-AGB tracks. However, in preliminary simulations we have found that neither third dredge up efficiency nor TP-AGB lifetimes play an important role in determining the location of post-AGB tracks. It remains to be seen if other physical assumptions like the overshooting efficiency during the core helium burning stage and early AGB (that also define the structure of the C-O core and are completely free parameters) may be playing a role in the location of post-AGB tracks. We are currently performing simulations of full stellar evolution sequences under different assumptions to clarify these issues." }, "0607/astro-ph0607224_arXiv.txt": { "abstract": "We present the results of numerical experiments, in which we study how the asphericities induced by the growth of the standing accretion shock instability (SASI) produce the gravitational waveforms in the postbounce phase of core-collapse supernovae. To obtain the neutrino-driven explosions, we parameterize the neutrino fluxes emitted from the central protoneutron star and approximate the neutrino transfer by a light-bulb scheme. We find that the waveforms due to the anisotropic neutrino emissions show the monotonic increase with time, whose amplitudes are up to two order-of-magnitudes larger than the ones from the convective matter motions outside the protoneutron stars. We point out that the amplitudes begin to become larger when the growth of the SASI enters the nonlinear phase, in which the deformation of the shocks and the neutrino anisotropy become large. From the spectrum analysis of the waveforms, we find that the amplitudes from the neutrinos are dominant over the ones from the matter motions at the frequency below $\\sim 100$ Hz, which are suggested to be within the detection limits of the detectors in the next generation such as LCGT and the advanced LIGO for a supernova at 10 kpc. As a contribution to the gravitational wave background, we show that the amplitudes from this source could be larger at the frequency above $\\sim$ 1 Hz than the primordial gravitational wave backgrounds, but unfortunately, invisible to the proposed space-based detectors. ", "introduction": "The gravitational astronomy is now becoming a reality. In fact, the ground-based laser interferometers such as TAMA300 \\citep{tama,tamanew} and the first LIGO \\citep{firstligo,firstligonew} are beginning to take data at sensitivities where astrophysical events are predicted. For the detectors including GEO600 and VIRGO, core-collapse supernovae especially in our Galaxy, have been supposed to be the most plausible sources of gravitational waves (see, for example, \\citet{new,kotake_rev} for review). Since the gravitational wave (plus neutrinos) is the only tool which gives us the information in the innermost part of evolved massive stars, the detection is important not only for the direct confirmation of gravitational waves but also for the understanding of supernova physics itself. So far, most of the theoretical predictions of gravitational waves from supernovae have focused on the bounce signal in the context of rotational \\citep{mm,ys,zweg,dimmel,fry,kotakegw,shibaseki,ott} and magnetorotational \\citep{kotakegwmag,obergaulinger} core collapse. In most of the previous studies, the iron core prior to core-collapse was assumed to rotate much more rapidly than predicted by the recent stellar evolution calculations \\citep{heger05}. Recently, the initial rotation periods were estimated to be larger than $\\sim$ 100 sec for the observed rotation periods of the radio pulsars \\citep{ott_birth}. In such a slowly rotating case, the bounce signal becomes too small to be detected even by the laser interferometers in the next generation for a galactic supernova, owing to the suppression of the rotation-induced deformation at core-bounce (see, e.g., \\citet{kotakegwmag}). Besides the rapid rotation of the cores, two other ingredients have been considered to be important in the much later phases after core bounce, namely convective motions and anisotropic neutrino emissions. Both of them contribute to the non-spherical parts in the energy momentum tensor of the Einstein equations, thus being the potential sources of the gravitational wave (see, \\citet{kotake_rev} for a review). One of the possibility as the origin of the asphericities may be large scale density inhomogeneities formed in the central core prior to collapse (e.g., \\citet{bazan,meakin}). \\citet{fryersingle} performed three dimensional SPH simulations and pointed out that the neutrino-originated gravitational waves, which dominate over the one from the convections, are within the detection limits for the advanced LIGO for the galactic supernova (see also, \\citet{burohey,fryersingle,fryer04,muyan97}). Another possibility to induce anisotropy is the (moderate) rotation of the core. \\citet{mueller} calculated the gravitational waves based on the two-dimensional (2D) Boltzmann transport simulations of slowly rotating core \\citep{buras} and found that the neutrino-originated gravitational waves exceed the bounce signal large enough to be detectable by the advanced LIGO with good signal-to-noise ratio for the galactic supernova (see, also \\citet{kotake_rev} for the properties of neutrino-originated gravitational waves in the rapidly rotating case). More recently, the new ingredient of the gravitational-wave emissions is reported \\citep{ott_new}, namely the g-mode excitations of the protoneutron stars, which was observed in the 2D approximate Boltzmann transport simulations at much later postbounce phase ($\\sim 600$ ms) \\citep{burr_new}. There is an another ingredient for producing large asphericity, to which much attention has been paid recently in the context of the studies about the explosion mechanisms, that is the so-called standing accretion shock instability (often called ``SASI''). In the numerical simulations by \\citet{blondin_03,scheck_04,blondin_05,ohnishi_1,ohnishi_2}, it was found that the standing shock wave is shown to be unstable to non-radial perturbations, and that the perturbations grow up to the non linear regime with clear low-mode ($\\ell=1,2$) dominance, leading to the global deformation of the shock wave later. Here $\\ell$ stands for the azimuthal index of the Legendre polynomials. The importance of SASI is also stressed by the recent studies, demonstrating that such an explosion is favorable to reproduce the observed synthesized elements of SN1987A \\citep{kifo} and also to explain the origin of the natal kicks of young pulsars \\citep{scheck_04}. These situations motivate us to study how the gravitational waveforms are originated from the asphericities by SASI. In this paper, we present the results of numerical experiments, in which we study how the asphericities induced by the growth of SASI produce the gravitational waveforms. To obtain the neutrino-driven explosions, we parameterize the neutrino fluxes emitted from the central protoneutron star and approximate the neutrino transfer by the light-bulb scheme. Based on the long-term two dimensional hydrodynamic results, we calculate the gravitational waveforms. By doing the spectrum analysis, we study the detectability of such signals from a nearby core-collapse supernova. It is noted that much attention has been paid recently to the core-collapse supernovae as one of the promising sources of the cosmological gravitational wave backgrounds (see, e.g., \\citet{buonanno} and references therein). Thus we calculate the SASI-induced gravitational wave backgrounds and discuss the detectability by the currently proposed space-based detectors such as LISA\\footnote{see http://lisa.jpl.nasa.gov/}, BBO\\footnote{see http://universe.nasa.gov/program/bbo.html}, and DECIGO \\citep{seto}. The plan of this paper is as follows: In Section \\ref{sec2}, we outline the initial models, the numerical methods, and shortly summarize the methods for calculating the waveforms. We show the main numerical results in Section \\ref{sec3}. We summarize and discuss our results in Section \\ref{sec4}. ", "conclusions": "} We presented the results of numerical experiments, in which we studied how the asphericities induced by the growth of the standing accretion shock instability could produce the gravitational waveforms in the postbounce phase of core-collapse supernovae. To obtain the neutrino-driven explosions, we parametrized the neutrino fluxes emitted from the central protoneutron star and approximated the neutrino transfer by the light-bulb scheme. By doing the spectrum analysis of the waveforms, we investigated the detectability of the signals from a single core-collapse supernova and the cosmological ones by the ground-based and space-based laser interferometers, respectively. Our main results can be summarized as follows. 1. The amplitudes of the gravitational waves from the anisotropic neutrino emissions are larger up to two orders of magnitudes than the ones from the matter motions during the SASI operations. It is found that the wave amplitudes from the neutrinos show the monotonic increase with time, regardless of the neutrino luminosities from the protoneutron star. We point out that this feature can be understood by the specific nature of SASI, which makes the deformation of the shock waves of $l=1,2$ modes dominant, leading to the enhanced neutrino emissions in the regions close to the symmetry axis. In fact, we show that the amplitudes become larger when the growth of the SASI enters the nonlinear phase, in which the deformation of the shocks and the neutrino anisotropy become large. 2. From the spectrum analysis of the waveforms, we find that the amplitudes from the anisotropic neutrino emissions are dominant over the ones from the matter motions at frequency $\\lesssim 100$ Hz. The detection of such signals from a galactic supernova may be marginal for the currently running detector of the first LIGO and promising for the detectors in the next generation such as LCGT and the advanced LIGO. 3. As for the background radiation, we indicate that the contribution of the gravitational signals considered here could be larger at frequency $\\gtrsim$ 1 Hz than the primordial gravitational wave backgrounds generated in the inflationary epoch. Unfortunately, however, it is found that this frequency range is just outside of the sensitivity of the proposed detectors, such as DECIGO. We give a brief comparison with recently published models. The monotonic increase with time in the wave amplitudes of the neutrino-originated gravitational waves is consistent with the model s15r of \\citet{mueller}, in which the operation of the SASI was seen. However, the amplitudes here are typically larger (up to 1 order of magnitude). This should be mainly because the neutrino luminosity here is taken to be higher than the one obtained in \\citet{mueller}, which is about $ \\sim 2 \\times 10^{52}$ erg/s during the SASI operation. As for the total gravitational-wave energy emission, typical values of the computed models here ($\\sim 10^{-10} M_{\\odot} c^2$, see Table 1) are one order magnitude smaller than the one in \\citet{mueller}. This should be owing to the excision of the protoneutron star, by which the contribution from the high frequency domain of the energy emissions are eliminated in this study. It should be noted that the larger oscillations of the protoneutron star in the postbounce phase \\citep{burr_new} and the resulting efficient gravitational emissions \\citep{ott_new} cannot be treated in principle here. The simulation highlighted here is nothing but an idealized study for the physical understanding of relation between the asphericities induced by the SASI and the resulting gravitational waves. Remembering the caveats about the assumptions of the artificially constructed initial condition, the fixed accretion rate, the absorbing boundary condition, and the fixed neutrino luminosity and energies, it is by no means definitive at all. Especially, much better neutrino transfer is indispensable for more reliable calculations of the neutrino-originated gravitational wave, which we only considered the radial transport. One more major deficit is the axial symmetry assumed in the present two-dimensional (2D) simulations. In three-dimensional environments, the pronounced dominance of $l =1,2$ along the symmetry axis, which is a coordinate singularity in the 2D computations, may become weaker, owing to the additional spatial degree of freedom in the azimuthal ($\\phi$) direction. In the 3D case, we think that the qualitative features of the plus mode waveform computed in the 2D case here will be unchanged, but quantitatively, we expect that the amplitudes become smaller owing to the reduced anisotropy along the symmetry axis. Thus the amplitudes calculated in this study could be an upper bound, in which the maximal anisotropy of the shock waves and thus neutrino emissions outside the neutrino sphere could be achieved. Furthermore we think that it is interesting to investigate the properties of the cross mode gravitational waves, which are of genuine 3D origin and could be possibly produced from the transfer of the $l=1,2$ modes to some modes with nonvanishing $m$ in $Y_{lm}$. This study is a prelude to the forthcoming 3D simulations to clarify those aspects, which will be presented elsewhere." }, "0607/astro-ph0607648_arXiv.txt": { "abstract": "We analyse a $z<0.1$ galaxy sample from the Sloan Digital Sky Survey focusing on the variation of the galaxy colour bimodality with stellar mass $\\mass$ and projected neighbour density $\\Sigma$, and on measurements of the galaxy stellar mass functions. The characteristic mass increases with environmental density from about $10^{10.6}\\Msun$ to $10^{10.9}\\Msun$ (Kroupa IMF, $H_0=70$) for $\\Sigma$ in the range 0.1--$10\\,\\perMpcsq$. The galaxy population naturally divides into a red and blue sequence with the locus of the sequences in colour-mass and colour-concentration index not varying strongly with environment. The fraction of galaxies on the red sequence is determined in bins of 0.2 in $\\log\\Sigma$ and $\\log\\mass$ ($12 \\times 13$ bins). The red fraction $f_r$ generally increases continuously in both $\\Sigma$ and $\\mass$ such that there is a unified relation: $f_r = F(\\Sigma,\\mass)$. Two simple functions are proposed which provide good fits to the data. These data are compared with analogous quantities in semi-analytical models based on the Millennium $N$-body simulation: the \\citet{bower06} and \\citet{croton06} models that incorporate AGN feedback. Both models predict a strong dependence of the red fraction on stellar mass and environment that is qualitatively similar to the observations. However, a quantitative comparison shows that the \\citeauthor{bower06}\\ model is a significantly better match; this appears to be due to the different treatment of feedback in central galaxies. ", "introduction": "\\label{sec:intro} Galaxies when characterised by their morphology or radial profiles, integrated or central colours, and total luminosity or stellar mass, exhibit a range of relationships. These include colour-morphology relations \\citep{Holmberg58,RH94}, and colour-magnitude relations separately for early-type galaxies \\citep{Faber73} and late-type galaxies \\citep{CR64}. While it was often considered that the natural dividing line was between spirals and ellipticals/lenticulars (e.g.\\ \\citealt*{TMA82}), it was not until the multi-wavelength Sloan Digital Sky Survey (SDSS) that the galaxy population was considered strongly bimodal in colour \\citep{strateva01}. Even when considering other galaxy properties such as radial profiles, the natural division is into two galaxy populations \\citep{hogg02red,ellis05,ball06}. Given that large automated imaging surveys are better at defining a galaxy's colour than morphology, it is more natural to describe a galaxy as being on the ``red sequence'' or ``blue sequence'' rather than being an ``early type'' or ``late type''. This interpretation also has the advantage that galaxy colours are directly related to the star formation, dust and metal enrichment history of the galaxy and can thus be more readily interpreted in theoretical models. A key goal of galaxy evolution theory is to explain the bimodality and the relationships within each sequence, and there has been considerable work in this area recently (e.g.\\ \\citealt{kang05,menci05,menci06,SDH05}; \\citealt{AvilaReese06,bower06,cattaneo06,croton06,DB06,perez06}). The key ingredient of many of these models is the inclusion of feedback from active galactic nuclei (AGN). Although AGN feedback (or its equivalent) is implemented in different ways in each of these models, the overall effect is to suppress cooling in massive halos. For example, in \\citet{bower06} the AGN feedback suppresses quasi-hydrostatic cooling flows, while \\citet{croton06} adopt a semi-empirical description for the AGN power related to the Bondi accretion rate. Bimodality of galaxy colours does not directly result from these schemes, however, the AGN feedback allows the star formation rate parameterisation to be adjusted in such a way as to simultaneously obtain a good description of the colour distribution at faint magnitudes and a good match to the shape of the luminosity function. In these models red galaxies at faint magnitudes are predominantly satellite galaxies of brighter systems, while at bright magnitudes the central galaxies are also red because of the AGN feedback (e.g.\\ fig.~3 of \\citeauthor{bower06}). In the real universe, however, the association between galaxy colours and their location in the halo is unlikely to be so simple \\citep{weinmann06}, and measurements of the dependence of galaxy colours on luminosity and redshift are an important constraint on the new generation of galaxy formation models. Analysis of the relationship between galaxy colours and environment will place important constraints on the processes defining galaxy evolution. This paper provides a detailed analysis of the variation of the bimodality with stellar mass and environment. The plan of the paper is as follows: in \\S\\,\\ref{sec:earlier-work}, we review a series of papers providing the buildup to this paper; in \\S\\,\\ref{sec:data}, we describe the data; in \\S\\,\\ref{sec:results}, we present the results; in \\S\\,\\ref{sec:discussion}, we discuss the implications and compare the data with models; and in \\S\\,\\ref{sec:summary}, we summarise the main results. Fits to the mass functions are presented in the Appendix. The data represented in this paper is available at http://www.astro.livjm.ac.uk/$\\sim$ikb/research/ or upon request. \\subsection{Previous work} \\label{sec:earlier-work} \\citet[hereafter Paper~I]{baldry04} characterised the volume-averaged colour-magnitude distribution of galaxies by fitting double Gaussian functions to the colour histograms in magnitude bins. This showed that the red sequence transitioned from a broader colour distribution at $M_r\\sim-19$ to a narrow distribution for massive early types at $M_r\\sim-21$ while the blue sequence become significantly redder over the range $-20$ to $-22$ (using centrally-weighted colours). This supported the suggestion of a transition in galaxy properties around $10^{10.5}\\,\\Msun$ \\citep{kauffmann03B} and the former is consistent with faint red-sequence galaxies forming more recently \\citep{delucia04,kodama04} but not necessarily in the richest clusters \\citep{Andreon06}. The difference in star formation history for galaxies below and above the transition mass may be related to the balance between hydrostatic versus rapid cooling \\citep{BD03,keres05,bower06,croton06}, or to the diversity of star formation histories of satellite galaxies at low masses \\citep{Bower91}. \\citet[hereafter Paper~II]{balogh04} analysed H$\\alpha$ emission strength as a function of galaxy environment. The H$\\alpha$ equivalent width (EW) distribution is bimodal with the distribution of the star forming population (blue sequence) not depending strongly on environment. The fraction of galaxies with ${\\rm EW}>4$\\AA\\ varied continuously with environmental density and there was no evidence of a break density that had been reported before \\citep{lewis02,gomez03}. This demonstrated the importance of describing the variation of the bimodal population by comparing the number of galaxies within each population rather than using a quantity averaged over both populations. In Paper~II, the results indicated a dependence of the star-forming fraction on scales of about 5\\,Mpc (after accounting for local environment). However, the interpretation is difficult because the small-scale measurement is noisy and the large-scale measurement could actually be adding information about the small-scale environment \\citep{blanton06}, and \\citet{kauffmann04} found no environmental dependence on scales larger than 1\\,Mpc. It is also known that for galaxies in clusters properties can depend on environment for scales smaller than 1\\,Mpc \\citep{Dressler80,WG91}. These argue for an $N$th nearest neighbour approach to measuring environmental density as the radius is smaller in high density regions and expands in low density regions where there are often insufficient galaxies on small scales (for a fixed luminosity cut). \\citet[hereafter Paper~III]{balogh04bimodal} extended the double Gaussian fitting of Paper~I to colour histograms across environment and luminosity. At fixed luminosity, the mean positions of the sequences become marginally redder with environmental density. For the red sequence, this can be explained by a small difference in age between low- and high-density environments of $\\sim2$\\,Gyr \\citep{thomas05} or less \\citep{hogg04,bernardi06}. In contrast, the fraction of red-sequence galaxies varied strongly with environment as measured by projected density. No effect related to velocity dispersion of a group or cluster was detected which is consistent with the morphology-density relation being similar in groups and clusters \\citep{PG84}. \\citet{baldry04conf} confirmed the shifts in colour of the red and blue sequences, 0.05 and 0.1 in $u-r$, respectively, over a factor of 100 in projected density. The red fraction varied from 0\\% to 70\\% for low-luminosity galaxies and 50\\% to 90\\% for high-luminosity galaxies (see also \\citealt{tanaka04}). The effects of environment and luminosity could be unified in that the fraction of red-sequence galaxies was related to a combined quantity: $\\Sigma_{\\rm mod} = (\\Sigma/\\perMpcsq) \\, + \\, (L_r / L_{r,{\\rm norm}}) \\,$ where $\\Sigma$ is the projected density and $L_{r,{\\rm norm}}$ is luminosity of a galaxy with $M_r=-20.2$. In this paper, this effect is explored in more detail using a larger sample and by converting luminosity to stellar mass. ", "conclusions": "\\label{sec:discussion} The advent of large-volume redshift surveys has greatly expanded the studies of galaxy populations as a function of environment that had traditionally been based around clusters \\citep{Dressler80} to the full range from `void' to cluster. There are many studies including Paper~II to this paper (\\S\\,\\ref{sec:earlier-work}) that focus on the variation of galaxy properties with local density \\citep{kauffmann04,tanaka04,KFS05,alonso06,haines06,mateus06,patiri06}. Other studies have analysed the average density as a function of galaxy property \\citep{hogg03,blanton05enviro}, the clustering properties of galaxies \\citep{wild05,zehavi05,li06}, and the relationship between galaxies and dark matter derived from galaxy-galaxy weak lensing \\citep{gray04,mandelbaum06}. The data available for $z<0.2$ is the most comprehensive, e.g., from SDSS and 2dFGRS data, but gains have been made at higher redshift \\citep{wilman05,yee05,cooper06,ilbert06}. While local density measurements illuminate various environmental trends in the data, galaxies in semi-analytical models are primarily associated with dark-matter halos. The main focus of this discussion is comparing models with the data by `measuring' the analogous quantity to projected galaxy density in the models. Our results show that the local galaxy population divides neatly into two types, and that the fraction of each type depends exclusively on a combination of stellar mass and local environment. Importantly, the environmental dependence is not a second-order effect, but is at least as important as stellar mass in determining the fraction of red galaxies in a population. By contrast, in simple galaxy formation models, the characteristics of a galaxy are usually determined primarily by the mass of its dark matter halo, which is closely related to the stellar mass of the galaxy. That is because the formation time, cooling rate and merging history of a halo are most strongly related to its present day mass. The most important environmental consideration in most of these models is that galaxies at the centre of a dark matter halo are treated differently from non-central, or satellite galaxies \\citep[e.g.][]{cole00}. Only the central galaxies are assumed to possess a halo of hot gas that can potentially cool and replenish the disk. In addition, there are second-order effects related to the large-scale density field that, qualitatively at least, can mimic some of the observed trends with environment \\citep{maulbetsch06}, depending on how the details of star formation and feedback are treated. Our results lead naturally to two questions in this context. One is, how does our measurement of environment, $\\Sigma$, relate to the dark matter density field $\\Delta\\rho/\\rho$? And the other is whether or not galaxy formation models that are successful in other respects are also able to reproduce the observed dependence of galaxy colour on environment. We will address these questions using the $z=0$ output of the Virgo Consortium's Millennium Simulation. The details of the dark matter are described in \\citet{springel05nat}, and we will compare with the galaxy formation models of both \\citet{croton06} and \\citet{bower06}. These models are improved over earlier efforts \\citep[e.g.][]{cole00} in that, by including a model of feedback from AGN, they are able to better match the observed colour distribution as a function of galaxy luminosity. Specifically, the inclusion of AGN feedback removes most of the bright blue galaxies, and increases the colour difference between the red and blue populations \\citep[see also][]{SDH05}. Although both models include feedback from radio-jets and AGN, the models use different schemes to implement the feedback. \\citet{croton06} compute an energy feedback rate based on a semi-empirical model involving the mass of the host halo and that of the central black hole. Their paper discusses how the expression they use is related to the Bondi accretion rate. In contrast, the \\citet{bower06} model assumes that AGN feedback will be self-regulating if the cooling time is long compared to the sound crossing time of the system so that the cooling of gas takes place from a quasi-hydrostatic atmosphere \\citep{Binney04}. The model also requires that central black hole is sufficiently massive to provide heating rate. In addition to the treatment of AGN jets, the models also differ in many details of the implementation of cooling galaxy merging, star bursts and many other factors. \\subsection{Galaxy versus dark matter densities in the models} To construct a mock-observational sample to compare with our data, we select galaxies at $z=0$ from the simulation. We restrict most of the analysis to galaxies with stellar masses $\\mass>10^{10}\\Msun$, which typically belong to dark matter halos with at least 100 particles (so their merger histories are reasonably well resolved). However, to define the local density we will select a luminosity-limited sample (see below) that includes lower-mass galaxies. The redshift of each galaxy is determined as $z=\\left(H_\\circ r_z+v_z\\right)/c$, where $r_z$ and $v_z$ are arbitrarily taken to be the position and peculiar velocity in the $z$-axis-coordinate of the simulation. The simulation box is 714\\,Mpc on a side, large enough that we can select the redshift range $0.01 0$. Yet, not all CCs that have $E_{k} - \\tau_{k} < 0$ are necessarily gravitationally bound, but are simply in the process of assembling from less dense gas by turbulent ram pressure. d) Despite their dynamical nature, some CCs are observed to have $W+\\Theta_{VT} \\sim 0$ ($\\Theta_{VT}=2~(E_{th}+E_{k}-\\tau_{th}-\\tau_{k})+E_{mag}+\\tau_{mag}$) and $|W|/|\\Theta_{VT}| \\sim 1$ which can lead them, erroneously, to be cataloged as being in a state of {\\it virial or magnetostatic equilibrium} or 'virialized', despite their dynamical nature. Additionally, we find, in all simulations, that the gravitational term must always be close to the gravitational energy and that the mass distribution outside the CCs does not have a significant contribution in distorting the CCs structure by external gravitational torques. e) There is no one-to-one correspondence between the state of gravitational boundedness of a CC as described by the virial balance analysis (i.e., gravity versus other virial energy terms) and as implied by the classical gravitational binding indicators $J_{c}$, $\\mu_{c}$, and $\\alpha_{vir}$. In general, from the virial analysis we observe that only the inner regions of the clumps (i.e., the dense cores selected at high density thresholds) are gravitationally bound, whereas Jeans number estimates of the same clumps tend to show that the objects are gravitationally bound at all threshold levels. On the other hand, the calculated $\\alpha_{vir}$ values not only shows that the clumps are more gravitationally bound at the lower threshold levels as the Jeans numbers, but also indicate a number of gravitationally bound objects always in excess of what is yielded by the virial analysis. Preliminary results by Dib \\& Kim (2006) show that the $J_{c}$ estimates and the corresponding energy ratio, when including the thermal surface energy ($E_{th}-\\tau_{th}/|W|$) are well correlated, and that the virial parameter $\\alpha_{vir}$ and the corresponding energy ratio $(E_{k}+E_{th}-\\tau_{k}-\\tau_{th})/|W|)$ or ($(E_{k}-\\tau_{k})/|W|$) are poorly correlated. A future study (Dib et al., in preparation) will investigate these relations in more detail for different models as well as the correlation between $\\mu_{c}$ and the magnetic energy to gravitational term ratio. f) The exponent, $\\epsilon$, of the virial parameter-mass relation, $\\alpha_{vir} \\propto M_{c}^{\\epsilon}$, shows a better agreement with observationally derived values for the case of the moderately supercritical cloud in which the most massive CCs are near critical ($\\mu_{c} \\sim 1$) , and a lesser good agreement for the subcritical (in which most CCs are subcritical) and strongly supercritical (in which the most massive object is supercritical) and non-magnetic cases. This result is not extremely conclusive yet because of the large uncertainties in the observational values of $\\epsilon$ and the limited statistics in our simulations. Yet, it might reflect an agreement with the recent predictions that most observed CCs have a mass-to magnetic flux ratio that is supercritical or subcritical by factors of 2 with the median value being around $\\mu_{c} \\sim 1$ (Crutcher \\& Troland 2006). g) In the non-magnetic simulation, we observe the formation of a core which possesses all the structural and dynamical properties of the Bok globule Barnard 68 (B68). This core is gravitationally bound. Such bound cores may survive the ionizing front of a \\ion{H}{2} region formed elsewhere in the cloud and which expands in the cloud's clumpy distribution (e.g., Mellema et al. 2006; Will Henney, private communication). The ionizing front can evacuate the gas around the core, and leave it confined by a surrounding warm gas such as in the case of B68. In the set of simulations presented in this work, we have used an isothermal equation of state to describe the gas physics. However, several authors have argued recently that deviations from isothermality in molecular clouds might lead them to have distinct structural and/or dynamical properties than isothermal clouds (e.g., Larson 2005). Li et al. (2003) and Jappsen et al. (2004) showed that the fragmentation process is dependent on the polytropic exponent which describes the equation of state. Dib et al. (2004) showed that the average density-size relation for non-isothermal clouds follows the observed $\\bar{n} \\propto R_{c}^{-1}$ Larson relations (Larson, 1981) unlike the isothermal clouds and their substructure is which this relation is not found (e.g., Ballesteros-Paredes \\& Mac Low 2002; Li et al. 2004). Hennebelle \\& Inutsuka (2006) discuss the possibility that warm gas can survive inside MCs. However, Pavlovski et al. (2006), showed that deviations from isothermality are only visible in terms of the high temperature zones behind shock waves. The detailed effects of a non-isothermal equation of state on the properties of individual cores has not been investigated in detail so far, and is left to a future work." }, "0607/astro-ph0607132.txt": { "abstract": "We report the results of H$^{13}$CO$^{+}$(1--0), CO(1--0), and 3.3 mm dust continuum observations toward one of the strongest mm-wave sources in OMC-3, MMS 7, with the Nobeyama Millimeter Array (NMA) and the Nobeyama 45 m telescope. With the NMA, we detected centrally-condensed 3.3 mm dust-continuum emission which coincides with the MIR source and the free-free jet. %8 $\\mu$m and 24 $\\mu$m SPITZER point sources and a 3.6 cm VLA source. The size and mass of the dusty condensation are 1500 $\\times$ 1200 AU (P.A. $\\sim$ 170$^{\\circ}$) and 0.36 - 0.72 M$_{\\rm{\\odot}}$ (for $T_{\\rm{dust}}$ = 26 - 50 K), respectively. Our combined H$^{13}$CO$^{+}$ observations with the 45 m telescope and the NMA have revealed a disk-like envelope around MMS 7 inside the H$^{13}$CO$^{+}$ core. The size and the mass of the disk-like envelope are 0.15 $\\times$ 0.11 pc and 5.1 - 9.1 M$_{\\rm{\\odot}}$ (for $T_{\\rm{ex}}$ = 26 - 50 K), respectively. The combined map also shows that the outer portion of the disk-like envelope has a fan-shaped structure which delineates the rim of the CO(1--0) outflow observed with the NMA. %along the bipolar outflow direction. The position-velocity (P-V) diagrams in the H$^{13}$CO$^{+}$ (1--0) emission show that the velocity field in the disk-like envelope is composed of a dispersing gas motion and a possible rigid-like rotation. The mass dispersing rate is estimated to be (3.4 - 6.0) $\\times$ 10$^{-5}$ M$_{\\rm{\\odot}}$ yr$^{-1}$, which implies that MMS 7 has an ability to disperse ${\\sim}$10 M$_{\\odot}$ during the protostellar evolutional time of a few $\\times$ $10^{5}$ yr. One of the probable dispersing mechanisms is the associated molecular outflow, and another the stellar wind which has enough power ($\\sim$76 L$_{\\odot}$) to drive the dissipation, (4.2 - 7.4) $\\times$ 10$^{-3}$ L$_{\\rm{\\odot}}$. The specific angular momentum of the possible rotation in the disk-like envelope is nearly two orders of magnitude larger than that in low-mass cores. %at the same size scale. The turn-over point of the power law of the angular momentum distribution in the disk-like envelope ($\\leq$ 0.007 pc), which is likely to be related to the outer radius of the central mass accretion, is similar to the size of the 3.3 mm dust condensation. We propose that the intermediate-mass protostar MMS 7 is in the last stage of the main accretion phase and that the substantial portion of the outer gas has already been dispersed, while the mass accretion may still be on-going at the innermost region traced by the dusty condensation. %the outer dense gas has been already being dispersed, while at %the innermost region traced by the dusty condensation mass accretion may be still ongoing. ", "introduction": "In the last two decades, developments of millimeter and submillimeter interferometers have enabled us to establish %make high-sensitivity observations. In low-mass star-forming regions, %it has been revealed that there are gas accretions with $\\dot{M}~{\\sim}~10^{-5} - 10^{-6}~\\rm{M_{\\odot}~\\rm{yr^{-1}}}$ %toward the central stars and circumstellar disks, %and active molecular outflows (e.g. Hayashi et al. 1993, Saito et al. 1996, Momose et al. 1998). a standard scenario of low-mass star formation via disk accretion processes (e.g. Hayashi et al. 1993, Saito et al. 1996, Myers et al. 2000). %In addition, there are samples of outflowing shells with wide opening angles in low-mass cores, %such as around IRAS 04368+2557 (L1527) \\citep{oha97a}. On the other hand, high- and intermediate-mass star-formation has not been clearly understood because of the complexity and the lack of observational studies, respectively. %Furthermore, the formation and evolution of intermediate-mass stars have not been well understood, %because there are limited observational studies. %Since intermediate-mass star formation is less %complicated, it would provide us a key to extend the low-mass scenario %and to bridge a missing link between the high-mass and low-mass star formation. %Interesting questions are if intermediate-mass protostars have higher mass-accretion or outflow %rates than those in low-mass cases. Since intermediate-mass star-formation is less complicated than high-mass star-formation, it would provide us a key to extend the low-mass scenario and to bridge a missing link between the high-mass and low-mass star formation. Particularly, it is important to verify the mechanisms and the differences of the destruction or dissipation of the natal dense cores between intermediate- and low-mass star formation, because these phenomena are related to the termination of the mass accretion which determines the final mass of the central star \\citep{nak95}. A key to these questions is investigation of structures and kinematics of dense cores around intermediate-mass protostars. For these purposes, we have been conducting survey observations toward protostellar cores in the Orion Molecular Cloud -2 and -3 region (OMC-2/3), relatively close ($d$ = 450 pc; Genzel and Stutzki 1989) and the most representative intermediate-mass star-forming region. We, here, report the results of one of the typical intermediate-mass protostars, MMS 7 in OMC-3, in the H$^{13}$CO$^{+}$(1--0), $^{12}$CO(1--0) and 3.3 mm continuum emission with the Nobeyama Millimeter Array (NMA) and the Nobeyama 45 m telescope. The H$^{13}$CO$^{+}$ (1--0) molecular line is one of the most appropriate tracers of dense gas, which has a high critical density ($n_{\\rm{crit}}~{\\sim}~10^5~\\rm{cm^{-3}}$) and is usually detected toward low- to intermediate-mass protostellar envelopes \\citep{tak00, sai01, fue05}. MMS 7 is %located at the northern part of an integral-shaped filament in the Orion A giant molecular cloud and is one of the Class 0 candidates identified by the 1.3 mm continuum observations toward OMC-3 \\citep{chi97}. This object is also identified as CSO12 by the 350 $\\mu$m continuum observations \\citep{lis98} and is associated with the IRAS point source 05329-0508. %H$^{13}$CO$^{+}$(1-0) observations were also carried out by Aso et al. (2000) toward the OMC-2/3 region with the Nobeyama %45 m telescope and substantial dense gas in MMS 7 was detected. The dust mass derived from the 1.3 mm continuum emission and the bolometric luminosity are 8 M$_{\\odot}$ and 76 L$_{\\odot}$, respectively \\citep{chi97}. This bolometric luminosity corresponds to $\\sim$ 3 M$_{\\odot}$ or a A0 star at the ZAMS. Dense molecular gas around MMS 7 is also detected by H$^{13}$CO$^{+}$(1--0) observations \\citep{aso00}. MMS 7 is also associated with a bright reflection nebula, Haro-5a/6a \\citep{har53}, and 2MASS and mid-infrared sources are located at the root of the eastern reflection nebula \\citep{nie03}. At MMS 7, a giant molecular outflow (0.86 pc; Aso et al. 2000) is also observed along the east-west direction, and 2.12 $\\mu$m H$_2 ~v$=1-0 $S$(1) knots are detected up to 1.45 pc to the western side \\citep{yu97, sta02}. In addition, there is a 3.6 cm continuum source elongated along the large-scale outflow, which traces a free-free jet from the protostar. The relative isolation of MMS 7 along with the above properties makes this source one of the most appropriate objects to investigate detailed spatial and velocity structures of a dense core around an intermediate-mass protostar. %The H$^{13}$CO$^{+}$ (1--0) molecular line is one of the most appropriate tracers of dense gas, %which has a high critical density ($n_{\\rm{crit}}~{\\sim}~10^5~\\rm{cm^{-3}}$) and %is usually detected toward low- to intermediate-mass protostellar envelopes \\citep{oni00, tak00, fue05}. %H$^{13}$CO$^{+}$(1-0) observations of dense gas in the Taurus Molecular Cloud complex have been carried out by Saito et al. (2001). %They revealed that the dense gas dissipates as protostars evolve and classified the evolutional phase from class A to C. %H$^{13}$CO$^{+}$(1-0) observations were also carried out by Aso et al. (2000) toward the OMC-2/3 region with the Nobeyama %45 m telescope and substantial dense gas in MMS 7 was detected. %Our observations with the higher spatial resolution and higher sensitivity enable us to discuss details of the structure and %kinematics of the dense gas around the intermediate-mass protostar of MMS 7 on 10$^3$-10$^4$ AU scale . % \\clearpage ", "conclusions": "Our new observations in the H$^{13}$CO$^{+}$ line have revealed systematic velocity structures in the disk-like envelope around the intermediate-mass protostar of MMS 7. In the subsequent sections, we will discuss the gas kinematics in the disk-like envelope and the difference from that in low-mass counterparts using position-velocity (P-V) diagrams. \\subsection{Radial Velocity Structure of the Disk-like Envelope} \\subsubsection{P-V diagram} We show a P-V diagram along the minor axis of the disk-like envelope in Figure \\ref{H13_PV}$b$. In the P-V diagram, there are at least two components; one is an eastern component at the velocity range of $V_{\\rm{LSR}}$=10.7 to 11.5 km s$^{-1}$ ((i) in Figure \\ref{H13_PV}$b$) and the other western-component at the velocity of $V_{\\rm{LSR}}$=9.6 to 10.9 km s$^{-1}$. The western component is associated with the 3 mm dust continuum emission (vertical solid line in Figure 7b). In this component, ``X-shape'' velocity structures are discerned in the P-V diagram. One of the X-shaped velocity structures as indicated by (ii) in Figure \\ref{H13_PV}$b$ shows blue-shifted emission ($V_{\\rm{LSR}}$=9.6 to 10.0 km s$^{-1}$) at the west of the protostar and red-shifted emission ($V_{\\rm{LSR}}$=10.6 to 10.9 km s$^{-1}$) at the east of the protostar. The sense of this velocity structure is opposite to that of the associated molecular outflow. On the assumption that the H$^{13}$CO$^{+}$ emission traces the flattened disk-like envelope which is perpendicular to the associated outflow, this velocity structure can be interpreted as an expanding motion in the flattened disk (e.g. Kitamura et al. 1996). The other velocity structure in the X-shape, with two components at the velocity range of $V_{\\rm{LSR}}$=9.9 to 10.4 km s$^{-1}$ (blueshifted; (iii)-1 in Figure \\ref{H13_PV}$b$) and $V_{\\rm{LSR}}$=10.5 to 11.0 km s$^{-1}$ (redshifted; (iii)-2 in Figure \\ref{H13_PV}$b$), has the same velocity sense as that of the associated outflow. These two components are located outside the central condensed gas, suggesting that these components are not associated with the central protostar. Thus, it is natural to interpret that these gas components trace swept-up dense gas by the associated outflow perpendicular to the disk-like envelope. A similar example has also been reported from observations in L1228 \\citep{taf94}. They detected a velocity shift in the high-density tracer of the C$_3$H$_2$ (2$_{12}$ - 1$_{01}$) emission, which has the same velocity sense as that of the associated CO bipolar outflow, and they interpreted this shift as an interaction between the dense core and the high-velocity outflow. The eastern component, indicated by (i) in Figure 7, is located at 25$''$ east of the continuum peak position and is a part of the fan-shaped structure seen in the combined total intensity map. We interpret that this component is remnant dense gas interacting with the associated molecular outflow. In summary, the P-V diagram can be interpreted as a mixture of the dispersing gas along the disk-like envelope and the interacting dense gas with the associated outflow perpendicularly to the disk-like envelope. Presumably due to the insufficient spatial resolution, an infall gas motion in the protostellar source of MMS 7 has not been clearly identified in our observations. \\subsubsection{Properties of the Dispersing Disk} The velocity structure of the disk-like envelope shows a dispersing gas motion as discussed in the last section. The virial mass of the disk-like envelope is estimated to be 23-30 M$_{\\odot}$ assuming $D_{\\rm{env}}$ = 0.15 pc, $T_{\\rm{env}}$ = 26 - 50 K, and $\\Delta v$ = 1.0 km s$^{-1}$, which is larger than the LTE mass of the disk-like envelope (5.1 - 9.1 M$_{\\odot}$). This result suggests that the disk-like envelope is gravitationally unbound, supporting our interpretation of the dispersing gas motion. We can estimate physical parameters of the dispersing motion, such as an expanding velocity ($V_{\\rm{exp}}$), momentum ($P_{\\rm{exp}}$), expanding energy ($E_{\\rm{exp}}$), and mechanical power ($L_{\\rm{exp}}$), as $V_{\\rm{exp}}({\\equiv}~V_{\\rm{max}})~=~1.2~\\rm{km~s^{-1}}$, $P_{\\rm{exp}}=MV_{\\rm{max}}~\\rm{M_{\\odot} km~s^{-1}}$ , $E_{\\rm{exp}}=MV_{\\rm{max}}^2/2~\\rm{M_{\\odot} km ^{2}~s^{-2}}$, and $L_{\\rm{exp}}=MV_{\\rm{max}}^3/2R~\\rm{L_{\\odot}}$. We adopt the disk inclination angle from the plane of the sky to be $i~\\sim$ 80$^{\\circ}$, which is consistent with the morphology of the CO outflow and the reflection nebula. Table \\ref{COMP_T} lists the estimated physical parameters of the dispersing motion. The dispersing envelope has also been observed in the $^{13}$CO(1-0) emission around a low-mass YSO of DG Tau \\citep{kit96}. We compare physical parameters specified in the dispersing process of MMS 7 to those of DG Tau in Table \\ref{COMP_T}. These parameters of MMS 7 are two orders of magnitude larger than those in DG Tau. The interferometric observations of DG tau may suffer from the problem of the missing flux, which prevents us from making a direct comparison of these parameters. However, even if the interferometric observations recover only 10 \\% of the total flux from DG Tau, these parameters of MMS 7 are still one order of magnitude larger than those in DG Tau. These results suggest that the intermediate-mass protostar of MMS 7 has more active dispersing processes than low-mass counterparts. There are recent studies on the dispersing gas motion around Herbig Ae/Be stars \\citep{fue02}. Their studies show that intermediate-mass stars disperse $\\gtrsim$ 90 \\% of the total mass of the parent core during the protostellar phase. Particularly, these objects which have a spectral type later than B6 evolve with the substantial dispersing material, 3 - 80 M$_{\\odot}$, in the protostellar phase. We, here, define the mass-dispersing rate (mass loss rate) which is $\\dot{M}_{\\rm{out}}\\sim M_{\\rm{env}}{\\cdot}V_{\\rm{exp}}/R_{\\rm{env}}=~(3.4 -6.0){\\times}10^{-5}~ \\rm{M_{\\odot}~yr^{-1}}$ . This value implies that MMS 7 has an ability to disperse a mass of 0.9-1.5 M$_{\\odot}$ during the outflow dynamical time scale of 2.5$\\times$10$^4$ yr, which is estimated from the ASTE data. We speculate that the intermediate-mass protostar of MMS 7 will keep dispersing substantial circumstellar material during the protostellar phase and will eventually be a Herbig Ae star in $\\tau \\leq$ a few$\\times$10$^{5}$ yr, as proposed by Fuente et al. (2002). \\subsubsection{Driving Mechanisms of the Dispersing Disk} What is the mechanism of the dispersing process in the disk-like envelope ? One of the probable candidates is the associated bipolar outflow, which is blowing away the material along the outflow axis and at the surface of the disk-like envelope. There is a twisting structure near the center of the reflection nebula Haro-5a/6a in the K's-band seen in Figure \\ref{INT}, and several knots (at least three components) in the red-shifted CO(3-2) emission in Figure \\ref{12CO_ASTE} (a),(b) denoted by the dashed-blue line. The connected trajectory of those emission peaks shows a wiggled structure. These results suggest that the precessing bipolar outflow from MMS 7 is able to destruct the surface of the disk-like envelope in the east-west direction with a wide opening angle. The importance of bipolar outflows in the dispersal of ambient gas around low-mass protostars has already been pointed out by previous studies \\citep{oha97a, vel98, tak03b}. Takakuwa et al. (2003b) has reported a dispersing low-mass protostellar core around IRAM 04191+1522. The detected blueshifted CH$_{3}$OH components (see BLUE1 and BLUE2 of Figure 3 in Takakuwa et al. 2003b) are most likely to be formed in consequence of the interaction between the outflow and the ambient dense gas surrounding the protostar and pushed away from the natal cloud core. In addition, there is a possibility that the stellar wind from the central protostar mainly contributes to the dispersion along the direction which is perpendicular to the outflow. A low-velocity wind ($\\sim$100 km s$^{-1}$) with a wide opening angle, ($\\sim$100$^{\\circ}$) is detected toward a low-mass protostar L1551 IRS 5 \\citep{pyo02, pyo05}. There are also studies on the stellar wind around Herbig Ae/Be stars, which show the derived mass-loss rate in the ionized gas $\\sim$10$^{-8}$ - 10$^{-7}$ M$_{\\odot}$ yr$^{-1}$ (e.g. Skinner et al. 1994). Therefore, it is possible that the stellar wind from MMS 7 drives the mass dispersion with the estimated rate of (3.4 - 6.0)$\\times$10$^{-5}$ M$_{\\odot}$ yr$^{-1}$. In fact, the bolometric power of MMS 7 is 76 $\\rm{L_{\\odot}}$, much larger than the mechanical luminosity of the dispersing motion, $L_{\\rm{exp}}=(4.2 - 7.4){\\times}10^{-3}~\\rm{L_{\\odot}}$ (see Table \\ref{COMP_T}), and is energetically sufficient to drive the expanding motion. \\subsection{Rotating Envelope} The P-V diagram along the major axis in Figure \\ref{H13_PV}$c$ shows that the velocity of the disk-like envelope increases with distance up to 5200 AU from the center. %One of the possible interpretations is the global motion of the OMC filament. %Because of systemic velocity of MMS 7, 10.6 km s$^{-1}$, shifted to the other OMC-3 sources in the filament, $\\sim$11.3 km s$^{-1}$. One of the possible (and simple) interpretations of this velocity gradient is rigid rotation in the disk-like envelope, although, it is difficult to clearly discern it owing to the mixture of the kinematics of the filament/core/envelope. %We, now, assume this velocity gradient as a rigid-rotation in following discussion. %This velocity structure is likely to suggest the presence of the rigid rotation in the disk-like envelope. Assuming the disk inclination of $80^{\\circ}$ and the radius of $5200$ AU, the timescale of the rigid rotation in this flattened disk is estimated to be $1.5 \\times 10^5$ yr, which is much longer than the outflow dynamical time scale of 2.5$\\times$10$^4$ yr. Red and blue plots in Figure \\ref{ang} show the local specific angular momentum in the disk-like envelope as a function of the radius around MMS 7. Our data demonstrate the angular momentum distribution on a size scale of $0.007 -0.02$ pc (corresponding to 1500 - 4700 AU) inside the single source. For comparison, we also plot the angular momentum of low-mass dense cores and circumstellar disks (Ohashi et al. 1997b, Goodman et al. 1993). The value of the specific angular momentum around the MMS 7 is nearly two orders of magnitude larger than that of the low-mass counterparts at the same size scale. The specific angular momentum distribution of the disk-like envelope around MMS 7 has a power-law index of $\\sim$ 1.8 which is closer to the index of rigid rotation, that is, 2.0. On the other hand, the specific angular momenta of low-mass NH$_{3}$ cores observed by Goodman et al. (1993) are fitted by a power-law index of $\\sim$ 1.6 as indicated by a dashed line in Figure {\\ref{ang}}. Furthermore, we note that there is no clear turn-over point of the power-law in the plot of MMS 7 as the low-mass case, which suggests that the turn-over point is smaller than that of low-mass NH$_{3}$ cores ($\\sim$ 0.03 pc) and is smaller than 0.007 pc, on the assumption that the value of the specific angular momentum at the innermost part of MMS 7 is the same as that of the low-mass value. The turn over point of the power law indicates the minimum size of the rotation where the specific angular momentum is not constant as a function of radius. At the dynamical collapsing region, the specific angular momentum should be constant in the entire radius. Therefore, the turn-over point can be considered as a starting point of the dynamical collapse toward the central protostar (Ohashi et al. 1997b). From these considerations, we suggest that the infalling radius in MMS 7 is smaller than that of low-mass counterparts. If dynamical collapse in MMS 7, if any, follows the inside-out collapse model \\citep{shu77}, the smaller infalling radius implies that the period passed after the accretion started is shorter than that of low-mass cores (a few $\\times$10$^{5}$ yr). %The other possibility to explain this velocity structure is the global velocity gradient in the cloud flament. %Because of systemic velocity of MMS 7, 10.6 km s$^{-1}$, shifted to the other OMC-3 sources in the filament, $\\sim$11.3 km s$^{-1}$. %In the subsesquent papers, we will also discuss the global velocity structure of OMC-2/3 filament with the extensive 45 m observations. \\subsection{Evolutional Stage of the Intermediate-mass Protostar MMS 7} The disk-like envelope, with the fan-shaped structure and the expanding motion around MMS 7, has a larger dispersing activity than that of low-mass cores (see section 4.1). MMS 7 has a large-scale (${\\sim}$1 pc) active CO outflow (Figure \\ref{12CO_ASTE}) which has a dynamical time scale of $\\tau_{\\rm{dyn}}{\\sim}2.5{\\times}10^{4}$ yr. These observational results suggest the presence of the significant mass-loss activity in the intermediate-mass protostar of MMS 7. Then, is there any mass supply onto the central star from the surrounding gas around MMS 7 ? In this subsection, we will discuss a possibility of the presence of mass accretion and the evolutionary stage of MMS 7. Although an infall motion is not clearly seen in our H$^{13}$CO$^{+}$ observations, the presence of the molecular outflow implies the existence of the mass accretion toward MMS 7 \\citep{bon96}. To make a 3 M$_{\\odot}$ protostar of MMS 7, a large accretion rate ($\\dot{M} \\sim 1.2 \\times 10^{-4}~\\rm{M_{\\odot}~yr^{-1}}$) is required over the outflow dynamical time (we, here, assume that $\\tau_{\\rm{dyn}}{\\sim} \\tau_{\\ast}$ and constant $\\dot{M}$). We found the smaller turn-over radius of the specific angular momentum distribution, which implies that the radius of the infalling region, $\\leq$ 0.007 pc (1500 AU), is smaller than low-mass counterparts. The size of the dusty condensation around MMS 7, 1500 $\\times$ 1200 AU, is similar to the inferred radius of the infalling region. Thus, it is possible that inside the dusty condensation, where the present spatial resolution is not high enough to trace the internal velocity structure, there is a substantial material which is accreting onto the central protostar. These considerations suggest that the intermediate-mass protostar MMS 7 is formed with higher mass-accretion and mass-loss rates than those of low-mass counterparts. The small infalling radius if present, short dynamical time of the outflow, and the presence of substantial circumstellar material, could imply that MMS 7 is in the early evolutional phase. On the other hand, our observational results show the energetic dispersing activity. Therefore, the early phase intermediate-mass protostar MMS 7 is already dispersing the surrounding envelope. Future interferometric observations with higher spatial resolution is required to investigate details of the innermost accreting region around MMS 7." }, "0607/astro-ph0607339_arXiv.txt": { "abstract": "We point out that the inclusion of a string component contributing around 5\\% to the CMB power spectrum amplitude on large scales can increase the preferred value of the spectral index $n_{s}$ of density fluctuations measured by CMB experiments. While this finding applies to any cosmological scenario involving strings, we consider in particular models of supersymmetric hybrid inflation, which predict $n_{s} \\stackrel{>}{{}_\\sim} 0.98$, in tension with the CMB data when strings are not included. Using MCMC analysis we constrain the parameter space allowed for $F$- and $D$-term inflation. For the $F$-term model, using minimal supergravity corrections, we find that $\\log\\kappa= -2.34\\pm 0.38$ and $M= (0.518\\pm 0.059)\\times 10^{16}{\\rm GeV}$. The inclusion of non-minimal supergravity corrections can modify these values somewhat. In the corresponding analysis for $D$-term inflation, we find $\\log\\kappa= -4.24\\pm 0.19$ and $m_{\\rm FI}= (0.245\\pm 0.031)\\times 10^{16}{\\rm GeV} $. Under the assumption that these models are correct, these results represent precision measurements of important parameters of a Grand Unified Theory. We consider the possible uncertainties in our measurements and additional constraints on the scenario from the stochastic background of gravitational waves produced by the strings. The best-fitting model predicts a $B$-mode polarization signal $\\approx 0.3 \\mu {\\rm K}$ rms peaking at $\\ell\\approx 1000$. This is of comparable amplitude to the expected signal due to gravitational lensing of the adiabatic $E$-mode signal on these scales. ", "introduction": "The publication of the most recent results from the Wilkinson Microwave Anisotropy Probe~\\cite{wmap} (WMAP) has focused attention in the direction of precision constraints on inflationary models~\\cite{PrecCon} believed to be the origin of the initial spectrum of density fluctuations~\\cite{DensityPerturbations}. By measuring the power spectrum of the $E$-mode polarization of the cosmic microwave background (CMB) on large-scales, the WMAP data constrains the optical depth to reionization, $\\tau_{\\rm R}=0.09\\pm 0.03$, with consequent improvement to parameters such as the spectral index of density fluctuations, $n_{s}=0.951^{+0.015}_{-0.019}$, which are degenerate with $\\tau_{\\rm R}$. Although ad-hoc inflationary models can be constructed with a range of values of $n_{s}$, some specific models predict very narrow ranges for $n_{s}$ making them vulnerable to the very tight constraints now available. One such class of models is supersymmetric (SUSY) hybrid inflation~\\cite{hybrid}, which has $F$-term~\\cite{CLLSW,DSS} and $D$-term~\\cite{DTerminf} variants. These scenarios are particularly attractive since the potential for the inflaton is flat at tree level. It acquires corrections from loop effects and when further tree-level, non-renormalisable Planck scale suppressed operators are added, and it can naturally meet the slow-roll conditions. Within certain minimal realisations, one is left with the choice of only one dimensionless coupling constant $\\kappa$ and a mass scale $M$. One can deduce the amplitude of curvature perturbations $P_{\\cal R}$, their spectral index $n_{s}$ and the dimensionless string tension $G\\mu$ from $\\kappa$ and $M$ (up to a weak dependence on the unknown reheat temperature, $T_{R}$). Therefore, these models may be considered as rather predictive when compared to other scenarios for inflation. One can make an analytic estimate of the scalar spectral index, which yields $n_{s} \\stackrel{>}{{}_\\sim} 1-{1/N_{\\rm e}}$, where $N_{\\rm e}$ is the number of e-foldings (measured from the end of inflation) when a particular observed scale leaves the horizon (or conversely when it comes back inside the horizon). For standard estimates of the reheat temperature, $N_{\\rm e}\\approx 50$, making $n_{s} \\stackrel{>}{{}_\\sim} 0.98$ a prediction of the simplest version of these models. If one takes the quoted observational constraint seriously, such a model would be excluded at around the $2-\\sigma$ level under the assumption that only adiabatic density fluctuations are created during inflation. However, in these models cosmic strings will be formed at the end of inflation, if the phase transition induces the spontaneous breakdown of a ${\\rm U}(1)$ symmetry by the waterfall fields. Since the predicted energy scale of these strings is around the grand unified (GUT) scale, they may contribute significantly to the observed perturbations~\\cite{jean97}, creating an interesting phenomenology~\\cite{WB,CHMb}. In this paper, we point out that the inclusion of a sub-dominant string contribution of around 5\\% to the large scale power spectrum amplitude of the CMB can increase the preferred value for the spectral index up to $n_s\\approx 0.98$ (and the maximum allowed value at $2-\\sigma$ level up to $n_s\\approx 1.02$), something which is a generic point valid for all models of inflation which produce cosmic strings. Naively, it may seem to be a grotesque violation of Occam's razor to have two sources of fluctuations with nearly equal amplitude; in no way do we claim that the data requires the additional string contribution in a Bayesian model selection sense. But as we shall describe, a string contribution of the required size arises very naturally in the class of models under consideration here. Note moreover, that these models constitute an attempt of fusing together the areas of particle physics and cosmology and may therefore be considered to be more attractive when put in the wider context. We then proceed to constrain the parameters $\\kappa$ and $M$ for specific realizations of $F$-term and $D$-term hybrid inflation models. Our results show that the inclusion of the string contribution is critical to determining the correct constraints on the parameters. We note that the upper bound on the string tension $\\mu$ has recently been discussed by a number of authors~\\cite{pog,fraisse05,fraisse06,slosel,bevis}. Their basic conclusion, using a variety of different methods, has been that there is a $2-\\sigma$ upper limit of $G\\mu < (2-3)\\times 10^{-7}$, something which we shall confirm. However, they have ignored the effect strings have on the preferred value of $n_{s}$. Qualitatively similar ideas were pointed out in ref.~\\cite{hind} in the context of constraints on the global texture model using the first year WMAP data. At that point in time, the constraint on $n_{s}$ was not as tight as is the case now and, therefore, the necessity of including the defect contribution was not so critical. ", "conclusions": "We have computed up-to-date precision constraints on the parameters $\\kappa$ and $M$ for $F$-term inflation, and $\\kappa$ and $m_{\\rm FI}$ for $D$-term inflation. Assuming that the data is correct and that one of these minimal SUSY hybrid inflation models is correct, we have shown that one can measure accurately parameters of supersymmetric models at Grand Unified scales. In particular, we have shown that the inclusion of the effects of strings is crucial to establish correct constraints. For comparison, we have performed most analyses also for the case where the string contribution is artificially excluded. Throughout our discussion we have also highlighted potential corrections to the constraints for the most minimal models, in particular we investigated the effect of the leading non-minimal SUGRA correction and of the tadpole induced by soft SUSY breaking. Without taking account of strings, one may infer from the WMAP3 data~\\cite{wmap} that the minimal $F$ term models, which predict $n_{s} \\approx 0.98$, are in tension with the data at $2-\\sigma$ level and that the small coupling domain of hybrid inflation, where $n_{s} \\approx 1$ is ruled out at $3-\\sigma$ level~\\cite{PrecCon,JP3}. The latter constraint would in fact rule out $D$-term inflation, which requires the parameter $\\kappa$ to be small in order not to have an excess contribution of strings to the perturbation spectrum. However, the self-consistent analysis performed here reveals that $D$-term inflation is not yet ruled out and, moreover, that the minimal six parameter $F$-term models with strings fit the data as well as the standard six parameter model. There are also some theoretical modelling uncertainties associated with establishing the cosmic string spectrum. We have shown that if $1.3\\le\\beta_r\\le 2.8$ then the effect on the string power spectrum is around $20\\%$, which is $\\sim 1\\%$ on the total. Sensible variations of $\\xi_{r}$ and $\\langle v^2\\rangle_{r}^{1/2}$ also induce variations of $<20\\%$. It appears that further refinement of the string power spectrum beyond this level of understanding is unnecessary for obtaining accurate constraints on hybrid inflation models. In addition to theoretical uncertainties there are a number of other issues associated with the use of the data which need to be carefully assessed: the discrepancy between the data and a SUSY hybrid inflation model with no strings and $n_{s}\\approx 0.98$ is only around $2-\\sigma$ (95\\% confidence level). Chief amongst the uncertainties is how the polarized foregrounds are extracted and how line-of-sight effects such as gravitational lensing~\\cite{CL} and the Sunyaev-Zeldovich (SZ) effect are dealt with. For example in the analysis performed by the WMAP team an SZ contribution was included in the analysis with an amplitude which was marginalized over. Since this contribution to the power spectrum will be increasing with $\\ell$, such an analysis will tend to lower the best fit value of $n_{s}$. Conversely, taking into account the gravitational lensing effect will have a tendency to increase the value of $n_{s}$. We have included neither in our analysis, presuming that they will cancel each other out. Once even higher precision data is available, from for example the PLANCK satellite, these effects may come to dominate the systematics. \\begin{figure}[htbp] \\epsfig{file=cls.eps,width=7.0cm} \\caption{Components of the temperature and polarization power spectra for the best-fitting $F$-term inflation model with an artificially large value of $r=0.1$ at $k=0.05{\\rm Mpc}^{-1}$. The solid lines are the adiabatic component, the dotted lines the string component and the dashed line is the $B$-mode component due to the gravitational lensing of the $E$-mode polarization of the adiabatic component. For the adiabatic and string components, the three curves (top to bottom) are the temperature power spectrum and that for the $E$- and $B$-mode polarization.} \\label{figure:bmode} \\end{figure} There is an important observational signature of these models which may be in reach in the near future. Inflationary models with non-zero $r$ predict that there will be $B$-mode polarization on large scales peaking around $\\ell\\approx 100$, but as we have already noted the SUSY hybrid inflation models predict $r<10^{-4}$ and a signal this weak is unlikely to ever be detected. However, as we have already pointed out the anisotropies created by cosmic strings create $B$-mode polarization since they do not distinguish between scalar, vector and tensor anisotropies. Fig.~\\ref{figure:bmode} shows both the adiabatic and cosmic string components to the temperature anisotropies and polarization for a model with an artificially large value of $r=0.1$ at $k=0.05{\\rm Mpc}^{-1}$. It can be seen that there is a $B$-mode polarization signal due to the cosmic string component, which has an amplitude of $\\approx 0.3\\mu{\\rm K}$ at around $\\ell\\approx 1000$. This has very different characteristics to that due to adiabatic tensor perturbations and is of similar amplitude to the contribution expected due to the conversion of $E$-mode polarization into $B$-mode by gravitational lensing. In ref.~\\cite{Seljak:2006hi} it was pointed out that string tensions of as low as $G\\mu\\approx 10^{-9}$ might be detectable in future CMB polarization missions if one is able to ``clean'' the lensing contribution using high resolution observations. Finally we should point out that a network of cosmic strings with $G\\mu\\sim 10^{-7}$ will lead to a number of other potentially observable effects. In particular, the decay of cosmic string loops could create a stochastic background of gravitational waves (see ref.~\\cite{Caldwell:1996en} and references therein) if the dominant decay channel for the strings is gravitation. Such a background is constrained by the lack of timing residuals in the observations of pulsars; the most recent constraint being that $\\Omega_{\\rm g}h^2<2\\times 10^{-9}$ at frequencies of $f\\approx 2\\times 10^{-9}{\\rm Hz}$~\\cite{pulsar}. If we assume that the absolute lower bound on the string spectrum is given by the ``red-noise'' spectrum generated by the decay of string loops in the radiation era, then one can use various parameters measured by string network simulations~\\cite{BB,AS} and the measured value of $\\Omega_{\\rm m}$ to compute a bound on $G\\mu$ as a function of the loop production size relative to the horizon, $\\alpha$. The results of doing this are presented in Fig.~\\ref{figure:pulsar}. We see that for small $\\alpha$ there is a plateau with $G\\mu<1.4\\times 10^{-7}$ and for larger values of $\\alpha$ (which are probably less likely) more tight constraints are possible. These results are considerably more stringent than those presented in ref.~\\cite{Caldwell:1996en} since at that time the limit $\\Omega_{\\rm g}h^2<9\\times 10^{-8}$ was used. Taken at face value these results appear to further constrain, but do not yet rule out, the $F$-term scenarios under consideration here to a narrow range of $\\log\\kappa\\approx -3$ and $M\\approx 4\\times 10^{15}{\\rm GeV}$, with a qualitatively similar situation in the $D$-term case. There are, however, numerous uncertainties, particularly in the details of string evolution which could substantially change the conclusions and therefore at this stage we feel that it would not be sensible to include these observations in our likelihood analysis. We have already noted that the anisotropy power spectrum that we would observe in the CMB is not that sensitive to the details of string evolution since it is sub-dominant and therefore it is unlikely that one would be able to unequivocally rule in or out these models on the basis of CMB measurements. Hence, it appears that improved observations of pulsar timing and a significantly better understanding of string network evolution, directed towards the pulsar bound, would be the best way of constraining these scenarios further. \\begin{figure}[htbp] \\epsfig{file=lim2.eps,width=7.0cm} \\caption{Constraints on $G\\mu$ from due to the absence of timing noise in pulsars against the loop production size $\\alpha$. The region above the solid line is ruled out at $2-\\sigma$ if $\\Omega_{\\rm g}h^2<2\\times 10^{-9}$ for $f\\approx 2\\times 10^{-9}{\\rm Hz}$. See note added in proof for the explanation of the dashed line.} \\label{figure:pulsar} \\end{figure}" }, "0607/astro-ph0607613_arXiv.txt": { "abstract": "{Model atmosphere fits to high-resolution optical spectra of \\wray\\ confirm the B~hypergiant classification of the massive companion to the X-ray pulsar \\object{GX301$-$2}. The models give a radius of 62~$R_{\\sun}$, an effective temperature of 18,100~K and a luminosity of $5 \\times 10^{5} \\, L_{\\sun}$. These values are somewhat reduced compared to the stellar parameters of \\wray\\ measured previously. The deduced mass-loss rate and terminal velocity of the stellar wind are $10^{-5}$~M$_{\\sun}$~yr$^{-1}$ and 305~\\kms, respectively. The interstellar $\\ion{Na}{i}$ D absorption indicates that \\wray\\ is located behind the first intersection with the Sagittarius-Carina spiral arm (1-2.5~kpc) and probably belongs to the stellar population of the Norma spiral arm at a distance of $3-4$~kpc. The luminosity derived from the model atmosphere is consistent with this distance ($3$~kpc). The luminosity of the wind-fed X-ray pulsar ($L_{X} \\sim 10^{37}$~erg~s$^{-1}$) is in good accordance with the Bondi-Hoyle mass accretion rate. The spectra obtained with UVES on the {\\it Very Large Telescope} (VLT) cover a full orbit of the system, including periastron passage, from which we derive the radial-velocity curve of the B hypergiant. The measured radial-velocity amplitude is $10 \\pm 3$~\\kms\\ yielding a mass ratio $q = M_{X}/M_{\\rm opt} = 0.046 \\pm 0.014$. The absence of an X-ray eclipse results in a lower limit to the mass of \\wray\\ of 39~M$_{\\sun}$. An upper limit of 68 or 53~M$_{\\sun}$ is derived for the mass of \\wray\\ adopting a maximum neutron star mass of 3.2 or 2.5~M$_{\\sun}$, respectively. The corresponding lower limit to the system inclination is $i > 44^{\\circ}$, supporting the view that the dip in the X-ray lightcurve is due to absorption by the dense stellar wind of \\wray\\ \\citep{Leahy02}. The ``spectroscopic'' mass of \\wray\\ is $43 \\pm 10$~M$_{\\sun}$, consistent with the range in mass derived from the binarity constraints. The mass of the neutron star is $1.85 \\pm 0.6$~M$_{\\sun}$. Time series of spectral lines formed in the dense stellar wind (e.g. $\\ion{He}{i}$ 5876~\\AA\\ and $\\rm H \\alpha$) indicate the presence of a gas stream trailing the neutron star in its orbit. The long-term behaviour of the $\\rm H \\alpha$ equivalent width exhibits strong variations in wind strength; the sampling of the data is insufficient to conclude whether a relation exists between wind mass-loss rate and pulsar spin period. ", "introduction": "\\wray\\ (BP~Cru) is the B-supergiant companion to the X-ray pulsar \\object{GX301$-$2}. Comparison of the mass functions derived for high-mass X-ray binaries (HMXBs) harbouring an X-ray pulsar shows that \\wray\\ is the most massive OB-star companion in these systems \\citep{Nagase89,Bildsten97}; HD153919, the O6.5 Iaf$^{+}$ companion to 4U1700-37, may have a higher mass ($M_{\\rm opt} = 58 \\pm 11$, \\citealt{Clark02}), but 4U1700-37 has not been proven to be an X-ray pulsar \\citep[e.g.][]{Hon04}. HMXBs are divided into two sub-groups: the Be/X-ray binaries (mainly X-ray transients) and OB-supergiant systems (of which \\wray\\ is a member). The transient character of the Be/X-ray binaries relates to ``outburst'' phases of the Be-type companion, which occur at irregular intervals, separated by years to decades. During such an outburst phase the Be star ejects matter which forms an equatorial disc. The crossing of the compact companion through the Be-star's equatorial disc leads to X-ray outbursts which recur with the orbital period. In OB-supergiant systems the X-ray source accretes from the strong stellar wind or is fed by Roche-lobe overflow \\citep[see][]{Kaper01}. In the latter case the much higher accretion rate results in an about 100 times higher X-ray luminosity ($\\sim 10^{38}$~erg~s$^{-1}$) and a short X-ray pulse period (seconds rather than minutes) in comparison to wind-fed systems. The mass function derived from pulse-timing analysis indicates that the mass of \\wray\\ is larger than 31.8~M$_{\\sun}$ \\citep{Sato86,Koh97}. Knowledge of the mass of \\wray\\ is important, because this information is used to determine the empirical lower mass limit for black-hole formation in a massive binary \\citep{VandenHeuvel84,Ergma98,Wellstein99}. The progenitor of \\object{GX301$-$2} was originally the most massive star in the system and left a neutron star (the X-ray pulsar) after the supernova explosion. The observed absence of X-ray eclipses sets an upper limit to the system's inclination, given an estimate of the radius of \\wray. The larger the radius, the lower the inclination, and the higher the lower limit to its (present) mass. Furthermore, the high mass of \\wray\\ may result in the formation of a black-hole -- neutron-star binary when \\wray\\ ends its life as a supernova or a gamma-ray burst. Based on a reclassification of its optical spectrum, \\citet{Kaper95} proposed that \\wray\\ is a B1~Ia+ hypergiant rather than a normal B supergiant \\citep{Parkes80}. Since a hypergiant is larger and more luminous than a supergiant of the same spectral type, this results in a larger mass (48~M$_{\\sun}$) and larger distance (5.3~kpc) of \\wray\\ than thought before. In fact, the radius adopted in \\citet{Kaper95} ($R_{\\star} = 87 \\, R_{\\sun}$) is larger than the Roche (and tidal) radius at periastron passage of this very eccentric ($e=0.462$) system. The interpretation of the pulse-period history (spin-up) of \\object{GX301$-$2} over the last decade \\citep{Pravdo95} and a reconsideration of the parameters of the binary system \\citep{Koh97} suggest that the distance, stellar radius and thereby the mass of \\wray\\ are less than proposed in \\citet{Kaper95}. From the cyclical occurence of X-ray flares, \\citet{Watson82} determined the orbital period of the system ($P_{\\rm orb} = 41.5$~d). It turns out that the periodic flare occurs just {\\it before} periastron passage \\citep[$\\sim 2$d, \\,][]{White84,Sato86,Chichkov95}. A similar X-ray lightcurve has been observed for 4U~1907+09 \\citep{IntZand98}; in this system the massive companion star also is a luminous supergiant with a dense stellar wind \\citep[O8-O9 Ia,][]{Cox05a}. Calculations by \\citet{Stevens88} of the dynamical effects of the neutron star on the stellar wind showed that a highly increased mass-loss rate from the primary can be expected at periastron passage. This provided the physical basis for the suggestion by \\citet{Haberl91} and \\citet{Leahy91} that the observed pre-periastron X-ray flares are due to the enhanced accretion rate during the passage of the neutron star through a gas stream in the stellar wind. \\citet{Pravdo95} found that also near apastron passage a periodic X-ray flare occurs; the gas-stream model cannot explain the apastron flare very well. They propose that both the asymmetry of the pre-periastron flare and the presence of the near-apastron flare can be explained by an equatorially enhanced stellar wind or a circumstellar disc around \\wray. The X-ray flares would occur when the X-ray source moves through the disc which is slightly inclined with respect to the orbital plane of the X-ray pulsar, as in the case of Be/X-ray binaries. The rapid spin-up episodes of \\object{GX301$-$2} discovered by \\citet{Koh97} suggest the formation of temporary accretion discs. The long-term spin-up trend of the X-ray pulsar observed since 1984 might be entirely due to such brief spin-up episodes. Numerical simulations carried out by \\citet{Layton98} confirm that tidal stripping in eccentric-orbit X-ray binaries can produce periodic flares. However, the tidally stripped mass only accretes when the supergiant is close to corotation at periastron, and the resulting X-ray flare occurs well {\\it after} periastron passage ($\\phi\\sim 0.2$) and not at apastron. The calculations further show that a transient accretion disc forms when the neutron star accretes from the tidal stream and persists for roughly half the binary period. This produces an extended epoch (many days) of spin-up reminiscent of the spin-up episodes observed by \\citet{Koh97}. \\citet{Layton98} remark that the X-ray flares are more likely to be due to enhanced accretion from an equatorial disc than to tidal stripping at periastion. \\begin{table}[!t] \\caption[]{Optical and near-infrared broad-band photometric parameters of \\wray\\ collected from literature; the infrared fluxes were observed with ISO. The fluxes are dereddened using the extinction law of \\citet{Mathis90} and $E(B-V)=1.9$. (1) \\citet{Bord76}; (2) \\citet{Hammerschlag76}; (3) \\citet{VanDishoeck89}; (4) \\citet{Coe97}; (5) \\citet{Glass79}; (6) ISO (see Sect.~\\ref{iso_observations}).} \\begin{center} \\begin{tabular}{llllll} \\hline \\hline Passband & $\\lambda$ & Mag & $f_{\\nu}$ & $A_{\\lambda}$ (mag) & $f_{\\nu}$ (cor) \\\\ & ($\\mu$m) & & (Jy) & ($R_{V}=3.1$) & (Jy) \\\\ \\hline U (1) & \\00.365 & 13.01 & 0.012 & 9.18 & 55.23 \\\\ B (1) & \\00.44 & 12.59 & 0.041 & 7.80 & 53.87 \\\\ V (2) & \\00.55 & 10.83 & 0.177 & 5.89 & 40.26 \\\\ I (3) & \\00.90 & \\07.6 & 2.049 & 2.82 & 27.52 \\\\ J (4) & \\01.25 & \\06.83 & 2.817 & 1.66 & 13.00 \\\\ H (4) & \\01.65 & \\06.11 & 3.526 & 1.04 & \\09.19 \\\\ K (4) & \\02.2 & \\05.72 & 3.194 & 0.63 & \\05.71 \\\\ L (5) & \\03.4 & \\05.25 & 2.304 & 0.30 & \\03.04 \\\\ LW2 (6) & \\06.75 & & 0.837 & 0.12 & \\00.94 \\\\ LW10 (6) & 11.5 & & 0.517 & 0.16 & \\00.60 \\\\ PHT03 (6) & 25.0 & & 0.089 & 0.08 & \\00.10 \\\\ \\hline \\end{tabular} \\end{center} \\label{tabphot} \\end{table} The main motivation of this paper is to provide a better and quantitative estimate of the stellar parameters of \\wray\\ based on model atmosphere fits to the optical line spectrum and the energy distribution. The mass ratio of the system is determined by measuring the radial-velocity orbit of \\wray. Further, we monitored \\wray\\ to search for the presence of a gas stream in the system. In the next section we describe the observations. Constraints on the distance and mass of \\wray\\ are discussed in Sect.~\\ref{constraints}. In Sect.~\\ref{modelling} we present the results of the spectrum modelling. In Sect.~\\ref{interaction} we study the interaction between the X-ray pulsar and the stellar wind, search for the presence of a gas stream and show that the observed X-ray flux is consistent with that expected from Bondi-Hoyle accretion. In the last section we summarize the conclusions and discuss the implications of the derived stellar parameters for the nature and evolutionary status of this binary system. ", "conclusions": "New optical spectra and infrared photometry strongly support the B hypergiant classification of \\wray\\ as proposed by \\citet{Kaper95}, even though the radius, distance and luminosity of \\wray\\ are somewhat reduced compared to the values listed in that paper: $R = 62 \\, R_{\\sun}$, $d \\sim 3$~kpc and $L = 5 \\times 10^{5} L_{\\sun}$. The radial-velocity curve, though hampered by the intrinsic scatter of individual measurements, clearly shows the orbital motion of \\wray\\ from which the radial-velocity amplitude is derived. Combined with the accurately determined orbital parameters of the X-ray pulsar, the mass ratio is set to $q = 0.046 \\pm 0.014$. The parameter missing to calculate the masses of the two stars is the orbital inclination. The system is not eclipsing, though the X-ray light curve shows a reduction in X-ray flux during the orbital phase interval when the X-ray source is behind the B hypergiant. That the orbital inclination indeed cannot be very low follows from the argument that the mass of the neutron star must be less than 3.2~M$_{\\sun}$ (causality limit) and likely even less than 2.5~M$_{\\sun}$ (maximum neutron star mass based on equation of state). As a consequence, the mass of \\wray\\ is less than 68 (53)~M$_{\\sun}$ and higher than 39~M$_{\\sun}$ (no X-ray eclipse). The lower limit on the mass of the neutron star is $1.85 \\pm 0.6$~M$_{\\sun}$, suggesting that the neutron star belongs to the high-mass peak in the bimodal neutron-star mass distribution as proposed by \\citet{Timmes96}, like Vela~X-1 \\citep{Barziv01} and probably 4U1700-37 \\citep{Clark02}. \\citet{VandenHeuvel84} propose that the system evolved from a 42+38 M$_{\\sun}$ binary, the lowest-mass realistic progenitor system fulfilling the condition that tandem evolution is avoided in a case~B scenario. \\citet{Wellstein99} propose that the system evolved in a case~A scenario, leading to a progenitor system of initially 25+24 M$_{\\sun}$. The phase of mass transfer has to be fully conservative in order to reproduce the current mass of \\wray\\ and its position in the Hertzsprung-Russell diagram. The surface chemical abundances of \\wray\\ derived from our spectra agree with the predictions of \\citet{Wellstein99}. Therefore, the progenitor mass of the neutron star may indeed be as low as 25~M$_{\\sun}$, rather than the 50~M$_{\\sun}$ proposed by \\citet{Kaper95}. As a consequence, the lower limit for black-hole formation in a massive binary, derived from this system, becomes 25~M$_{\\sun}$. The orbital modulation of spectral lines formed in the stellar wind of \\wray, such as $\\rm H \\beta$ and $\\ion{He}{i}$ 5876~\\AA, indicate the presence of a gas stream in the system. Such a gas stream was proposed to explain the peaks in the X-ray light curve near apastron and periastron. We find no observational evidence for an extended equatorial disc surrounding \\wray. The B hypergiant is not in corotation with the X-ray pulsar during periastron passage, which may be a complicating factor for models explaining the presence of a gas stream in the system originating from tidal interaction \\citep{Layton98}. However, it may well be that \\wray\\ exceeds its tidal lobe during periastron passage, enabling the formation of a gas stream in the system. The decrease in X-ray pulse period of \\object{GX301$-$2} from 1984 to 1991 may be due to an increase in the mass and angular momentum accretion rate related to a higher wind mass-loss rate of \\wray." }, "0607/astro-ph0607049_arXiv.txt": { "abstract": "If inflation was preceded by a radiation era then at the time of inflation there will exist a decoupled thermal distribution of gravitons. Gravitational waves generated during inflation will be amplified by the process of stimulated emission into the existing thermal distribution of gravitons. Consequently the usual zero temperature scale invariant tensor spectrum is modified by a temperature dependent factor. This thermal correction factor amplify the $B$-mode polarization of the CMB by an order of magnitude at large angles, which may now be in the range of observability of WMAP. ", "introduction": " ", "conclusions": "" }, "0607/astro-ph0607563_arXiv.txt": { "abstract": "Inverse Compton (IC) scattering by relativistic electrons produces a major component of the diffuse emission from the Galaxy. The photon fields involved are the cosmic microwave background and the interstellar radiation field (ISRF) from stars and dust. We expect which could be detectable by instruments such as GLAST. Even individual nearby luminous stars could be detectable assuming just the normal cosmic-ray electron spectrum. We present the basic formalism required and give possible candidate stars to be detected and make prediction for GLAST. Then we apply the formalism to the OB associations and the Sun, showing that the IC emission produced is not negligible compared to the sensitivity of current or coming detectors. We estimate that the gamma-ray flux from the halo around the Sun contributes to the diffuse background emission at few percent level. ", "introduction": "\\label{intro} In the early 90s there was already the idea to consider the IC gamma-ray emission generated by electrons accelerated by shocks in winds around hot stars (e.g.\\cite{Ref0}). In the present work, for the first time, we show that even the gamma-ray emission by the ambient cosmic-ray electrons via IC scattering of the stellar radiation field could be detected by GLAST. We begin with the simplest possible rough estimate to show that the IC emission from luminous stars could be visible. The optical luminosity of the Galaxy is about 3$\\times$10$^{10}$ L$_\\odot$, and a typical O star has 10$^{5}$ L$_\\odot$ i.e. about 10$^{-5}$ of the Galaxy. Consider such a star at a distance of 100 pc: compared to the entire Galaxy (distance to centre = 8.5 kpc) this inverse Compton source is on average about a factor 100 closer and hence the IC is 10$^{-5}$ $\\times$ 100$^2$ of the Galactic IC, suggesting it is significant. The IC luminosity $L_{IC}$ within a volume surrounding a star is proportional to the radius $r$ around the star times the optical luminosity of the star: $L_{IC}\\propto r~L_{STAR}$. The flux depends on the star's distance d: $flux_{IC}\\propto L_{IC}/d^{2}$ and for an angle $\\alpha \\propto ~r/d$, we obtain $flux_{IC}\\propto L_{STAR} ~\\alpha/d$. A more precise formulation is given in the next section. ", "conclusions": "\\label{sec:7} We have estimated the gamma-ray emission by IC scattering of cosmic-ray electrons with the radiation field around stars. We find that the contribution of the most luminous stars is non-negligible and even individual luminous stars could be detectable by GLAST. Moreover OB associations can contribute to the clumpiness of the emission. The same model applied to the Sun \\footnote{When this work had already been completed we learned about work by Moskalenko et al. (2006) on the Sun \\cite{Ref16}.} shows that the IC emission produced is significant and should be accounted for in diffuse background studies." }, "0607/astro-ph0607080_arXiv.txt": { "abstract": "We present results of N-body simulations aimed at understanding the dynamics of young stars near the Galactic center. Specifically, we model the inspiral of a cluster core containing an intermediate mass black hole and $N \\sim 50$ cluster stars in the gravitational potential of a supermassive black hole. We first study the elliptic three-body problem to isolate issues of tidal stripping and subsequent scattering, followed by full N-body simulations to treat the internal dynamics consistently. We find that our simulations reproduce several dynamical features of the observed population. These include the observed inner edge of the claimed clockwise disk, as well as the thickness of said disk. We find that high density clumps, such as that claimed for IRS13E, also result generically from our simulations. However, not all features of the observations are reproduced. In particular, the surface density profile of the simulated disk scales as $\\Sigma \\propto r^{-0.75}$, which is considerably shallower than that observed. Further, at no point is any significant counter-rotating population formed. ", "introduction": "In the past decade, observations have conclusively established the presence of a supermassive black hole (SMBH) at the Galactic Center (GC) (Sanders 1992, Haller et al. 1996, Ghez et al. 2003). Efforts to measure the mass of the SMBH have led to a large body of evidence regarding the stellar kinematics and mass distribution in this region (Genzel et al. 2003, Ghez et al. 2003, Sch\\\"odel et al. 2003, Paumard et al. 2005, Maillard et al. 2004, Eisenhauer et al. 2005, Ghez et al. 2005). The most surprising findings include the discovery of a number of young, massive stars closely orbiting the central object Sgr A*, young stellar populations in two possibly counter-rotating stellar disks orbiting Sgr A* further out, and small stellar clumps or associations (so-called ``comoving groups'') within these disks (Levin \\& Beloborodov 2003, Genzel et al. 2003, Lu et al. 2005, Maillard et al. 2004, Sch\\\"odel et al. 2005). The origin and peculiar kinematics of these structures beg theoretical explanation. Formation of stars near the SMBH is problematic because the strong tidal field is likely to shear and disrupt normal molecular clouds well before they can gravitationally collapse. Other formation scenarios include molecular cloud collisions, star-star collisions, and giants whose envelopes have been tidally stripped, but none of these is entirely satisfactory. Current opinion falls into one of two classes -- either formation of stars by gravitational instability in an AGN-like accretion disk (Kolykhalov \\& Sunyaev 1980, Shlosman \\& Begelman 1989, Morris 1996, Sanders 1998, Goodman 2003, Levin \\& Beloborodov 2003, Nayakshin \\& Cuadra 2005) or rapid inward transport (due to dynamical friction) and subsequent tidal disruption of a star cluster that formed at larger radii (Gerhard et al. 2001). Early numerical simulations of the latter scenario revealed that the cluster would not survive the infall to small radii unless extraordinary demands were placed on the mass and stellar density (McMillan \\& Portegies Zwart 2003, Kim \\& Morris 2003). An enhancement of this idea was the inclusion of an intermediate mass black hole (IMBH) at the center of the cluster, which served to both maintain the cluster potential well and slow internal relaxation (Hansen \\& Milosavljevi\\a'c 2003). Further direct simulations, including this refinement, again placed stringent demands on the cluster initial conditions, although the required core density decreased (Kim, Figer, \\& Morris 2004). Recent Monte Carlo and N-body simulations of this process have demonstrated that $\\sim 100$ stars can be transported via dynamical friction to about $1\\pc$ from the GC, with an IMBH formed naturally through a runaway process of stellar merger during the migration (Baumgardt et al. 2004, G\\\"urkan \\& Rasio 2005). However, these simulations could not accurately follow the further evolution of this cluster core because of algorithmic limitations. It is the principal goal of this paper to follow the physics of this process further inwards to determine to what extent this scenario may reproduce the observed features of the young star distribution. To date, the reliability of existing simulations has effectively ended at about $1\\pc$. A chief cause of failure of these codes is the presence of very strong tidal fields due to the nearby SMBH, which causes normally simple algorithms for energy conservation and treatment of close encounters to become delicate and highly complex, resulting in occasional failure. Since the majority of the interesting and puzzling observations have been made interior to this radius, simulations that might illuminate answers to these riddles are necessary. This paper discusses simulations of the dynamics of remnant cluster cores as they sink towards the GC, specifically focusing on the region interior to $1\\pc$. The paper is organized as follows: \\S~\\ref{sec:nummeth} describes the numerical methods of simulating the inspiral of general, three-body and $\\sim 50-$body systems, including the implementation of dynamical friction that creates the inspiral. \\S~\\ref{sec:3body} describes the three-body simulations, and \\S~\\ref{sec:nbody} the N-body simulations, and compares and contrasts the two regimes. Finally \\S~\\ref{sec:disc} concludes with a discussion of how these results can be used to understand dynamics at the Galactic Center, with particular regard to the curious observed structures such as the S-stars and the comoving groups IRS13E and IRS16SW. ", "conclusions": "\\label{sec:disc} The goals of this study were twofold: to study the effectiveness with which an IMBH can transport a handful of stars deep into a SMBH potential well, and in the process, gain insight into the dynamics of young stars there. The results of the three-body integrations suggest that the principal factors that influence the eventual deposition of stars are the IMBH orbital eccentricity and the inspiral speed. Paumard et al. (2006) provide a review of particular features of the observations, and we compare our results to this. \\begin{itemize} \\item The observations show two counter-rotating disks, oriented at large angles with respect to one another, with inner radii of about $0.05\\pc$. The cluster core simulations produce a single, pronounced disk of stars, but in no case is there any significant counter-rotating disk formed. In order to explain the claimed multiple contemporaneous disks, one would need to invoke the almost simultaneous infall of two distinct clusters. \\item The disks are observed to have a well-defined inner radius, or edge. The edge of the more populous disk (the clockwise disk) has an edge at 1$''$. Such edges also emerge naturally from the cluster core simulations. They result from the internal dynamical relaxation within the cluster, which sets up a Bahcall-Wolf density cusp. While individual stars may be sufficiently tightly bound to be transported to smaller radii, this edge represents the limit to which any significant cusp containing several stars may be transported. \\item Paumard et al. (2006) find that both disks have surface density profiles that scale as $\\Sigma \\propto r^{-2}$. Our simulations produce disks with surface density $\\Sigma \\propto r^{-0.75}$. This result appears to be robust with respect to changes in the initial cluster density law. We return to this below. \\item There is an absence of stars on larger scales as might be expected from tidal stripping of less tightly bound material (e.g., Kim \\& Morris 2003, Kim, Figer, Morris 2004). Our simulations are devoted to the cluster core alone and do not address this directly. Simulations of larger clusters suggest that this lack of observed massive stars could be a significant constraint on the cluster scenario (e.g., Portegies Zwart \\& McMillan 2003). However, the observations are limited to massive stars and this constraint may be avoided if the massive stars in the cluster are initially centrally concentrated (G\\\"{u}rkan \\& Rasio 2005), in which case they would be stripped only at small radii. \\item The disks are observed to have moderate thickness ($14\\dgrs \\pm 4\\dgrs$ for the clockwise disk), which is well produced by our cluster core simulations. \\item Most of the stars in the clockwise disk are on low eccentricity orbits, while those in the counter-clockwise system are mostly on eccentric orbits. Taken together, this is consistent with the two rings forming from entirely separate clusters, since the stripped disk stars tend to have eccentricities similar to those of the parent cluster orbit. The fact that the more eccentric orbits lie further out is also consistent with the ability of circular IMBH orbits to transport stars deeper into the potential well. \\item Paumard et al. (2006) confirm earlier claims that the IRS13E 'clump' corresponds to a real overdensity. Such core remnants emerge naturally from our simulations as well. Further, our simulations mirror the observations in that such associations reside in the tidal tail of stripped stars. \\end{itemize} Thus, the cluster core simulations demonstrate an encouraging ability to reproduce several of the principal dynamical features of the observed disks. In particular, the observed inner edge appears naturally and the presence of dynamically long-lived clumps appears generic, although the endurance of such clumps is constrained by high IMBH mass and low eccentricity. Furthermore, we find that the cluster scenario naturally produces the observed thickness of the disks and results in similar eccentricities amongst stars in a given disk. Nevertheless, there are several observed features which remain elusive. In no case do we produce a significant second disk -- if both disks are the result of cluster infall, then there must have been two separate clusters. Furthermore, the resulting surface density profile appears both robust and at odds with the observations. It is a consequence of the internal dynamics of the cluster, and not a result of initial conditions specific to any particular formation scenario. Internal dynamical relaxation leads to a cluster density profile given by the Bahcall \\& Wolf (1976) law. Tidal stripping of this cusp results in remnant disks with inevitable surface densities of $\\Sigma \\propto r^{-0.75}$. The same internal density profile also explains the inner edge we observe. The number of surviving cluster stars, assuming a fixed Bahcall-Wolf density profile, scales as $N \\propto r^{5/4}$. Thus, if the cluster core contained 100 stars at 1~pc, the radius at which a single star remains bound to the IMBH is $r \\sim (1/100)^{4/5} \\sim 0.025$~pc. Thus, our simulations show that the gross features of the resulting population can be understood in large part by approximating the internal structure of the cluster as a relaxed Bahcall-Wolf cusp and treating the stripping with a simple Roche lobe criterion. In addition to the disk stars, there is also the cluster of B-type stars close to the SMBH, the S-star cluster. While our simulations do suggest that an IMBH can carry one or two stars very deep into the potential well, in no case do we find transport of a sufficient number of stars to explain the S-stars as the remnant of a single inspiral. Further, the observed S-stars have varying semimajor axes and eccentricities. These simulations show that at the end of its inspiral, the IMBH orbit has circularized, and stars acquire orbital parameters similar to that of the IMBH upon stripping. This implies that the large variation in the orbital parameters of the S-stars cannot be accounted for simply by their deposition by a cluster infall. Levin \\& Beloborodov (2003) argue that Lens-Thirring precession may account for some randomization of the orbits, but this can only be effective for stars with initially small semimajor axes and large eccentricities, like SO-2; otherwise, the precession timescales are far too long, compared to the stellar lifetimes. Infalling clusters likely do not survive to small radii, and are efficiently disrupted by strong peribothric tidal stresses when in highly eccentric initial orbits. Thus, the origin and subsequent evolution of the S-stars remains unclear from these simulations, although a suite of scattering experiments is currently underway which may provide some insight. Finally, the inclusion of a dark-mass cusp, perhaps representing a population of small black holes (Miralda-Escud\\'e \\& Gould 2000), significantly decreased the inspiral time of the IMBH. This somewhat loosens the constraints on the infalling cluster scenario's ability to explain the peculiar dynamical environment at the Galactic Center, because ultimately, the scenario is constrained by the ages of the OB stars observed in the inner few arcseconds. Consequently, given such a cusp, the delivery of very young OB stars to the Galactic Center can occur over a longer range of timescales than previously considered, since their $\\sim 10\\Myr$ measured ages allow them to have spent longer times outside the central parsec. The outer portion of the inspiral thus may proceed more slowly, allowing the cluster time to further relax and mass-segregate. \\subsection{Comparison to other work} A recent paper by Levin, Wu, \\& Thommes (2005) reports on the simulations of a three-body problem similar to that treated in \\S~\\ref{sec:3bresults}. They investigate the problem using a symplectic integrator in extended phase space, with an ad hoc treatment of close-encounters, which are the bane of many symplectic algorithms. The IMBH orbit analytically decays, and all stars are set to be massless with the same initial Jacobi radius. Further, their choice of inspiral timescale is in the range $1000-10000$ IMBH orbits. They find that for circular IMBH inspirals, stars can achieve significant eccentricities but only low inclinations, even for slower inspiral. These results are mirrored by those reported in \\S~\\ref{sec:3bresults}. For eccentric inspirals, their results show two groups of stars, one in a thin disk of half-opening angle $10^{\\circ}$, and the other with inclinations of $10^{\\circ}2$ \\cite{2003Ap&SS.283..347C}. For $N=2$ Newtonian systems are regular, but post Newtonian dynamics\\cite{1915SPAW.......778E,1915SPAW.......799E,1916AnP....49..769E} can already reveal chaotic behaviour \\cite{1978MitAG..43Q.121B,2003PhRvL..90m4101B} by $N=2$. Chaoticity in a self gravitating systems with $N>2$ can reveal itself on a very short time scales, of the order of orbital periods. On the other hand, in some systems like the planetary system around the Sun, chaoticity only reveals itself on a time scale of billions of orbits \\cite{1860MNRAS..21...60D,1860MNRAS..20..240D}. In a stellar cluster which does not contain a dominant central mass, orbits can be chaotic on a much smaller time scale. In such an environment seemingly regular orbits can suddenly become highly irregular, to return later to regular again \\cite{1962AJ.....67..591K,1979IAUS...81..231K}. In this study we focus on the characterisation of orbits which show irregular behaviour on a short (dynamical) time scale, and not on a long (relaxation) time scale. Star clusters, with contain between about 100 and {\\large O($10^7$)} stars, are governed by microscopic few body interactions \\cite{1996magr.meet..167K}. Analysis of such systems is often hampered by this internal microscopic physics and a qualitative indicator of chaotic behaviour could enormously assist in the understanding of large scale gravitational $N$-body simulations. In particular, since it is thought that stars on irregular orbits in self gravitating $N$-body systems have an important effect of the bulk properties of such systems \\cite{2002SSRv..102..115M}. We present a transparent diagnostic to qualify chaotic behaviour of individual trajectories in $N$-body systems. In addition, the method has some quantitative qualities. The application of this method ranges from 3-body interactions to simulations of entire star clusters and galaxies. For clarity, we define irregular orbits as orbits which are deterministic though sensitive to the initial conditions and which cannot be described as a sum of periodic motions. The motion of a star in an \\emph{N}-body simulation can change from regular to irregular (and vice versa) due to gravitational interactions with other stars. Irregularity of an orbit then is a local quantity, and this forces us to deviate from the conventional methods based on Lyapunov \\cite{Lyapunov:1901} numbers such as explained in \\cite{1997CeMDA..67...41F,Brasser:2004,Sandor:2004}. In addition, the calculation of Lyapunov exponents is costly, may require reruns and are therefore less suited for a direct diagnostic, whereas we are predominantly interested in a diagnostic that can be evaluated at runtime. The Lyapunov indicator, however, is well suited for discriminating between ordered and weak chaotic motion, like planetary systems, whereas we are in particular interested in very chaotic systems \\cite{1997CeMDA..67...41F}. Others methods for detecting chaos in Hamiltonian systems, such as SALI \\cite{Kalapotharakos:2004,Skokos:2004}, Fourier Transform \\cite{Aguilar:1998,Laskar:1998,Meritt:1998}, Poincar\\'e Section \\cite{Poincare:1892}, the zero-one method \\cite{Gottwald:2002} and the geometric indicator \\cite{Cipriani:2002} are also less suited for analysing gravitational $N$-body simulations during run time, since they also are not practical in providing an instantaneous quantification of the chaoticity of the system we are interested in here. In \\S\\,\\ref{sec:method} we discuss the basics of the Continuous Wavelet Transform (CWT), which we present in a package called CWaT. We then proceed by applying this method to several \\emph{N}-body simulations in \\S\\,\\ref{sec:application}, to conclude in \\S\\,\\ref{Sect:Conclusions}. ", "conclusions": "\\label{Sect:Conclusions} We present a qualitative diagnostic to analyze irregular behaviour in particle based gravitational $N$-body simulations. The method, dubbed CWaT, is based on continuous wavelet transforms and provides a direct and instantaneous diagnostic for irregular behaviour. In our analysis we demonstrate that CWaT is suitable for analyzing self gravitating $N$-body simulations, as it does not require a long time series for the analysis and it provides an instantaneous indicator for the degree of chaos, allowing a quantification of the number of irregular orbits at any moment in time in a large $N$-body simulation. We applied CWaT successfully to analyse gravitational 3-body interactions. The results of the CWaT method for 3-body systems compares well with Poincar\\'e sections. Here CWaT provides a qualification of the irregularity of the orbit, whereas for Poincar\\'e sections such an analysis is technically much harder to perform. Eventually we applied the CWaT method to several N-body systems with 100 particles in which all stars had the same mass, and another set of simulations in which a mass spectrum was adopted. The multi mass system systematically produced fewer stars on irregular orbits, compared to the equal mass system. In addition, we noticed that during a deep core collapse the number of irregular orbits dropped even further, whereas in the relatively low density inter-core collapse stages the number of irregular orbits tends to increases again." }, "0607/astro-ph0607669_arXiv.txt": { "abstract": "We explore the nature of variations in dust emission within an individual galaxy using 3.6 - 160~$\\mu$m {\\it Spitzer} Space Telescope observations and 450 and 850~$\\mu$m James Clerk Maxwell Telescope observations of the edge-on Sd spiral galaxy NGC~4631 with the goals of understanding the relation between polycyclic aromatic hydrocarbons (PAH) and dust emission, studying the variations in the colors of the dust emission, and searching for possible excess submillimeter emission compared to what is anticipated based on dust models applied to the mid- and far-infrared data. PAH emission at 8~$\\mu$m is found to correlate best with hot dust emission at 24~$\\mu$m on kiloparsec scales, although the relation breaks down on scales equal to hundreds of parsecs, possibly because of differences in the mean free paths between the photons that excite the PAHs and heat the dust and possibly because the PAHs are destroyed by the hard radiation fields in the centers of some star formation regions. The ratio of 8~$\\mu$m PAH emission to 160~$\\mu$m cool dust emission appears to be a function of radius. The 70/160 and 160/450~$\\mu$m flux density ratios are remarkably constant even though the surface brightness varies by a factor of 25 in each wave band, which suggests that the emission is from dust heated by a nearly-uniform radiation field. Globally, we find an excess of 850~$\\mu$m emission relative to what would be predicted by dust models. The 850~$\\mu$m excess is highest in regions with low 160~$\\mu$m surface brightness, although the strength and statistical significance of this result depends on the model fit to the data. We rule out variable emissivity functions or $\\sim4$~K dust as the possible origins of this 850~$\\mu$m emission, but we do discuss the other possible mechanisms that could produce the emission. ", "introduction": "} Before the launch of the {\\it Spitzer} Space Telescope \\citep{wetal04}, research into the 1 - 1000~$\\mu$m spectral energy distributions (SEDs) of nearby galaxies was greatly restricted by the angular resolution and signal-to-noise levels of far-infrared data. Most of the {\\it Infrared Astronomical Satellite} (IRAS) and {\\it Infrared Space Observatory} (ISO) data were only usable as global flux density measurements. In the few cases where galaxies were resolved at far-infrared wavelengths, the spatial information was limited. With {\\it Spitzer} data, it is now possible to examine 1 - 1000~$\\mu$m SEDs for kiloparsec-sized regions in galaxies at distances up to 10~Mpc. These higher-resolution, higher-sensitivity data allow for new observational investigations into the properties of the polycyclic aromatic hydrocarbons (PAHs) and other dust across this large wavelength regime. In this paper, we will use {\\it Spitzer} data to address two major issues related to dust emission that were not completely resolved using IRAS and ISO data. The first major question that can be asked is how best to describe the SEDs of discrete regions within individual nearby spiral galaxies, particularly in the context of describing the dust emission between 60 and 2000~$\\mu$m. Studies combining IRAS and ISO results with submillimeter or millimeter data suggested that a range of simple one or two-component blackbodies modified with different emissivity laws could describe the data. The exact descriptions, however, were confusing and contradictory. The major point of debate was how to describe the emission at submillimeter or millimeter wavelengths that exceeded what is expected when extrapolating from the 60 - 200$\\mu$m regime using one or a series of $\\sim20$ - 30~K blackbodies modified by $\\lambda^{-2}$ emissivity functions. Some studies have suggested that the power law describing the emissivity function varies among nearby galaxies \\citep[e.g][]{deeiac00, betal03}. Other studies found that the submillimeter or millimeter excess could be described by a cold blackbody modified by a $\\lambda^{-2}$ emissivity law. This cold dust temperature component could have a temperature as low as $\\lesssim10$~K \\citep[e.g][]{kszc98, skc99, gmjwbl03, gmjwb05}. \\citet{rtbetal04} were among the first to examine this issue using {\\it Spitzer} data. They examined the SEDs of the central starburst ring and outer disk of NGC~7331 and found the 70 - 850~$\\mu$m dust emission from the ring was consistent with dust in the $\\sim20$ - 25~K range with an emissivity law of $\\lambda^{-2}$. However, a limitation of the study by Regan et al. was the lack of significant submillimeter emission at a high signal-to-noise level outside of the ring. They therefore could not determine whether the description of the dust emission in the center of NGC~7331 applies throughout the galaxy's disk. This issue is particularly important in terms of modeling the dust emission and estimating the dust mass within nearby galaxies. If dust as cold as $\\sim6$~K is present, it will only be a significant source of emission at submillimeter wavelengths, but it may contribute substantially to the mass of the interstellar medium. If the dust emissivity is variable then dust models need to be adjusted to account for this variability. Without these results, the full picture of dust emission in nearby galaxies is incomplete. The other major question is to determine how PAH emission, particularly the 7.7~$\\mu$m feature, is correlated with dust emission on scales of hundreds of parsecs. Results from ISO gave differing results. \\citet{frsc04}, among others, found a strong correlation between 7 and 15~$\\mu$m emission in galaxy disks, thus demonstrating a correlation between PAH and hot ($\\sim100$~K) dust emission (which is commonly associated with transiently-heated very small grain emission but which may also represent emission from grains that are in thermal equilibrium at $\\sim100$~K). \\citet{hkb02}, however, found that 8~$\\mu$m PAH emission was correlated more closely with 850~$\\mu$m emission from cool ($\\sim25$~K) grains than with the 15~$\\mu$m emission from hot dust. Using {\\it Spitzer} observations of NGC~300, \\citet{hraetal04} found that 8~$\\mu$m emission originated from the rims of star formation regions within the galaxy whereas the 24~$\\mu$m emission was more strongly peaked in the centers of the star formation region, which suggested a correlation on kiloparsec scales but not on smaller scales. \\citet{detal05}, however, found that 8~$\\mu$m emission from kiloparsec sized regions within M~81, M~51, and NGC~7331 appeared to be more closely correlated with the total infrared flux than with the 24~$\\mu$m flux density. In a separate analysis of the {\\it Spitzer} data for M~81, P\\'erez-Gonz\\'alez et al. (2006, in press) also found that the 8~$\\mu$m luminosity was closely correlated with the total infrared luminosity on kiloparsec scales, although they did not examine the relation between 8 and 24~$\\mu$m emission. Identifying how PAH emission is correlated with dust emission will lead to a better understanding of how it can be used as a star formation indicator, as has been suggested by the results of \\citet{rsvb01} and \\citet{frsc04} (although the contradictory results of \\citet{cetal05} should also be noted). If PAH emission is correlated with 24~$\\mu$m hot dust emission on scales of hundreds of parsecs or kiloparsecs, then it should be as effective a star formation indicator as 24~$\\mu$m emission on those spatial scales. However, if PAH emission is more closely correlated with 160~$\\mu$m cool dust emission, which may trace extended cirrus emission unrelated to star formation \\citep[e.g.][]{hrgetal04}, then the PAH emission should be used with greater caution as a star formation tracer. Among the data available in the {\\it Spitzer} Infrared Nearby Galaxy Survey \\citep[SINGS;][]{ketal03}, the data for \\object{NGC 4631} are optimal for studying the complete SED of the dust emission from mid-infrared to submillimeter wavelengths. NGC~4631 is a nearly edge-on \\citep[inclination $85\\deg$;][]{t88} Sd galaxy at a distance of 9.0~Mpc with an optical disk of $15^\\prime.5 \\times 2^\\prime.7$ \\citep{ddcbpf91}. The edge-on nature of the galaxy makes it a high surface brightness source that is easier to detect at all wavelengths, particularly submillimeter wavelengths. However, the edge-on orientation adds a geometry effect to the observed surface brightness variations; regions may appear bright because they are intrinsicly luminous or because they represent regions with high column densities. Nonetheless, in terms of this study, the advantages of the edge-on orientation outweigh the disadvantages. In addition to the Infrared Array Camera \\citep[IRAC;][]{fetal04} and Multiband Imaging Photometer for {\\it Spitzer} \\citep[MIPS;][]{ryeetal04} data taken as part of the SINGS legacy program, deep archival SCUBA \\citep{hetal99} 450 and 850~$\\mu$m data covering most of the optical disk \\citep[previously published in][]{adb99, sag05} are available for this galaxy. These data almost completely cover the extended dust emission in this galaxy and allow for extraction of 3.6-850~$\\mu$m SEDs from regions throughout the disk. Mid-infrared-to-submillimeter SEDs of this galaxy have been studied before \\citep{betal03, dkw04, sag05}. However, these previous studies have relied on IRAS and ISO data that have lacked the resolution, the sensitivity, or the spatial coverage to study anything more than the global SED or the SED of the center of the object. With these new data, we can study infrared color variations across the entire optical disk of the galaxy and examine the SEDs within discrete regions throughout the disk. These results will lead to answers as to how to characterize far-infrared to submillimeter dust emission throughout a spiral galaxy and how PAH and dust emission at different wavelengths are interrelated. In Section~\\ref{s_obs}, we present information on the observations and data reduction for this galaxy. We briefly discuss the images in Section~\\ref{s_images}. In Section~\\ref{s_color}, we present simple color information from flux densities integrated over set apertures distributed throughout the plane of the galaxy. Then, in Section~\\ref{s_sed}, we present the 3.6 - 850~$\\mu$m spectral energy distributions of the galaxy as a whole and of discrete regions within the galaxy. We then summarize the results in Section~\\ref{s_conclusions}. ", "conclusions": "} The primary results from this research can be summarized as follows: 1. PAH emission at 8~$\\mu$m is closely related to hot dust emission at 24~$\\mu$m. The relation holds on 1.7~kpc scales but begins to break down at 650~pc scales. Furthermore, the PAH emission is not as centrally peaked as the hot dust emission on scales of hundreds of pc. 2. Variations in the strength of PAH emission relative to cool dust emission and in the fraction of dust mass in PAHs depend primarily on radius, not infrared surface brightness. 3. The 70 - 450~$\\mu$m color temperature does not appear to vary with surface brightness in this galaxy, even though the 24/70~$\\mu$m and 24/160~$\\mu$m flux density ratios do vary with surface brightness. This implies that a substantial part of the dust emission in this wavelength regime may originate from a cool, diffuse cirrus component. 4. The 850 and 1230~$\\mu$m emission in this galaxy is found to exceed what is anticipated from either the 23~K blackbody modified by a $\\lambda^{-2}$ emissivity law or the semi-empirical model that describes the emission shortward of 850~$\\mu$m. The physical dust models marginally describe the data, but they leave open the possibility that excess emission at 850~$\\mu$m is present. If present, the 850~$\\mu$m excess is highest in regions of moderate infrared brightness and constitutes the greatest fraction of the 850~$\\mu$m emission in infrared-faint regions. Two major implications for using PAH emission as a star formation indicator arise from this research. First, these results demonstrate that the spatial correlation between PAH emission and 24~$\\mu$m dust emission breaks down on scales of hundreds of parsecs. This implies that PAH emission cannot be used as an accurate star formation tracer on such scales (if 24~$\\mu$m emission primarily traces star formation regions). Second, these results suggest that the ratio of PAH to dust mass varies radially, although the reason for the radial variation needs further study. If PAH emission is to be used as a tracer of dust or star formation, then these radial effects need to be taken into account. Further comparisons using 8, 24, and 160~$\\mu$m data of nearby galaxies are needed to understand how PAH emission relates to dust emission in other galaxies and to understand the radial variations in the PAH emission. The color variations and SEDs found within NGC~4631 have some major implications for dust modeling. The absence of variations in the 70/160~$\\mu$m and 160/450~$\\mu$m flux density ratios in this galaxy but the presence of variations in the 70/160~$\\mu$m ratio in other galaxiesneeds to be studied further so as to understand the conditions that cause the ratios to vary or remain constant. Second, dust models may need to be adjusted to better account for the excess emission at 850~$\\mu$m if it is present. Further analysis with {\\it Spitzer} 70 and 160~$\\mu$m data as well as 850~$\\mu$m data with comparable resolution are needed not only to confirm the presence of the 850~$\\mu$m excess emission in other galaxies but to provide better descriptions of where the excess can be found. Furthermore, additional observations of nearby galaxies at multiple wavelengths between 160 and 850~$\\mu$m are needed so as to better constrain the shape of the SED in this wavelength regime." }, "0607/astro-ph0607175_arXiv.txt": { "abstract": "{ Since the discovery of the first giant planet outside the solar system in 1995 (Mayor \\& Queloz 1995), more than 180 extrasolar planets have been discovered. With improving detection capabilities, a new class of planets with masses 5-20 times larger than the Earth, at close distance from their parent star is rapidly emerging. Recently, the first system of three Neptune-mass planets has been discovered around the solar type star HD69830 (Lovis et al. 2006). Here, we present and discuss a possible formation scenario for this planetary system based on a consistent coupling between the extended core accretion model and evolutionary models (Alibert et al. 2005a, Baraffe et al. 2004,2006). We show that the innermost planet formed from an embryo having started inside the iceline is composed essentially of a rocky core surrounded by a tiny gaseous envelope. The two outermost planets started their formation beyond the iceline and, as a consequence, accrete a substantial amount of water ice during their formation. We calculate the present day thermodynamical conditions inside these two latter planets and show that they are made of a rocky core surrounded by a shell of fluid water and a gaseous envelope. ", "introduction": "The three Neptune-mass planetary system orbiting HD69830, a 4-10 Gyr old nearby star with a mass estimated at $0.86 \\pm 0.03 \\msun$, has been discovered through high precision measurements obtained with the HARPS spectrograph installed at La Silla, Chile (Lovis et al. 2006). The three planets, planets b,c and d, are located at 0.0785, 0.186 and 0.63 AU from the central star, and their minimum masses are equal to 10.2, 11.8, 18.1 $\\mearth$ respectively. This system, with three sub-Neptune mass planets within 1 AU, represents a considerable challenge for planet formation models, namely the disk instability (DI) model and the core-accretion (CA) model. ", "conclusions": "We have presented calculations which provide a fully consistent scenario for the formation and evolution of the planetary system around HD69830. From the calculations presented we can infer the following general scenario for the formation of the system. All three planets start by accreting planetesimals and very little gas as they migrate inwards until they reach a region depleted in solids either by the passage of a previous planet or because of too high temperatures. The main heating source being suppressed they essentially accrete gas at a rate given by their Kelvin-Helmholtz (KH) timescale (Ida \\& Lin 2004). To remain of Neptune-mass without requiring unlikely timing with the disapearence of the disk, a given planet must enter this depleted region when its KH timescale is of the order of the lifetime of the disk which corresponds to a mass of order ~8-12 $\\mearth$ (Ida \\& Lin 2004). For the three planets to collect this mass of heavy elements implies a significant amount of migration of the growing cores. A question which naturally arises with such a planet formation model, is the degree of fine-tuning of the initial conditions needed to produce a planetary system with similar properties (an exact match is meaningless). In this regard, the protoplanetary disk mass and lifetime we require are typical of observed values (Haisch et al. 2001, Beckwith \\& Sargent 1996). In fact, the major constraint comes from the fact that the planets, in order to remain of small mass, must enter the planetesimal depleted region of the disk at a time when their accretion timescale (roughly the core's Kelvin-Helmholtz timescale) is comparable to the lifetime of the disk. For the three planets (b,c and d), this timescale is around 1 Myr, 0.6 Myr and 0.2 Myr. In our simulation the lifetime of the disk is of order 2 Myr. Hence, it is only for the second and third planet (c and d) that this requirement is really limiting the possibilities, but it is certainly not fine-tuning. For the first time, our consistent formation/evolution calculations lead to the determination of the bulk composition and the inner structure of three Neptune-mass planets: the innermost one consists of a rocky core, with possibly a tiny gaseous envelope, whereas the two outermost planets are made of a central rocky core, a shell of super-critical fluid water and a gaseous envelope. A clear test of the present formation and evolution scenario could be achieved by the determination of the mean density of the planets. This would only be possible if the system is seen edge-on and transits are detected so as to measure the radius of the planets. While difficult from the ground, such observations are within reach of HST, COROT or KEPLER. Even if the present system does not lead to observable transits, it is likely that similar, transiting Neptune-mass systems will be discovered in a near future. Confrontation of the present theory with such observations will improve dramatically our understanding of planet formation. Finally, Spitzer observations of the HD69830 system have revealed the presence of micron sized dust at distances lower than 1 AU from the central star, that could result from the presence of an asteroid belt (Beichman et al. 2005). Preliminary order of magnitude estimates have shown that the passage of the two inner planets during their formation may significantly but not completely deplete the asteroid belt. Hence, the belt, if present prior to the formation of the planets, would be able to survive at least in part. Interestingly, we note that dust is observed in regions in mean motion resonances with the outermost planet (1:2 and 1:3), that may excite the asteroids, leading to collisions and dust production." }, "0607/astro-ph0607496_arXiv.txt": { "abstract": "I review the observational prospects to constrain the equation of state parameter of dark energy and I discuss the potential of future imaging and redshift surveys. Bayesian model selection is used to address the question of the level of accuracy on the equation of state parameter that is required before explanations alternative to a cosmological constant become very implausible. I discuss results in the prediction space of dark energy models. If no significant departure from $w=-1$ is detected, a precision on $w$ of order 1\\% will translate into strong evidence against fluid--like dark energy, while decisive evidence will require a precision of order $10^{-3}$. ", "introduction": "One of the most fundamental problems of contemporary physics is to elucidate the nature of the ``Dark Sector'' of the Universe. A wealth of cosmological observations seem presently to point to a concordance cosmological model where ``normal'' (i.e.\\ baryonic) matter accounts for a mere 4\\% of the matter--energy contents of the cosmos. The remaining 96\\% makes up the so--called ``Dark Sector'', with about 19\\% of cold dark matter (CDM) and 77\\% of ``dark energy''. The details of this cosmic budget vary somewhat depending on the data sets used and the assumptions one makes, but the errors on the different components are below 10\\% (for details, see e.g.~\\cite{Efstathiou,Percival,Seljak,Sanchez,Spergel:2006hy}). Guidance as to the nature of dark energy requires stronger observational proof of its properties today and in the past. A first important step is to discriminate between an evolving dark energy (whose energy density changes with cosmic time) and a cosmological constant of the form proposed by Einstein in the 1910s. A handle on this question is offered by the equation of state parameter, $w$, that measures the ratio of pressure to energy density of dark energy. Current data are consistent with $w=-1$ out to a redshift of about $1$, with an uncertainty of order $5-10\\%$, by using all of the available data sets (see e.g.~\\cite{Seljak:2006bg}). However, one must be very careful when assessing the combined constraining power of different data sets whenever each one of them does not provide strong constraints when taken alone. Combination of mutually inconsistent data can potentially lead to unwarranted conclusions on the dark energy parameters. We first briefly review the observational prospects for constraining the dark energy equation of state parameter, referring the reader to \\cite{Trotta:2006gx} for a more detailed discussion. In section \\ref{sec:bayes} we present and discuss some results on the required accuracy on $w$ from the perspective of Bayesian model selection and conclude in section \\ref{sec:conclusion}. ", "conclusions": "\\label{sec:conclusion} We have argued that the most promising methods for dark energy investigation are weak lensing and acoustic oscillations, because of their statistical accuracy (weak lensing) and robustness to systematic errors (acoustic oscillations). Weak lensing has the potential of achieving 1\\% accuracy on $\\weff$ but this precision requires an exquisite control of various systematic errors. Observations of baryonic oscillations with a spectroscopic survey have less statistical power than weak lensing (roughly a factor of 5), but are less prone to systematic errors due to the characteristics of the acoustic signature. The above goals could be reached within the next decade thanks to a vigorous observational campaign, involving collaborations such as DES, darkCAM, Pan--STARRS and WFMOS. We have shown that Bayesian model selection can offer a guidance as to the level of precision required on $w=-1$ before explanations alternative to a cosmological constant appear extremely unlikely in terms of posterior odds of models. We have found that phantom models where one can have $\\weff \\ll -1$ are strongly disfavoured by present--day data. Gathering decisive evidence against a fluid--like model for dark energy will however require a precision of order $10^{-3}$ on $\\weff$. In conclusion, the observational study of dark energy is a crucial area of cosmological research. Thanks to a host of ambitious proposals and a strong support by several funding bodies, key advances are likely to be made within the next decade both from the observational and the theoretical points of view." }, "0607/nucl-th0607039_arXiv.txt": { "abstract": "The nucleus of ${}^{208}$Pb --- a system that is 18 order of magnitudes smaller and 55 orders of magnitude lighter than a neutron star --- may be used as a miniature surrogate to establish important correlations between its neutron skin and several neutron-star properties. Indeed, a nearly model-independent correlation develops between the neutron skin of $^{208}$Pb and the liquid-to-solid transition density in a neutron star. Further, we illustrate how a measurement of the neutron skin in ${}^{208}$Pb may be used to place important constraints on the cooling mechanism operating in neutron stars and may help elucidate the existence of quarks stars. ", "introduction": "\\label{Sec:Introduction} It is an extrapolation of 18 orders of magnitude from the neutron radius of a heavy nucleus --- such as $^{208}$Pb with a neutron radius of $R_{n}\\!\\approx\\!5.7$~fm --- to the approximately 10~km radius of a neutron star. Yet both radii depend on our incomplete knowledge of the equation of state of neutron-rich matter. That strong correlations arise among objects of such disparate sizes is not difficult to understand. Heavy nuclei develop a neutron-rich skin as a result of its large neutron excess ({\\it e.g.,} $N/Z\\!=\\!1.54$ in $^{208}$Pb) and because the large Coulomb barrier reduces the proton density at the surface of the nucleus. Thus the thickness of the neutron skin depends on the pressure that pushes neutrons out against surface tension. As a result, the greater the pressure, the thicker the neutron skin~\\cite{Brown:2000}. Yet it is this same pressure that supports a neutron star against gravitational collapse~\\cite{Lattimer:2000nx,Steiner:2004fi}. Thus models with thicker neutron skins often produce neutron stars with larger radii~\\cite{Horowitz:2001ya}. The above discussion suggests that an accurate and model-independent measurement of the neutron skin of even a single heavy nucleus may have important implications for neutron-star properties. Attempts at mapping the neutron distribution have traditionally relied on strongly-interacting probes. While highly mature and successful, it is unlikely that the hadronic program will ever attain the precision status that the electroweak program enjoys. This is due to the large and controversial uncertainties in the reaction mechanism~\\cite{Ray:1985yg,Ray:1992fj}. The mismatch in our knowledge of the proton radius in ${}^{208}$Pb relative to that of the neutron radius provides a striking example of the current situation: while the charge radius of ${}^{208}$Pb is known to better than 0.001~fm~\\cite{Fricke:1995}, realistic estimates place the uncertainty in the neutron radius at about 0.2 fm~\\cite{Horowitz:1999fk}. The enormously successful parity-violating program at the Jefferson Laboratory~\\cite{Aniol:2005zf,Aniol:2005zg} provides an attractive electroweak alternative to the hadronic program. Indeed, the Parity Radius Experiment (PREX) at the Jefferson Laboratory aims to measure the neutron radius of $^{208}$Pb accurately (to within $0.05$~fm) and model independently via parity-violating electron scattering~\\cite{Horowitz:1999fk}. Parity violation at low momentum transfers is particularly sensitive to the neutron density because the $Z^0$ boson couples primarily to neutrons. Moreover, the parity-violating asymmetry, while small, can be interpreted with as much confidence as conventional electromagnetic scattering experiments. PREX will provide a unique observational constraint on the thickness of the neutron skin of a heavy nucleus. We note that since first proposed in 1999, many of the technical difficulties intrinsic to such a challenging experiment have been met. For example, during the recent activity at the Hall A Proton Parity Experiment (HAPPEX), significant progress was made in controlling helicity correlated errors~\\cite{Michaels:2005}. Other technical problems are currently being solved --- such as the designed of a new septum magnet --- and a specific timeline has been provided to solve all remaining problems within the next two years~\\cite{Michaels:2005}. Our aim in this contribution is to report on some of our recent results that examine the correlation between the neutron skin of ${}^{208}$Pb and various neutron-star properties~\\cite{Horowitz:2000xj,Horowitz:2001ya, Horowitz:2002mb}. In particular, we examine the consequences of a ``softer'' equation of state that is based on a new accurately calibrated relativistic parameter set that has been constrained by both the ground state properties of finite nuclei and their linear response. Further, results obtained with this new parameter set --- dubbed ``FSUGold''~\\cite{Todd-Rutel:2005fa} --- will be compared against the NL3 parameter set of Lalazissis, Konig, and Ring~\\cite{Lalazissis:1996rd,Lalazissis:1999} that, while highly successful, predicts a significantly stiffer equation of state. ", "conclusions": "\\label{Sec:Conclusions} In conclusion, a new accurately calibrated relativistic model (``FSUGold'') has been fitted to the binding energies and charge radii of a variety of magic nuclei. In this regard, the new parametrization is as successful as the NL3 set which has been used here as a useful paradigm. In particular, symmetric nuclear matter saturates at a Fermi momentum of $k_{\\rm F}\\!=\\!1.30~{\\rm fm}^{-1}$ (corresponding to a baryon density of $0.15~{\\rm fm}^{-3}$) with a binding energy per nucleon of $B/A\\!=\\!-16.30$~MeV. Further, by constraining the FSUGold parameter set by a few nuclear collective modes, we obtain a nuclear-matter incompressibility of $K\\!=\\!230$~MeV and a neutron skin thickness in ${}^{208}$Pb of $R_{n}-R_{p}\\!=\\!0.21$~fm. While the description of the various collective modes imposes additional constraints on the EoS at densities around saturation density, the high-density component of the EoS remains largely unconstrained. Thus, we made no attempts at constraining the EoS at the supranuclear densities of relevance to neutron-star physics. Rather, we simply explored the consequences of the new parametrization on a variety of neutron star observables and eagerly await high-quality data that will constrain the high-density component of the EoS. In particular, we found a limiting neutron-star mass of $M_{\\rm max}\\!=\\!1.72 M_{\\odot}$, a radius of $R\\!=\\!12.66$~km for a $M\\!=\\!1.4 M_{\\odot}$ neutron star, and no direct URCA cooling in neutron stars with masses below $M\\!=\\!1.3 M_{\\odot}$. It is interesting to note that recent observations of pulsar-white dwarf binaries at the Arecibo observatory suggest a pulsar mass for PSRJ0751+1807 of $M\\!=\\!2.1^{+0.4}_{-0.5} M_{\\odot}$ at a 95\\% confidence level~\\cite{Nice:2005fi}. If this observation could be refined, not only would it redefine the high-density behavior of this (and many other) EoS, but it could provide us with a precious boost in our quest for the equation of state." }, "0607/astro-ph0607482_arXiv.txt": { "abstract": "Properties of groups of galaxies depend sensitively on the algorithm for group selection, and even the most recent catalogs of groups built from redshift-space selection should suffer from projections and inclusion of infalling galaxies. The cosmo-dynamical evolution of groups from initial Hubble expansion to collapse and virialization leads to a fundamental track in virial-theorem estimated $M/L$ vs crossing time. The increased rates of mergers, both direct and after orbital decay by dynamical friction, in (low velocity dispersion) groups relative to clusters, explain the higher fraction of elliptical galaxies at given local number density in X-ray selected groups, relative to clusters, even when the hierarchical evolution of groups is considered. Galaxies falling into groups and clusters should later travel outwards to typically 2 virial radii, which is close to but somewhat less than the outermost radius where galaxy star formation efficiencies are observed to be enhanced relative to field galaxies of same morphological type. An ongoing analysis of the internal kinematics of X-ray selected groups suggests that the radial profiles of line of sight velocity dispersion are consistent with isotropic NFW distributions for the total mass density, with higher concentrations in massive groups than $\\Lambda$CDM predictions and lower concentrations in low mass groups. The critical mass, at $M_{200} \\approx 10^{13} M_\\odot$ is consistent with possible breaks in the X-ray luminosity-temperature and Fundamental Plane relations. The internal kinematics of groups indicate that the $M-T$ relation of groups should agree with that extrapolated from clusters with no break at the group scale. The analyses of observed velocity dispersion profiles and of the fundamental track both suggest that low velocity dispersion groups (compact and loose, X-ray emitting or undetected) are quite contaminated by chance projections. ", "introduction": "The attractive nature of gravity tends to assemble galaxies together in groups. With typical grouping algorithms, roughly half of all galaxies reside in groups. A smaller fraction of galaxies live in virialized groups of at least 4 bright galaxies, and a considerably smaller fraction live in the more massive virialized clusters. Given typical scaling relations, defining the virial radius of groups where the mean density is 200 times the critical density of the Universe, groups of galaxies have ranges of mass within the virial radius, virial and turnaround radii (all assuming $H_0 = 70 \\,\\rm km \\,s^{-1} \\, Mpc^{-1}$), velocity dispersion and temperature shown in Table~\\ref{mamon:grouppars}. \\begin{table}[ht] \\centering \\caption{Typical scales of groups} \\begin{tabular}{lccccr} \\hline & $\\log M_{200}$ & $r_{200}$ & $r_{\\rm ta}$ & $\\sigma_v$ & \\multicolumn{1}{c}{$kT$} \\\\ & ($M_\\odot$) & (Mpc) & (Mpc) & ($\\, \\rm km \\, s^{-1}$) & (keV) \\\\ \\hline Minimum & 12.5 & 0.3 & 1.1 & 140 & 0.2\\ \\ \\ \\\\ Maximum & 14.0 & 1.0 & 3.4 & 450 & 2\\ \\ \\ \\\\ \\hline \\end{tabular} \\label{mamon:grouppars} \\end{table} More massive objects can be called clusters. Of course, the limiting mass between groups and clusters is arbitrary and historical. The more massive groups have properties (e.g. $L_X-T$ \\cite{mamon:OP04} and Fundamental Plane \\cite{mamon:SMCB93} relations) expected from the extrapolation of clusters, while the less massive groups do not appear to follow such extrapolations, with the separation between massive cluster-like and low mass groups occurring at $M_{200} \\approx 10^{13}\\,M_\\odot$. Groups of galaxies thus provide an important laboratory to understand how the density of the environment affects the properties of galaxies. In turn, the modulation with environment of galaxy properties serves as an important constraint for (semi-)analytical models of galaxy formation. This review focusses on several dynamical and cosmological aspects of the evolution of groups and of their constituent galaxies. ", "conclusions": "Our understanding of the evolution of groups and of galaxies therein is making rapid progress thanks to the advent of 1) large galaxy surveys such as the 2dFGRS and SDSS, 2) multi-wavelength observations of groups, and 3) high resolution cosmological $N$ body simulations. Many of the results presented here need to be confirmed with these large data sets and simulation outputs. I am grateful to my collaborator, Andrea Biviano, for allowing me to mention our work in progress. I also thank Ivo Saviane, Valentin Ivanov and Jordanka Borissova for organizing a very exciting and high-level meeting and for being extremely patient with the manuscript." }, "0607/astro-ph0607161_arXiv.txt": { "abstract": "\\noindent We investigate the global transition from a turbulent state of superfluid vorticity (quasi-isotropic vortex tangle) to a laminar state (rectilinear vortex array), and vice versa, in the outer core of a neutron star. By solving numerically the hydrodynamic Hall-Vinen-Bekarevich-Khalatnikov equations for a rotating superfluid in a differentially rotating spherical shell, we find that the meridional counterflow driven by Ekman pumping exceeds the Donnelly-Glaberson threshold throughout most of the outer core, exciting unstable Kelvin waves which disrupt the rectilinear vortex array, creating a vortex tangle. In the turbulent state, the torque exerted on the crust oscillates, and the crust-core coupling is weaker than in the laminar state. This leads to a new scenario for the rotational glitches observed in radio pulsars: a vortex tangle is sustained in the differentially rotating outer core by the meridional counterflow, a sudden spin-up event (triggered by an unknown process) brings the crust and core into corotation, the vortex tangle relaxes back to a rectilinear vortex array (in $\\lsim 10^5$ {\\rm s}), then the crust spins down electromagnetically until enough meridional counterflow builds up (after $\\lsim 1$ {\\rm yr}) to reform a vortex tangle. The turbulent-laminar transition can occur uniformly or in patches; the associated time-scales are estimated from vortex filament theory. We calculate numerically the global structure of the flow with and without an inviscid superfluid component, for Hall-Vinen (laminar) and Gorter-Mellink (turbulent) forms of the mutual friction. We also calculate the post-glitch evolution of the angular velocity of the crust and its time derivative, and compare the results with radio pulse timing data, predicting a correlation between glitch activity and Reynolds number. Terrestrial laboratory experiments are proposed to test some of these ideas. ", "introduction": "Timing irregularities in a rotation-powered pulsar, such as discontinuous glitches \\citep{lss00,zwwmwz04,s05} and stochastic timing noise \\citep{hthesis02,sfw03}, provide an indirect probe of the internal structure of the star. The physical processes usually invoked to explain these phenomena are (un)pinning of Feynman-Onsager vortices in the crystalline inner crust \\citep{ai75}, starquakes \\citep{r76}, and thermally driven vortex creep \\citep{apas84,leb93}. Less attention has been directed at the {\\it global hydrodynamics} of the superfluid, except within the context of the spin-up problem in cylindrical geometry \\citep{aprs78, r93, cls00}. The importance of the global hydrodynamics was demonstrated by \\citet{tsatsa80}, who simulated pulsar rotational irregularities in the laboratory by impulsively accelerating rotating containers of He II, obtaining qualitative agreement with radio timing data (e.g. glitch amplitudes and post-glitch relaxation times). In this paper, we examine how the global flow pattern of superfluid in the outer core of a neutron star affects the rotation of the star. We focus on the outer core for simplicity: vortex pinning is thought to be weak or non-existent there \\citep{als84,dp03,dp04}, and the fluid is mainly isotropic ($^1S_0$ Cooper pairing) \\citep{ss95,yls99}, reducing the problem to a hydrodynamic one in a spherical shell. Even with this simplification, the calculation remains numerically challenging: the spherical Couette problem for a superfluid was solved for the first time only recently \\citep{pmgo05a}, generalizing previous work on the cylindrical Taylor-Couette problem for a superfluid \\citep{hbj95,hb04} and the spherical Couette problem for a classical viscous fluid \\citep{mt87a,dl94}. An isotropic ($^1S_0$-paired) neutron superfluid is described by the two-fluid Hall-Vinen-Bekarevich-Khalatnikov (HVBK) model. In this model, the viscous normal fluid and inviscid superfluid components feel a mutual friction force whose magnitude and direction depends on the distribution of Feynman-Onsager vortices \\citep{hv56a,bk61}. The vortices are organized in a rectilinear array if the flow is strictly toroidal, but they evolve into a tangle of reconnecting loops when the counterflow along the rotation axis exceeds a threshold, exciting the Donnelly-Glaberson instability (DGI) \\citep{gjo74,sbd83,tab03}. \\citet{pmgo05a} showed that the DGI is excited in a neutron star under a wide range of conditions, driven by the meridional component of the spherical Couette flow (SCF) in the interior. The mutual friction force changes dramatically during transitions between a vortex array and a vortex tangle \\citep{gm49,vinen57c,sbd83,s85}, affecting the rotational evolution of the star. In this paper, we propose a phenomenological model for timing irregularities in radio pulsars based on the creation and destruction of a vortex tangle --- {\\it superfluid turbulence} \\citep{bsd95,vinen03} --- in the outer core of a rotating neutron star. In our scenario, a glitch comprises the following sequence of events. (i) Differential rotation between the outer core and crust of the star, built up over time through electromagnetic spin down, generates a meridional Ekman counterflow in the outer core. We show that the axial counterflow exceeds the DGI threshold, creating a vortex tangle throughout the outer core. The mutual friction in this turbulent state takes the isotropic Gorter-Mellink (GM) form and is much weaker than the mutual friction associated with a rectilinear vortex array. (ii) When the glitch occurs, triggered by an unknown mechanism, the outer core and inner crust suddenly come into corotation and the vortex tangle decays, ultimately converting into a rectilinear vortex array. The decay process lasts $\\lsim 10^5$ {\\rm s}, depending on the drag force acting on the vortex rings, after which the mutual friction takes the anisotropic Hall-Vinen (HV) form \\citep{hv56a,hv56b} and increases by $\\sim 5$ orders of magnitude, precipitating a ``torque crisis\". (iii) After the glitch, differential rotation builds up again between the outer core and the crust due to electromagnetic spin down. When the axial counterflow exceeds the DGI threshold, after $\\lsim 1$ {\\rm yr}, the vortex array breaks up again into a tangle and the mutual friction drops sharply. Similar transitions from turbulent to laminar flow in a superfluid have been observed in laboratory experiments where He II, cooled to a few mK, flows around an oscillating microsphere \\citep{nks02,s04}. The paper is organized as follows. HVBK theory is briefly reviewed in \\S\\ref{sec:HVBKT}, together with the pseudospectral numerical method which we use to solve the HVBK equations in a differentially rotating spherical shell. The physics of the turbulent-laminar transition in a generic glitch scenario is elaborated in \\S\\ref{sec:VTMODEL}. The response of the stellar crust to a turbulent-laminar transition in the outer core is calculated numerically in \\S\\ref{sec:rotevol}. Finally, the results are summarized and applied to observational data in \\S\\ref{sec:summary}; a fuller observational analysis will be carried out in a future paper. ", "conclusions": "\\label{sec:summary} In this paper, we investigate how transitions between turbulent and laminar states of superfluid vorticity alter the standard theoretical picture of pulsar rotational irregularities like glitches and timing noise. (i) Most of the time, except in the immediate aftermath of a glitch, differential rotation in the outer core drives a nonzero, poloidal counterflow which continuously excites the DGI. A vortex tangle is thereby maintained in the outer core. The mutual friction in this regime, which is of GM form, couples the normal and superfluid components loosely. (ii) Immediately after a glitch, the differential rotation ceases, as does the poloidal counterflow. The vortex tangle decays over the mean life-time of its constituent vortex rings, $\\tau_{\\rm d} = 7.6 \\times 10^{5} (\\alpha/10^{-7})^{-1} (\\dot{\\Omega}_*/10^{-13} \\, {\\rm rad} \\, {\\rm s}^{-2})^{-2} ( t/1 \\, {\\rm yr})^{-2}$ {\\rm s}. A rectilinear vortex array develops, and the mutual friction switches to HV form, coupling the normal and superfluid components much more strongly. (iii) After $t_{\\rm tan} = 0.49 \\, (\\Omega_*/10^2 \\, {\\rm rad} \\, {\\rm s}^{-1})^{1/2}(\\dot{\\Omega}_*/10^{-13} \\, {\\rm rad} \\, {\\rm s}^{-2})^{-1}$ {\\rm yr}, electromagnetic spin down builds up the differential rotation sufficiently to drive a poloidal counterflow that exceeds the DGI threshold. A vortex tangle forms again in a time $\\tau_{\\rm g} = 4.9 \\times 10^{4} (\\alpha/10^{-7})^{-1}(\\dot{\\Omega}_*/10^{-13} \\, {\\rm rad} \\, {\\rm s}^{-2})^{-1}$ {\\rm s}, and the mutual friction reverts to GM form. Note that vortex pinning provides the boundary conditions for the superfluid SCF but does not occur within the outer core itself \\citep{dp03}. {\\it Therefore our new phenomenological picture is not a complete model for glitches}. It merely clarifies the vorticity state of the outer core before and after a glitch as an input into future models than incorporate the full glitch dynamics, including trigger mechanisms related to pinning in the inner crust. We draw together the strands of the model in Figure \\ref{fig:fig9}, which displays the evolution of the torque and the regions where the DGI is active during the following numerical experiment: we fix $\\Delta \\Omega = 0.1$ until $t=20$, accelerate the outer sphere instantaneously to corotation at $t=20$, then decelerate the outer sphere according to $\\Omega_2(t) = 1 - 0.001(t - 20)$ for $t>20$. This mimics the situation in a real pulsar, where we have $t_{\\rm E} \\ll t_{\\rm tan}$, i.e. Ekman pumping brings the fluid into corotation {\\it before} the DGI gradually starts being reexcited throughout the outer core. To make the experiment as realistic as possible, we do not assume that the mutual friction takes the same form everywhere in the outer core, but rather choose GM or HV friction at each point according to whether $(\\vv_{ns})_z$ is greater or less than $v_{\\rm DG}$ locally. In this comparison, we approximate $\\vv_{ns}$ by $\\vv_n$, as in equation (\\ref{eq:vns_ep}), because $\\vv_s$ can become very complicated (e.g. Figure \\ref{fig:fig1}), creating numerical difficulties. In order to satisfy $t_{\\rm E} \\ll t_{\\rm tan}$ while keeping $\\Delta \\Omega$ large enough so that the spheres ``lap\" each other several times, we are forced by computational exigencies to adopt a relatively low Reynolds number $\\Rey = 100$, shortening the Ekman time ($t_{\\rm E} \\sim 10 \\Omega_1^{-1} \\ll t_{\\rm tan}$), and to artificially boost $v_{\\rm DG}$, so that $|(\\vv_{n})_z/v_{\\rm DG}|$ does not exceed $\\sim 3$ throughout the computational domain. The results of the above numerical experiment are presented in Figure \\ref{fig:fig9}. Contours of $|(\\vv_{n})_z/v_{\\rm DG}|$, before and after the spin up at $t=20$, are plotted in Figures \\ref{fig:fig9}a--\\ref{fig:fig9}e. Shaded regions indicate where the DGI is active, i.e. $|(\\vv_{n})_z/v_{\\rm DG}| > 1$. Just before the glitch (Figure \\ref{fig:fig9}a), $32$ \\% of the superfluid is in a turbulent state, with the DGI active close to the inner sphere and at intermediate latitudes where meridional circulation is significant. After the differential rotation shuts off at $t=20$, the DGI initially spreads through the shell as transient axial flows increase, occupying $39$ \\% of the volume at $t=22$. However, the flow quickly settles down to a state of near-corotation during the time interval $ 24 \\lsim t \\lsim 50$, HV friction dominates, and the torque decays exponentially, with time constant $\\sim 10$ (followed by a linear decay). At $t=50$, the DGI slowly begins to assert itself again, starting from the inner sphere. As for a classical viscous fluid, when the outer sphere spins down, it pumps fluid radially inward and along the axis of rotation \\citep{vanyo}, so the axial flow speed is greatest near the inner sphere. By $t=120$, when $\\Delta \\Omega = 0.1$, the vorticity state is similar to that at $t=20$, before spin up. One might wonder if, in a realistic neutron star, the superfluid ever exits the turbulent state and becomes laminar. For the simulations in this paper, which have $\\Rey \\leq 3 \\times 10^4$, the answer is yes. Figure \\ref{fig:fig9} shows that, when the glitch occurs and the spheres come into corotation, $|(\\vv_{n})_z|$ falls below the DGI threshold after a time $t \\sim t_{\\rm E}$. The vortex tangle is then guaranteed to decay on a time-scale given by (\\ref{eq:decay_tangle1}) and (\\ref{eq:decay_tangle2}), as observed in terrestrial experiments. However, for more realistic neutron star Reynolds numbers ($\\Rey \\geq 10^{8}$), which are too challenging to simulate at present, the turbulent eddies in the normal fluid decay more slowly when the spheres come into corotation, so that $|(\\vv_{n})_z|$ remains above the DGI threshold for longer. If this happens, the vortex tangle may persist until the next glitch occurs, so that the superfluid never exits the turbulent state. What are the implications of Figure \\ref{fig:fig9} and the results in \\S\\ref{sec:VTMODEL} and \\S\\ref{sec:rotevol} for observations of glitches? Before considering this question, we emphasize again that the results in this paper do not constitute a theory of glitches, because important questions regarding the glitch trigger remain unresolved. Nevertheless, some general remarks can be made. First of all, it is clear that transitions between flow (and vorticity) states in superfluid SCF are caused by changes in $\\Rey$, and that such transitions become more frequent and complicated as $\\Rey$ increases \\citep{ybm77,je00}. This is compatible with the observation that adolescent pulsars ($\\sim 10^4$ {\\rm yr} old, like Vela) glitch most actively \\citep{lss00}. In younger pulsars (age $\\lsim 10^4$ {\\rm yr}), $T$ and hence $\\nu_n$ are relatively high, so $\\Rey$ is low. In older pulsars (age $\\gsim 10^4$ {\\rm yr}), $\\Omega$ and hence $\\Rey$ are low following electromagnetic spin down [although this trend is not straightforward and can be masked by localized heating from differential rotation between the superfluid and the crust \\citep{g75,ll99} or crust cracking \\citep{lfe98,fle00}]. A systematic statistical study of glitch activity versus $\\Rey$ will be published elsewhere \\citep{mpw06}, but preliminary estimates of $T$ and hence $\\Rey$ from cooling curves \\citep{tsuruta74,tsuruta98,plps04} including superfluidity \\citep{fi76,fi79,acg05} give Reynolds numbers in the range $10^{8} \\lsim \\Rey \\lsim 10^{12}$ for glitching pulsars. Two of the most active glitchers, the Crab and Vela, have $\\Rey \\sim 10^{9}$ and $\\Rey \\sim 10^{10}$ respectively. One expects that, at such high $\\Rey$, the fluid is turbulent, with the kinetic energy concentrated at large scales \\citep{yav78,yb86}, as for a classical viscous fluid \\citep{sdgv93,bsbd97}. This suggests that superfluid turbulence in pulsar interiors is an important factor in glitch dynamics. If it is true that the vorticity in the outer core exists in a turbulent state before a glitch, as postulated in our model, then $t_{\\rm tan}$ represents a lower bound on the time between glitches. In testing whether this bound is respected by the glitching pulsars currently known, we are hampered by the fact that most of these objects have only glitched once. Nevertheless, for all the $28$ pulsars that have glitched repeteadly, we find that the minimum inter-glitch time interval $t_{\\rm min}$ is greater than $t_{\\rm tan}$, as the theory predicts \\citep{mpw06}. The object PSR $2116+1414$ approaches the bound most closely, with $t_{\\rm min}=3.9$ {\\rm yr} and $t_{\\rm tan} = 2.0$ {\\rm yr}. This is encouraging, because the $28$ objects cover five decades in $t_{\\rm tan}$ and three decades in $t_{\\rm min}$, and the theoretical expression (\\ref{eq:reform}) for $t_{\\rm tan}$ contains zero free parameters. Note that $t_{\\rm tan}$ is proportional to the characteristic age ($=\\Omega_*/2\\dot{\\Omega}_*$) divided by $\\Omega_*^{1/2}$. Note also that the activity parameter defined by \\citet{ml90} involves glitch amplitudes (which are highly variable) as well as mean recurrence times, so we do not predict a correlation between the activity parameter and $t_{\\rm tan}$. It is harder to test the theoretical decay time-scale of the vortex tangle, as predicted by (\\ref{eq:decay_tangle2}), because it remains unclear what observable features are engendered by the decay process. The observed exponential post-glitch relaxation is of viscous origin and occurs on a time-scale much larger than $\\tau_{\\rm d}$. On the other hand, the decay of the tangle is accompanied by a large increase in mutual friction (GM $\\rightarrow$ HV), which may be connected with the rapid jump in $\\Omega$ during a glitch. The jump in $\\Omega$ has never been resolved in time, in pulsars which are nearly constantly monitored, consistent with the predictions of (\\ref{eq:decay_tangle2}) for the Crab ($\\tau_{\\rm d} = 3 \\times 10^{-4}$ {\\rm s}) and Vela ($\\tau_{\\rm d} = 0.2$ {\\rm s}). However, equation (\\ref{eq:decay_tangle2}) predicts that it may be possible to resolve the $\\Omega$ jump in older pulsars, provided that the time between glitches does not increase faster than $\\dot{\\Omega}_*$. In making these estimates we assume the canonical value $\\alpha = 10^{-7}$ for every object, in the absence of a microscopic theory, yet this is clearly an oversimplification because $\\alpha$ is sensitively temperature dependent. Oscillations in $\\dot{\\Omega}_*$ were observed before (period $\\sim 10$ {\\rm d}) and after (period $\\sim 25$ {\\rm d}) the Vela Christmas glitch, with $\\Delta \\dot{\\Omega}/\\dot{\\Omega}_* \\approx 0.17$ \\citep{mhmk90}. In our numerical simulations, persistent torque oscillations are always observed when the outer core rotates differentially, as occurs before a glitch. They are also observed after a switch from GM to HV friction, as occurs after a glitch. By comparing the dashed and solid curves in Figure \\ref{fig:fig8}b, we see that the oscillations are sustained by HV mutual friction. The oscillation period in our simulations is much shorter than in pulsar data, because we are restricted to $\\Rey \\leq 3 \\times 10^{4}$. An alternative explanation is that vortices in the inner crust oscillate relative to the normal fluid in the core \\citep{ssa95}. Several of the effects explored in this paper have been studied in terrestrial laboratories. Our results will motivate new experiments of this sort, cf. \\citet{alpar78} and \\citet{aprs78}. Although it is hard to access the neutron star regime $\\rho_n \\ll \\rho_s$ in He II, where interatomic forces are appreciable, transitions between turbulent and laminar superfluid vorticiy have been observed in experiments with microspheres immersed in $^4$He at {\\rm mK} temperatures \\citep{nks02}. Promising results on the relaxation of rotating He II were obtained by \\citet{tsatsa80}, but again these results are for $\\rho_n \\lsim \\rho_s$ and hollow spheres rather than a differentially rotating shell. We propose to extend these experiments in two directions: (i) by investigating low-Rossby-number ($\\Delta \\Omega/\\Omega \\ll 1$), high-Reynolds-number ($\\Rey \\gg 10^5$) SCF with He II at the temperature which minimizes $\\rho_n/\\rho_s$; and (ii) by repeating (i) with a nonideal dilute-gas Bose-Einstein condensate confined in a differentially rotating magneto-optical trap, in order to probe the stability of a vortex lattice to Kelvin wave excitations in the regime $\\rho_n \\ll \\rho_s$ \\citep{pa05}. The presence of a vortex tangle in He II can be detected by standard second-sound absorption techniques \\citep{hv56a,sbd83}, and the torque in experiment (i) can be monitored to look for oscillations when a change from HV to GM friction (or vice versa) is triggered by the DGI. In classical Navier-Stokes fluids, injection of vorticity into a metastable laminar state can trigger turbulence, e.g. seed vortices injected into a cylindrical vessel containing $^3$He-B (with $T \\leq 0.6 T_c$) generate a vortex tangle that eventually decays into a rectilinear vortex array \\citep{finne03}. Unlike He II, the normal component in $^3$He-B is laminar in these experiments and does not participate in the turbulent dynamics, due to its comparatively high viscosity [$\\nu_n \\sim 1$ {\\rm cm}$^2$ {\\rm s}$^{-1}$ $\\gg \\nu_s$, cf. $\\nu_n \\sim \\nu_s$ in He II; see \\citet{fbek04}]. Standard glitch theories assume that the normal component is tightly coupled to the crust by the external magnetic field \\citep{as88,jm98}. This suggest a second possible SCF experiment with $^3$He-B (at $T \\leq 0.6 T_c$) in a hollow spherical container in which the normal fluid corotates with the container, and therefore does not participate in the DGI, while the superfluid is free to become turbulent. Such an experiment can probe what aspects of the turbulent-laminar transition are caused by the normal and superfluid components respectively. Transitions between a vortex tangle and a rectilinear array can be detected using non-invasive nuclear magnetic resonance techniques \\citep{finne03}." }, "0607/astro-ph0607357_arXiv.txt": { "abstract": "We present polarization maps of G30.79 FIR 10 (in W43) from thermal dust emission at 1.3 mm and from CO J=$2 \\rightarrow 1$ line emission. The observations were obtained using the Berkeley-Illinois-Maryland Association array in the period 2002-2004. The G30.79 FIR 10 region shows an ordered polarization pattern in dust emission, which suggests an hourglass shape for the magnetic field. Only marginal detections for line polarization were made from this region. Application of the Chandrashkar-Fermi method yielded $B_{pos} \\approx 1.7$ mG and a statistically corrected mass to magnetic flux ratio $\\lambda_{C} \\approx 0.9$, or essentially critical. ", "introduction": "The star formation process involves a number of physical parameters, of which the magnetic field is the least observed. Magnetic field observations are divided into measurements of the Zeeman effect (in order to obtain the magnetic field strength in the line of sight), and linear polarization observations of dust and spectral-line emission. Polarization of dust emission is believed to be perpendicular to the magnetic field under most conditions \\citep{Lazarian2003}; hence, polarization of dust emission has been used as a major probe for the magnetic field geometry. In order to efficiently map the polarization of dust emission and infer information about the magnetic field morphology, high resolution observations are required. The BIMA millimeter interferometer has been used previously to obtain high-resolution polarization maps in several star forming cores \\citep{Rao1998,Girart1999b,Lai2001, Lai2002,Lai2003}. These results show fairly uniform polarization morphologies over the main continuum sources, suggesting that magnetic fields are strong, and therefore cannot be ignored by star formation theory. However, the number of star formation regions with maps of magnetic fields remains small, and every new result is significant. Spectral line linear polarization has been suggested to arise from molecular clouds under anisotropy conditions \\citep{Goldreich1981}. The prediction suggests that a few percent of linearly polarized radiation should be detected from molecular clouds and circumstellar envelopes in the presence of a magnetic field. It is also predicted that the molecular line polarization will be either parallel or perpendicular to the magnetic field, depending on the angles between the line of sight, the magnetic field, and the anisotropic excitation direction \\citep{Goldreich1982}. This process is known as the Goldreich - Kylafis effect. In order to use these techniques, we mapped the massive star forming region G30.79 FIR 10 with the BIMA array. We measured continuum polarization at 1.3 mm and CO $J=2\\rightarrow 1$ line polarization obtaining high resolution interferometric maps for both measurements. The remainder of this paper is divided in five major sections. Section 2 reviews information about the source, section 3 describes the observation procedure. Section 4 presents the results, section 5 gives the discussion, and section 6 the conclusions and summary. ", "conclusions": "We observed G30.79 FIR 10 and successfully mapped CO $J=2 \\rightarrow 1$ line and 1.3 mm dust continuum polarized emission with a resolution of $4^{\\prime \\prime}$. G30.79 FIR 10 is not a well studied region; however, there is evidence that points toward an early stage of development of high-mass star formation. We found a remarkably uniform pattern in our polarized dust emission map, which suggests an hourglass magnetic field morphology. Using the Chandrasekhar-Fermi method, we inferred a plane-of-the-sky magnetic field strength 1.7 mG, which yielded a geometry-corrected mass to magnetic flux ratio of 0.9 to respect to critical. This result is similar to those found in many other regions of low-mass and high-mass star formation. Polarized line emission was also detected in this region; these results pose questions about the sources of anisotropy, which will require more detailed modeling. This research was partially funded by NSF grants AST 02-05810 and 02-28953." }, "0607/astro-ph0607027_arXiv.txt": { "abstract": "Narrow depolarized canals are common in maps of the polarized synchrotron emission of the Milky Way. Two physical effects that can produce these canals have been identified: the presence of Faraday rotation measure ($\\RM$) gradients in a foreground screen and the cumulative cancellation of polarization known as differential Faraday rotation. We show that the behaviour of the Stokes parameters $Q$ and $U$ in the vicinity of a canal can be used to identify its origin. In the case of canals produced by a Faraday screen we demonstrate that, if the polarization angle changes by $90\\degr$ across the canal, as is observed in all fields to-date, the gradients in $\\RM$ must be discontinuous. Shocks are an obvious source of such discontinuities and we derive a relation of the expected mean separation of canals to the abundance and Mach number of supernova driven shocks, and compare this with recent observations by \\citet{Haverkorn03}. We also predict the existence of less common canals with polarization angle changes other than $90\\degr$. Differential Faraday rotation can produce canals in a uniform magneto-ionic medium, but as the emitting layer becomes less uniform the canals will disappear. We show that for moderate differences in emissivity in a two-layer medium, of up to $1/2$, and for Faraday depth fluctuations of standard deviation $\\lesssim 1\\,\\mathrm{rad}$, canals produced by differential rotation will still be visible. ", "introduction": "Polarized (synchrotron) radio emission is a rich source of information about the relativistic and thermal plasmas and magnetic fields in the interstellar medium (ISM). Recent observations have revealed an abundance of unexpected features that arise from the propagation of the emission through the turbulent ISM \\citep{Wieringa93, Uyaniker98, Duncan99, Gray99, Haverkorn00, Gaensler01, Wolleben06}. Arguably the most eye-catching is the pattern of depolarized canals: a random network of dark narrow regions, clearly visible against a bright polarized background. These canals evidently carry information about the ISM, but it is still not quite clear how this information can be extracted. Two theories for the origin of the canals have been proposed; both attribute the canals to the effects of Faraday rotation, but one invokes steep gradients of the Faraday rotation measure ($\\RM$) across the telescope beam, in a Faraday screen \\citep*{Haverkorn00, Haverkorn04}, whereas the other relies on the line-of-sight effects producing differential Faraday rotation \\citep{Beck99, Shukurov03}. In order to use the canal properties to derive parameters of the ISM, one must correctly identify their origin. In this letter we briefly discuss a few important aspects of the two theories describing the origin of canals. Detailed discussion of relevant depolarization mechanisms is presented in \\citet[][and in preparation]{Fletcher06a}. Canals produced by Faraday rotation measure gradients are discussed in Section~\\ref{sec:screen} where we show that these canals require a \\emph{discontinuous\\/} distribution of free electrons and/or magnetic field; we further suggest an interpretation of the discontinuities in terms of interstellar shocks. The case of differential Faraday rotation is discussed in Section~\\ref{sec:diffrot}; in Section~\\ref{sec:origin} we suggest an observational test that can be used to identify the specific mechanism that produces a given canal. The defining features of the canals are as follows: \\begin{enumerate} \\item \\label{itemi} the observed polarized emission $P$ approaches the polarization noise level $\\sigma_P$, $P\\la\\sigma_P$; \\item the canal is about one beam wide; \\item the canal passes through a region of significant polarized intensity, say $P\\ga3\\sigma_P$; \\item the canal is not related to any structure in total intensity, and so cannot be readily explained by, e.g., an intervening gas or magnetic filament. \\end{enumerate} Polarized intensity vanishes when emission within the telescope beam consists of two equal parts with mutually orthogonal polarization planes. Thus, feature \\ref{itemi} most often arises because the polarization angle $\\ps$ changes by $90\\degr$ across the canal. However, in Section~\\ref{sec:screen} we predict a new type of canal across which the observed polarization angle does not change. We only consider canals occurring in properly calibrated maps: \\citet{Haverkorn04} argue that the observations we discuss in Section~\\ref{sec:meansep} do not suffer from missing large-scale structure; \\citet{Reich06} discusses the calibration of radio polarization maps in depth. Polarized radiation is commonly described in terms of the complex polarization, \\begin{equation} \\cP = p\\exp{(2\\mi\\ps)}\\;, \\label{eq:compp} \\end{equation} where $p$, the degree of polarization, is the fraction of the radiation flux that is polarized. When polarized emission passes through magnetized and ionised regions, the local polarization angle $\\psi$ (at position $\\bmath{r}$) changes by an amount depending on the wavelength $\\lambda$ due to the Faraday effect \\[ \\psi(\\bmath{r})=\\pso(\\bmath{r})+\\phi(\\bmath{r}); \\hspace{5mm} \\phi(\\bmath{r})=\\lam^2 K\\int^{\\infty}_{z}\\nel B_z \\dif z^\\prime, \\] where $K=0.81\\rad\\m^{-2}\\cm^3\\mkG^{-1}\\p^{-1}$ is a constant, $\\nel$ is the number density of free thermal electrons, $B_z$ is the component of the magnetic field along the line of sight (here aligned with the $z$-axis), and the observer is located at $z\\to\\infty$. $\\phi(z)$ is known as the \\emph{Faraday depth to a position $z$} and gives the change in polarization angle of a photon of wavelength $\\lam$ as it propagates from $z$ to the observer. The maximum amount of Faraday rotation in a given direction is called the \\emph{Faraday depth} \\footnote{This terminology may cause confusion: many authors, including \\citet{Burn66} and \\citet{Sokoloff98}, define the Faraday depth as $F/\\lambda^2$, in our notation. However, it is more convenient, and physically better motivated, to define the Faraday depth, similarly to the optical depth, as a dimensionless quantity, as used by \\citet{Spangler82} and \\citet{Eilek89}} \\[ F=\\phi(-\\infty)=\\lambda^2 K \\int^{\\infty}_{-\\infty}\\nel B_z \\,\\dif z^\\prime\\;. \\] The observed amount of Faraday rotation, determined by the rotation measure $\\RM=\\dif\\ps/\\dif\\lam^2$, cannot exceed $F$, i.e., $|\\RM|\\leq |F|\\lambda^{-2}$. The value of $\\RM$ is related to $F$, but often in a complicated manner \\citep[see, e.g.,][]{Burn66, Sokoloff98}. Simplest is the case of a Faraday screen, where the source of synchrotron emission is located behind a magneto-ionic region (e.g., because relativistic and thermal electrons occupy disjoint regions): then $\\RM=F\\lambda^{-2}$. In a homogeneous region, where relativistic and thermal electrons are uniformly mixed, $\\RM=0.5 F\\lambda^{-2}$. Observations of linearly polarized emission provide the Stokes parameters $I$, $Q$, $U$ which are related to $p$ and $\\ps$ via $\\cP=(Q+\\mi U)/I$: \\begin{eqnarray} p & = & \\frac{(Q^2+U^2)^{1/2}}{I} \\label{eq:qup}\\;,\\\\ \\ps & = & \\label{eq:qupsi} \\sfrac{1}{2}\\left[\\arctan{\\frac{U}{Q}}-\\sfrac{1}{2}\\pi(\\sign Q-1)\\sign U\\right]\\;, \\end{eqnarray} and the polarized intensity is $P=(Q^2 + U^2)^{1/2}=pI$. The complex polarization can be written in terms of $\\phi$, \\begin{equation} \\cP = \\frac{p_0}{I}\\int_V W(\\bmath{r}_\\perp)\\epsilon(\\bmath{r}) \\exp\\left\\{2\\mi\\left[\\psi_0(\\bmath{r})+\\phi(\\bmath{r})\\right]\\right\\}\\dif V\\;, \\label{cPF} \\end{equation} where integration extends over the volume of the telescope beam $V$, $W(\\bmath{r}_\\perp)$ defines the shape of the beam, a function of position in the sky plane $\\bmath{r}_\\perp=(x,y)$, and $\\epsilon(\\bmath{r})$ is the synchrotron emissivity. The total intensity is similarly given by $I=\\int_V W(\\bmath{r}_\\perp)\\epsilon(\\bmath{r})\\,\\dif V$; the Faraday depth is a function of $\\bmath{r}_\\perp$, $F=F(\\bmath{r}_\\perp)$. ", "conclusions": "\\label{sect:summary} The main points of this paper can be summarized as follows: \\begin{enumerate} \\item The behaviour of the Stokes parameters $Q$ and $U$ in the vicinity of a canal allows one to identify whether a foreground Faraday screen or differential Faraday rotation is the cause of the canal (Section~\\ref{sec:origin}). \\item A foreground Faraday screen can produce canals with any polarization angle change across the canal, $0<\\Delta\\ps<90\\degr$. However, discontinuous jumps in the Faraday depth will only produce canals with $\\Delta\\ps\\simeq 90\\degr$ (Section~\\ref{sec:screen}). \\item If shocks produce the discontinuities in a foreground Faraday screen that generate canals, the mean separation of the canals can provide information about the Mach number and separation of shocks in the screen (Section~\\ref{sec:meansep}). \\item Canals produced by differential Faraday rotation are sensitive to non-uniformity in the medium along the line of sight, systematic or random. However, they will remain recognisable if the synchrotron emissivity varies by less than a factor of about $2$ or if the standard deviation of the Faraday depth is $\\sigma_{F}<1$ (Section~\\ref{subsec:diffrot}). \\end{enumerate}" }, "0607/astro-ph0607211_arXiv.txt": { "abstract": "We present a comprehensive description of the theory and practice of opacity calculations from the infrared to the ultraviolet needed to generate models of the atmospheres of brown dwarfs and extrasolar giant planets. Methods for using existing line lists and spectroscopic databases in disparate formats are presented and plots of the resulting absorptive opacities versus wavelength for the most important molecules and atoms at representative temperature/pressure points are provided. Electronic, ro-vibrational, bound-free, bound-bound, free-free, and collision-induced transitions and monochromatic opacities are derived, discussed, and analyzed. The species addressed include the alkali metals, iron, heavy metal oxides, metal hydrides, $H_2$, $H_2O$, $CH_4$, $CO$, $NH_3$, $H_2S$, $PH_3$, and representative grains. Once monochromatic absorption cross sections for all consistuents have been derived, chemical abundances have to be obtained before the resulting product can be summed to obtain total opacities. Hence, we include a review of the thermochemistry, techniques, and databases needed to derive equilibrium abundances and provide some sample results. ", "introduction": "\\label{intro} The discoveries of extrasolar giant planets and brown dwarfs in 1995 have opened up exciting new fields in astrophysics. These enable model atmospheres to be directly tested against observations for atmospheric temperatures and pressures at which a large number of molecules that had not been considered before in stellar astrophysics reside. Moreover, many of the extrasolar planets, including the first to be discovered, are much closer to their parent star than Mercury is to the Sun so the atmospheric conditions of these planets are totally unlike any previously investigated. In such situations, many of these molecules are subjected to a large ultraviolet flux from the parent star. The absorption of radiation by molecules is generally much more complicated than by atoms. Since many of the molecules in the atmospheres of brown dwarfs and extrasolar giant planets had not been studied in detail in an astrophysical context before 1995, their properties are still poorly known. Even those that had been investigated, had not been for the temperatures (100 - 3000 K) and pressures (10$^{-6}$ - 100 atmospheres) found in the atmospheres of such objects. Moreover, an additional factor of considerable importance is the formation of grains. In this paper, we address the spectroscopy and opacities of the molecules and atoms central to an understanding of substellar objects. Such opacities are required for models of their spectra and evolution. Another consequence of the emergence of this young field is the extension of the spectral classification system beyond late-type M-dwarf stars to include two new spectroscopic classes at lower temperatures, L and T dwarfs. The criteria for these new classes depend upon the strengths of atomic and molecular spectral lines, which in turn are governed by absorption cross sections and the abundances of various species. In this paper, we discuss the methods used to calculate opacities from extant line lists or by directly calculating the energy levels. The absorption due to spectral lines is the most important source of opacity, as well as the most difficult to calculate, and this is considered in \\S\\ref{lines}. After dealing with the background theory for calculating line strengths and broadening, and how to handle a frequency grid that samples the lines, we review the chemical species used for calculating substellar atmospheres. These are grouped together when they share important similarities, such as being metal hydrides or absorbing in the same general part of the spectrum. Table 1 lists the species considered in this work and gives the sections where the species are discussed in detail, including several condensates addressed with the Mie theory. In \\S\\ref{hitran}, we discuss the species with data obtained from the HITRAN (Rothman et al. 2003, 2005) and GEISA (Jacquinet-Husson et al. 1999, 2003, 2005) databases, together with other sources. The spectral range covered is the visible and infrared. The molecules considered in detail are $H_2O$ (Partridge \\& Schwenke 1997; Barber et al. 2006), $CH_4$ (Strong et al. 1993; Karkoscka 1994; Borysow et al. (2003); Brown et al. 1997), $NH_3$ (Brown el al. 2000), $CO$ (Goorvitch, 1994). Of these, $CH_4$ and $NH_3$ are important at the lower temperatures of interest, $CO$ is important at the upper temperatures of interest, and $H_2O$ is very important over the whole temperature range. In \\S\\ref{oxides}, we discuss the two heavy metal oxides, $TiO$ (Allard, Hauschildt \\& Schwenke 2000) and $VO$ (Plez 1998). Of the heavy metal oxides, these are the most important sources of opacity at the higher end of the temperatures of interest. Both have very similar properties. This is followed by \\S\\ref{hydrides}, where we discuss the metal hydrides $TiH$ (Burrows et al. 2005), $CrH$ (Burrows et al. 2002), $FeH$ (Dulick et al. 2003), $MgH$ (Skory et al. 2003; Weck et al. 2003a), and $CaH$ (Leininger \\& Jeung 1995), which all have similar properties. As much as possible, the most extensive and up-to-date spectroscopic data available are considered. In the case of $H_2S$ and $PH_3$, tables of precomputed opacities can be used, so calculations are restricted to interpolating in the ranges of temperature and pressure in the tables, but since these species have low abundances, they are relatively unimportant. Ultraviolet opacities are important when a planet is irradiated by the star it orbits, and in \\S\\ref{uvopac} we discuss these opacities for atomic hydrogen (Menzel 1969; Bethe \\& Salpeter 1957; Carson 1988a), atomic iron (Kurucz 1995), and the molecules $H_2$, $CO$, $SiO$, $H_2O$ (Kurucz 1993), and $H_2S$ (Lee, Wang \\& Suto 1987). The available data vary considerably, and this dictates the methods used to calculate the opacities. The absorption of the alkali elements in monatomic form: $Li$, $Na$, $K$, $Rb$, and $Cs$ using the Vienna Atomic Line Data, (VALD - Piskunov 1994) is considered in \\S\\ref{vald}. The line strengths and widths are all calculated in a uniform way, except for the wings of the resonance lines of $Na$ and $K$ (Burrows \\& Volobuyev 2003; Allard et al. 2003; Zhu, Babb, \\& Dalgarno 2006), for which a separate treatment should be applied. In addition to lines, the absorption due to the underlying continuum is calculated, as discussed in \\S\\ref{continuous}. Rosseland and other harmonic mean opacities are divergent if the absorption drops to zero at any point in the frequency region considered, and a realistic truncation of the wings of spectra lines is required. The four processes discussed are free-free, bound-free, collision-induced absorption in the gas phase, and grain absorption by condensed phases. For the temperatures and pressures of interest, free-free opacity sources are the least important, so are only briefly covered in \\S\\ref{freefree}, since they can make at most a minor contribution at the upper end of the temperature range of interest. We discuss a number of free-free opacity sources, but in practice the only ones of even minor importance are $H_{ff}$ (Carson 1988a), $H^-_{ff}$ (Wishart 1979), $He^-_{ff}$ (Bell, Berrington \\& Croskery 1982), and $H_{2ff}^-$ (Bell 1980). We adopt the convention of including the free electron in the net charge of the whole system undergoing photon absorption, so here $H_{ff}$ refers to an electron moving in the field of a proton, and $H^-_{ff}$ refers to an electron moving in the field of a neutral hydrogen atom, and likewise for the other species. Of considerably greater importance is bound-free absorption, as discussed in \\S\\ref{boundfree}. We calculate the contribution due to $H$ with Gaunt factors from Carson (1988a), its negative ion $H^-$ (Wishart 1979; Bell \\& Berrington 1987), and the atomic species $Na$ (Cunto \\& Mendoza 1992; Cunto et al. 1993), $K$ (Verner \\& Yakovlev 1995), and $Fe$ (Kurucz 1995). The ion $H^-$ is important only when there is a supply of free electrons, and this is the case only when the species with the lowest ionization potentials, namely the alkali elements, become ionized. At high pressures, collision-induced absorption, as discussed in \\S\\ref{ciaopac}, can be very important in the infrared. In an astrophysical mixture for the temperatures of interest to us here, the two most important contributions are due to $H_2-H_2$ (Borysow et al. 1985; Zheng \\& Borysow 1995a; Zheng \\& Borysow 1995b; Borysow \\& Frommhold 1990; Lenzuni, Chernoff, \\& Salpeter 1991; Guillot et al. 1992), and $H_2-He$ collisions (Borysow, Frommhold, \\& Birnbaum 1988; Borysow, Frommhold, \\& Moraldi 1989; Borysow \\& Frommhold 1989). Although the contribution to the opacity due to $H_2-CH_4$ collisions is much weaker, $CH_4$ is an important molecule and the data are available. The detailed discussion of opacity calculations is completed in \\S\\ref{grainopac} by calculating Mie scattering due to grains formed by the condensation of species out of the gas phase (Van de Hulst 1957; Sudarsky 2002). Both the real and imaginary components of the refractive indices are needed and used. Once all the individual opacity sources have been calculated, they must be weighted by abundances and combined into monochromatic opacities or suitable mean opacities, as discussed in \\S\\ref{total}. For the abundances calculated in \\S\\ref{atmosabun}, suitable thermodynamic data are required, and methods of handling them are discussed. Once the abundances are used to derive the total monochromatic opacity, Rosseland and Planck mean opacities can be calculated, if desired. ", "conclusions": "\\label{conclusions} In this work, we have discussed and detailed the main sources of opacity in the cool atmospheres of brown dwarfs and extrasolar giant planets. Since these objects have lower atmospheric temperatures than stars, a number of diatomic and polyatomic molecules are present, which are not found in abundance in most stellar types. Such molecules can have a very complicated spectrum with a large number of lines. Because of the complex spectrum of molecules, and because they are such important sources of absorption, a considerable amount of effort is required to calculate their contribution to the opacity. Detailed calculations for many millions of lines over an extensive frequency grid are involved. Most other papers dealing with the opacities in the atmospheres of substellar objects, consider only one or a few specific molecules, often in a restricted wavelength region. This is one of the first papers to provide in one place a comprehensive discussion of the important molecular and atomic opacity sources, covering a broad range of wavelengths from the near ultraviolet through to the far infrared. In practical opacity calculations, once the monochromatic absorption of each species has been considered, their abundances have to be calculated before they can be combined to obtain total opacities. The minimization of the free energy of the system can be used to determine equilibrium abundances, and we have provided a comprehensive summary of the techniques needed to derive such abundances and results of such calculations. Once the complete monochromatic opacities are obtained, mean opacities weighted by abundances can be calculated. In order to model substellar objects correctly, a knowledge of how the atmosphere absorbs radiation is necessary. This paper is meant to provide a compendium of approaches, references, and results to aid the student interested in substellar dense atmospheres and the associated spectroscopy and chemistry." }, "0607/hep-th0607140_arXiv.txt": { "abstract": "To alleviate the black-hole (BH) information problem, we study a holographic-principle-inspired nonlocal model of Hawking radiation in which radiated particles created at different times all have the same temperature corresponding to the instantaneous BH mass. Consequently, the black hole loses mass not only by continuously radiating new particles, but also by continuously warming previously radiated particles. The conservation of energy implies that the radiation stops when the mass of the black hole reaches the half of the initial BH mass, leaving a massive BH remnant with a mass much above the Planck scale. ", "introduction": " ", "conclusions": "" }, "0607/hep-th0607006_arXiv.txt": { "abstract": " ", "introduction": "Recent Astrophysical Data, from either studies of distant supernovae type Ia~\\cite{snIa}, or precision measurements of temperature fluctuations in the cosmic microwave background radiation from the WMAP satellite~\\cite{wmap}, point towards a current-era acceleration of our Universe, as well as a very peculiar energy budget for it, 70\\% of the energy density of which consists of an unknown energy substance, termed dark Energy. In fact, global best-fit models of a compilation of all the available data at present are provided by simple Einstein-Friedman Universes with a (four space-time dimensional) positive {\\it cosmological constant} $\\Lambda$, whose value saturate the Newtonian upper limit obtained from galactic dynamics, namely in order of magnitude \\begin{equation} \\Lambda \\sim 10^{-122}M_P~\\qquad (M_P = 10^{19}~{\\rm GeV}~). \\label{lambdaval} \\end{equation} Although, as a classical (general relativistic) field theory, such a model is fairly simple, from a quantum theory view point it appears to be the less understood at present. The reason is simple: Since in cosmology~\\cite{cosmology} the radiation and matter energy densities scale with inverse powers of the scale factor, $a^{-4}$ and $a^{-3}$ respectively, in a Universe with positive cosmological constant $\\Lambda$, the vacuum energy density remains constant and positive, and eventually dominates the energy budget. The asymptotic (in time) Universe becomes a {\\it de Sitter} one, and in such a Universe the scale factor will increase exponentially, \\begin{equation} a(t) = a_0 e^{\\sqrt{\\frac{\\Lambda}{3}}t}~, \\label{inflation} \\end{equation} thereby implying that the Universe will eventually enter an inflationary phase again, and in fact it will accelerate eternally, since ${\\ddot a} > 0$, where the overdot denotes derivative with respect to the Robertson-Walker cosmic time, $t$, defined by: \\begin{equation} ds_{RW}^2 = -dt^2 + a^2(t)ds_{\\rm spatial}^2 \\label{rw} \\end{equation} In such de Sitter Universes there is unfortunately a {\\it cosmic horizon} \\begin{equation} \\delta \\propto \\int_{t_0}^{t_{End}} \\frac{cdt}{a(t)} < \\infty \\label{horizon} \\end{equation} where $t_{End}$ indicates the end of time. For a closed Universe $t_{End} < \\infty$, but for an open or flat Universe $t_{End} \\to \\infty$. The Cosmic Microwave (CMB) data of WMAP and other experiments at present indicate that our Universe is spatially {\\it flat}, and hence $t_{End} \\to \\infty$. The presence of a cosmic horizon implies that it is not possible to define pure state vectors of quantum asymptotic (in time) states. Therefore, the entire concept of a well-defined and gauge invariant Scattering matrix $S$ breaks down in quantum field theories defined on such de Sitter space time backgrounds. For string theory this is bad news, because precisely by construction~\\cite{strings}, perturbative string theory is based on the well-defined nature of scattering amplitudes of various excitations, and hence on a well-defined S-matrix~\\cite{smatrix}. This is a challenge for string theory, and certainly one of the most important issues I would like to discuss in this brief review. A straightforward way out, would be {\\it quintessence}-like scenaria for dark energy~\\cite{quint}, according to which the latter is due to a potential of a time dependent scalar field, which has not yet reached its equilibrium point. If, then, the asymptotic value of the dark energy vanishes in such a way so as not to have a cosmic horizon, then the model could be accommodated within string-inspired effective field theories, and could thus characterise the low-energy limit of strings, given that an asymptotic S-matrix could be defined in such a case. However, this does not mean that de Sitter Universes {\\it per se} cannot be accommodated somehow into a (possibly non perturbative) string theory framework. Their anti-de-Sitter (AdS, negative cosmological constant) counterparts certainly do, and in fact there have been important development towards a holographic property of quantum field theories in such Universes, due to the celebrated Maldacena conjecture~\\cite{malda}, concerning quantum properties of (supersymmetric) conformal field theories on the boundary of AdS space time. As we shall discuss in the next section, similar conjectures~\\cite{strom} may characterise their de Sitter counterparts, and this may be a way forward to accommodate such a space time into string theory. Finally, a more straightforward (perturbative) approach to discuss de Sitter and inflationary scenaria in string theories, will be to use the so-called non-critical (or Liouville) string framework~\\cite{ddk}, dealing with a mathematically consistent way of discussing strings propagating in non-conformal backgrounds, such as the de Sitter space time. This theory, however, at least as far as computation of the pertinent correlation functions are concerned, has not been developed to the same level of mathematical understanding as the critical strings. A crucial ingredient in this approach is the identification of the Liouville mode with the target time~\\cite{emn}, which allows for some non-conformal backgrounds in string theory, including de Sitter space times and accelerated Universes, to be accommodated in a mathematically consistent manner. We should stress at this point, that the above considerations, regarding S-matrix amplitudes in de Sitter Universes, refer to pure perturbative string theories. In the modern approach to string theory, where membrane (D-brane) structures~\\cite{membranes} also appear as mathematically consistent entities, the presence of a dark energy on the string theory on the brane is unavoidable, unless extreme conditions on unbroken supersymmetry and static nature of brane worlds are imposed. However, in brane cosmology one needs moving branes, in order to obtain a cosmological space time~\\cite{binetruy}, and in this case, target space time supersymmetry breaks down, due to the brane motion, resulting in non-trivial vacuum energy contributions on the brane~\\cite{maartens}. The structure of the article will be the following: in section 2, I will deal with mathematical properties of de Sitter space times: after reviewing briefly basic features of this geometry, I will describe modern approaches to the issue of placing a quantum field theory in de Sitter space times, by discussing briefly a holographic conjecture, put forward by Strominger~\\cite{strom}, according to which a quantum field theory on the single boundary of de Sitter space can be related to a classical theory in the bulk, in a way not dissimilar to the celebrated Maldacena conjecture~\\cite{malda} for anti-de Sitter spaces (negative cosmological constant space times). In section 3, I will discuss the issue of cosmic horizons in perturbative string theory, and give further arguments that consistent perturbative strings cannot be characterized by such horizons. In section 4, I will discuss quintessence scenaria in strings, where the dilaton behaves as the quintessence field, responsible for the current acceleration of the Universe. I will discuss two opposite examples, a pre Big-Bang scenario~\\cite{veneziano}, in which the string coupling increases at late times, with string loop corrections playing a dominant r\\^ole, and another scenario~\\cite{emn,dgmpp}, in which the string coupling becomes more and more perturbative as the time passes, leading asymptotically to a vanishing dark energy, in such a way that S-matrix states can be defined. In this second scenario the current-era acceleration parameter turns out to be proportional to the square of the string coupling, which at present enjoys perturbative values compatible with particle physics phenomenology. I will briefly discuss predictions of such models in the context of recent data, but also unresolved problems. I will not discuss the issue of dark energy in brane cosmologies in this article, as this is a topic covered by other lecturers in the school~\\cite{maartens}. Conclusions and directions for future research in the issue of Dark Energy in Strings will be presented in section 5. ", "conclusions": "In this work we have reviewed various issues related to the consistent incorporation of Dark Energy in string theory. We have discussed only traditional string theory and did not cover the modern extension, including membranes. This topic has been covered by other lecturers in the School. One of the most important issues concerns de Sitter space, and in general space-times with horizons in string theory. We have studied general properties, including holographic scenaria, which may be the key to an inclusion of such space times in the set of consistent (possibly non perturbative) ground states of strings. We have also seen that perturbative strings are incompatible with space times with horizons, mainly due to the lack of a scattering matrix. However, non-critical strings may evade this constraint, and we have discussed briefly how accelerating universes can be incorporated in non critical (Liouville) strings. The use of Liouville strings to describe the evolution of our Universe is natural, since non-critical strings are associated with non-equilibrium situations which undoubtedly occurred in the early Universe. The dilaton played an important r\\^ole in string cosmology, and we have seen how it can act as a quintessence field, responsible for the current-era acceleration of the Universe. There are many phenomenological tests of this class of cosmologies that can be performed, which the generic analysis presented here is not sufficient to encapsulate. Tensor perturbations in the cosmic microwave background radiation is one of them. The emission of gravitational degrees of freedom from the hot brane to the cold bulk, during the inflationary and post-inflationary phases in models involving brane worlds is something to be investigated in detail. A detailed knowledge of the dependence of the equation of state on the redshift is something that needs to be looked at in the context of specific models. Moreover, issues regarding the delicate balance of the expansion of the Universe and nucleosynthesis, which requires a very low vacuum energy, must be resolved in specific, phenomenologically semi-realistic models, after proper compactification to three spatial dimensions, in order that the conjectured cosmological evolution has a chance of success. Finally, the compactification issue \\emph{per se} is a most important part of a realistic stringy cosmology. In our discussion above, we have assumed that a consistent compactification takes place, leading to effective four-dimensional string-inspired equations of motion. In realistic scenaria, however, details of how the extra dimensions are compactified play a key r\\^ole in issues like supersymmetry breaking. In this review I did not discuss higher-curvature modifications of the low-energy Einstein action, which characterise all string-inspired models, including brane worlds scenaria. Such terms may play an important r\\^ole in Early Universe cosmology. For instance, they may imply initial singularity-free string cosmologies~\\cite{tamvarizos}, or non-trivial black hole solutions with (secondary) dilaton hair~\\cite{kanti}, which can play a r\\^ole in the Early universe sphaleron transitions. So, before closing the lecture, I will devote a few words on their form. In ordinary string theory, which is the subject of the present lecture, such higher-order terms possess ambiguous coefficients in the effective action. This is a result of local field redefinitions, which leave the (low-energy) string scattering amplitudes invariant, and hence cannot be determined by low energy considerations. In ordinary string theory~\\cite{strings}, with no space-time boundaries in (the low-energy) target space time, such ambiguities imply that the so-called ghost-free Gauss-Bonnet combination $\\frac{1}{g_s^2}\\left(R_{\\mu\\nu\\rho\\sigma}^2 - 4 R_{\\mu\\nu}^2 + R^2\\right)$, with $g_s = e^\\Phi$ the string coupling and $\\Phi$ the dilaton field, can always be achieved for the quadratic curvature terms in the string-inspired low-energy effective action, which constitutes the first non trivial order corrections to Einstein term in bosonic and heterotic string effective actions. However, in the case of brane worlds, with closed strings propagating in the bulk, things are not so simple. As discussed in \\cite{mp}, field redefinition ambiguities for the bulk low-energy graviton and dilaton fields, that would otherwise leave bulk string scattering amplitudes invariant, induce brane (boundary) curvature and cosmological constant terms, with the unavoidable result of ambiguities in the terms defining the Einstein and cosmological constant terms on the brane. This results in (perturbative in $\\alpha '$) ambiguities in the cross-over scale of four-dimensional brane gravity, as well as the brane vacuum energy. It is not clear to me, however, whether these ambiguities are actually present in low-energy brane world scenaria. I believe that these bulk-string ambiguities can be eliminated once the brane effective theory is propertly defined, given that closed and open strings also propagate on the brane world hypersurfaces, and thus are characterised by their own scattering amplitudes. Matching these two sets of scattering amplitudes properly, for instance by looking at the conformal theory describing the splitting of a closed-string bulk state, crossing a brane boundary, into two open string excitations on the brane, may lead to unambiguous brane cross-over and cosmological constant scales, expressed in terms of the bulk string scale and coupling. These are issues that I believe deserve further investigation, since they affect early Universe cosmologies, where such higher-curvature terms are important. I will not, however, discuss them further here. I would like to close this lecture with one more remark on the non-equilibrium Liouville approach to cosmology advocated in \\cite{emn,emn04}, and discussed last in this article. This approach is based exclusively on the treatment of target time as an irreversible dynamical renormalization-group scale on the world sheet of the Liouville string (the zero mode of the Liouville field itself). This irreversibility is associated with fundamental properties of the world-sheet renormalization group, which lead in turn to the loss of information carried by two-dimensional degrees of freedom with world-sheet momenta beyond the ultraviolet cutoff~\\cite{zam} of the world-sheet theory. This fundamental microscopic time irreversibility may have other important consequences, associated with fundamental violations of CPT invariance~\\cite{mavrodecoh} in both the early Universe and the laboratory, providing other tests of these ideas. \\noindent {\\bf Acknowledgements} \\\\ It is my pleasure to thank the organisers, and especially E. Papantonopoulos, for the invitation to lecture in this very interesting school and workshop. This work is partially supported by funds made available by the European Social Fund (75\\%) and National (Greek) Resources (25\\%) - (EPEAEK II) - PYTHAGORAS." }, "0607/astro-ph0607094_arXiv.txt": { "abstract": "We have studied several sources of systematic uncertainty in calculating the aperture of the High Resolution Fly's Eye experiment (HiRes) in monocular mode, primarily as they affect the HiRes-II site. The energy dependent aperture is determined with detailed Monte Carlo simulations of the air showers and the detector response. We have studied the effects of changes to the input energy spectrum and composition used in the simulation. A realistic shape of the input spectrum is used in our analysis in order to avoid biases in the aperture estimate due to the limited detector resolution. We have examined the effect of exchanging our input spectrum with a simple E$^{-3}$ power law in the ``ankle'' region. Uncertainties in the input composition are shown to be significant for energies below $\\sim$ 10$^{18}$ eV for data from the HiRes-II detector. Another source of uncertainties is the choice of the hadronic interaction model in the air shower generator. We compare the aperture estimate for two different models: QGSJet01 and SIBYLL 2.1. We also describe the implications of employing an atmospheric database with hourly measurements of the aerosol component, instead of using an average as has been used in our previously published measurements of the monocular spectra. ", "introduction": "\\label{intro} The High Resolution Fly's Eye experiment consists of two air fluorescence detectors (``HiRes-I'' and ``HiRes-II'') located in the desert of Utah. HiRes observes ultra-high energy cosmic rays indirectly through extensive air showers, i.e. cascades of secondary charged particles, which are caused by interactions of the primary cosmic ray particles with the earth's atmosphere. In the wake of the air shower, excited nitrogen molecules emit fluorescence light in the ultraviolet, which is collected by mirrors and projected onto clusters of photomultiplier tubes. Detailed descriptions of the detectors can be found in~\\cite{hr_det1} and~\\cite{hr_det2}. The HiRes experiment aims at measuring the arrival directions, composition, and flux of the most energetic cosmic rays. The two detectors allow stereoscopic observation of air showers, which yields the best resolution in shower geometry and cosmic ray energy. An advantage of data analysis in monocular mode, i.e. separate analyses of the data from each of the two detectors, lies in the higher statistics that can be achieved at the high energy end due to the longer lifetime of the HiRes-I detector, which started operation two years before HiRes-II, in 1997. Monocular analysis also allows an extension of the observed energy range down to energies as low as $\\sim$10$^{17}$ eV due to the larger elevation coverage and better time resolution of the HiRes-II detector, and also due to the necessity of triggering only one detector. The differential flux or ``energy spectrum'' observed in monocular mode by HiRes shows a hardening in the flux at around $10^{18.5}$ eV, known as the ``ankle'', and a suppression of the flux near $10^{19.8}$ eV, at the expected energy of the GZK flux suppression~\\cite{gr,zk}. These results have been published in~\\cite{hr_ankle}. The main systematic uncertainties that are introduced in the UHECR spectrum measurement with the HiRes detectors have been reported in~\\cite{hr_prl}. They are uncertainties in the absolute phototube calibration ($\\pm$10\\%), the fluorescence yield ($\\pm$10\\%) and the correction for ``missing energy'' ($\\pm$5\\%). The latter refers to the energy component that is channeled mainly into neutrinos and does not contribute to the ionization process. Not taking into account atmospheric effects, the uncertainty in the energy scale is $\\pm$15\\%, which results in a systematic uncertainty in the flux $J$ of $\\pm$27\\%. The effect on the energy scale of a variation of the average vertical aerosol optical depth (VAOD) by $\\pm$1 RMS value, from 0.04 to 0.06 and 0.02, has also been described in~\\cite{hr_prl}. It is not larger than 9\\% on the average. This results in a total uncertainty in the energy scale of $\\pm$17\\%. The effect of the same VAOD variation on the aperture leads to an average atmospheric uncertainty in the flux $J$ of $\\pm$15\\%. The total systematic uncertainty in the measured flux adds up to $31 \\%$ for each of the two monocular spectrum measurements. In this paper, we examine additional systematic uncertainties that may affect the calculation of the HiRes aperture in monocular mode. Since the aperture of an air fluorescence detector is a function of the energy of the observed cosmic rays, it has to be modeled carefully with detailed Monte Carlo (MC) simulations. The HiRes MC simulation programs use libraries of air shower profiles, generated at different energies with the air shower simulation program CORSIKA~\\cite{corsika} and the hadronic interaction code QGSJet~\\cite{qgsjet}, for a realistic representation of the fluctuations in the observed charged particle profiles. A detector response MC program simulates the light emission process along the shower and traces the photons through the atmosphere to the telescopes of the two detectors, taking into account all relevant atmospheric effects. The detector optics, electronics and trigger system are modeled in great detail using databases that record variable detector settings, as well as density fluctuations of aerosols in the atmosphere. After performing extensive comparisons between simulated events and data, which allow us to verify the quality of our simulations (see~\\cite{hr_hr2}), we estimate the detector acceptance using the ratio of accepted MC events ($\\nu^{MC}$) to generated MC events ($\\mu^{MC}$) in each energy bin. To correctly simulate effects stemming from the finite energy resolution of the detectors and their limited elevation coverage, we use a continuous energy spectrum and a bi-modal composition based on previous measurements as inputs to our simulation programs. The differential flux $J$ in each energy bin is calculated as: \\begin{equation} J(E_i)= N(E_i)\\cdot \\frac{1}{\\Delta E} \\cdot \\frac{1}{C_i \\cdot A \\Omega \\cdot t} \\label{eq:espec} \\end{equation} where $N(E_i)$ is the number of observed events in the energy bin and $\\Delta E$ is the bin-width. The geometrical aperture (area $\\times$ solid angle) used in generating MC events is noted by $A \\Omega$, and $t$ is the detector live-time. Through our use of a continuous and realistic input energy spectrum, the finite energy resolution of the detectors is taken into account in the acceptance $C_i$ = $\\frac{\\nu_i^{MC}}{\\mu_i^{MC}}$. This will be explained in the next chapter. In the following, we refer to the product of the constant $A \\Omega$ and the acceptance as (instantaneous) aperture. We will first consider the effects of varying the input energy spectrum on the calculated aperture and thus on the measured spectrum, in Section~\\ref{einput}. In Section~\\ref{hadmod}, we examine the implications of exchanging the hadronic interaction model in the air shower generator. For this study, we replace the QGSJet model, which is used in our standard spectrum measurement, with the SIBYLL model~\\cite{sibyll}. The effect of a variation of the assumed input composition on the measured spectrum is presented in Section~\\ref{compinput}. Another systematic uncertainty that can affect the aperture estimate of the experiment is that due to variations in the aerosol component of the atmosphere. For the analysis of the monocular spectra, we used an average atmospheric description based on measurements with laser systems that are installed at each detector site~\\cite{hr_atmos1}. In Section~\\ref{atmos}, we re-analyze the HiRes-II monocular data with a database containing hourly measurements of the aerosol component of the atmosphere and compare it to the average description in our standard analysis. Although the systematic studies presented in this paper have been carried out with simulation and reconstruction tools of the HiRes-II analysis, their results are applicable to the HiRes-I spectrum measurement as well. ", "conclusions": "\\label{conclusions} None of the sources of possible systematic uncertainties we have studied here contribute significantly to our published estimate of the systematic uncertainty. The bias introduced by using an E$^{-3}$ power law instead of a more realistic spectral shape is not very significant in the ``ankle'' region. Our calculated aperture is sensitive to the assumed input composition for energies below $\\sim$10$^{18}$ eV for HiRes-II. By using a measured composition as an input to our simulation programs, our analysis does not depend on the assumed hadronic interaction model. For the 17 month period tested here, the description of the aerosol density using an hourly database does not cause any significant differences in the spectrum, when compared with an average atmosphere. We also found no significant changes in the reconstructed energies for the time period under study." }, "0607/astro-ph0607577_arXiv.txt": { "abstract": "The statistical expectation values of the temperature fluctuations and polarization of cosmic microwave background (CMB) are assumed to be preserved under rotations of the sky. We investigate the statistical isotropy (SI) of the CMB maps recently measured by the Wilkinson Microwave Anisotropy Probe (WMAP) using the bipolar spherical harmonic formalism proposed in Hajian \\& Souradeep 2003 for CMB temperature anisotropy and extended to CMB polarization in Basak, Hajian \\& Souradeep 2006. The {\\em Bipolar Power Spectrum (BiPS)} had been measured for the full sky CMB anisotropy maps of the first year WMAP data and now for the recently released three years of WMAP data. We also introduce and measure directional sensitive {\\em reduced Bipolar coefficients} on the three year WMAP ILC map. Consistent with our published results from first year WMAP data we have no evidence for violation of statistical isotropy on large angular scales. Preliminary analysis of the recently released first WMAP polarization maps, however, indicate significant violation of SI even when the foreground contaminated regions are masked out. Further work is required to confirm a possible cosmic origin and rule out the (more likely) origin in observational artifact such as foreground residuals at high galactic latitude. ", "introduction": "In standard cosmology, CMB anisotropy signal is expected to be statistically isotropic, i.e., statistical expectation values of the temperature fluctuations $\\Delta T(\\hat q)$ are preserved under rotations of the sky. In particular, the angular correlation function $C(\\hat{q},\\, \\hat{q}^\\prime)\\equiv\\langle\\Delta T(\\hat q)\\Delta T(\\hat q^\\prime)\\rangle$ is rotationally invariant for Gaussian fields. In spherical harmonic space, where $\\Delta T(\\hat q)= \\sum_{lm}a_{lm} Y_{lm}(\\hat q)$ the condition of {\\em statistical isotropy} (SI) translates to a diagonal $\\langle a_{lm} a^*_{l^\\prime m^\\prime}\\rangle=C_{l} \\delta_{ll^\\prime}\\delta_{mm^\\prime}$ where $C_l$, is the widely used angular power spectrum of the CMB anisotropy. Statistical isotropy of the CMB sky is essential for $C_l$ to be a complete description of (Gaussian) CMB anisotropy and, hence, an adequate measure for carrying out cosmological parameter estimation of `standard' (SI) model. Hence, it is crucial to be able to determine from the observed CMB sky whether it is a realization of a statistically isotropic process, or not. The detection of statistical isotropy (SI) violations in the CMB signal can have exciting and far-reaching implications for cosmology and the cosmological principle. For example, a generic consequence of cosmic topology is the breaking of statistical isotropy in characteristic patterns determined by the photon geodesic structure of the manifold as probed by the CMB photons traveling to us from the surface of last scattering over a distance comparable to the cosmic horizon, $R_H$~\\cite{bps98,bps00a,bps00b}. Mildly anisotropic cosmological models predict charateristic patterns hidden in the CMB sky. On the other hand, SI violation could also arise from foreground contamination, or, artifacts of observational and analysis techniques. The CMB measurements of the {\\it Wilkinson Microwave Anisotropy Probe} ({\\it WMAP}) are consistent with predictions of the concordance $\\Lambda$CDM model with (nearly) scale-invariant and adiabatic fluctuations expected to have been generated during an inflationary epoch in the early universe~\\cite{hin_wmap03, kogut_wmap03, sper_wmap03, page_wmap03, peiris_wmap03,sper_wmap06,hin_wmap06,pag_wmap06}. After the first year of {\\it WMAP} data, the SI of the CMB anisotropy ({\\it i.e.} rotational invariance of n-point correlations) has attracted considerable attention. Tantalizing evidence of SI breakdown (albeit, in very different guises) has mounted in the {\\it WMAP} first year sky maps, using a variety of different statistics~\\cite{erik04a, Copi:2003kt, Schwarz:2004gk, Hansen:2004vq, angelwmap, Land:2004bs,Land:2005ad, Land:2005dq, Land:2005jq,Land:2005cg, Bielewicz:2004en, Bielewicz:2005zu,Copi:2005ff, Copi:2006tu, Naselsky:2004gm, Prunet:2004zy,Gluck:2005td, Stannard:2004yp, Bernui:2005pz, Bernui:2006ft, Freeman:2005nx, Chen:2005ev,Wiaux:2006zh}. The three-year WMAP maps are consistent with the first-year maps up to a small quadrupole difference. The two additional years of data and the improvements in analysis has not significantly altered the low multipole structures in the maps~\\cite{hin_wmap06}. Hence, `anomalies' are expected to persist at the same modest level of significance but are now less likely to be artifacts of noise, uncorrected systematics, or the analysis of the first year data. The cosmic significance of these `anomalies', however, remains debatable because of the aposteriori statistics employed to ferret them out of the data. {\\em More importantly, what is missing is a common, well defined, mathematical language to quantify SI (as distinct from non Gaussianity) and the ability to ascribe statistical significance to the anomalies unambiguously.} We employ a well defined, mathematical language of Bipolar harmonic decomposition of the underlying correlation to quantify SI that can ascribe statistical significance to the anomalies unambiguously. The observed CMB sky is a single realization of the underlying correlation, hence the detection of SI violation or correlation patterns pose a great observational challenge. For statistically isotropic CMB sky, the correlation function \\begin{equation} C(\\hat{n}_1,\\hat{n}_2)\\equiv C(\\hat{n}_1\\cdot\\hat{n}_2) = \\frac{1}{8\\pi^2}\\int d{\\mathcal R} C({\\mathcal R}\\hat{n}_1,\\, {\\mathcal R}\\hat{n}_2), \\label{avg_cth} \\end{equation} where ${\\mathcal R}\\hat{n}$ denotes the direction obtained under the action of a rotation ${\\mathcal R}$ on $\\hat{n}$, and $d{\\mathcal R}$ is a volume element of the three-dimensional rotation group. The invariance of the underlying statistics under rotation allows the estimation of $C(\\hat{n}\\cdot\\hat{n}')$ using the average of the temperature product $\\widetilde{\\Delta T}(\\hat n) \\widetilde{\\Delta T}(\\hat n')$ between all pairs of pixels with the angular separation $\\theta$. In the absence of statistical isotropy, $C(\\hat{n},\\hat{n}')$ is estimated by a single product $\\widetilde{\\Delta T}(\\hat n)\\widetilde{\\Delta T}(\\hat n')$ and hence is poorly determined from a single realization. Although it is not possible to estimate each element of the full correlation function $C(\\hat{n},\\hat{n}')$, some measures of statistical anisotropy of the CMB map can be estimated through suitably weighted angular averages of $\\widetilde{\\Delta T}(\\hat n)\\widetilde{\\Delta T}(\\hat n')$. The angular averaging procedure should be such that the measure involves averaging over sufficient number of independent `measurements', but should ensure that the averaging does not erase all the signature of statistical anisotropy. We proposed the Bipolar Power spectrum (BiPS) $\\kappa_\\ell$ ($\\ell=1,2,3, \\ldots$) of the CMB map as a statistical tool of detecting and measuring departure from SI~\\cite{us_apjl,us_pascos} and reviewed in this article in sec.~\\ref{bips}. The non-zero value of the BiPS spectrum imply the break down of statistical isotropy \\begin{equation} {\\mathrm {\\Huge STATISTICAL\\,\\,\\,\\, ISOTROPY}} \\,\\,\\,\\,\\,\\,\\, \\Longrightarrow \\,\\,\\,\\,\\,\\,\\, \\kappa_\\ell\\,=\\,0 \\,\\,\\,\\,\\,\\,\\, \\forall \\ell \\ne 0. \\label{bipsstate} \\end{equation} The BiPS is sensitive to structures and patterns in the underlying total two-point correlation function \\cite{us_apjl, us_pascos}. The BiPS is particularly sensitive to real space correlation patterns (preferred directions, etc.) on characteristic angular scales. In harmonic space, the BiPS at multipole $\\ell$ sums power in off-diagonal elements of the covariance matrix, $\\langle a_{lm} a_{l'm'}\\rangle$, in the same way that the `angular momentum' addition of states $l m$, $l' m'$ have non-zero overlap with a state with angular momentum $|l-l'|<\\ell$=$-$2.4), typical of tailed radio sources in clusters of galaxies; \\item the significantly flatter spectrum of the component associated with BG2 with respect to the extended emission ($<\\alpha>=-0.4$ and $<\\alpha>=-2.6$). \\end{itemize} \\noindent We therefore consider that R5 is a NAT radio source associated with BG2, with a core component in its Northernmost side, and an extended ($\\sim170{\\times}160$~\\h~kpc) tail. This latter is distorted towards NW, possibly due to a lower local density of the intra-cluster gas revealed by the XMM observations of B05. The tail morphology is very likely explained by ram-pressure effects due to the relative motion between the ICM and BG2, in turn related to the merging event of the host sub-cluster (A3921-B). The total radio power of R5 at 1.344~GHz is Log $(P{\\rm (W/Hz)})$=24.0. Its total flux density is dominated by the extended emission located outside the optical galaxy, with a flux density ratio of $\\sim$18 between the extended and the core component ($\\sim41.4$ and $\\sim2.3$ mJy respectively). R5 does not show visible jets and the emission in the tail is very steep ($\\alpha{\\simeq}-2{\\div}-3.5$), suggesting that electrons are suffering strong radiation losses. The diffuse component of R5 would therefore be a `dying' tail, which has been partly detached from an earlier period of activity of the BG2 radio galaxy and whose final evolution is dominated by synchrotron losses. The extended emission is now presumably supported and distorted by the ICM pressure. The strong polarization at the extreme of the NW component of the tail suggests that the radio jet has encountered a denser (or higher pressure) region leading to magnetic field compression." }, "0607/astro-ph0607261_arXiv.txt": { "abstract": "{} {We present the results of a detailed analysis of 452 ground-based high-resolution high S/N spectroscopic measurements spread over 4.5 years for $\\beta$~Canis Majoris with the aim to determine the pulsational characteristics of this star, and to use them to derive seismic constraints on the stellar parameters.} {We determine pulsation frequencies in the Si\\,III 4553~\\AA\\ line with Fourier methods. We identify the $m$-value of the modes by taking into account the photometric identifications of the degrees $\\ell$. To this end we use the moment method together with the amplitude and phase variations across the line profile. The frequencies of the identified modes are used for a seismic interpretation of the structure of the star.} {We confirm the presence of the three pulsation frequencies already detected in previous photometric datasets: $f_1 = 3.9793~\\cpd$ ($46.057~\\muhz$), $f_2 = 3.9995~\\cpd$ ($46.291~\\muhz$) and $f_3 = 4.1832~\\cpd$ ($48.417~\\muhz$). For the two modes with the highest amplitudes we unambiguously identify $(\\ell_1,m_1) = (2,2)$ and $(\\ell_2,m_2) = (0,0)$. We cannot conclude anything for the third mode identification, except that $m_3 > 0$. We also deduce an equatorial rotational velocity of $31 \\pm 5~\\kmps$ for the star. We show that the mode $f_1$ must be close to an avoided crossing. Constraints on the mass ($13.5 \\pm 0.5 \\msun$), age ($12.4 \\pm 0.7$~Myr) and core overshoot ($0.20 \\pm 0.05\\,H_P$) of \\bcma\\ are obtained from seismic modelling using $f_1$ and $f_2$.} {} ", "introduction": "\\label{sec:intro} Many breakthroughs have recently been achieved in the field of asteroseismology of \\bcep\\ stars. The observation of a few pulsating modes led to constraints not only on global stellar parameters but also on the core overshoot parameter and on the non-rigid rotation of several \\bcep\\ stars. In particular modelling has been performed for HD~129929 \\citep{aerts03a,dupret04} and $\\nu$~Eri \\citep{pamyatnykh04,ausseloos04}. Our aim is to add other \\bcep\\ stars to the sample of those with asteroseismic constraints. The B\\,1\\,II-III bright \\bcep\\ star $\\beta$~Canis Majoris (HD\\,44743, HR\\,2294, $V_\\mathrm{mag} = 1.97$) is particularly interesting to study. Indeed, earlier photometric and spectroscopic data revealed that this object exhibits multiperiodicity with rather low frequencies in comparison with the frequencies of other \\bcep\\ stars, which would indicate that \\bcma\\ is either a reasonably evolved star or oscillates in modes different from the fundamental. The variability of \\bcma\\ has been known for one century and the star has been extensively studied. We refer to \\cite{albrecht08}, \\citet{henroteau18}, \\citet{meyer34} and \\citet{struve50} for the first spectroscopic measurements of \\bcma. Later \\citet{shobbrook73} found three pulsation frequencies from extensive photometric time series. The same three frequencies were recently confirmed by \\citet{shobbrook06} who analysed photometric measurements of a multisite campaign dedicated to the star. \\citet{aerts94} collected spectroscopic data in order to identify the modes of the known frequencies of \\bcma. In this paper, we present a similar analysis but based on a much larger number of spectra and using the version of the moment method improved by \\citet{briquet03}. We then construct stellar models which show oscillations in accordance with our unique identification of the modes of $\\beta$~Canis Majoris. The paper is organised as follows. Section~\\ref{sec:spec} describes the results from our spectroscopic observations, including data reduction, frequency analysis and mode identification. In Sect.~\\ref{sec:model} we present our seismic interpretation of \\bcma. We end the paper with a discussion about our results in Sect.~\\ref{sec:concl}. ", "conclusions": "\\label{sec:concl} $\\beta$~Canis Majoris is one of the \\bcep\\ stars whose variability has been observed and analysed for one century. It was discovered that this star pulsates with three frequencies rather low in comparison with other known stars of its type. For this reason \\bcma\\ is an important target for asteroseismology purposes. However, so far, no definite mode identification had been achieved for this star so that no modelling could be attempted. Our aim was to increase the number of known pulsating frequencies and mostly to provide a unique identification of the modes of \\bcma. Our study was based on 452 ground-based high-resolution high S/N spectroscopic measurements spread over 4.5 years. We used the Si\\,III 4553~\\AA\\ line to derive the pulsation characteristics of \\bcma. Our dataset unfortunately suffers from strong aliasing but the three established frequencies of the star were confirmed in the first three velocity moments of the line and in the spectra themselves. They are $f_1 = 3.9793~\\cpd$ ($f_1 = 46.057~\\muhz$), $f_2 = 3.9995~\\cpd$ ($f_2 = 46.291~\\muhz$) and $f_3 = 4.1832~\\cpd$ ($f_3 = 48.417~\\muhz$). Unfortunately no new frequencies were discovered neither in our spectroscopic observations nor in the recent multisite photometric measurements led by \\citet{shobbrook06}. The important result of the combination of both intensive campaigns is an identification of the two main modes of \\bcma, which is a strong constraint for further asteroseismic modelling of the star. The photometric identification by \\citet{shobbrook06} yielded $\\ell_1 = 2$ and $\\ell_2 = 0$. Our spectroscopic data could corroborate that the mode with $f_2$ is radial. We adopted the photometric identification of $\\ell_1$ and spectroscopic techniques allowed us to derive the $m$-value of the main mode. The application of the moment method gave a preference to $m_1 = 2$. Because moment solutions could not definitely exclude $m_1 = 1$ we made use of the behaviour of the amplitude distributions across the line profile for the best parameter sets given by the moment method. In this way we could conclude without any doubt that $(\\ell_1,m_1) = (2,2)$. For the third mode nothing could be concluded. In addition we derived a stellar equatorial rotational velocity of $31 \\pm 5 \\kmps$. The definite identification of two of the observed modes and a much improved estimate of the rotation velocity of \\bcma\\ allowed us to attempt the first seismic modelling of this star. Although it is not realistic to hope for a unique model to fit just two frequencies, we have thoroughly explored the stellar parameter space to derive reasonable constraints for the mass, age, and core overshoot. The most significant aspect of the seismic analysis is the fact that we could assert that the non-radial mode, $f_1$, is close to an avoided crossing. This implies a very strong constraint on the stellar parameters, especially the age of the star. At the same time, it rules out the possibility of the radial mode, $f_2$, being the fundamental mode. This makes \\bcma\\ one more \\bcep\\ star known to have a dominant radial overtone mode of pulsation \\citep[cf.][]{aerts06}. Our best fit models indicate that \\bcma\\ has a mass of $13.5\\pm 0.5 \\msun$, an age of $12.4\\pm 0.7$~Myr ($\\xc = 0.126\\pm 0.003$) and core overshoot of $\\dov = 0.20\\pm 0.05$. No satisfactory model can be found if core overshoot is absent. A small overshoot parameter is possible only for a higher mass along with high metallicity (and proportionally higher helium content). On the other hand, higher core overshoot is required if the star is metal-poor ($\\mbyh < -0.05$). However, \\citet{morel06} found the composition of \\bcma\\ not much different from that of the Sun, making such possibilities unlikely. Therefore, it is safe to conclude that the models with $\\dov=0.20\\pm 0.05$ are the most likely ones. If the chemical composition of \\bcma\\ could be known to higher accuracy independently, one would be able to constrain the other parameters even more. All the solutions turn out to have effective temperatures close to the cooler edge of the adopted errorbox on the \\hrd. This is consistent with the recent estimates of \\teff\\ \\citep[e.g.,][]{morel06}. We note that we had deliberately chosen a very conservative errorbox for effective temperature and luminosity; a stricter limit on these parameters would rule out some of our possible models. In retrospect, one can also try to identify the mode of oscillation for the third frequency, $f_3$, by comparing it with the theoretical model frequencies. In this comparison, we allowed for different rotational splitting values, $m_3$, for each non-radial mode with the restriction $m_3>0$, as we found in Sect.~\\ref{sec:spec}. This leads us to only one possibility for $f_3$, for all the models which match $f_1$ and $f_2$: $\\ell_3 = 3, m_3 = 2, n_3=-1$. We hope that this identification of $f_3$ can be checked through future observations. However, we cannot place further constraints on the models at this stage using this frequency; all the models with stellar parameters in the range restricted by the first two frequencies also match the third frequency with the identification given above within the uncertainty associated with the rotational velocity. One would need a more precise estimate of the rotational velocity to distinguish between these models. The rotation period may be calculated from our estimate of the equatorial velocity (Sect.~\\ref{sec:spec_vel}) and the radius of our best models which indeed lie in a narrow range (Table~\\ref{tab:bestmodels}). We estimate the rotation period to be $18.6 \\pm 3.3$~days, which indicates that \\bcma\\ is indeed a slow rotator; therefore, our assumption in neglecting higher order terms of the rotation velocity while calculating the frequency splitting stands justified. Despite the knowledge of only two frequencies for this star, the occurrence of the avoided crossing goes a long way towards constraining most of the stellar parameters. While one cannot expect to be so lucky for every star, we have shown that the identification of an avoided crossing might help us to extract a lot more information about the star than any normal mode." }, "0607/astro-ph0607507_arXiv.txt": { "abstract": "We present {\\it Spitzer Space Telescope} observations of the well-studied extremely red objects (EROs) \\hr\\ and \\lbds\\ from 3.6~$\\mu$m to 160~$\\mu$m. These galaxies are the prototypes of the two primary classes of EROs: dusty starbursts and old, evolved galaxies, respectively. Both galaxies, as well as \\lbdsb, another example of an old, quiescent galaxy, are well-detected out to 8 $\\mu$m. However, only the dusty starburst \\hr\\ is detected in the far-infrared. All three EROs have stellar masses of a few $\\times\\, 10^{11} M_\\odot$. Using evolutionary model fits to their multiband photometry, we predict the infrared colors of similar EROs at $1 < z < 2$. We find that blueward of observed 10~$\\mu$m, the two ERO classes are virtually indistinguishable photometrically. Deep spectroscopy and 24~$\\mu$m data allow the classes to be separated. ", "introduction": "} It has now been nearly twenty years since the first near-infrared surveys identified an intriguing population of optically faint galaxies with surprisingly bright near-infrared magnitudes \\markcite{elston:88, elston:89}(Elston, Rieke, \\& Rieke 1988, 1989). Initially proposed to be extremely distant ($z \\sim 10$) sources whose optical emission was absorbed by intergalactic hydrogen \\markcite{Mobasher:05}(\\eg Mobasher {et~al.} 2005), the ``extremely red object'' (ERO) population\\footnote{EROs are typically selected to have extreme optical-minus-infrared colors, such as $(R-K)_{\\rm Vega}>6.0$.} is now instead recognized to be primarily comprised of two interesting galaxy populations: (1) old, evolved galaxies whose red colors are caused by a dearth of young, luminous, massive stars, and (2) dusty starburst galaxies whose red colors are caused by dust absorption of the bluer light in actively star-forming regions. The prototype of each class is, respectively, the ``old, dead, and red'' radio galaxy \\lbds\\ at $z=1.552$ \\markcite{dunlop:96, spinrad:97}(Dunlop {et~al.} 1996; Spinrad {et~al.} 1997) and the dusty starburst galaxy \\hr\\ at $z=1.44$ \\markcite{hu:94, graham:96, dey:99}(Hu \\& Ridgway 1994; Graham \\& Dey 1996; Dey {et~al.} 1999). The weak radio galaxy \\lbdsb\\ at $z=1.432$ \\markcite{nolan:03, dey:06}(Nolan {et~al.} 2003; Dey {et~al.} 2006) is a second well-studied example of the old, evolved ERO class. EROs have a surface density of approximately 1000\\,deg$^{-2}$ for $R-K>6$ and $K<20$ (Vega), and comprise $\\sim 10\\%$ of all $K$-selected galaxies to $K\\simlt20$ \\markcite{thompson:99, Moustakas:04}(Thompson {et~al.} 1999; Moustakas {et~al.} 2004). For a recent comprehensive review of EROs see \\markcite{mccarthy:04}McCarthy (2004). With redshifts $1 \\simlt z \\simlt 2$ \\markcite{cimatti:03}(\\eg Cimatti {et~al.} 2003), clustering properties similar to elliptical galaxies \\markcite{daddi:03, moustakas:02}(Daddi {et~al.} 2003; Moustakas \\& Somerville 2002), and substantial stellar masses that contribute significant fractions of the global stellar mass density at the redshifts they are found \\markcite{kong:06}(\\eg Kong {et~al.} 2006), EROs clearly are an important population in the context of understanding galaxy formation and evolution \\markcite{somerville:04, nagamine:05}(\\eg Somerville {et~al.} 2004; Nagamine {et~al.} 2005). Volume-limited galaxy surveys of the local universe find a clear bimodality in galaxy properties, with the ``red sequence'' generally populated by the more massive and established galaxies \\markcite{kauffmann:03, blanton:05}(\\eg Kauffmann {et~al.} 2003; Blanton {et~al.} 2005). There is increasing evidence that this sequence is securely in place not only out to $z\\sim1$ \\markcite{bell:04, faber:06}(Bell {et~al.} 2004; Faber {et~al.} 2006), but possibly to much higher redshifts \\markcite{RoccaVolmerange:04, Labbe:05, papovich:06}({Rocca-Volmerange} {et~al.} 2004; Labb\\'e {et~al.} 2005; Papovich {et~al.} 2006). The processes by which these massive galaxies form must happen early. This is consistent with measurements of the evolution of the global stellar mass density $\\Omega_{*}(z)$ in field surveys. While between half and three-quarters of the present-day stellar mass is in place by $z\\sim1$ \\markcite{dickinson:03, rudnick:03, fontana:04, drory:04}(Dickinson {et~al.} 2003; Rudnick {et~al.} 2003; Fontana {et~al.} 2004; Drory {et~al.} 2004), only $\\sim 10\\%$ of the present-day stellar mass is in place at $z \\simgt 3$ \\markcite{drory:05}(Drory {et~al.} 2005). The evolution of $\\Omega_{*}$ is dramatic: either the high-redshift accountings are incomplete (\\eg because of cosmic variance effects, unaccounted for obscured populations, or significant population synthesis deficiencies), or probes in that redshift range will see the most dramatic epoch of galaxy assembly in progress. Indeed, studies of infrared selected samples of distant galaxies in the Great Observatories Origins Deep Survey \\markcite{Giavalisco:04}(GOODS; Giavalisco {et~al.} 2004) find significant evidence of massive galaxy assembly at $z \\sim 1 - 3$ \\markcite{Caputi:06, papovich:06}({Caputi} {et~al.} 2006; Papovich {et~al.} 2006). Understanding ERO demographics and energetics will provide key insight into all of these fundamental questions. EROs are represented by two different galaxy populations. With enough dedicated time and effort, young, dusty starbursts and old, evolved galaxies may be distinguished through spectroscopic features in the restframe UV and optical \\markcite{Cimatti:02}(\\eg Cimatti {et~al.} 2002a). A photometric technique proposed by \\markcite{Pozzetti:00}Pozzetti \\& Mannucci (2000) to distinguish between ERO types using combinations of optical and near-infrared colors was designed to work for objects around $z\\approx1.5$ \\markcite{Mannucci:02}(\\eg Mannucci {et~al.} 2002). \\markcite{Dorman:03}Dorman, {O'Connell}, \\& Rood (2003) suggest an alternate discriminating technique using mid-ultraviolet colors. A more detailed look at the actual $\\lambda_{\\rm rest}\\sim0.1-1.1\\, \\mu$m spectral energy distributions (SEDs) of EROs by \\markcite{Moustakas:04}Moustakas {et~al.} (2004) demonstrates that a larger baseline, extending longward of the observed $K$-band, is needed to differentiate all but the most dramatic EROs. In this paper we present 3.6\\,$\\mu$m to 160\\,$\\mu$m observations of the archetypical EROs obtained with the {\\it Spitzer Space Telescope} \\markcite{werner:04}(Werner {et~al.} 2004). Section~2 briefly summarizes previous studies of these three EROs. Section~3 describes the {\\it Spitzer} observations and data reduction, followed by an analysis of the broadband SEDs of these galaxies in \\S~4. Surprisingly, the two ERO populations remain nearly indistinguishable out to observed 8\\,$\\mu$m (rest-frame 3.2\\,$\\mu$m) and it is only at longer wavelengths that the populations separate (\\S~5). We adopt the concordance cosmology and use Vega-system magnitudes unless stated otherwise. ", "conclusions": "} We present {\\it Spitzer} IRAC and MIPS observations of three well-studied EROs. \\hr\\ is a dusty starburst class of EROs, while the weak radio sources \\lbds\\ and \\lbdsb\\ are evolved, quiescent EROs. The IRAC data show the $1.6 \\mu$m stellar peak in all three sources, providing accurate derivations of basic properties of their stellar populations. Consistent with the bright near-infrared magnitudes, we find all three galaxies are massive, with stellar masses $M_* = (3 - 6) \\times 10^{11} M_\\odot$. We also derive ages of a few Gyr for all three galaxies, though \\hr\\ is found to be dustier and to have proportionally more recent star formation. Only \\hr\\ is detected by MIPS, illustrating that $24 \\mu$m photometry provides a robust discriminant between the two primary classes of EROs. Blueward of $10 \\mu$m, we find that both ERO classes are similar; identifying the relative fractions of massive galaxies at $z \\sim 1.5$ which are old or dusty and starforming requires either longer wavelength data or deep spectroscopy. From a sample of EROs in a 64 square arcmin area of the {\\it Spitzer} First Look Survey, \\markcite{Yan:04c}Yan {et~al.} (2004b) find that approximately half of the ERO population is detected to a $24 \\mu$m flux limit of $40 \\mu$Jy (3$\\sigma$), suggesting a roughly even distribution of starbursts and quiescent galaxies in the ERO classification. This breakdown is consistent with the courageous, yet incomplete, spectroscopic studies \\markcite{Cimatti:02, Yan:04b, Doherty:05}(\\eg Cimatti {et~al.} 2002a; Yan {et~al.} 2004a; Doherty {et~al.} 2005). Since EROs are the likely progenitors of massive early-type galaxies in the early universe, this suggests both that substantial formation of early-type galaxies occurs at $z \\simgt 2$, but that many massive galaxies are also still actively forming stars at $z \\simlt 2$." }, "0607/astro-ph0607394_arXiv.txt": { "abstract": "We investigate structural properties of dark matter halos of disk galaxies in hierarchical $\\Lambda$CDM cosmology, using a well-defined sample of 81 disk-dominated galaxies from the SDSS redshift survey. We model the mass-velocity (TF) and fundamental plane (FP) relations of these galaxies, which are constructed from the galaxy stellar mass, disk scale length, and optical H$\\alpha$ rotation velocity at 2.2 scale lengths. We calculate a sequence of model galaxy populations, defined by the distribution of the stellar disk-to-total mass fraction, $m_d$. We include the effect of adiabatic contraction of dark matter halos in response to condensation of baryons. We find that models with constant $m_d$ under-predict the intrinsic scatter of the TF and FP relations and predict an (unobserved) strong correlation between TF residuals, even with the full range of halo concentration scatter. Introducing a scatter of disk mass fractions and allowing the mean value $\\bar{m}_d$ to scale with the stellar surface density significantly improves observational match of both the slope and intercept of the model TF relation and reduces the predicted residual correlation enough to be statistically consistent with the data. The distribution of angular momentum parameters $\\lambda_d$ required to match the observed disk scale lengths is significantly narrower than that predicted for halo spin parameters. Our best-fit models with a Kroupa stellar IMF over-produce the galaxy stellar mass function and predict the virial $r$-band mass-to-light ratios, $M_{\\rm vir}/L_r$, systematically lower than those inferred from galaxy-galaxy weak lensing and satellite dynamics. We investigate three possible solutions to these problems: (1) ignoring the effects of adiabatic contraction, (2) adopting a ``light'' stellar IMF with $M_*/L$ lower than the Kroupa IMF by 0.15 dex, or (3) considering the lower halo concentrations predicted for a low cosmological power spectrum normalization $\\sigma_8 \\approx 0.74$. In combination with our proposed correlation of $\\bar{m}_d$ with stellar surface density, any of these solutions yields acceptable residual correlations and relieves most of the observational tension between the TF relation and the galaxy stellar mass function. ", "introduction": "\\label{sec:intro} In the standard theoretical framework of galaxy formation, baryons cool, condense, and form stars in the centers of dark matter halos \\citep{white_rees78}. Observed disk galaxies obey a tight correlation between luminosity and rotation speed, known as the Tully-Fisher (hereafter TF) relation \\citep{tully_fisher77}. The slope, intercept, and scatter of the TF relation are critical constraints on galaxy formation models \\citep[e.g.,][]{cole_kaiser89, kauffmann_etal93, cole_etal94, eisenstein_loeb96, somerville_primack99, steinmetz_navarro99}. These constraints can be characterized clearly within the dissipative collapse modeling framework developed by \\citet{fall_efstathiou80} and \\citet{gunn83}, updated to the cold dark matter (CDM) scenario by \\citet{dalcanton_etal97} and \\citet{mo_etal98}. In this framework, the disk rotation velocity is determined by the ratio, $m_d$, of the disk mass to the total halo mass, by the halo density profile, and by the angular momentum parameter, $\\lambda_d$, which sets the disk scale length. In analogy to ``fundamental plane'' studies of the elliptical galaxy population \\citep{djorgovski_davis87, dressler_etal87}, one can consider disk scale length as an additional parameter in galaxy scaling relations \\citep{shen_etal02}. Since $V^2 \\sim GM(R)/R$, ``maximal disk'' models in which baryons dominate the observed rotation curve predict a strong anti-correlation between TF residual and disk size \\citep{courteau_rix99}. In this paper, we derive empirical constraints on the distributions of $m_d$ and $\\lambda_d$ by modeling the sample of disk-dominated galaxies from \\citet{pizagno_etal05}. This sample is comprised of 81 galaxies with H$\\alpha$ rotation curves selected from the Sloan Digital Sky Survey \\citep[SDSS;][]{york_etal00} main galaxy redshift sample \\citep{strauss_etal02} in the absolute magnitude range $-18 > M_{r} > -23$, for which disk-bulge decomposition yields a best-fit $i$-band bulge fraction $\\leq 0.1$. \\citet{pizagno_etal05} estimate galaxy stellar masses $M_*$ from the observed luminosities and colors, using the population synthesis models of \\citet{bell_etal03} to relate the color to the mean stellar mass-to-light ratio. They characterize rotation velocities by the amplitude of the observed rotation curve at 2.2 disk scale lengths, $V_{2.2}$. We construct model galaxy populations with different $P(m_d)$ distributions and test their ability to reproduce the parameters of the TF relation and its residual correlations, after imposing the observed $M_* - R_d$ distribution as a constraint. Our approach differs from previous studies by using a new, homogeneous data set with well-defined selection criteria and small observational errors, by working with estimated stellar masses instead of luminosities, and by deriving the $P(\\lambda_d)$ distribution empirically from the data instead of imposing it {\\it a priori} from theory. While the distribution of the {\\it halo} angular momentum parameter $P(\\lambda)$ has been well studied with $N$-body simulations \\citep[e.g.,][]{barnes_efstathiou87, bullock_etal01b}, the distribution of disk angular momenta could be different because baryons and dark matter exchange angular momentum, because a biased subset of halo baryons settle into the disk, or because disk-dominated galaxies form in a biased subset of dark matter halos. An empirical determination of $P(\\lambda_d)$ is therefore a valuable diagnostic of galaxy formation physics. For our purposes, we define a model of the disk galaxy population by the probability distribution of $m_d$, which may depend on stellar mass and scale length, $P(m_d | M_*,R_d)$. We assume that the density distribution of dark matter halos is given by the NFW profile \\citep{navarro_etal97} with a concentration parameter $c$. The distribution function $P(c)$ and mass dependence of the halo concentrations have been carefully studied and are no longer a systematic uncertainty of the models, but they influence the scatter and slope of the predicted TF relation. For each observed galaxy, we draw random values of $m_d$ and $c$ consistent with the assumed $P(m_d)$ and theoretically estimated $P(c)$, and calculate the halo mass as $M = M_*/m_d$. Note that our definition of $m_d$ includes only the stellar mass and {\\it does not} include the contribution of cold gas in the disk. We make this choice of necessity because we do not have gas mass measurements; in the range of our sample, gas fractions are typically $\\sim 15\\%$. We determine the disk angular momentum parameter $\\lambda_d$ of each galaxy by matching the observed disk scale length $R_d$. We judge the acceptability of a model by how well it reproduces the observed joint distributions of $M_*, R_d$, and $V_{2.2}$. Since we use the values of stellar masses and disk sizes as input to the models, the bivariate $M_* - R_d$ distribution is reproduced by construction. The slope, amplitude, and scatter of the $M_* - V_{2.2}$ relation and the correlation of residuals of the $M_* - R_d$ and $M_* - V_{2.2}$ relations ($\\Delta R_d$ and $\\Delta V_{2.2}$, respectively) serve as tests. Alternatively, we consider a ``fundamental plane''-type (FP) relation for all three variables, $M_*$, $R_d$, and $V_{2.2}$. After finding the distributions $P(m_d | M_*,R_d)$ that give acceptable fits to the observed relations, we examine the ability of these models to fit the observed baryon mass function of galaxies. We also make predictions for extended mass distributions around the late-type galaxies, which are measured by weak lensing and satellite dynamics. We include in our models the effect of adiabatic contraction (AC) of dark matter halos in response to central condensation of baryons. The deeper potential well created by the cooling and compression of gas attracts more dark matter to the central regions of the galaxy and increases its density. In the analytic literature on disk galaxy modeling, this effect has sometimes been omitted, or its inclusion treated as a ``free parameter'', on the suspicion that the spherical symmetry and adiabatic growth assumptions used in the standard derivation of \\citet{blumenthal_etal86} might make the results invalid in the hierarchical merging paradigm of galaxy formation. Recently, the AC effect has been tested and quantified for galaxies forming by hierarchical merging, using ultrahigh resolution gasdynamics cosmological simulations with cooling and star formation \\citep{gnedin_etal04}. The effect is not as strong as predicted by the original \\citet{blumenthal_etal86} model, mainly because the dark matter particle orbits are elongated rather circular, but the enhancement of dark matter density in the luminous parts of galaxies is robust. The number of galaxies studied at this level of detail remains small, but all available analytical and numerical studies show that enhanced density of dark matter halos by dissipative baryons is a generic result, which is not sensitive to the assumption of slow, smooth growth of the central galaxy component \\citep{jesseit_etal02, sellwood_mcgaugh05, choi_etal06, maccio_etal06, weinberg_etal06}. We are not aware of any hydrodynamic simulations of galaxy formation that do not show this effect. The inclusion of AC amplifies the effects of disk gravity, making it difficult to reproduce the observed {\\it lack} of correlation between disk size and rotation speed \\citep{courteau_rix99}. By boosting disk rotation speeds relative to halo virial velocities, AC also makes it difficult to reconcile the observed Tully-Fisher relation with the observed galaxy luminosity function, given the halo population of typical CDM models; this reconciliation is a long-standing challenge to semi-analytic models of galaxy formation \\citep[e.g.,][]{kauffmann_etal93, cole_etal94} because the conflict emerges from a relatively simple, halo-counting argument. These observational challenges make it tempting to simply omit AC when creating models of the galaxy population \\citep[e.g.,][]{somerville_primack99, dutton_etal05, dutton_etal06}. Given the numerical simulation results cited above, however, we regard omitting AC as a radical departure from a well established element of galaxy formation physics. It is a departure that must be considered if the observations require it, but one should consider comparably radical changes (e.g., to cosmological parameters or the stellar initial mass function) on an equal basis. Throughout this paper, we compare models with AC (computed using the modified model of \\citealt{gnedin_etal04}) and without AC, but we regard the former as ``standard galaxy formation physics'' and the latter as a speculative scenario. ", "conclusions": "We have constrained the structural properties of dark matter halos of disk-dominated galaxies using a well-defined sample of 81 late-type galaxies with H$\\alpha$ rotation curves, selected from the SDSS redshift survey \\citep{pizagno_etal05}. We model the Tully-Fisher and fundamental plane relations constructed from the galaxy stellar mass, $M_*$, disk scale length, $R_d$, and optical H$\\alpha$ rotation velocity at $2.2 \\, R_d$, $V_{2.2}$. The stellar mass is determined from the $i$-band luminosity and a stellar mass-to-light ratio estimated from the galaxy's $g-r$ color. The normalization of these stellar masses depends on the IMF, for which we generally adopt the form of \\citet{kroupa01}. The maximum likelihood fit to the TF relation for our sample is $\\log{V_{2.2}} = 0.29 (\\log{M_*} - 10.5) + 2.23$. The FP relation is $\\log{V_{2.2}} = 0.27 (\\log{M_*} - 10.5) + 0.085 (\\log{R_d} - 0.65) + 2.23$. Residuals from the TF relation are weakly, and positively, correlated with residuals from the $R_d - M_*$ relation. The weakness of this correlation, and the small $R_d$ coefficient in the FP relation, imply that the TF relation is a nearly edge-on view of the fundamental plane of disk galaxies. We define models of the disk galaxy population by the probability distribution function $P(m_d | M_*,R_d)$ of the disk stellar mass fraction $m_d = M_*/M_{\\rm vir}$, where $M_{\\rm vir}$ is the virial mass of the host dark matter halo, assumed to have an NFW profile with the concentration predicted by $N$-body studies. For each sample galaxy, we draw $m_d$ from the model distribution, then find the disk angular momentum parameter $\\lambda_d$ required to reproduce the galaxy's observed disk scale length. Our standard calculations include the impact of adiabatic contraction (AC) on the halo profile \\citep[e.g.,][]{blumenthal_etal86, gnedin_etal04}, but we also consider models without AC. Models with a constant value of $m_d$ require $m_d \\approx 0.08$ with AC ($m_d \\approx 0.04$ without AC) to reproduce the TF intercept and slope. These models under-predict the scatter in the TF relation, even when we include the full range of concentration scatter found in $N$-body simulations of dark matter halos. More importantly, they predict a strong anti-correlation of TF and $R_d - M_*$ residuals, with more compact galaxies rotating faster at fixed $M_*$. Equivalently, they predict an $R_d$ coefficient in the FP relation that has the wrong sign and is incompatible with the observed value. Adding random scatter to $V_{2.2}$ values to match the observed intrinsic scatter does not resolve this conflict with the observed residual correlations. The conflicts are more severe for the models with AC because of the stronger impact of disk gravity. The AC models also predict a TF slope that is slightly too steep. We conclude that a scatter of $m_d$ values, with dispersion in $\\ln{m_d}$ of $\\sigma_m \\approx 0.25$, is required to reproduce the observed scatter of the TF relation, and that models with AC require higher disk mass fractions in higher mass galaxies to reproduce the observed slope. These findings confirm the more qualitative arguments of \\citet{pizagno_etal05} based on the same data set. We can obtain consistency with the weak TF residual correlation and the small $R_d$ FP coefficient by tying the mean disk mass fraction to the stellar surface density. The best-fit model with AC has $\\bar{m}_d \\propto \\Sigma_*^{0.65}$ and a mean fraction $\\bar{m}_d = 0.1$ at $M_* \\, R_d^{-2} = 10^{9.2}\\, \\Msun\\, {\\rm kpc}^{-2}$. Without AC, the best-fit model has a shallower slope, 0.2 instead of 0.65, and a lower normalization, $\\bar{m}_d = 0.04$ at $10^{9.2}\\, \\Msun\\, {\\rm kpc}^{-2}$. The $\\bar{m}_d - \\Sigma_*$ correlation counteracts the effect of stronger disk gravity in more compact galaxies. It produces a major improvement for the AC model, and a modest improvement for the no--AC model. The derived $\\lambda_d$ distributions for these two models are significantly narrower than the $\\lambda$ distribution for dark matter halos in $N$-body simulations: $\\sigma_{\\lambda d} \\approx 0.25$ (AC) or 0.39 (no AC) vs. $\\sigma_\\lambda \\approx 0.56$. The mean $\\bar{\\lambda}_d$ depends on the typical $m_d$ values, which determine halo virial radii: we find $\\bar{\\lambda}_d \\approx 0.054$ (AC) or $\\bar{\\lambda}_d \\approx 0.036$ (no AC) vs. $\\bar{\\lambda} \\approx 0.045$ for $N$-body halos. The systematic difference between the $\\lambda_d$ and $\\lambda$ distributions implies either that disk baryons have systematically different angular momentum from the dark matter in their host halo or that only a subset of halos form disk-dominated galaxies like those in our sample. We consider the $\\bar{m}_d - \\Sigma_*$ correlation a physically plausible mechanism for bringing theoretically motivated disk formation models, which include adiabatic contraction, into agreement with the slope and residual correlation of the observed TF relation. However, our best-fit AC model suffers the long-standing problem of over-producing the galaxy stellar mass function given the halo population of a $\\Lambda$CDM cosmological model with observationally favored parameter values and the concentration-mass relation of \\citet{bullock_etal01}. It also predicts virial mass to $r$-band light ratios lower than some recent estimates from galaxy-galaxy lensing and satellite dynamics. Lowering $\\Omega_{\\rm m}$, and thus the halo mass function, can ameliorate the abundance problem, but the required value, $\\Omega_{\\rm m} \\approx 0.15$, seems implausibly low and the mass-to-light ratio problem remains in any case. Eliminating AC can resolve both problems, at the expense of dropping a well-tested ingredient of galaxy formation theory. We suggest an alternative solution: adopting a ``light'' stellar IMF with $M_*/L$ lower than the \\citet{kroupa01} IMF by about 0.15 dex. This model yields acceptable residual correlations with a modest trend $\\bar{m}_d \\propto \\Sigma_*^{0.4}$, predicts virial mass-to-light ratios in accord with recent measurements, and reduces but does not entirely remove the tension with the galaxy baryon mass function. We do not know of other observational evidence to support such a light IMF, but it is perhaps within the uncertainties of existing observational constraints. Systematic uncertainties in the \\citet{bell_etal03} population synthesis models, which we use to compute $M_*/L$ from $g-r$ color, could account for some of the required change even without altering the IMF itself. We also show that the lower $\\bar{c}(M)$ relation predicted for $\\sigma_8 = 0.74$ by Bullock et al.'s (2001b) analytic model relieves most of the tension between the TF relation and the mass function and $M_{\\rm vir}/L_r$ constraints, even in the case of a Kroupa IMF. A full assessment of this solution awaits better numerical confirmation of the analytic model and better observational constraints on $\\sigma_8$. Other changes to halo profiles that reduce the fraction of mass at small radii would have similar effects. Our observational findings are consistent with those of \\citet{dutton_etal06}, which is reassuring given the different characteristics of the data sets. Our theoretical inferences are consistent but not identical. We are more reluctant than \\citet{dutton_etal06} to abandon adiabatic contraction, given the numerous simulations in which it occurs and the absence of any simulations in which it does not. The $\\bar{m}_d - \\Sigma_*$ correlation can remove the $\\Delta R_d - \\Delta V_{2.2}$ anti-correlation that otherwise plagues models with AC, though it leaves problems with the baryon mass function and weak lensing mass-to-light ratios. The addition of a light IMF or the adoption of lower halo concentrations largely resolve these difficulties as well, and we consider some combination of these effects to be a plausible resolution of the observational tensions. Modeling of rotation curve data at larger and smaller separations and galaxy-galaxy lensing measurements for more finely divided galaxy sub-classes may provide further insight that could help discriminate among these solutions. The fact that disk formation models based on the most straightforward theoretical and observational inputs do not reproduce all aspects of the current data suggests that interesting surprises may still lie ahead." }, "0607/astro-ph0607441_arXiv.txt": { "abstract": "The double-peaked broad emission lines are usually thought to be linked to accretion disks, however, the local viscous heating in the line-emitting disk portion is usually insufficient for the observed double-peaked broad-line luminosity in most sources. {It was suggested that the X-ray radiation from an ion-supported torus in the inner region of the disk can photo-ionize the outer line-emitting disk region. However, our calculations show that only a small fraction ($\\la$ 2.3 per cent) of the radiation from the radiatively inefficient accretion flow (RIAF) in the inner region of the disk can photo-ionize the line-emitting disk portion, because the solid angle of the outer disk portion subtended to the inner region of the RIAF is too small.} We propose a physical model for double-peaked line emitters, in which only those AGNs with sufficient matter above the disk (slowly moving jets or outflows) can scatter enough photons radiated from the inner disk region to the outer line-emitting disk portion extending from several hundred to more than two thousand gravitational radii. Our model predicts a power-law $r$-dependent line emissivity $\\epsilon^{{\\rm H}\\alpha}\\propto R_{\\rm d}^{-\\beta}$, where $\\beta\\sim 2.5$, which is consistent with $\\beta\\sim 2-3$ required by the model fittings for double-peaked line profiles. Using a sample of radio-loud AGNs with double-peaked emission lines, we show that the outer disk regions can be efficiently illuminated by the photons scattered from slow or mild relativistic electron-positron jets with $\\gamma_{\\rm j}\\la 2$. It is consistent with the fact that no double-peaked emission line is present in strong radio quasars with relativistic jets. For radio-quiet double-peaked line emitters, slow outflows with Thomson scattering depth $\\sim 0.2$ instead of jets can scatter sufficient photons to (illuminate) the line-emitting regions. This model can therefore solve the energy budget problem for double-peaked line emitters. ", "introduction": "Only a small fraction of active glactic nuclei (AGNs) exhibit double-peaked broad-line profiles, for example, the largest sample available so far is the 116 double-peaked Balmer line AGNs, which are found from an initial sample consisting of 5511 broad line AGNs with $z<0.5$ observed by the Sloan Digital Sky Survey(SDSS) \\citep{s03}. Some previous authors mainly focused on the double-peaked broad lines in radio-loud(RL) AGNs \\citep*[e.g.,][]{p88,chf89,ch89,eh94,eh03}. A complete survey on RL AGNs finally resulted in 20 double-peaked line emitters \\citep{eh94,eh03}. It is still a mystery why (only) a small fraction of AGNs exhibit double-peaked broad-line profiles. Some different scenarios were suggested for the origin of double-peaked emission lines, namely, (1) emission from the accretion disk \\citep*[e.g.,][]{chf89,ch89}, (2) emission from a binary broad-line region (BLR) in a binary massive black hole system \\citep*[e.g.,][]{g83,g88}, (3) emission from the bipolar outflows \\citep*[e.g.,][]{zsb90}, (4) emission from a spherically symmetric BLR illuminated by an anisotropic ionizing radiation source \\citep*[e.g.,][]{gw96}. The accretion disk model is the most favorable one among them \\citep*[see detailed comparison between these different scenarios in][and references therein]{eh03}. The widths of double-peaked lines range from several thousand to nearly 40000 km~s$^{-1}$ \\citep*[e.g.][]{w05}. In the accretion disk model, the double-peaked emission lines are radiated from the disk region between around several hundred gravitational radii to more than two thousand gravitational radii, and their profiles can be well fitted by the accretion disks with a power-law line emissivity \\citep*[e.g.,][]{chf89,ch89,eh03}. However, none of these scenarios can answer why only a small fraction of AGNs have been detected as double-peaked line emitters. For example, in the accretion disk model, one may expect to observe double-peaked lines in most AGNs, as the AGNs with higher inclination angles have more chances to be observed, until their accretion disks are obscured by the putative tori \\citep*[e.g.,][]{a93}, if the orientations of AGNs are isotropically distributed in space. Another difficulty for the disk model is the ``energy budget\" problem, i.e., the local viscously dissipated power in the line-emitting disk portion is usually insufficient for the observed double-peaked broad-line luminosity in most sources, and the temperatures of the line-emitting disk regions are too low to produce the observed H$_\\alpha$ lines \\citep{chf89,ch89,eh94,eh03}. The X-rays from a hot ion-supported torus in the inner disk region are assumed to irradiate the outer line-emitting disk region and then to solve the ``energy budget\" problem \\citep*[e.g.,][]{ch89}. {We discuss the illumination of the outer disk regions by the inner ion-supported tori in \\S 2.} In this paper, we suggest a physical model for these double-peaked line emitters. In this model, a fraction of photons from the accretion disk are scattered by the electrons in the jet/outflow back to photo-ionize the outer line-emitting region, and then to produce the observed emission lines. We describe our model in \\S 3. The cosmological parameters $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0=70~ {\\rm km~s^{-1}~Mpc^{-1}}$ have been adopted in this work. ", "conclusions": "In the ion-supported torus scenario, the torus is truncated to a thin disk at a radius $R_{\\rm d,tr}$, and the X-ray photons from the torus are thought to illuminate the outer line-emitting disk region. The most favorable candidate for such an ion-supported hot torus is the RIAF \\citep{ny94,ny95}. However, some double-peaked emitters have very luminous X-ray emission, which cannot be reproduced solely by the RIAFs in these sources if a reasonable viscosity parameter is adopted (see discussion in \\S 2). It is quite doubtful whether RIAFs are in these X-ray luminous double-peaked line emitters. In Fig. 1, we find that most energy is radiated in the inner edge of the RIAF (more than 70 per cent of the total radiation of the RIAF is emitted from the region $R_{\\rm d}\\la 0.3 R_{\\rm d,tr}$). Our calculation shows that less than 2.3 per cent radiation from the RIAF in the inner region of the disk can illuminate the outer disk region with $R_{\\rm d}\\ge R_{\\rm d,tr}$ (see \\S 2), which implies that the radiation of RIAFs is unable to solve the energy budget problem for most sources even if all the observed luminous X-ray emission in some double-peaked line emitters can be attributed to RIAFs. Comparison between jet power estimated from the extended radio emission and kinetic luminosity of the jet requires $\\gamma_{\\rm e,min}\\sim 1$ for electron-positron jets, and $\\gamma_{\\rm e,min}\\sim 100$ for electron-proton jets \\citep*[e.g.,][]{cf93}. For electron-proton jets, the inverse-Compton scattered photons by the electrons in the jet are in hard X-ray/$\\gamma$-ray bands, which are inefficient for photo-ionizing the line-emitting disk region. In our present model, the line-emitting disk region is assumed to be irradiated by the scattered X-ray photons from the electron-positron jet. Most photons Compton up-scattered by the electrons in the electron-positron jets are in the soft X-ray band, which can efficiently photo-ionize the outer line-emitting disk region. Our calculations show that the jet can be Compton thin down to $r_{\\rm j,tr}$ if the jet kinetic luminosity $L_{\\rm kin}/L_{\\rm Edd}$ is low, while it can be Compton thick even at very large radii for high values of $L_{\\rm kin}/L_{\\rm Edd}$, because much matter is loaded by the jet with high $L_{\\rm kin}/L_{\\rm Edd}$. In our calculations, the minimal radius of the jet $r_{\\rm j,min}=10$ is adopted. In principle, we can calculate the cases for a jet down to a radius lower than $r_{\\rm j,min}=10$, though it is still unclear whether the K\\\"onigl's conical jet is still valid to a very small radius near the black hole. Fortunately, the photons scattered near the bottom of the jet cannot photo-ionize the outer line-emitting disk region efficiently, because of the small solid angle subtended to the line-emitting region, which is similar to the RIAF case as discussed in \\S 2. Thus, the final results have not been affected much even if a value of $r_{\\rm j,min}<10$ is adopted. In Fig. \\ref{fig3}, we find that the emissivities $\\epsilon^{{\\rm H}\\alpha}$ have similar power-law $r$-dependence for different Lorentz factors $\\gamma_{\\rm j}$ adopted. We find $\\epsilon^{{\\rm H}\\alpha}$ varies with $r$ nearly $\\propto r^{-2.5}$ along the disk radius for all cases. The model fittings on the observed double-peaked line profile requires a power-law line emissivity with an index of $\\sim 2-3$ for different sources \\citep*[e.g.,][]{ch89,chf89,eh94,eh03}, which is consistent with our model calculations. It is found that the emissivity decreases rapidly with increasing the Lorentz factor $\\gamma_{\\rm j}$. For given jet kinetic luminosity, a higher $\\gamma_{\\rm j}$ leads to lower electron density (see Eqs. \\ref{lkin1} and \\ref{mdotjet}), which decreases the inverse-Compton scattered photons from the jet, and then the emissivity $\\epsilon^{{\\rm H}\\alpha}$. The ratios $L_{{\\rm H}\\alpha}/L_{\\rm disk}$ varying with $L_{\\rm kin}/L_{\\rm Edd}$ for different values of $\\gamma_{\\rm j}$ are plotted in Fig. \\ref{fig4}. We plot the data of the sample in Fig. \\ref{fig4} to test our model by comparison with the model calculations. We find that the Lorentz factors of the jets $\\gamma_{\\rm j}\\la 2$ are required to explain the observed data for most sources in this sample. The apparent angular velocity of the brightest jet component in 3C390.3 measured by VLBI is $\\sim$0.54 mas/year \\citep{k04}. The Lorentz factor of the jet in this source $\\gamma_{\\rm j}\\simeq 2.1$ is therefore derived from its inclination angle $i=26^\\circ$, which is estimated from the fitting of its double-peaked line profile by \\citet{eh94}. In Fig. \\ref{fig4}, one can find that a jet with $\\gamma_{\\rm j}\\sim 1.5-2$ is required by the model calculation for this source. For this source, $W_{\\rm d}\\simeq 4.6 L_{\\rm H\\alpha}$ \\citep{eh03}, which means that the Blamer line emission may be partly attributed to the local viscous heating of the line-emitting disk region, so our result is roughly consistent with that derived from the VLBI observations. Only one source, Pictor A, a jet with relatively larger Lorentz factor $\\gamma_{\\rm j}>2$ is required to illuminate the line-emitting region. This is a powerful radio galaxy \\citep*[e.g.,][]{w01}, and X-ray jet emission is detected in this source \\citep{h05}. The X-ray emission from the jet is important for this source, and the accretion disk luminosity derived from its X-ray luminosity is obviously over-estimated, so the location of this source in Fig. \\ref{fig4} should be shifted upwards. A mild relativistic jet may probably provide sufficient scattered photons for photo-ionizing the line-emitting disk region in this source, if the X-ray emission is dominantly from the jet. {For Arp 102B, a very low $\\gamma_{\\rm j}\\sim 1.01$ ($\\sim 0.14c$) is required (see Fig. \\ref{fig4}), i.e., a slowly moving conical outflow is present in this source, which is consistent with its compact radio structure \\citep{b81,p86,c01}. } Another special source, IRAS 0236.6-3101, which is difficult to be explained by this model (see Fig. \\ref{fig4}). This is a star-forming galaxy with very weak radio emission (9.5 mJy at 1.4 GHz), and it probably has no jet. The origin of its double-peaked lines may be similar to that of radio-quiet double-peaked line emitters, which will be addressed in the last paragraph of this section. For the seven sources with only upper limits on black hole masses (lower limits on the ratios $L_{\\rm kin}/L_{\\rm Edd}$), they are similar to others and the jets with moderate Lorentz factors $\\gamma_{\\rm j}\\la 2$ can solve the ``energy budget\" problem. As the detailed model fittings for individual sources indicate that the line-emitting regions in these sources have different sizes \\citep*[see][for details]{eh03}, our calculations have been carried out for different values of $\\xi_1$ and $\\xi_2$(see different line types in Fig. \\ref{fig4}), which have not altered our main conclusion. There are 13 sources in this sample with $W_{\\rm d}/L_{{\\rm H}\\alpha}>1$. Although it means the local viscous power in the line-emitting disk regions may be sufficient for the line emission for these 13 sources, the temperatures of the line-emitting disk portions are still too low for all the observed H$_\\alpha$ line emission. The external illuminating from the jets is still necessary for these sources. In our model, an electron-positron jet is required to scatter disk photons back to photo-ionize the outer disk surface to produce Balmer lines. We suggest that future VLBI polarization observations on these radio-loud double-peaked line emitters can test this model. In this paper, some quantities, $M_{\\rm bh}$, $L_{\\rm bol}$, and $L_{\\rm kin}$, are derived from observed quantities by using the conventional approaches reported in literature, of which the uncertainties might affect our results. However, the qulititative conclusion that slow jets/outflows are required in double-peaked emitters would not be altered (see Fig. 4), for typical errors for these quantities, i.e., a factor of $\\sim 3$ for $M_{\\rm bh}$, and the derived kinetic luminosity $L_{\\rm kin}$ is accurate at an order of magnitude. {In our present calculation of $L_{\\rm H\\alpha}$, we have not included the thermalization od the disk. If the thermalization is considered, the resulted $L_{\\rm H\\alpha}$ will become lower, so a lower bulk Lorentz factor $\\gamma_{\\rm j}$ of the jet than our present model calculation is required (see Fig. \\ref{fig4}). } Besides radio-loud AGNs in this sample, several ten radio-quiet double-peaked line emitters have been discovered \\citep{s03}. Although no jet is present in these radio-quiet AGNs, we speculate that slow moving outflows may be in these sources, which play the same role of the jets in their radio-loud counterparts. If the electrons in the outflows are non-relativistic, a fraction of disk power radiated from its inner region can be Thomson scattered to the outer line-emitting portion of the disk. If the Thomson scattering depth is $\\sim 0.2$, the outflows can scatter nearly one-tenth of the disk radiation to illuminate the outer line-emitting disk portion. If this is the case, the absorption by the outflows in these radio-quiet AGNs is expected to be detected by X-ray observations, which can be a test on our model. The lack of detected double-peaked broad emission lines in strong radio-loud quasars with relativistic jets is a natural prediction of our model, i.e., the Compton scattering power is too weak to irradiate the outer line-emitting portion for fast moving jets. For radio-quiet AGNs, only those having outflows with suitable scattering depth can scatter disk photons back to illuminate the outer line-emitting portion, which leads to double-peaked broad emission lines. Recent X-ray observations showed that a fraction of quasars and Seyfert galaxies have X-ray absorption with column density $\\ga 10^{23} {\\rm cm}^{-3}$ \\citep*[e.g.,][]{row03,y05,p05}, which implies that the AGNs having outflows with suitable scattering depth $\\sim 0.2$ are not common among all AGNs. {However, it does not mean that all AGNs with suitable scattering depth will exhibit double-peaked lines. Some other factors, such as, the viewing angle and the properties of the accretion disk, may play important roles in producing double-peaked lines.} The quantitative model calculations to fit the line profiles for individual AGNs are beyond the scope of this work, which will be reported in the forthcoming paper. {In this work, we suggest that the soft photons scattered by electrons in electron-position jets for radio-loud double-peaked line emitters (or in slow outflows with suitable scattering depth for radio-quiet counterparts) can efficiently ionize the line-emitting disk regions. Another possibility is that slow outflows play the same role in radio-loud double-peaked line emitters as radio-quiet sources, though it is still unclear whether relativitic jets can co-exist with such cold outflows in radio-loud double-peaked line emitters. However, it is unable to explain the fact that no double-peaked emission line is present in strong radio quasars with relativistic jets, while this is a natural prediction of the electron-positron jet model proposed in this work.}" }, "0607/astro-ph0607488_arXiv.txt": { "abstract": "High-resolution mid-infrared images have been obtained in N-band and Q-band for the proto-planetary nebula IRAS 16594$-$4656. A bright equatorial torus and a pair of bipolar lobes can clearly be seen in the infrared images. The torus appears thinner at the center than at the edges, suggesting that it is viewed nearly edge-on. The infrared lobes correspond to the brightest lobes of the reflection nebula seen in the Hubble Space Telescope ({\\it HST}) optical image, but with no sign of the point-symmetric structure seen in the visible image. The lobe structure shows a close correspondence with a molecular hydrogen map obtained with {\\it HST}, suggesting that the dust emission in the lobes traces the distribution of the shocked gas. The shape of the bipolar lobes shows clearly that the fast outflow is still confined by the remnant circumstellar envelope of the progenitor asymptotic giant branch (AGB) star. However, the non-detection of the dust outside of the lobes suggests that the temperature of the dust in the AGB envelope is too low for it to be detected at 20 $\\mu$m. ", "introduction": "Proto-planetary nebulae (PPNe) are the long-sought-after missing link between the end of the asymptotic giant branch (AGB) phase and the beginning of planetary nebula phase of stellar evolution. After the {\\it Infrared Astronomical Satellite} ({\\it IRAS}) mission, a number of objects were proposed as candidate PPNe based on their infrared colors and other spectral properties. These are typically stars of G to B spectral type with significant infrared excesses due to the remnant circumstellar dust shell ejected in the AGB phase. Of particular interest among these candidates are a number of carbon-rich objects whose abundances show a strong enhancement of s-process elements, as expected from the dredge-up of material in thermal pulses during the AGB evolution \\citep{kwok93, vanwinckel03}. For other candidates there is some possibility of confusion with massive supergiants, but those in this carbon-rich group are almost certainly bona-fide PPNe. One of these objects is IRAS 16594$-$4656. It is a bright mid-infrared source which has typical colors of a PPN \\citep{volk89}. Optically it is associated with a southern emission-line star. It was found to be of spectral type B7 with V magnitude 14.6, subject to about 7.5 magnitudes of visual extinction \\citep{steene00}. It was not clear from the original {\\it IRAS} spectral observations whether the object was oxygen-rich or carbon-rich, but subsequent {\\it Infrared Space Observatory} ({\\it ISO}) spectral observations showed that it has carbon-based dust features, including the 21 $\\mu$m feature \\citep{garcia99}. Optical images obtained with the {\\it Hubble Space Telescope} showed that the star has a surrounding reflection nebula with a complex structure \\citep{hrivnak99}. The relative faintness of the reflection nebula compared to the star, together with the optical morphology, led \\citet{hrivnak99} to conclude that the nebula is intrinsically bipolar (or multi-polar) viewed at an intermediate angle to the bipolar axis. Both optical and near-infrared spectral observations show emission lines, which are thought to be shock-excited rather than radiatively excited by the star \\citep{garcia99,steene03}. No radio emission has been detected from IRAS 16594$-$4656 \\citep{steene93}. The distance to this object is uncertain. The dust shell model of \\citet{hrivnak00} suggests a distance of 2.6 kpc if the total luminosity of the star is 10$^4$ $L_\\odot$. A slightly smaller estimate of (2.2$\\pm$0.4) kpc is given by \\citet{steene03} for the same assumed total luminosity. Most of the luminosity is being emitted in the infrared where extinction effects are smaller than in the optical, so these estimates are not strongly affected by the non-spherical morphology of the dust shell. The optical morphology of the nebula in IRAS 16594$-$4656 appears complex, with what appear to be pairs of symmetric structures at multiple position angles, for which it was named the ``Water Lily Nebula'' \\citep{hrivnak99}. The optical images also showed several concentric arcs centered on the star \\citep{hrivnak01}. However, near-infrared observations with the {\\it NICMOS} instrument on {\\it HST} showed a somewhat simpler morphology, although it was not clear whether this was simply an effect of reduced dynamic range compared to the optical observations \\citep{su03}. An initial N-band observation of IRAS 16594$-$4656 with the TIMMI2 camera on the ESO 3.6m telescope failed to resolve the dust shell \\citep{steene00}, while more recent TIMMI2 observations did marginally resolve the structure \\citep{garcia04}. In this paper we present new, higher sensitivity and higher angular resolution mid-infrared images of IRAS 16594$-$4656 which resolve the dust shell in thermal emission. We present the observations in section 2. We then derive a dust color temperature map in section 3, and compare the morphology as observed in the mid-infrared to that seen at other wavelengths in section 4. We give a brief discussion of the results in the final section. ", "conclusions": "Figure \\ref{fig3} shows the deconvolved Si5 image with overlaid contour plots in the optical I-band (0.8 $\\mu$m; Su, Hrivnak, \\& Kwok 2001) and near-infrared H$_2$ filter images (2.12 $\\mu$m; Hrivnak, Kelly, \\& Su 2004). This allows direct comparison of the morphology of the nebula in these different wavelengths. The H$_2$ image was matched to the Si5 filter image since they were immediately seen to have very similar morphologies. This was used to obtain a reasonably accurate estimate of the central star position in the T-ReCS image. That position was then used to match the optical image to the mid-infrared image. The Si5 image is plotted in absolute surface brightness units, Jy/square arc-second. As shown in the upper right panel of Figure \\ref{fig3}, there is a close match of features in the lobes between the {\\it NICMOS} H$_2$ filter image and the T-ReCS image. The bright lobe edges are regions of strong H$_2$ emission. The cap in the east lobe is clearly visible in this panel, and it also corresponds to a region of strong H$_2$ emission. It is more difficult to determine if any of the mid-infrared bright waist structure is also detected in the H$_2$ image, because the star is saturated in the H$_2$ image, but it does not look as if anything but the edges of the lobes are detected in the H$_2$ image. Since the H$_2$ emission is mainly shock excited \\citep{steene03}, this raises the possibility that the dust emission from these regions are partially shock-excited. Another possibility that could explain why the dust emission region so closely matches the shock is that the dust grains may be much smaller downstream from the shock than they are before the shock, and that as a result the small grains are transiently heated to relatively high temperatures. Comparison of the T-ReCS images with the {\\it HST} optical images \\citep{hrivnak99} shows that the two lobes seen in the T-ReCS images correspond closely to the central brightest part of the reflection nebula. The total size of the optical reflection nebula is 12.3$^{\\prime\\prime}$ by 8.8$^{\\prime\\prime}$, which is much larger than the size of the N-band or Q-band images. The optical image overlay in the lower left panel of Figure \\ref{fig3} uses logarithmically spaced contours ranging from 0.0085\\% of the stellar peak brightness up to the peak brightness, with each contour at a level 2.7 times the previous contour. The optical reflection nebula is at most about 0.75\\% of the stellar peak brightness. The lowest contours show what has been suggested to be point-symmetric morphology, with three pairs of oppositely directed features at position angles of about 40$^\\circ$, 57$^\\circ$, and 87$^\\circ$ east of north as can be measured from the I-band image presented in \\citet{hrivnak99}. The optical lobes have been suggested to be caused by a rapidly precessing, columnated high-speed wind from the star \\citep{garcia99}. The T-ReCS image shows just the two lobes at position angle 75$^\\circ$. While there seems to be some very faint mid-infrared emission detected in the Si5 filter on a size scale roughly three times that of the main mid-infrared lobes, over-plotting this with the optical image does not show any correspondences with the faint extended optical structure. In particular, none of this emission is seen just outside the bright torus either to the north or to the south. Kinematic observations of the individual optical lobes is needed to determine if they are distinct structures or not. If these point symmetric features are not distinct kinematic structures, then perhaps there are just two lobes but the optical appearance is due to structure in the walls of these lobes which make them look more complex. The optical emission is due to reflection by dust, and so the larger size of the optical images compared to the mid-infrared images shows that the dust shell is much larger than is obvious from the N-band or Q-band images. The dust outside the central bright region delineated by the shocked H$_2$ emission is clearly much colder than that inside the shocks. At the ends of mid-infrared torus, in particular, there must be a large discontinuity in temperature and optical depth so that there is a large change in mid-infrared surface brightness but a much smaller change in the optical brightness of the scattered light. These mid-infrared images suggest that the torus is perpendicular to the plane of the sky. This is consistent with recent observations of the bipolar lobes, which indicate that they are nearly in the plane of the sky. Unpublished high resolution long-slit spectra in the near-infrared \\citep{hrivnak06} have been obtained with the Phoenix spectrograph on Gemini South. Analysis of these spectra, which map the molecular hydrogen line at 2.12 $\\mu$m for cuts at different position angles through the star and the lobes, shows similar velocities for the two lobes, indicating that they are oriented very close to the plane of the sky. This was also concluded by \\citet{ueta05} based upon near-infrared polarimetry along with dust shell modeling of the spectral energy distribution of the object. Since the evidence indicates that the nebular axis is oriented very close to the plane of the sky, the east lobe must actually be smaller than the west lobe. The T-ReCS images show that the central bright region of the circumstellar dust shell is quite different than that of other well-known bipolar nebulae IRAS 17150$-$3224, IRAS 17441$-$2411, Roberts 22, and Hen 3-401, in all of which the dust emission is strongly peaked at the stellar position. For IRAS 16594$-$4656 the mid-infrared brightness is at a minimum near the stellar position and is much higher to the north and the south along the central waist. Possibly the torus is of much lower optical depth in this object than in the others, or it is highly asymmetric with a low optical depth along our line of sight. The latter suggestion is consistent with the visibility of the central star. Certainly for IRAS 17150$-$3224 and IRAS 17441$-$2411 the spectral energy distribution implies a relatively high optical depth along our line of sight to the star, and the star is not seen in visible light. The spectral type of the star in IRAS 16594$-$4656 is much earlier than that in the latter two objects, so it may simply be more evolved and the torus may have had more time to disperse. Unlike the hour-glass or open lobes observed in most bipolar planetary nebulae (e.g. NGC 6302), the bipolar lobes of IRAS 16594$-$4656 as shown in Figure \\ref{fig2} are closed and resemble the lobes seen in the PPN IRAS 17106$-$3046 and the young planetary nebula Hen 2-320. The morphology of the lobes clearly shows that the lobes are confined by the circumstellar medium, and the fast collimated outflow which creates the bipolar lobes has not yet broken through the stellar wind of the AGB progenitor. The interaction between the fast and slow winds is clearly delineated by dust distribution in the lobes. We also note that there is \"bulge\" at the tip of the western lobe, which suggests that the high-velocity flow is on the verge of breaking out. In contrast to the cap at the tip of the eastern lobe (which represents a pile-up at the wind interface), the western lobe may represent a slightly more advanced stage of the breakout, and therefore explains the difference in sizes between the two lobes. In a few hundred years, we expect that both lobes will open up into butterfly morphology. We are therefore witnessing a critical phase of morphological transformation of PNs. We have successfully detected the bipolar lobes and a central bright waist in mid-infrared images of IRAS 16594$-$4656. The bright waist suggests that we are seeing a central torus nearly edge-on. While this is consistent with published polarization and unpublished kinematic results, it differs from earlier published interpretations of the visible image that concluded that the lobes are seen at an intermediate orientation. This emphasizes the need for multi-wavelength observations to confidently understand the structure of proto-planetary nebulae. This result may well be applicable to our understanding of other bipolar phenomenon such as YSOs and AGNs." }, "0607/hep-ph0607032_arXiv.txt": { "abstract": "The early stages of the universe evolution are discussed according to the hot big bang model and the grand unified theories. The shortcomings of big bang are summarized and their resolution by inflationary cosmology is sketched. Cosmological inflation, the subsequent oscillation and decay of the inflaton field, and the resulting reheating of the universe are studied in some detail. The density perturbations produced by inflation and the temperature fluctuations of the cosmic microwave background radiation are introduced. The hybrid inflationary model is described. Two natural variants of the supersymmetric version of this model which avoid the disaster encountered in its standard realization from the overproduction of magnetic monopoles are presented. ", "introduction": "\\label{sec:introduction} \\par The discovery of the cosmic microwave background radiation (CMBR) together with the observed Hubble expansion of the universe had established hot big bang as a viable model of the universe (for a textbook treatment of this model, see e.g. Ref.~\\cite{wkt}). The success of nucleosynthesis (see e.g. Ref.~\\cite{bbn}) in reproducing the observed abundance of light elements in the universe and the proof of the black body character of the CMBR then imposed hot big bang as the standard cosmological model. This model combined with the grand unified theories (GUTs) \\cite{ggps} of strong, weak, and electromagnetic interactions provides the scientific framework for discussing the early stages of the universe evolution. \\par The standard big bang (SBB) cosmological model, despite its great successes, had some long-standing shortcomings. One of them is the so-called horizon problem. The CMBR received now has been emitted from regions of the universe which, according to this model, had never communicated before sending this radiation to us. The question then arises how come the temperature of the black body radiation from these regions is so finely tuned as the measurements of the cosmic background explorer (COBE) \\cite{cobe} and the Wilkinson microwave anisotropy probe (WMAP) \\cite{wmap1,wmap3} satellites show. Another important issue is the so-called flatness problem. It is a fact \\cite{wmap3} that the present universe appears to be almost flat. This means that, in its early stages, the universe must have been flat with a great accuracy, which requires an extreme fine tuning of its initial conditions. Also, combined with GUTs which predict the existence of superheavy magnetic monopoles \\cite{monopole}, the SBB model leads \\cite{preskill} to a catastrophe caused by the overproduction of these monopoles. Finally, the model has no explanation for the small density perturbations which are required for explaining the structure formation in the universe (for a pedagogical discussion, see e.g. Ref.~\\cite{structure}) and the generation of the observed \\cite{cobe,wmap1,wmap3} temperature fluctuations in the CMBR. \\par It is clear that cosmological inflation \\cite{guth} offers an elegant solution to all these problems of the SBB model (for a textbook introduction or previous reviews on inflation, see e.g. Refs.~\\cite{book,inflation}). The idea is that, in the early universe, a real scalar field (the inflaton) was displaced from its vacuum value. If the potential energy density of this field happens to be quite flat, the roll-over of the field towards the vacuum can be very slow for a period of time. During this period, the energy density is dominated by the almost constant potential energy density of the inflaton. As a consequence, the universe undergoes a period of quasi-exponential expansion, which can readily solve the horizon and flatness problems by stretching the distance over which causal contact is established and reducing any pre-existing curvature in the universe. It can also adequately dilute the GUT magnetic monopoles. Moreover, it provides us with the primordial density perturbations which are needed for explaining the large scale structure in the universe \\cite{structure} as well as the temperature fluctuations observed in the CMBR. Inflation can be easily incorporated in GUTs. It can occur during the GUT phase transition at which the GUT gauge symmetry breaks by the vacuum expectation value (VEV) of a Higgs field, which also plays the role of the inflaton. \\par After the termination of inflation, the inflaton field starts performing oscillations about the vacuum. These oscillations are damped because of the dilution of the field energy density by the cosmological expansion and the decay of the inflaton into light particles. The resulting radiation energy density eventually dominates over the field energy density and the universe returns to a normal big bang type evolution. The temperature at which this occurs is historically called `reheat' temperature although there is neither supercooling nor reheating of the universe \\cite{reheat} (see also Ref.~\\cite{dilution}). \\par The early realizations of inflation share the following important disadvantage. They require tiny parameters in order to reproduce the COBE or WMAP measurements on the CMBR. In order to solve this naturalness problem, hybrid inflation has been introduced \\cite{hybrid}. The basic idea was to use two real scalar fields instead of one that was customarily used. One field may be a gauge non-singlet and provides the `vacuum' energy density which drives inflation, while the other is the slowly varying field during inflation. This splitting of roles between two fields allows us to reproduce the temperature fluctuations of the CMBR with natural (not too small) values of the relevant parameters. Hybrid inflation, although it was initially introduced in the context of non-supersymmetric GUTs, can be naturally incorporated \\cite{lyth,dss} in supersymmetric (SUSY) GUTs. \\par It is unfortunate that the magnetic monopole problem reappears in hybrid inflation. The termination of inflation, in this case, is abrupt and is followed by a `waterfall' regime during which the system falls towards the vacuum manifold and performs damped oscillations about it. If the vacuum manifold happens to be homotopically non-trivial, topological defects will be copiously formed \\cite{smooth} by the Kibble mechanism \\cite{kibble} since the system can end up at any point of this manifold with equal probability. Therefore, a cosmological disaster is encountered in the hybrid inflationary models which are based on a gauge symmetry breaking predicting magnetic monopoles. \\par One way \\cite{smooth,jean,shi,talks} to solve the magnetic monopole problem of SUSY hybrid inflation is to include into the standard superpotential for hybrid inflation the leading non-renormalizable term. This term cannot be excluded by any symmetries and, if its dimensionless coefficient is of order unity, can be comparable with the trilinear coupling of the standard superpotential (whose coefficient is typically $\\sim 10^{-3}$). Actually, we have two options. We can either keep \\cite{jean} both these terms or remove \\cite{smooth,shi} the trilinear term by imposing a discrete symmetry and keep only the leading non-renormalizable term. The pictures emerging in the two cases are quite different. However, they share an important common feature. The GUT gauge group is spontaneously broken already during inflation and, thus, no topological defects can form at the end of inflation. So, the magnetic monopole problem is solved. ", "conclusions": "\\label{sec:conclusions} \\par We summarized the shortcomings of the SBB cosmological model and their resolution by inflation, which suggests that the universe underwent a period of exponential expansion. This may have happened during the GUT phase transition at which the relevant Higgs field was displaced from the vacuum. This field (inflaton) could then, for some time, roll slowly towards the vacuum providing an almost constant vacuum energy density. Inflation generates the density perturbations needed for the large scale structure of the universe and the temperature fluctuations of the CMBR. After the end of inflation, the inflaton performs damped oscillations about the vacuum, decays, and reheats the universe. \\par The early inflationary models required tiny parameters. This problem was solved by hybrid inflation which uses two real scalar fields. One of them provides the vacuum energy density for inflation while the other one is the slowly rolling field. Hybrid inflation arises naturally in many SUSY GUTs, but leads to a disastrous overproduction of magnetic monopoles. We constructed two extensions of SUSY hybrid inflation which do not suffer from this problem." }, "0607/astro-ph0607455_arXiv.txt": { "abstract": "We have performed $N$-body simulations of the formation of hyper-velocity stars (HVS) in the centre of the Milky Way due to inspiralling intermediate-mass black holes (IMBHs). We considered IMBHs of different masses, all starting from circular orbits at an initial distance of 0.1\\,pc. We find that the IMBHs sink to the centre of the Galaxy due to dynamical friction, where they deplete the central cusp of stars. Some of these stars become HVS and are ejected with velocities sufficiently high to escape the Galaxy. Since the HVS carry with them information about their origin, in particular in the moment of ejection, the velocity distribution and the direction in which they escape the Galaxy, detecting a population of HVS will provide insight in the ejection processes and could therefore provide indirect evidence for the existence of IMBHs. Our simulations show that HVS are generated in short bursts which last only a few Myrs until the IMBH is swallowed by the supermassive black hole (SMBH). HVS are ejected almost isotropically, which makes IMBH induced ejections hard to distinguish from ejections due to encounters of stellar binaries with a SMBH. After the HVS have reached the galactic halo, their escape velocities correlate with the distance from the Galactic centre in the sense that the fastest HVS can be found furthest away from the centre. The velocity distribution of HVS generated by inspiralling IMBHs is also nearly independent of the mass of the IMBH and can be quite distinct from one generated by binary encounters. Finally, our simulations show that the presence of an IMBH in the Galactic centre changes the stellar density distribution inside $r<0.02$\\,pc into a core profile, which takes at least 100 Myrs to replenish. ", "introduction": "\\label{sec:intro} Hills (1988) was the first to show that the ejection of stars with velocities $>1000\\kms$ is a natural consequence of galaxies hosting supermassive black holes. He named these stars 'hyper-velocity stars' (HVS). Recently, several HVS have been discovered in the Galactic halo \\citep{b05, hi05, b06a, b06b}. Except for one star which might have been ejected from the LMC \\citep{edel05}, the travel times of all HVS are short enough that the stars could have been ejected from the Galactic centre within the lifetimes of the stars, confirming Hills' predictions. The exact formation mechanism of HVS is however still a matter of debate. The ejection of stars by supernova explosions in close binary systems \\citep{b61} and dynamical encounters \\citep{p67} cannot produce main-sequence stars with velocities exceeding a few hundred kilometers per second \\citep{gps05}, leaving the interaction of stars in galactic nuclei around super-massive black holes (SMBHs) as the only possible source for HVS. Yu \\& Tremaine (2003) considered three processes which could eject stars from the vicinity of SMBHs: (1) close encounters between two single stars, (2) encounters between stellar binaries and the central SMBH and (3) encounters between single stars and a massive black hole binary. In the case of the SMBH in the Galactic centre, they found that close encounters between single stars can eject stars with a rate of only $10^{-11}$\\,yr$^{-1}$ which would create less than one HVS during the lifetime of the Milky Way. The other two processes were found to eject stars with rates of up to $10^{-4}$\\,yr$^{-1}$, sufficiently high to explain the observed number of HVS in the halo of the Milky Way. Similar results were also obtained by \\citet{gps05}, who studied the ejection of stars by SMBHs by means of scattering experiments and found that HVS formation is possible by both processes. The tidal disruption of binaries was found to create HVS with higher velocities while the ejection from binary black hole systems creates HVS with a higher rate. Similarly, \\citet{gl06} found that the tidal breakup of stellar binaries can create HVS with velocities significantly higher than what has been found so far. A distinction between the two scenarios might come from a detailed analysis of the spatial and kinematical distribution of HVS: the ejection of stars due to the interaction of stellar binaries with an SMBH should be nearly constant with time since the reservoir of binary stars in the Galactic centre is depleted only slowly. Furthermore, the distribution of binary orbits should be nearly isotropic in sufficiently relaxed nuclei, implying that the resulting HVS distribution will also be isotropic. In contrast, the ejection of stars from an SMBH-IMBH binary should show characteristic variations with time and spatial direction: HVS are mainly ejected when the inspiralling IMBH reaches the centre since the density of stars is highest in the centre and the velocity dispersion is also highest close to the SMBH. In addition, escaping stars acquire their extra velocities mainly in the direction of motion of the IMBH, introducing a spatial anisotropy in the HVS distribution. These considerations were confirmed by \\citet{l05}, who studied analytically the distribution of escapers created by an SMBH-IMBH pair in a dense stellar cusp. He found that the ejection of stars in case of a black hole binary is happening mainly in a burst which lasts a few dynamical friction timescales and that if the IMBH is initially in a nearly circular orbit, the velocity vectors of the ejected stars also cluster around the orbital plane. If the IMBH moves in an eccentric orbit, stars are ejected in a broad jet roughly perpendicular to the Runge-Lenz vector of the IMBHs orbit. Similarly, \\citet{shm06} found through three-body scattering experiments that eccentric black hole binaries eject stars along a broad jet perpendicular to the semimajor axis of the binary. IMBHs could form in galaxies through runaway collisions of stars in star clusters, giving rise to ultra-luminous X-ray sources \\citep{pbhmm04}. They could later be brought into the centres of galaxies through dynamical friction of the star clusters \\citep{pbmmhe06}. Their merger with the central SMBHs would be an important source of gravitational waves, detectable with the next generation of gravitational wave detectors, like e.g. {\\it LISA}. IMBHs might also be an important contribution for the growth of SMBHs in the early universe \\citep{ebi01}. In the present paper we have therefore performed collisional $N$-body simulations of the dynamics of inspiralling IMBHs in stellar cusps around supermassive black holes. The aim of our simulations is to study whether the ejection of stars by IMBHs leads to observable consequences in the distribution of HVS which might help to distinguish between different ejection scenarios and which could give an indirect hint for the presence of one or more IMBHs in the centre of the Milky Way. ", "conclusions": "\\label{sec:concl} We have performed simulations of the inspiral of massive black holes into the centres of galaxies and of the subsequent ejection of hyper-velocity stars. We found that the spatial distribution of HVS is nearly isotropic and would be difficult to distinguish from a HVS distribution created by interactions of stellar binaries with an SMBH if only few HVS were found, as is presently the case. A better indication comes from the escape times of HVS: our simulations confirm that most HVS are ejected in a short burst, lasting only a few Myrs for typical IMBH masses, as soon as the IMBH reaches the galactic centre. The ejection ends when the IMBH merges with the central SMBH, which should take less than 10\\,Myrs. Even if merging can be avoided, the ejection rate of HVS is a factor of 30 to 100 lower than during the burst maximum. The currently observed HVS show a broad distribution of escape times \\citep{b05}, which argues against ejection due to a single IMBH, but would still be consistent with the inspiral of several IMBHs. The evidence is not conclusive yet and requires more HVS to be found, which should be possible with future astrometric surveys like e.g. {\\it GAIA}. In case HVS are ejected by an IMBH, we also expect a strong correlation of escape velocity with galactocentric distance in the sense that the fastest HVS can be found at the largest distances. \\begin{figure} \\begin{center} \\includegraphics[width=8.3cm]{escrate_vgl.eps} \\end{center} \\caption{Rate of stars ejected from the central cusp as a function of time for an IMBH mass of 3000 $M_\\odot$ but three different stellar masses. As expected from theoretical arguments, the maximum ejection rate drops linearly with the number of stars while the escape rate at later times is nearly independent of the average stellar mass.} \\label{fig:eject_scal} \\end{figure} Another prediction from our runs is that IMBHs deplete the central region in stars so that an initial cusp profile is turned into a nearly constant density core with core radius $r=0.02$\\,pc. If an IMBH was present in the Galactic centre within the last 100 Myrs, such a core should still be visible in the stellar density distribution. If on the other hand the cusp profile observed at larger radii continues all the way down to the centre, this would be evidence against the presence of an IMBH in the Galactic centre within the last $T \\approx 100$ Myrs." }, "0607/astro-ph0607663_arXiv.txt": { "abstract": "Synchrotron is considered the dominant emission mechanism in the production of gamma-ray burst photons in the prompt as well as in the afterglow phase. Polarization is a characteristic feature of synchrotron and its study can give a wealth of information on the properties of the magnetic field and of the energy distribution in gamma-ray burst jets. In this paper I will review the theory and observations of gamma-ray bursts polarization. While the theory is well established, observations have prove difficult to perform, due to the weakness of the signal. The discriminating power of polarization observations, however, cannot be overestimated. ", "introduction": "The ``standard'' model to interpret GRB emission implies dissipation of bulk kinetic energy via collisionless shocks~\\cite{Rees92,Meszaros93,Piran99}. Magnetic fields are generated in the process and highly relativistic electrons are accelerated in a power-law distribution of energies (or Lorentz factors)~\\cite{Silva03,Hededal04}. In these conditions, radiation is emitted through the synchrotron process, giving rise to a broad band radiation source~\\cite{Meszaros97,Piran99}. Synchrotron radiation from a coherent magnetic field is known to be polarized. As a consequence, theoretical and observational efforts were produced to understand if and how much polarization should be observed in the various phases of the GRB process. In this paper we divide the GRB in two distinctive phases, the prompt and the afterglow. For each phase we consider both theoretical predictions for linear polarization (circular polarization is supposed to be very small, especially in the afterglow phase~\\cite{Matsumiya03}) and observations. In both cases we will emphasize how the understanding of the polarization properties of GRBs can shed light on other important issues such as the release mechanism of GRB jets, the micro-physics of collisionless shocks, the energy distribution of jets and the properties of dust along the line of sight to the GRB. ", "conclusions": "" }, "0607/astro-ph0607349_arXiv.txt": { "abstract": "We investigate the relative importance of starbursts and AGN in nuclear activities of ultra-luminous infrared galaxies (ULIRGs) based on chemodynamical simulations combined with spectrophotometric synthesis codes. We numerically investigate both the gas accretion rates (${\\dot{m}}_{\\rm acc}$) onto super massive black holes (SMBHs) and the star formation rates (${\\dot{m}}_{\\rm sf}$) in ULIRGs formed by gas-rich galaxy mergers and thereby discuss what powers ULIRGs. Our principal results, which can be tested against observations, are as follows. (1) ULIRGs powered by AGN can be formed by major merging between luminous, gas-rich disk galaxies with prominent bulges containing SMBHs, owing to the efficient gas fuelling (${\\dot{m}}_{\\rm acc} > 1 {\\rm M}_{\\odot}$ yr$^{-1}$) of the SMBHs. AGN in these ULIRGs can be surrounded by compact poststarburst stellar populations (e.g., A-type stars). (2) ULIRGs powered by starbursts with ${\\dot{m}}_{\\rm sf} \\sim 100 {\\rm M}_{\\odot}$ yr$^{-1}$ can be formed by merging between gas-rich disk galaxies with small bulges having the bulge-to-disk-ratio ($f_{\\rm b}$) as small as 0.1. (3) The relative importance of starbursts and AGN can depend on physical properties of merger progenitor disks, such as $f_{\\rm b}$, gas mass fraction, and total masses. For example, more massive galaxy mergers are more likely to become AGN-dominated ULIRGs. (4) For most models, major mergers can become ULIRGs, powered either by starbursts or by AGN, only when the two bulges finally merge. Interacting disk galaxies can become ULIRGs with well separated two cores ($>$ 20kpc) at their pericenter when they are very massive and have small bulges. (5) Irrespective of the choice of model, interacting/merging galaxies show the highest accretion rates onto the central SMBHs, and the resultant rapid growth of the SMBHs occur when their star formation rates are very high. Based on these results, we discuss an evolutionary link between ULIRGs, QSOs with poststarburst populations, and ``E+A'' galaxies. We also discuss spectroscopic properties (e.g., H$\\beta$ luminosities and line ratio of \\oiii/H$\\beta$) in galaxy mergers with starbursts and AGN. ", "introduction": "The observation that ultra-luminous infrared galaxies (ULIRGs, defined as those with infrared luminosities greater than $10^{12}$ $ L_{\\odot}$) show a mixture of two distinct types of nuclear activities, namely starbursts and active galactic nuclei (AGN), has led to many observational studies of their formation and evolution processes (e.g., Sanders et al. 1988; Solomon et al. 1992; Soifer et al. 1986; Clements et al. 1996; Murphy et al. 1996; Sanders \\& Mirabel 1996; Gao \\& Solomon 1999; Trentham et al. 1999; Veilleux et al. 1999; Scoville et al. 2000; Surace et al. 2000; Bushouse et al. 2002; Tacconi et al. 2002; Farrah et al. 2003; Armus et al. 2004; Imanishi \\& Terashima 2004; Colina et al. 2005; Iwasawa et al. 2005). For example, Sanders et al. (1988) proposed that ULIRGs formed by gas-rich galaxy mergers can finally evolve into QSOs after the removal of dust surrounding QSO black holes. Spectroscopic properties of ULIRGs have been extensively discussed in terms of the relative importance of starbursts and active galactic nuclei (AGN) in the energy budget of ultra-luminous infrared galaxies (e.g., Genzel et al. 1998; Lutz et al. 1998). These observations have so far raised many questions, the most significant being: (1) whether all ULIRGs evolve into QSOs, (2) what mechanisms are responsible for triggering starbursts and AGN obscured heavily by dust in ULIRGs, (3) what determines the relative importance of starbursts and AGN in spectral energy distributions (SEDs) of ULIRGs, (4) whether there is an evolutionary link between starbursts and AGN in ULIRGs, and (5) whether there can be physical relationships between low redshift (low-$z$) ULIRGs and high-$z$ dust-enshrouded starbursts and AGN at intermediate and high redshifts recently revealed by SCUBA (Submillimeter Common-User Bolometer Array) (e.g., Barger et al. 1998; Smail et al. 1997, 1998, 1999; Blain et al. 1999). Morphological studies of ULIRGs revealed that they show strongly disturbed morphologies indicative of violent galaxy interaction and merging. Previous theoretical studies have tried to answer the above five questions in the context of gas fuelling to the central region of galaxy mergers (See Shlosman et al. 1990 for more general discussions on fuelling mechanism in galaxies). Physical mechanisms responsible for the formation of starbursts in galaxy mergers have been investigated by many authors (e.g., Olson \\& Kwan 1990; Barnes \\& Hernquist 1991; Noguchi 1991; Mihos \\& Hernquist 1994, 1996; Gerritsen \\& Icke 1997). For example, Olson \\& Kwan (1990) suggested that high velocity disruptive cloud-cloud collisions, which are more prominently enhanced in mergers, are responsible for the observed high star formation rates in galaxy mergers. Although these previous numerical studies provided some theoretical predictions on star formation rates (SFRs) and their dependence on the initial physical parameters of galaxy merging (e.g., bulge-to-disk-ratio and gas mass fraction), they did not investigate both SFRs and accretion rates (ARs) onto the central super-massive black holes (SMBHs) simultaneously. Therefore, they did not provide useful theoretical predictions on the formation and evolution of AGN, or on a possible evolutionary link between starbursts and AGN in ULIRGs. Physical processes of gas fuelling to the central SMBHs in galaxy mergers have been investigated by a number of authors (Bekki \\& Noguchi 1994; Bekki 1995; Di Matteo et al. 2005; Springel et al. 2005a, b). Using dynamical simulations with rather idealized modeling of gas dynamics and star formation, Bekki \\& Noguchi (1994) first investigated both SFRs and ARs in merging galaxies and found that SFRs become very high at the epoch of the coalescence of the cores of two merging galaxies, whereas ARs attain their maxima only after the coalescence. Recently, Springel et al (2005a) have performed more sophisticated, high-resolution SPH simulations including feedback effects of AGN on the interstellar medium (ISM), and thereby demonstrated that AGN feedback can be quite important for global photometric properties of elliptical galaxies formed by major galaxy merging. These previous models however did not discuss the latest observational results of ULIRGs, partly because their model do not allow authors to investigate photometric and spectroscopic properties of dusty starbursts and AGNs in galaxy mergers. The purpose of this paper is thus to investigate simultaneously both SFRs and ARs of merging galaxies in an self-consistent manner and thereby try to address the aforementioned questions related to the origin of ULIRGs. We particularly try to understand (1) physical conditions required for galaxy mergers to evolve into ULIRG with AGN (or starbursts), (2) key factors which determine the relative importance of starbursts and AGN, and (3) epochs when mergers become ULIRGs with AGN. We develop a new model in which the physics of star formation (including gas consumption and supernovae feedback by star formation), the time evolution of accretion disks around SMBHs, and the growth of SMBHs via gas accretion from the accretion disks are included. By using this new model, we show (1) how SFRs and ARS in merging galaxies evolve with time, (2) how they depend on galactic masses, mass ratios of two merging spirals, and bulge-to-disk-ratios of the merger progenitor spirals, and (3) how SMBHs grow in the central regions of starbursting mergers. We also show emission line properties of galaxies with starbursts and AGNs by combining the results of the simulated SFRs and ARs with spectral evolution codes. Although previous numerical simulations combined with spectrophotometric synthesis codes have already derived SEDs of {\\it purely starburst} galaxies obscured by dust (Bekki et al. 1999; Bekki \\& Shioya 2000, 20001; Jonsson et al. 2005), they did not discuss at all the spectrophotometric properties of galaxies {\\it where starbursts and AGN coexist}. Therefore our new way of deriving spectral properties based on simulation results enables us to answer some key questions raised by recent large, systematic survey of AGN (e.g., Kauffmann et al. 2003), such as why a significant fraction of high-luminosity AGN have the Balmer absorption lines. Previous one-zone spectroscopic models discussed what controls emission and absorption line properties of galaxies with starbursts and AGN (Baldwin, Phillips \\& Terlevich 1981; Veilleux \\& Osterbrock 1987; Kewley et al. 2001; Dopita et al. 2006). The present simulations allow us to discuss this point based on the results of SFRs and ARs derived by chemodynamical simulations with growth of SMBHs. The plan of the paper is as follows: In the next section, we describe our numerical model for calculating SFRs and ARs in merging galaxies. In \\S 3, we present the numerical results on the time evolution of SFRs and ARs and its dependences of model parameters. In this section, we also show emission line properties of galaxies mergers with starbursts and AGN. We discuss the present results in terms of formation and evolution of ULIRGs and QSOs in \\S 4. We summarise our conclusions in \\S 5. \\begin{table*} \\centering \\begin{minipage}{185mm} \\caption{Model parameters} \\begin{tabular}{ccccccccl} {Model no. } & $M_{d}$ ($\\times$ $10^{10} {\\rm M_{\\odot}} $) & {$f_{\\rm g}$% \\footnote{initial gas mass fraction}} & {$f_{\\rm b}$% \\footnote{mass ratio of bulge to disk}} & {$m_{\\rm 2}$% \\footnote{mass ratio of merging two disks}} & {orbital type} & {$m_{\\rm sf,max}$% \\footnote{maximum star formation rate (${\\rm M}_{\\odot}$ yr$^{-1}$)}} & {$m_{\\rm acc,max}$% \\footnote{maximum accretion rate (${\\rm M}_{\\odot}$ yr$^{-1}$)}} & Comments \\\\ M1 & 6.0 & 0.2 & 0.5 & 1.0 & FI & $2.6 \\times 10^0$ & $2.5\\times 10^{0}$ & standard \\\\ M2 & 6.0 & 0.2 & 0.5 & 1.0 & FI & $6.4 \\times 10^2$& $0\\times 10^{0}$ & no accretion onto SMBHs \\\\ M3 & 0.15 & 0.2 & 0.5 & 1.0 & FI & $1.3 \\times 10^{-1}$ & $4.0 \\times 10^{-4}$ & \\\\ M4 & 3.0 & 0.2 & 0.5 & 1.0 & FI & $5.0\\times 10^0$ & $7.0 \\times 10^{-1}$ & \\\\ M5 & 30.0 & 0.2 & 0.5 & 1.0 & FI & $3.0 \\times 10^2$ & $9.2\\times 10^1$ & \\\\ M6 & 0.15 & 0.2 & 0.1 & 1.0 & FI & $2.1 \\times 10^{-1}$ & $6.9 \\times 10^{-5}$ & \\\\ M7 & 3.0 & 0.2 & 0.1 & 1.0 & FI & $5.1\\times 10^0$ & $8.2 \\times 10^{-3}$ & \\\\ M8 & 6.0 & 0.2 & 0.1 & 1.0 & FI & $9.5\\times 10^0$ & $3.0 \\times 10^{-2}$ & \\\\ M9 & 30.0 & 0.2 & 0.1 & 1.0 & FI & $9.8 \\times 10^1$ & $9.5 \\times 10^{-1}$ & \\\\ M10 & 6.0 & 0.2 & 0.5 & 1.0 & HI & $2.2 \\times 10^1$ & $2.3\\times 10^0$ & \\\\ M11 & 6.0 & 0.2 & 0.5 & 1.0 & RR & $2.2 \\times 10^1$ & $3.6\\times 10^0$ & \\\\ M12 & 6.0 & 0.2 & 0.5 & 1.0 & BO & $3.3 \\times 10^1$ & $4.1\\times 10^{0}$ & \\\\ M13 & 6.0 & 0.2 & 0.5 & 0.1 & BO & $6.0 \\times 10^0$ & $1.0\\times 10^{-2}$ & LSB minor merger \\\\ M14 & 6.0 & 0.2 & 0.5 & 0.3 & BO & $6.8 \\times 10^0$ & $1.0\\times 10^{-1}$ & unequal-mass merger \\\\ M15 & 6.0 & 0.2 & 0.5 & 0.1 & BO & $8.5 \\times 10^0$ & $2.0\\times 10^{-1}$ & HSB minor merger \\\\ M16 & 6.0 & 0.02 & 0.5 & 1.0 & FI & $6.8 \\times 10^{-1}$ & $4.9 \\times 10^{-3}$ & gas poor \\\\ M17 & 6.0 & 0.05 & 0.5 & 1.0 & FI & $1.8 \\times 10^0$ & $1.0 \\times 10^{-1}$ & \\\\ M18 & 6.0 & 0.1 & 0.5 & 1.0 & FI & $8.3 \\times 10^0$ & $1.6 \\times 10^0$ & \\\\ M19 & 6.0 & 0.2 & 0.02 & 1.0 & FI & $1.1 \\times 10^2$ & $2.9 \\times 10^{-3}$ & smaller bulge \\\\ M20 & 6.0 & 0.2 & 1.0 & 1.0 & FI & $6.0 \\times 10^1$ & $1.6 \\times 10^0$ & bigger bulge\\\\ M21 & 6.0 & 0.05 & 1.0 & 1.0 & FI & $3.1 \\times 10^0$ & $6.1 \\times 10^{-1}$ & bigger bulge, gas poor \\\\ M22 & 30.0 & 0.2 & 0.1 & 1.0 & TI & $5.7 \\times 10^1$ & $2.4 \\times 10^0$ & tidal interaction \\\\ \\end{tabular} \\end{minipage} \\end{table*} \\begin{figure*} \\psfig{file=f1.eps,width=18.0cm} \\caption{ Mass distributions projected onto the $x$-$y$ plane for the standard model. For convenience, stellar particles (old stars) and gaseous ones are shown in magenta (i.e.\\ dark matter halo particles are not shown). Big green dots represent the locations of SMBH1 and 2. Time ($T$), SFRs (${\\dot{m}}_{\\rm sf}$ in units of ${\\rm M}_{\\odot} {\\rm yr}^{-1}$), ARs (${\\dot{m}}_{\\rm ac}$ in units of ${\\rm M}_{\\odot} {\\rm yr}^{-1}$) and the simulation scale are shown at upper left, upper right, lower right, and lower left, respectively, for each frame. Here time $T$ represents the time that has elapsed since the simulation starts. Note that ARs can become very high ($> 1{\\rm M}_{\\odot} {\\rm yr}^{-1}$) when the two bulges finally merge (i.e.\\ when the two SMBHs become very close with each other). } \\label{Figure. 1} \\end{figure*} ", "conclusions": "We have numerically investigated both SFRs and ARs in forming ULIRGs via gas-rich galaxy merging in an self-consistent way. Dependences of the time evolution of SFRs and ARs on model parameters are mainly investigated. We summarize our principle results as follows: (1) ULIRGs powered by AGN can be formed by major merging between luminous, gas-rich disk galaxies with prominent bulges containing SMBHs owing to the efficient gas fuelling (${\\dot{m}}_{\\rm acc} > 1 {\\rm M}_{\\odot}$ yr$^{-1}$) to the SMBHs. AGN in these ULIRGs can be surrounded by compact poststarburst stellar populations (e.g., A-type stars). These results suggest that ULIRGs and QSOs can show strong Balmer absorption lines. (2) ULIRGs powered by starbursts with ${\\dot{m}}_{\\rm sf} \\sim 100 {\\rm M}_{\\odot}$ yr$^{-1}$ can be formed by merging between gas-rich disk galaxies with small bulges having the bulge-to-disk-ratio ($f_{\\rm b}$) as small as 0.1. As long as the accretion radii ($R_{\\rm B}$) of SMBHs are proportional to the masses of the SMBHs, galaxy mergers with smaller bulges are more likely to become starburst-dominated ULIRGs (i.e., they can not show AGN activity owing to a smaller amount of gas accretion onto the SMBHs). (3) The relative importance of starbursts and AGN can depend on physical properties of merger progenitor disks, such as $f_{\\rm b}$, gas mass fraction, and total masses. For example, more massive galaxy mergers are more likely to become AGN-dominated ULIRGs. (4) For most models, major mergers can become ULIRGs powered either by starbursts or by AGN, when the two bulges finally merge. Interacting disk galaxies can become ULIRGs with well separated two cores ($>$ 20kpc) at their pericenter only when they are very massive and have small bulges. These suggest that it is highly unlikely for interacting/merging pair of galaxies to become ULIRGs with double/multiple nuclei. We note, however, the results of Veilleux et al (2002), who found that about 7\\% of ULIRGs in their sample have nuclear separations in excess of 20kpc. This may suggest that ULIRGs can be formed via alternate routes to the major mergers examined herein. (5) Irrespectively of models, interacting/merging galaxies show the highest accretion rates onto the central SMBHs and the resultant rapid growth of the SMBHs, when their star formation rates are very high. (6) ARs can become high ($1 {\\rm M}_{\\odot} {\\rm yr}^{-1}$) enough to show QSO-like activities ($L_{\\rm bol} \\approx 10^{12} {\\rm L}_{\\odot}$) mostly in major mergers between massive disk galaxies with remarkable bulges. ARs however can not reach the required rates for QSOs (${\\dot{m}}_{\\rm acc} \\approx 0.7{\\rm M}_{\\odot} {\\rm yr}^{-1}$) in minor and unequal-mass mergers that form S0s. These results therefore imply that only forming elliptical via major mergers can show QSO-like activities whereas forming S0s (or early-type spirals with big bulges) via minor and unequal-mass merging show low luminosity AGN (e.g., type 1/2 Seyfert). (7) Maximum ARs (${\\dot{m}}_{\\rm acc,max}$) can correlate with maximum SFRs (${\\dot{m}}_{\\rm sf,max}$) in the sense that galaxy mergers with higher ${\\dot{m}}_{\\rm sf,max}$ are likely to show higher ${\\dot{m}}_{\\rm acc,max}$. This suggests that mergers and ULIRGs with more pronounced AGN activities are likely to show stronger starburst components in their nuclei. The correlations can be discussed in the context of recent observational results (e.g., Goto 2005) on correlations between infrared luminosities of ULIRGs, star formation rates, and AGN luminosities (measured from [OIII] emission lines). (8) The ratio of ${\\dot{m}}_{\\rm acc,max}$ to ${\\dot{m}}_{\\rm acc,sf}$ can correlate with ${\\dot{m}}_{\\rm acc,max}$ in the sense that galaxy mergers with higher ${\\dot{m}}_{\\rm acc,max}$ are likely to show higher ${\\dot{m}}_{\\rm acc,max}/{\\dot{m}}_{\\rm sf,max}$. This implies that merger and ULIRGs with higher AGN (thus total) luminosities are likely to be dominated by AGN rather than by starbursts. This result can be also consistent with recent results on AGN fraction as a function of infrared luminosities of galaxies (e.g., Goto 2005). (9) There could be evolutionary links between ULIRGs, Q+A's, QSOs, and E+A's. Galaxy mergers between less massive disks are more likely to evolve from starburst-dominated ULIRGs into E+As without experiencing QSO phases, whereas those between more massive disks with prominent bulges can evolve from AGN-dominated ULIRGs, to QSOs (and/or Q+A's), and finally to E+A's, if the lifetimes of QSOs are as short as $\\sim 0.1$ Gyr. Removal of gas reservoir for star formation via supernovae and AGN feedback could be essentially important for the above evolutionary links. (10) Time evolution of emission line properties of galaxies with starbursts and AGNs is investigated based on SFRs and ARs derived from chemodynamical simulations. For example, simulated mergers are demonstrated to evolve from those with smaller \\oiii/H$\\beta$ (starburst-dominated) to those with larger \\oiii/H$\\beta$ (AGB-dominated). It is suggested that strong Balmer absorption lines are more likely to be detected in type 2 Seyfert than in type 1 owing to the less amount of dilution of stellar light by emission from hidden broad line regions of type 2 Seyfert. Direct comparison between the predicted spectrophotometric properties of galaxy mergers with dusty starbursts and AGNs and the corresponding observations will be done in our forthcoming papers." }, "0607/astro-ph0607513_arXiv.txt": { "abstract": "We present 426 epochs of optical monitoring data spanning 1000 days from December 2003 to June 2006 for the gravitationally lensed quasar SDSS J1004+4112. The time delay between the A and B images is $\\Delta t_{\\mathrm{BA}}=38.4\\pm2.0$ days ($\\Delta\\chi^2=4$) in the expected sense that B leads A and the overall time ordering is C-B-A-D-E. The measured delay invalidates all published models. The models failed because they neglected the perturbations from cluster member galaxies. Models including the galaxies can fit the data well, but strong conclusions about the cluster mass distribution should await the measurement of the longer, and less substructure sensitive, delays of the C and D images. For these images, a delay of $\\Delta t_{\\mathrm{CB}} \\simeq 681 \\pm 15$~days is plausible but requires confirmation, while delays of $\\Delta t_{\\mathrm{CB}} > 560$~days and $\\Delta t_{\\mathrm{AD}} > 800$~days are required. We clearly detect microlensing of the A/B images, with the delay-corrected flux ratios changing from $m_B-m_A=0.44\\pm0.01$ mag in the first season to $0.29\\pm0.01$ mag in the second season and $0.32\\pm0.01$ mag in the third season. ", "introduction": "The wide-separation lensed quasar SDSS J1004+4112 was discovered in the Sloan Digital Sky Survey search for lenses (Inada et al. 2003; Oguri et al. 2004; Sharon et al. 2005; Wambsganss 2003). The lens consists of at least four images of a redshift $z_s=1.734$ quasar whose $\\sim 15\\farcs0$ Einstein ring diameter is created by a redshift $z_l=0.68$ cluster. The cluster has been characterized with X-ray observations (Ota et al. 2006) and there are additional multiply imaged arcs formed from still higher redshift background galaxies (Sharon et al. 2005). There is also strong evidence for a fifth, lensed image of the quasar located near the center of the brightest cluster galaxy (Inada et al. 2005), which in combination with a future velocity dispersion measurement for the galaxy will strongly constrain the central mass distribution of the lens (e.g. Sand et al.~2002, but see Dalal \\& Keeton 2003). Thus, it is not only feasible to cleanly compare X-ray and lensing mass distributions in this galaxy cluster, but it may also be possible to test the cosmological model by measuring the increase of the Einstein radius with source redshift due to the $D_{LS}/D_{OS}$ distance ratio scaling of the lens deflection (Soucail et al. 2004). That the source is a time-variable quasar offers further and unique opportunities for this cluster lens. First, the time delay between the quasar images can be measured as a constraint on the mass distribution. In theory, the time delays determine the mean surface density near the images for which the delay is measured (Kochanek 2002), so the mass sheet ($\\kappa$) degeneracy of most cluster lensing measurements can be broken under the assumption that the Hubble constant is well-determined by other means. Several theoretical studies of the time delays in SDSS J1004+4112 (Oguri et al. 2004; Williams \\& Saha 2004; Kawano \\& Oguri 2006) have explored their dependence on the mean mass profile of the cluster, finding a broad range of potential delays. As we shall see, all these models are incorrect in their details because they neglected cluster member galaxies whose deflection scales are larger than the positional constraints on the quasar images used in the models (see the discussion in Keeton et al. 2000 on the failure of similar models for the cluster lens Q0957+561 and the general discussion in Kochanek 2005). Nonetheless, all these models indicate that the delay between the A and B images is relatively short (weeks) and that its value should indicate the magnitude of the much longer (years) delays of the C and D images. The second unique property of the lens is that microlensing of the quasar accretion disk by any stars in the cluster halo or small satellites near the images can be used as an added probe of the structure of the cluster (see Wambsganss 2006). Because the cluster has a higher velocity dispersion ($700$~km/s) than a typical galaxy lens ($\\sim 200$~km/s), the microlensing time scales in this system may also be shorter than for a lens by about a factor of 3. There is already evidence for microlensing from the time variability of the \\ion{C}{4} 1549\\AA\\ line in image A that is not observed in image B (Richards et al. 2004, Lamer et al. 2006, G\\'omez-\\'Alvarez et al. 2006), although recently Green (2006) has suggested that this could also be due to time variable absorption in the source quasar. For three years we have conducted an optical monitoring campaign to measure the optical variability of this system. This has proved more challenging than desired because the quasars are somewhat faint for monitoring with available telescopes and modest exposure times. However, we have succeeded both in measuring the A/B time delay and clearly detecting microlensing of the optical continuum of the quasar. In \\S\\ref{obsis} we present the data from the monitoring campaign for the four bright lensed quasar images. In \\S\\ref{abdelay} we determine the A/B time delay, discuss the presence of microlensing in the system, and place constraints on the long delays between the close image pair A and B and the fainter images C and D. In \\S\\ref{models} we discuss the failure of existing models for the system and introduce a simple successful model that includes the perturbations of cluster galaxies, and we conclude in \\S\\ref{discuss}. ", "conclusions": "\\label{discuss} We have measured the A/B time delay of SDSS J1004+4112 to be $(38.4\\pm 2.0)$~days ($\\Delta\\chi^2=4$), which fixes the overall time ordering of the images to be C-B-A-D-E. While this is the time ordering predicted in published models (Oguri et al. 2004, Williams \\& Saha 2004, Kawano \\& Oguri 2006) it is significantly longer than the delays predicted by these models. The cause of the discrepancy is that the previously published models overly simplified the mass distribution by neglecting the deflections generated by the cluster member galaxies. Models including the eleven most important galaxies can simultaneously fit the A--E image positions and the measured A/B time delay with reasonable parameter values. Modelers of this system need to remember the lesson of Q0957+561 -- model constraints that are applied more tightly than the deflection scale of the most massive, neglected components of the lens lead to incorrect results (Keeton et al. 2000). We note that Sharon et al. (2005) also needed to include some of the member galaxies in order to model the higher redshift lensed arcs, but made no predictions for the time delay. Fortunately, the A/B delay should be the most sensitive of the delays to the effects of cluster galaxies because it is a merging image pair. The longer delays for the C and D images relative to A and B should be less affected by substructure, so their measurement should provide constraints on the cluster halo properties that are less sensitive to the member galaxies. At present we cannot claim a measurement of these longer delays. A lower bound on the delay $\\Delta t_{\\mathrm{AD}} > 800$~days is consistent with our models, which predict delays of 5--7~years for this image pair. The shorter C/B delay is at least $\\Delta t_{\\mathrm{CB}} > 560$~days but there is a possible delay of $\\Delta t_{\\mathrm{CB}} \\simeq 680\\pm15$~days that should be confirmed or rejected during the next observing seasons and is consistent with our models. We have also clearly detected microlensing variability in the A/B images, with changes of order $0.15$~mag in the A/B flux ratio over the course of the three observing seasons. This result provides strong evidence that the differential changes in the A/B emission line profiles are also due to microlensing (Richards et al. 2004, Lamer et al. 2006, G\\'omez-\\'Alvarez et al. 2006) rather than variable absorption in the source (Green 2006). The microlensing time scales in SDSS J1004+4112 should be relatively shorter than in most single galaxy lenses because the internal velocities of the cluster are about 3 times higher than those of a galaxy. While the flux ratio changes in the optical continuum are modest, we would expect to find significantly larger effects at shorter wavelengths where the source size should be more compact. There is already some evidence for this from the X-ray flux ratios measured by Ota et al. (2006) and Lamer et al. (2006). A campaign to monitor this system in X-rays would both allow us to study the size of the X-ray emission region and provide the added data on the emission from the cluster needed to provide a precision comparison of the mass distributions estimated using X-ray data and lens models. Such careful tests will be essential if measurements of the increase of the Einstein radius of the cluster with source redshift based on the surrounding multiply imaged arcs are to be used as a new test of the cosmological model as proposed by Soucail et al. (2004) and Sharon et al. (2005)." }, "0607/astro-ph0607039_arXiv.txt": { "abstract": "We conduct a detailed comparison of broad-band spectral energy distributions of six $z\\ga 5.5$ galaxies against galaxies drawn from cosmological hydrodynamic simulations. We employ a new tool called \\spoc, which constrains the physical properties of observed galaxies through a Bayesian likelihood comparison with model galaxies. We first show that \\spoc~self-consistently recovers the physical properties of a test sample of high-redshift galaxies drawn from our simulations, although dust extinction can yield systematic uncertainties at the $\\approx50\\%$ level. We then use \\spoc~to test whether our simulations can reproduce the observed photometry of six $z>5.5$ galaxies drawn from the literature. We compare physical properties derived from simulated star formation histories (SFHs) versus assuming simple models such as constant, exponentially-decaying, and constantly rising. For five objects, our simulated galaxies match the observations at least as well as simple SFH models, with similar favored values obtained for the intrinsic physical parameters such as stellar mass and star formation rate, but with substantially smaller uncertainties. Our results are broadly insensitive to simulation choices for galactic outflows and dust reddening. Hence the existence of early galaxies as observed is broadly consistent with current hierarchical structure formation models. However, one of the six objects has photometry that is best fit by a bursty SFH unlike anything produced in our simulations, driven primarily by a high $K$-band flux. These findings illustrate how \\spoc~provides a robust tool for optimally utilizing hydrodynamic simulations (or any model that predicts galaxy SFHs) to constrain the physical properties of individual galaxies having only photometric data, as well as identify objects that challenge current models. ", "introduction": "Over the last few years, observations of galaxies at $z\\sim 6$ have opened up a new window into the reionization epoch, turning it into the latest frontier both for observational and theoretical studies of galaxy formation. Planned~\\citep{gon05} and existing wide-area narrowband searches for $z\\ga 5.5$ objects such as the Subaru Deep Field~\\citep{aji06,shi06}, the Large Area Lyman-Alpha Survey~\\citep{rho01,mal04}, the Chandra Deep Field-South~\\citep{wan05,mal05}, and the Hubble Ultra Deep Field~\\citep{mal05} are now combining with Lyman-alpha dropout searches~\\citep{dic04,bou04a,bou04b,mob05,bou06,eyl06,lab06}, targeted searches near lensing caustics in galaxy clusters~\\citep{kne04, hu02} and occasionally serendipity~\\citep{ste05} to uncover star-forming galaxies from the reionization epoch in significant numbers \\citep[see] [for a listing of spectroscopically-confirmed $z>5$ galaxies]{ber06}. A question immediately raised by this new stream of observations is, what are the physical properties of these early galaxies? Optimally, one would determine properties such as the stellar mass, star formation rate, and metallicity directly from high-quality spectra, but at present this is infeasible for such faint systems. Hence properties must be inferred from photometry alone, occasionally augmented by emission line information. This requires making some poorly constrained choices for the intrinsic galaxy properties. A commonly applied method known as Spectral Energy Distribution (SED) fitting involves generating an ensemble of population synthesis models under a range of assumptions for the intrinsic nature of the object, and then finding the set of assumptions that best reproduces a given galaxy's observed photometry~\\citep[e.g.][]{ben00,kau03}. The physical properties that yield the lowest $\\chi^2$ model are then forwarded as the most probable values, sometimes with little attention to statistical uniqueness or robustness~\\citep[see][for a nice exploration of such issues]{sch05}. Amongst the various assumptions used in SED fitting, the one that is often least well specified and produces the widest range in final answers is the galaxy's star formation history (SFH). With no prior information, common practice is to use simple SFHs with one free parameter such as constant, single-burst, or exponentially-decaying, which in aggregate are assumed to span the range of possible SFHs for a given galaxy. Indeed, in most cases all one-parameter SFHs yield plausible results, though the parameters obtained and quality of fits in each case can vary significantly. If it were possible to narrow the allowed range of SFHs through independent considerations, physical parameters could in principle be more precisely determined. One approach for constraining SFHs a priori is to incorporate information from currently favored hierarchical structure formation models. As we will discuss in this paper, hydrodynamic simulations tend to produce a relatively narrow range of star formation histories for early galaxies. Their galaxies' SFHs tend to follow a generic {\\it form} at these early epochs, best characterized as a constantly-rising SFH. This form is broadly independent of cosmology, feedback assumptions, or other ancillary factors, and is furthermore distinct from any one-parameter models commonly used today. A primary aim of this paper is to test whether this relatively generic SFH form is consistent with observations, and if so, what the implication are for the physical properties of high-redshift galaxies. Despite impressive recent successes in understanding cosmology and large-scale structure in our Universe~\\citep[e.g.][]{spe06,spr06}, many uncertainties remain in our understanding of galaxy formation. Several recent papers have tested models of high-$z$ galaxy formation by comparing them to observed bulk properties such as luminosity functions at rest-frame UV and \\lya wavelengths. These comparisons have shown that such models are broadly successful at reproducing observations, under reasonable assumptions for poorly constrained parameters such as dust extinction~\\citep{som01,idz04,nig05,fin06,dav06a}. While this broad success is encouraging, it is subject to some ambiguousness in interpretation, because the properties of individual galaxies are not being compared in detail. One could envision situations in which a model reproduces an ensemble property of galaxies but not the detailed spectra of individual objects. As an example, it was forwarded by \\citet{kol99} that Lyman break galaxies at $z\\sim 3$ are actually merger-driven starbursts, in contrast to many other models predicting them to be large quiescent objects. Despite quite different SFHs, both models reproduced many of the same bulk properties such as number densities and clustering statistics. For $z\\ga 6$ galaxies where statistics are currently poor, such degeneracies can hamper interpretations of bulk comparisons of observations to models. A complementary set of constraints on galaxy formation models may be obtained by comparing models to the individual spectra of observed galaxies. In practice, for high-$z$ systems, photometry over a reasonably wide set of bands must substitute for detailed spectra. Such comparisons of models to data would move towards more precise and statistically robust analyses that do not rely on having a large ensemble of objects. This last aspect is critical, because the very earliest observed objects that may provide the greatest constraints on models will in practice always be few in number and detected only at the limits of current technology. In short, what is desireable would be a tool to compare models and observations of high-redshift galaxies that (1) employs reasonably generic predictions of current galaxy formation models; (2) provides a quantitative and robust statistical assessment of how well such models reproduce observations; (3) yields information on the physical properties of galaxies under various assumptions; (4) obtains such information based solely on observed photometry; and (5) does all this on a galaxy-by-galaxy basis rather than relying on having a large statistical sample of observed galaxies. In this paper we introduce such a tool, called \\spoc~(Simulated Photometry-derived Observational Constraints). \\spoc\\ takes as its input the photometry (with errors) of a single observed galaxy along with an ensemble of model spectra drawn either from simulations or generated using one-parameter SFHs. The output is probability distributions of physical parameters derived using a Bayesian formalism, along with goodness-of-fit measures for any given model. The probability distributions give quantitative constraints on the galaxy's physical properties, while the goodness-of-fit can be used to discriminate between models and determine whether a given model (be it simulated or one-parameter SFHs) is able to provide an acceptable fit to that galaxy's photometry. After introducing and testing \\spoc, we apply it to a sample of six $z>5.5$ galaxies from the literature that have published near-infrared photometry. We show that in five of six cases, the simulated galaxies fit observations at least as well as one-parameter SFHs. Since there is no guarantee that simulations produce galaxy SFHs that actually occur in nature, the fact that good fits are possible shows that the existence of the majority of observed $z\\ga 5.5$ galaxies is straightforwardly accommodated in current galaxy formation models. However, in one case, we find that simulated galaxies provides a much poorer fit than can be obtained with one-parameter SFHs, as burstier SFHs provide a much better fit than can be obtained from any simulated galaxies. At face value, this implies that our simulations cannot yet accommodate the full range of observed galaxies, and that some physical process may be missing, although we will explore alternate interpretations. For each galaxy we also present the best-fit physical parameters, with uncertainties, obtained using each model SFH. The simulations provide significantly tighter constraints than the full range of one-parameter SFHs, as expected based on their relatively small range of SFHs produced. These values can therefore be regarded as predictions of our simulations that may be tested against future observations. \\S~\\ref{sec:spoc} introduces \\spoc, detailing our Bayesian formalism and discussing systematic uncertainties. \\S~\\ref{sec:models} presents the simulations and the one-parameter models that will be used as the template library for \\spoc. \\S~\\ref{sec:performance} discusses what drives the inferred physical properties in the context of our simulations, and shows that~\\spoc~accurately recovers the physical properties of simulated galaxies. \\S~\\ref{sec:kesr} explores the best-fit parameters of one observed reionization-epoch galaxy in detail, and compares with results from traditional one-parameter SFH models. \\S~\\ref{sec:otherGals} repeats the previous comparison for a larger set of observed galaxies, highlighting the variety of interesting results that \\spoc\\ obtains. Finally, in \\S~\\ref{sec:summary} we present our conclusions. ", "conclusions": "\\label{sec:summary} In this paper we present a Bayesian SED-fitting engine called \\spoc, which provides constraints on the physical properties of galaxies from photometric data. \\spoc\\ takes as input a galaxy's photometry and a set of model galaxy spectra, obtained either by assuming an analytic form for the galaxy's star formation history or from numerical simulations of galaxy formation, and outputs probability distributions for the physical properties of individual galaxies based on those model priors. Here we compare and explore implications for different models, with an eye towards better constraining the physical properties of high-redshift galaxies. Because \\spoc\\ is intended to test whether predicted SFHs match those inferred from data, it provides constraints on models of galaxy formation that complement comparisons to bulk properties such as galaxy luminosity functions or clustering. A successful model must not only reproduce observed bulk statistics, but must also reproduce the colors of individual galaxies across all observed bands. Galaxies with photometry that cannot be well fit in a given model provide insights into model failings, and ultimately into the physics that drives galaxy formation. Using \\spoc\\ it is possible to quantitatively determine how well a particular model matches an individual galaxy. Such a methodology is particularly useful at epochs when only a small number of galaxies are detectable, such as for the earliest galaxies known. We show that \\spoc\\ accurately recovers the input physical parameters of model galaxies when fitted with SFH derived from the simulations themselves, with typical systematic errors less than the formal $1\\sigma$ fitting errors. When a simulated galaxy is fit with one-parameter star formation histories or the incorrect extinction model, then larger deviations can occur. Since the true SFH and extinction law are unknown, these deviations can be regarded as systematic errors on the determination of physical parameters. In most cases, such systematic errors are less than 50\\% in stellar mass, star formation rate, and $A_V$, and $\\la 2\\%$ in redshift. We then apply \\spoc\\ to six galaxies at $z\\ga 5.5$ that have published photometry spanning the 4000\\AA\\ break (which is necessary in order to obtain meaningful constraints on physical properties). We begin with a more in-depth study of Abell 2218 KESR ($z\\approx 6.7$), since at present it is probably the best-studied reionization-epoch system, and then apply \\spoc\\ to five more galaxies in order to investigate a wider variety of galaxy properties. Our main conclusions are summarized as follows: \\begin{itemize} \\item The physical parameters derived for Abell 2218 KESR using our library of simulated galaxies are in line with previous determinations, and those employing various one-parameter SFHs. The formal uncertainties are smaller when using simulated galaxies, because the simulations intrinsically produce a relatively narrow range of SFHs for galaxies at these epochs. The existence of an object with the properties of Abell 2218 KESR at these epochs does not pose a significant challenge to current models of galaxy formation. \\item All inferred physical parameters for Abell 2218 KESR except metallicity are remarkably insensitive to the explored choices for superwind feedback model, dust extinction law, and cosmology. The fundamental reason for this is that regardless of such choices, simulations produce a similar {\\it form} for the SFH of early galaxies. Disappointingly, this means that photometry alone cannot constrain such modeling uncertainties. On the other hand, this means that simulation predictions of physical parameters using \\spoc\\ are quite robust. Discrimination between, say, superwind feedback models can be obtained through comparisons with bulk properties such as luminosity functions \\citep[as shown in][]{dav06a}. \\item Exploring a set of six $z\\ga 5.5$ galaxies, we find that for five of the systems the simulated galaxies and one-parameter models produce overlapping probability contours and best-fit physical parameters that are formally consistent, and the minimum reduced $\\chi^2$ is similar between all models. This shows that simulations are usually but not always able to produce galaxies with properties similar to observed $z\\sim 6$ systems. \\item The spectra of simulated galaxies always show a significant Balmer break, despite the fact that their SFHs are best characterized as constantly-rising (not decaying). The median stellar age is $\\approx 120-250$~Myr ($2\\sigma$) for all six objects, spanning a range in stellar masses from $\\sim 5\\times 10^8M_\\odot \\la M_*\\la 2\\times 10^{10}M_\\odot$. Hence the existence of somewhat older stellar populations in these early systems is consistent with simulation predictions. \\item The object GLARE\\#3001 is an outlier that is poorly fit by our simulations, though in fairness the fits with one-parameter models are not optimal either. The best fits tend to favor models with younger stellar ages, lower masses, and higher SFRs than any simulated galaxies, perhaps indicating that this is a galaxy undergoing a burst. However, as discussed in the text, these fits are mainly driven by a high $K_s$-band flux. If that data point were confirmed it would suggest that the simulations may be missing some physical process(es) that governs a small fraction of galaxies at this epoch. \\end{itemize} Our hydrodynamic simulations of galaxy formation make some clear predictions for the star formation histories of galaxies at these early epochs. In particular, they predict that the stellar mass, star formation rate, and metallicity are all tightly correlated~\\citep{fin06,dav06a}. Indeed such trends are generically seen in most hydrodynamic simulations of galaxy formation, though not necessarily in semi-analytic models. Our preliminary comparisons to $z\\sim 6$ galaxies show that, while the simulations are not necessarily statistically favored over other classes of SFHs, the available data at least do not rule out these simulation predictions. Together with the fact that such simulations can reproduce key bulk properties of early galaxies~\\citep{fin06,dav06a}, this suggests that models of early galaxy formation are able to reproduce a range of observed properties of galaxies at these epochs. In this work we have performed the simplest possible spectral synthesis calculations, primarily because our focus has been a comparison between different SFH models. In the future we plan to investigate different population synthesis models~\\citep[e.g.,][]{mar06}, IMFs~\\citep[e.g.][]{far06}, and nonstellar contributions to the observed fluxes such as emission lines from HII regions. Eventually we will apply \\spoc\\ to a larger sample of galaxies, which will enable a number of interesting investigations. First, the bulk statistics (e.g. stellar mass or star formation rate functions) derived using \\spoc\\ can be compared against that directly produced in simulations. While this is partly a circular comparison because simulated galaxies are being used to infer the physical properties, in practice there is no guarantee that consistency will be achieved because the simulated SFHs are quite generic; hence this should provide a stringent test of models of galaxy formation. Second, \\spoc\\ can be used to identify populations of galaxies that deviate dramatically from simulation predictions (such as GLARE\\#3001), in order to isolate which physical processes may be missing in models. In principle this could quantify the contributions from nonstellar emission-dominated sources such as AGN and extremely dusty objects. Finally, the redshift evolution of the galaxy population provides a strong test of models, particularly investigating the dramatic change in the nature of massive galaxies that appears to occur at $z\\sim 1-2$ \\citep{pap05}. \\spoc\\ is a general-purpose code, in principle able to utilize any type of model that produces detailed galaxy SFHs, be it a hydrodynamic, semi-analytic, or analytic model. In principle there is nothing that restricts \\spoc\\ to high redshift use, but in practice other physical phenomena such as AGN contamination may become more important at lower redshifts. \\spoc\\ could also provide a useful tool to identify galaxy classes that are not readily reproduced in current models. We are currently planning to make \\spoc~and our latest library of simulated galaxies publicly available upon publication of this paper. We hope that \\spoc\\ will be a useful and flexible tool for conducting detailed comparisons between simulations and observations, as are critical for advancing our understanding of galaxy formation." }, "0607/astro-ph0607569_arXiv.txt": { "abstract": "{ The extreme synchrotron BL Lac object H\\,2356$-$309, located at a redshift of $z = 0.165$, was observed from June to December 2004 with a total exposure of $\\approx$40\\,h live-time with the \\hess\\ (High Energy Stereoscopic System) array of atmospheric-Cherenkov telescopes (ACTs). Analysis of this data set {yields, for the first time, a} strong excess of \\excess\\ $\\gamma$-rays (10 {standard deviations above background}) from H\\,2356$-$309, corresponding to an observed integral flux above 200\\,GeV of I($>$200\\,GeV) = \\iflux\\ {(statistical error only)}. The differential energy spectrum of the source between {200\\,GeV} and 1.3\\,TeV is well-described by a power law with a normalisation (at 1 TeV) of N$_0$ = \\normalisation\\ and a photon index of $\\Gamma$ = \\spectralindex . H\\,2356$-$309 is one of the most distant BL Lac objects detected at very-high-energy $\\gamma$-rays so far. Results from simultaneous observations from ROTSE{-III} (optical), RXTE (X-rays) and NRT (radio) are also included {and used together with the \\hess\\ data} to constrain a single-zone homogeneous synchrotron self-Compton (SSC) model. {This model provides an adequate fit to the \\hess\\ data when using a reasonable set of model parameters.} ", "introduction": "\\label{intro} The Spectral Energy Distribution (SED) of Active Galactic Nuclei (AGN) spans the complete electromagnetic spectrum from radio waves to very-high-energy (VHE; E\\,$>$\\,100\\,GeV) $\\gamma$-rays. In the widely-accepted unified model of AGN \\citep[e.g.][]{rees:1984a,urry:1995a}, the ``central engine'' of these objects consists of a super-massive black hole (up to 10$^9$\\,M$_\\odot$) surrounded by a thin accretion disk and a dust torus. In {some} radio-loud AGN, i.e. objects with a radio to B-band flux ratio F$_\\mathrm{5GHz}$/F$_\\mathrm{B}$\\,$>$\\,10, two relativistic plasma outflows (jets) {presumably} perpendicular to the plane of the accretion disk have been observed. AGN are known to be VHE $\\gamma$-ray emitters since the detection of Mrk\\,421 above 300\\,GeV by the Whipple group \\citep{punch:1992a}, {{who} pioneered the imaging atmospheric-Cherenkov technique.} At very high energies, a number of AGN ($\\approx$10) were subsequently detected by different groups using a similar technique. Almost all these objects {are {BL Lacertae (BL Lac)} objects, {belonging to the class of Blazars (BL Lac objects and Flat Spectrum Radio Quasars),}} i.e. AGN having their jet pointing at a small angle to the line of sight. The only confirmed VHE detection of an extragalactic object not belonging to the BL Lac class is the giant radio galaxy M\\,87 \\citep{aharonian:2003b,aharonian:2005f}. Two broad peaks are present in the observed SED of AGN. The first peak is located in the {radio, optical, and X-ray bands}, the second peak is found at higher energies and can extend to the VHE band. The observed broad-band emission from AGN is commonly explained by two different model types. In leptonic models, the lower-energy peak is explained by synchrotron emission of relativistic electrons and the high-energy peak is assumed to result from inverse Compton (IC) scattering of electrons off a seed-photon population, see e.g. \\citet[][]{sikora:2001a} and references therein. In hadronic models, the {emission is} assumed to be produced via the interactions of relativistic protons with matter \\citep{pohl:2000a}, ambient photons \\citep{mannheim:1993a} or magnetic fields \\citep{aharonian:2000c}, or {via the interactions of relativistic protons with} {photons and magnetic fields} \\citep{muecke:2001a}. The observed $\\gamma$-ray emission {from BL Lac objects} shows high {variability} ranging from short bursts of sub-hour duration to long-time activity of the order of months. Detailed studies of variability of BL Lac type objects can contribute to the understanding of their intrinsic acceleration mechanisms \\citep[e.g.][]{krawczynski:2001a,aharonian:2002e}. Additionally, observations of distant objects in the VHE band provide an indirect measurement of the SED of the Extragalactic Background Light (EBL), see e.g. \\citet{stecker:1992a,primack:1999a} and references therein. Due to the absorption of VHE $\\gamma$-rays via e$^+$e$^-$ pair production with the photons of the EBL, the shape of the observed VHE spectra is distorted as compared to the intrinsically emitted spectra. Using a given spectral shape of the EBL, the observed AGN spectrum can be corrected for this absorption. The resulting intrinsic (i.e., corrected) spectrum can then be compared to basic model assumptions on the spectral shape of the $\\gamma$-ray emission, thereby constraining the applied shape of the EBL. In this context, it is especially important to detect AGN at higher redshifts but also to study the spectra of objects over a wide range of redshifts, in order to disentangle the effect of the EBL from the intrinsic spectral shape of the objects. To date, the redshifts of VHE emitting {BL Lac objects} {with measured spectra} range from $z = 0.033$ to $z = 0.129$. The high frequency peaked BL Lac object (HBL) H\\,2356$-$309, identified in the optical by \\citet{schwartz:1989a}, is hosted by an elliptical galaxy located at a redshift of $z = 0.165$ \\citep{falomo:1991a}. The object was first detected in X-rays by the satellite experiment UHURU \\citep{forman:1978a} and subsequently by the Large Area Sky Survey experiment onboard the HEAO-I satellite \\citep{wood:1984a}. The spectrum of H\\,2356$-$309 as observed by BeppoSAX \\citep{costamante:2001a} is not compatible with a single power law model, indicating that the peak of the synchrotron emission lies within the energy range {of BeppoSAX}. A broken power law fit yields a synchrotron peak around 1.8\\,keV, with a detection of the source up to 50\\,keV. These observations qualified the object as an \\emph{extreme synchrotron blazar}. A selection of TeV candidate BL Lac objects was proposed by \\citet{costamante:2002a}. The objects were selected from several BL Lac samples and using information in the radio, optical and X-ray bands. VHE predictions for the selected objects were given by the authors based on a parametrisation proposed by \\citet{fossati:1998a}, suitable for predictions of high state flux of an average source population. {The authors also gave VHE flux predictions based on} a simple one-zone homogeneous SSC model \\citep{ghisellini:2002a}, appropriate for a quiescent state of the specific VHE source candidate. H\\,2356$-$309 is included in this list and the predicted integral flux values above 300\\,GeV for H\\,2356$-$309 are 8.4$\\times$\\ifunit\\ for the parametrisation and 1.9$\\times$\\ifunit\\ for the SSC model. {It should be noted that no absorption due to the EBL was taken into account in these calculations.} In this paper the discovery of VHE $\\gamma$-rays from H\\,2356$-$309 with the \\hess\\ Cherenkov telescopes in 2004 is reported. With a redshift of $z = 0.165$, \\mbox{H\\,2356$-$309} is one of the most distant AGN detected at VHE energies so far. { \\mbox{H\\,2356$-$309} was observed by \\hess\\ from June to December 2004 (see sections~\\ref{hessobs} and~\\ref{hessres}). Simultaneous observations were carried out with RXTE {(Rossi X-ray Timing Explorer)} in X-rays on 11th of November 2004 (see section~\\ref{rxtesection}), with the Nan\\c{c}ay decimetric radio telescope (NRT) between June and October 2004 (see section~\\ref{nrtsection}) and with ROTSE{-III} (see section~\\ref{rotsesection}) in the optical, covering the whole 2004 \\hess\\ observation campaign. } ", "conclusions": "The high frequency peaked BL Lac object H\\,2356$-$309, located at a redshift of $z = 0.165$, was discovered in the VHE regime by the \\hess\\ Cherenkov telescopes. Two different reconstruction and analysis methods were applied to the data both yielding consistent results. No strong evidence for variability in the VHE band is found within the \\hess\\ observations. The same holds true in the X-ray band, where the object {does} not show any strong flux variability, neither in the ASM nor in the pointed observations. Additionally, the RXTE flux, observed simultaneously to the \\hess\\ observations, is lower than the previously-measured BeppoSAX flux. This might indicate that our observations took place during a relatively low state of emission. For the first time, an SED comprising simultaneous radio, optical, X-ray and VHE measurements was made. {A simple one-zone SSC model, taking into account absorption by the EBL \\citep{aharonian:2005a}, {provides} a satisfactory description of these data}. Given the high redshift of the object, the observed \\hess\\ spectrum provides strong constraints on the density of the EBL \\citep{aharonian:2005a}. Future observations of H\\,2356$-$309 with \\hess\\ will improve the accuracy of the spectral measurement and might also allow an extension of the observed spectrum to higher energies. This will provide further constraints on the absorption of $\\gamma$-rays by the EBL." }, "0607/astro-ph0607043_arXiv.txt": { "abstract": "Key targets for gravitational wave (GW) observatories, such as LIGO and the next generation interferometric detector, Advanced LIGO, include core-collapse of massive stars and the final stage of coalescence of compact stellar remnants. The combined GW signal from such events occurring throughout the Universe will produce an astrophysical GW background (AGB), one that is fundamentally different from the GW background by very early Universe processes. One can classify contributions to the AGB for different classes of sources based on the strength of the GW emissions from the individual sources, their peak emission frequency, emission duration and their event rate density distribution. This article provides an overview of the detectability regimes of the AGB in the context of current and planned gravitational wave observatories. We show that there are two important AGB signal detection regimes, which we define as `continuous' and `popcorn noise'. We describe how the `popcorn noise' AGB regime evolves with observation time and we discuss how this feature distinguishes it from the GW background produced from very early Universe processes. ", "introduction": "Astronomy has at its disposal a powerful set of technological tools to explore the cosmos, enabling new windows to the Universe to be opened. During the 19th and early part of the 20th century no one envisaged the scope of astrophysical phenomena that the electromagnetic (EM) spectrum would reveal. The optical window into the Universe has been extended to include radio, infrared, ultraviolet, x-ray and $\\gamma$-ray components of the EM spectrum, and neutrino detectors have been constructed and brought into successful operation. These technology-based advances have led to many dramatic discoveries, including the cosmic microwave background radiation and exotic objects such as neutron-stars and super-massive black-holes. Gravitational wave (GW) astronomy, a new window, offers the possibility of dramatically extending our present understanding of the Universe. This new spectrum is presently totally unexplored. Three sets of long baseline laser-interferometer GW detectors have been, or are nearly, constructed. The US LIGO (Laser Interferometer Gravitational-wave Observatory) has started its 4th science run (S4), using two 4-km arm detectors situated at Hanford, Washington, and Livingston, Louisiana; the Hanford detector also contains a 2-km interferometer. The Italian/French VIRGO project is commissioning a 3-km baseline instrument at Cascina, near Pisa. There are detectors being developed at Hannover (the German/British GEO-600 project with a 600-m baseline, commissioning completed May 2006) and near Perth (the Australian International Gravitational Observatory, AIGO, initially with an 80-m baseline). A detector at Tokyo (TAMA-300, 300-m baseline) has been in operation since 2001. The astrophysical detection rates are expected to be low for the current interferometers, such as `Initial LIGO', but second-generation observatories with high optical power are in the early stages of development; these `Advanced' interferometers have target sensitivities that are predicted to provide a practical detection rate. The currently operating detectors are searching for both astrophysical sources of GWs, such as merging binary neutron-star and black-hole systems and core-collapse supernovae, and stochastic backgrounds composed of many individually unresolved astrophysical sources at cosmological distances \\citep{fer99,cow01,cow02}. In addition to these sources, detectors may also be sensitive to a stochastic background of gravitational radiation from the early Universe, a quest that is considered by many to be the `holy grail' of GW searches. We provide here a brief sketch of several key ideas underpinning models for the primordial background, but our main focus is on the other background--a stochastic background produced by astrophysical sources distributed throughout the Universe. This GW signal, as well as being of profound fundamental interest, provides an interesting comparison to that expected from the primordial Universe. \\section[Primordial GW background]{Cosmological background of primordial GWs} Detection of a GW background from the early Universe would have a profound impact on early-Universe cosmology and on high-energy physics, opening up a new window to explore both the Planck epoch at $10^{-43}$ seconds into the evolution of the Universe and physics at correspondingly high energies that will never be accessible by any other means. Relic GWs must carry unique information from the primordial plasma at a graviton decoupling time when the Universe had a temperature $\\sim 10^9$ K, providing a `snapshot' of the state of the Universe at that epoch: cosmological GWs could probe to great depth in the very early Universe. As currently understood, mechanisms for generating GWs in the primordial Universe can best be described as speculative. Four popular mechanisms are mentioned briefly here; see Maggiore (2000, Sects. 8, 9 \\& 10) for a comprehensive account. \\noindent{\\bf Vacuum fluctuations.} In inflationary models, as the Universe cooled, it passed through a phase in which its energy density was dominated by vacuum energy, and the scale factor increased exponentially fast. The spectrum of GWs that might be observed today from this period results from adiabatically-amplified zero-point energy fluctuations. The spectrum is expected to be extremely broad, covering $10^{-18}-10^8$ Hz. \\noindent{\\bf Phase transitions.} Phase transitions in the early Universe, particularly at the electro-weak scale, are potential sources of GWs. The Standard Model of particle physics predicts that there will be a smooth cross-over between states rather than a phase transition, but supersymmetric theories predict first-order phase transitions, hence allowing for possible strong GW emission in the $10^{-5}-1$ Hz band. \\noindent{\\bf Cosmic strings.} Grand unified theories predict topological defects; among these, vibrating cosmic strings, characterized by a mass per unit length, could potentially be strong sources of GWs. The spectrum has a narrow band feature peaking at $10^{-12}$ Hz and a broad band component in the $10^{-10}-10^{10}$ Hz band. \\noindent{\\bf Coherent excitations of the Universe}, regarded as a `brane' or surface defect in a higher-dimensional universe. The excitations could be of a `radion' field that controls the size or curvature of the additional dimensions, or they could be of the location and shape of our Universe's brane in the higher dimensions. If there was an equipartition of energy between these excitations and other forms of energy in the very early Universe, then the excitations could produce GWs strong enough for detection by advanced LIGOs. They could probe one or two additional dimensions of size $\\sim 10^{-13}$ -- $10^{-16}$ m. If the number of extra dimensions is more than 2, much smaller scales could be reached (Cutler \\& Thorne 2002; Hogan 2000). \\subsection{Characterizing the background}\\label{chartheback} A cosmological stochastic background of GWs is expected to be isotropic, stationary and unpolarized. The intensity of a stochastic background of GWs is conventionally characterized by the dimensionless `closure density', $\\Omega _{GW}(f)$, defined as the fraction of energy density of GWs per logarithmic frequency interval normalized to the cosmological critical energy density $\\rho_c c^2$ required to close the Universe: \\begin{equation} \\Omega_{GW}(f)=\\frac{1}{\\rho_c c^2}\\frac{{\\rm d}\\rho_{\\rm GW}}{{\\rm d}\\log f}, \\end{equation} where $\\rho_{\\rm GW}(f)$ is the energy density of the background at frequency $f$. In terms of the present value of the Hubble constant $H_0$, written as $h_0\\times100$ km s$^{-1}$ Mpc$^{-1}$: \\begin{equation} \\rho_c =\\frac{3H_0^2}{8\\pi G}\\approx 1.9 \\times 10^{-26} h_0^2 \\hspace{0.25cm} \\mathrm{kg\\hspace{0.15cm} m^{-3}}. \\end{equation} \\indent There are several other ways of describing the spectrum, including the spectral density (density in frequency space) of the ensemble average of the Fourier component of the metric\\footnote{see Maggiore 2000, Sect 2.2 for a complete derivation of $S_h(f)$}, $S_h(f)$, with units of Hz$^{-1}$; this allows the experimentalist to express the sensitivity of the detector in terms of a strain sensitivity, $S_h^{1/2}(f)$, with dimension Hz$^{-1/2}$. Also, the spectrum can be expressed as a dimensionless characteristic amplitude\\footnote{note that the Fourier convention used here is $\\tilde{g}(f)=\\int_{-\\infty}^{\\infty} {\\rm d}t \\,{\\rm exp}(2 \\pi i f t) g(t)$ so that $g(t)=\\int_{-\\infty}^{\\infty} {\\rm d}f \\,{\\rm exp}(-2 \\pi i f t) \\tilde{g}(f)$}. The relationships between $h_c(f)$, $S_h(f)$ and $h_0^2\\Omega_{GW}(f)$, as derived by e.g. Maggiore (2000, Sect. 2.2), are: \\begin{equation}\\label{1hc} h_c(f)= \\sqrt{2fS_h(f)}, \\end{equation} \\begin{eqnarray}\\label{1omega1} h_0^2\\Omega_{GW}(f)=(2\\pi^2/3) f^2h_c^2(f)=(4\\pi^2/3) f^3S_h(f). \\end{eqnarray} \\subsection[Observational bounds]{Observational bounds on $\\Omega_{GW}$} The integrated closure density of GWs from all frequency bands is expressed in dimensionless form as \\begin{equation} \\Omega_{GW} \\equiv \\int{\\Omega_{GW}(f) d(\\ln f)}. \\end{equation} The current best limit for $h_0^2\\Omega_{GW}$ is $6\\times 10^{-6}$ (Cutler \\& Thorne 2002); a value larger than this would have meant the Universe has expanded too rapidly through the era of primordial nucleosynthesis (Universe age $\\sim $ a few minutes), distorting the universal abundances of light elements away from their observed values. Pulsars are natural GW detectors. As a GW passes between us and the pulsar, the time of arrival of the pulse will fluctuate. The bound on $\\Omega_{GW}(f)$ from pulsar timing is (Maggiore 2000, Sect. 7.3): \\begin{equation} h_0^2 \\Omega_{GW}(f) < 10^{-8},\\hspace{0.4cm} f \\sim 10^{-8}\\hspace{0.1cm} \\mathrm{Hz}\\hspace{0.1cm}.\\\\ \\end{equation} A background of GWs will also cause fluctuations in the temperature of the cosmic microwave background. The bound from anisotropy measurements (Maggiore 2000, Sect. 7.2) is: \\begin{equation} h_0^2 \\Omega_{GW}(f) < 7\\times 10^{-11}(H_0/f)^2,\\hspace{0.5cm} 1 10^4 \\, {\\rm K/hr}$) at a temperature range of $1273 - 1473 \\, {\\rm K}$, in which the isotopic fractionation should occur associated with evaporation of troilites (Tachibana \\& Huss 2005). Below $1273 \\, {\\rm K}$, troilites is solid state and it is assumed that no isotopic fractionation occurs when solid troilite evaporates because the time scale of sulfur diffusion in FeS is much larger than that of evaporation (the evaporation P\\'{e}clet number for troilites at temperatures close to the eutectic point is about 100 for $50 \\, {\\rm \\mu m}$-sized grain). On the contrary, above $1473 \\, {\\rm K}$, the melting troilite grains would have been surrounded by melted silicate. In this case, evaporation of sulfur would be controlled by diffusion through the surrounding silicate melt. However, since sulfur can hardly dissolve to the chondrule silicate melt, the troilite would not evaporate and a large degree of isotopic fractionation is not expected. Dust aggregates that have never melted before are thought to be fluffy. When such aggregates are heated, the isotopic fractionation of troilites should occur associated with the evaporation from inside of fluffy aggregates. If the heating rate is slow, it is impossible to suppress the isotopic fractionation, because the duration of the evaporation becomes long enough to produce certain amount of isotopic fractionation. On the other hand, since once molten dust particles are not fluffy, they do not have to be heated rapidly to prevent from isotopic fractionation because FeS in the dust particle is completely covered by silicate components. Thus, we do not have to take into account the isotopic fractionation during the cooling phase. Namely, chondrules have to be heated rapidly at least in the first melting heating event. The isotopic fractionation can be suppressed due to not only the rapid heating but also the presence of back reaction from evaporated sulfur gas if the dust-to-gas mass ratio is large enough. However, in our situation, this effect can be negligible. Tachibana \\& Huss (2005) also calculated the degree of isotopic fractionation of sulfur under conditions of sufficiently high dust-to-gas mass ratio in a closed system, and concluded that the required dust-to-gas mass ratio to suppress the isotopic fractionation by the back reaction is higher than about ten thousands times the solar value. It means that the dust-to-gas mass ratio should be greater than about 100. It has not been well investigated whether the shock-wave heating (the gas drag heating) works well in such extremely high dust-to-gas mass ratio environment or not, and to answer this problem is beyond the scope of this paper. Thus, we do not take into account the back reaction in this study. Tachibana {\\it et al.} (2004) investigated the heating rate of chondrules in the framework of the shock-wave heating model and concluded that the gas drag heating in the post-shock region can heat chondrules rapidly enough to suppress the isotopic fractionation. However, it is known that chondrules are also heated in the pre-shock region due to the radiation emitted by ambient dust particles (DC02, CH02). The effect is well known as the blanket effect, which was not taken into account in the study by Tachibana {\\it et al.} (2004). Results by DC02, in which the dust thermal radiation is taken into consideration as the radiation source, showed that the heating speed due to the radiation is too slow to suppress the isotopic fractionation ($\\sim 300 \\, {\\rm K/hr}$). CH02 also performed numerical simulation taking the transfer of the dust thermal continuum emission into consideration. In contrast with DC02, results of CH02 showed that the heating rate of chondrules is large enough to suppress the isotopic fractionation even if the pre-shock region is dusty environment (dust-to-gas mass ratio is about 1.5, which is a few hundreds times larger than that of the solar abundance). In previous studies of shock-wave heating model taking into account the radiation transfer (DC02 and CH02), only the dust thermal continuum was taken into consideration as the radiation source. However, there is the other radiation source, the line emission of gas molecules. DC02 and CH02 neglected the line cooling because they assumed that the shock region is optically too thick to its own line emission to lose gas thermal energy effectively. On the contrary, MN05 showed an estimation that the post-shock gas is not so optically thick and the post-shock gas in a few $100-1000 \\, {\\rm km}$ behind the shock front can cool. Therefore, the line cooling should be taken into account. Moreover, numerical simulations in DC02 and CH02 have done for the limited shock conditions (the pre-shock gas number density is $n_0 \\simeq$ a few $10^{14} \\, {\\rm cm^{-3}}$ and the shock velocity is $v_{\\rm s} \\simeq 7 - 9 \\, {\\rm km \\, s^{-1}}$). These shock conditions are classified into the high-density and low-velocity shock waves in the appropriate shock condition for chondrule formation (INSN). Recently, a new shock wave generation mechanism which are sufficient to account for melting of dust particles was proposed (Nakamoto {\\it et al.} 2005). In the case, X-ray flares generated at close to the protosun induce the shock waves in the upper solar nebula where the gas density is low. The typical shock condition of the new mechanism is estimated as $n_0 \\sim 10^{11} \\, {\\rm cm^{-3}}$ and $v_{\\rm s} \\sim 40-60 \\, {\\rm km \\, s^{-1}}$. There is no study for the low-density and high-velocity shock waves taking into account the thermal radiation from dust particles. The purpose of this paper is to develop a new shock-wave heating code taking into account the radiation transfer of both the line emission of gas molecules and the dust thermal radiation and to investigate how the thermal history of dust particles in the shock waves is affected by the optical properties of the pre-shock region. Especially, we focus the heating rate of precursor dust particles in a temperature range of $1273 - 1473 \\, {\\rm K}$, in which the isotopic fractionation should occur. The optical properties of the flow depend mainly on the dust size distribution and the dust-to-gas mass ratio. Moreover, we perform the simulation with various shock conditions (the pre-shock gas number density and the shock velocity) in order to discuss which shock generating mechanism is appropriate for chondrule formation. We describe details of our model and basic equations in section 2. We show the calculation results in section 3 and discussion in section 4. Finally, we summarize our study in section 5. ", "conclusions": "We numerically simulated the shock-wave heating for chondrule formation taking into account the radiation transfer of the line emission of gas molecules and the dust thermal continuum emission. Regarding the line cooling due to the gas molecules, we estimated the cooling rate using the photon escape probability method in order to reflect the column density of gas molecules. We focused the dust thermal histories in the pre-shock region, especially the heating rate of chondrules $R_{\\rm heat}$, which has to be large enough to suppress the isotopic fractionation. In order to clear the condition in which the rapid heating constraint is satisfied, we performed simulations for various shock conditions, the pre-shock gas number density ($n_0 = 10^{11} - 10^{14} \\, {\\rm cm^{-3}}$) and the shock velocity ($v_{\\rm s} = 10 - 55 \\, {\\rm km \\, s^{-1}}$), and the dust models, the initial dust size distribution (the power-law size distribution and the lognormal one) and the dust-to-gas mass ratio ($C_{\\rm d} = 0.01$, $0.03$, and $0.10$). We found the following results: \\begin{enumerate} \\item The line emission can remove the post-shock gas thermal energy away even if the pre-shock gas number density is relatively high ($n_0 = 10^{14} \\, {\\rm cm^{-3}}$). Therefore, the line emission plays an important role as the gas cooling mechanism. Moreover, it indicates that it is also an important radiation source term. \\item When the optical depth of the pre-shock region increases, it results into the higher dust temperatures in the pre-shock region. If the dust temperatures just in front of the shock front exceed $1273 \\, {\\rm K}$, the heating rate of the dust particles in a temperature range of $1273 - 1473 \\, {\\rm K}$ is too slow to prevent the isotopic fractionation. On the contrary, the dust temperatures just in front of the shock front are lower than $1273 \\, {\\rm K}$, the dust particles are heated by the gas frictional heating in the post-shock region so rapidly that the isotopic fractionation is prevented. Therefore, the condition to prevent the isotopic fractionation is that the dust temperatures just in front of the shock front do not exceed $1273 \\, {\\rm K}$. \\item We analytically derived the dust temperature just in front of the shock front using the theory of the radiative diffusion. The analytic solution well explains the results of our numerical simulations for optically thick case (optical depth of the pre-shock region $\\tau_{\\rm pre} \\ga 1$). \\item There is the upper limit of the optical depth of the pre-shock region above which the isotopic fractionation will occur. The value of the upper limit depends on the shock condition. For the low-velocity and high-density shock waves, the upper limit is about unity. For the high-velocity and low-density shock waves, it is about 10. \\item The fundamental factors to determine the pre-shock dust thermal histories are the gas energy flux flowing into the shock front and the optical depth of the pre-shock region. Even if the spacial dimension of the shock-wave heating region is different, the dust temperatures just in front of the shock front become almost the same as long as the pre-shock regions take the similar values of the optical depths. \\item We also evaluated the cooling rate of the precursor dust particles $R_{\\rm cool}$ at a crystallization temperature. It was found that the dust thermal histories obtained by our simulations are roughly classified into two categories: the first is the slow heating ($R_{\\rm heat} = 10^2 - 10^3 \\, {\\rm K \\, hr^{-1}}$) and slow cooling ($R_{\\rm cool} = 5 - 20 \\, {\\rm K \\, hr^{-1}}$), and the second is the rapid heating ($R_{\\rm heat} \\sim 10^6 \\, {\\rm K \\, hr^{-1}}$) and the rapid cooling ($R_{\\rm cool} = 10^3 - 10^5 \\, {\\rm K \\, hr^{-1}}$). If supporting an opinion that the slow cooling rate is appropriate for producing chondrules, it seems difficult to meet two conditions simultaneously in the framework of a one-dimensional, plane-parallel, and steady shock wave model. In order to satisfy the rapid heating constraint and the appropriate slow cooling constraint simultaneously, a non-steady shock wave model or a multiple heating scenario might be needed. \\end{enumerate}" }, "0607/astro-ph0607532_arXiv.txt": { "abstract": "A survey of linearized cosmological fluid equations with a number of different matter components is made. To begin with, the one-component case is reconsidered to illustrate some important mathematical and physical points rarely discussed in the literature. The work of some previous studies of two-component systems are examined and re-analyzed to point out some deficiencies of solutions, and further solutions and physical interpretation are then presented. This leads into a general two-component model with variable velocity dispersion parameters and mass density fractions of each component. The equations, applicable to both hot dark matter (HDM) and cold dark matter (CDM) universes are solved in the long wavelength limit. This region is of interest, because some modes in this range of wavenumbers are Jeans unstable. The mixture Jeans wavenumber of the two-component system is introduced and interpreted, and the solutions are discussed, particularly in comparison to analogous solutions previously derived for plasma modes. This work is applicable to that region in the early Universe ($20 < z < 140$), where large scale structure formation is thought to have occurred. ", "introduction": "The theory of structure formation in the Universe has become one of the most popular and intensely studied topics in modern cosmology. Throughout the past century there has been an accumulating volume of work on the analytic investigation of the cosmological structure formation equations. The various approaches include both fluid and kinetic theory formulations. They principally consider the gravitational interaction of components of the cosmological medium, though sometimes other forms of interaction such as magnetic fields are also included (for some standard examples see e.g. \\cite{padman,peebles}). The analysis of these equations has employed ever more diverse and complicated techniques and approximation schemes to model increasingly realistic physical situations. This has been comprehensively supported and now superseded by large N-body simulations. The algorithms which govern these large numerical studies have grown progressively more refined and subtle, and are now producing very accurate and realistic results, which can be directly compared with observations (e.g.\\cite{virgo}). Despite the current trends in modern cosmological structure formation theory, much can still be learned from relatively simple analytic models. We consider such models, in the face of modern computing power, to analyze at a fundamental level some of the basic physical processes which cause the clustering observed in the Universe. This helps to isolate physical mechanisms difficult to discern numerically. In this paper our interest will focus on the linearized cosmological fluid equations. These equations have been used to build up the components of the cosmological density perturbation power spectrum, and must be evolved through the various stages of cosmological evolution, and over a large range of physical scales. There are several reasons for taking such an approach. The equations may be solved numerically to give detailed power spectra for the various cosmological models currently viable. The power spectra may then be used as initial data to evolve the large N-body simulations, which are ultimately compared to observations. The equations may also be used to build up a semi-quantitative picture of the evolution of the power spectrum. This can show how the various sized perturbations scale with respect to the Friedmann expansion parameter $a$ during different epochs of the Universe, and give a direct insight into some of the fundamental physical processes operating to produce structure in the Universe (see e.g. \\cite{padman}). The evolution of perturbation modes with wavelength greater than the Hubble radius may be studied through a relativistic formulation of the perturbation equations, whereas for modes with wavelengths smaller than this radius, a Newtonian formalism suffices. There is a range of physical parameters and variety of differential equations describing the evolution processes of density fluctuations in the early Universe. This involves such elements as equations of state for energy constituents, and specification of the expansion parameter $a$ by the Friedmann cosmological equations. As a consequence, there is a large wealth of literature on this subject, and most of the currently important techniques and results have been collected in some well known textbooks \\cite{padman,peebles,texts}. Some of the relatively complicated systems of equations have also been studied in the literature, and it is our goal to both review many of these studies, and to extend them in new directions, achieving some new unique results. In this paper we will only study the Newtonian limit of the linearized cosmological perturbation equations, valid for density fluctuations on scales well within the Hubble radius. Our main concern is with some mathematically more complicated multi-component models, which although not usually considered in standard power spectrum analysis, have realistic and interesting physical meaning. The primary concept of the Jeans gravitational instability \\cite{jeans} has been investigated in a static universe for multi-component models, to reveal the more complicated structure of modes possible \\cite{russians1,carvalho}. This provides some interesting qualitative ideas about the possible mechanisms for structure formation, but the lack of an expanding background spacetime in the models leads to unrealistic solutions, exponential in form. The inclusion of cosmological expansion in the equations leads to the more realistic power law and logarithmic solutions, familiar from the standard power spectrum analysis. Previous work in this area has focused both on some particular models \\cite{russians2}, and on a more general classification of the equations and solutions for a range of parameter values (some of them only of mathematical interest) and physical contexts \\cite{fargion,haubold1,haubold2}. Analytic solutions for some of the most general cases of the equations considered above, which often have significant physical interest, have not been achieved. It is our aim here to rectify this situation and investigate a system of equations modeling a two-component fluid in the matter dominated post-recombination era of an Einstein--deSitter universe. One of these components consists of baryons, and the other some form of nonrelativistic dark matter particles. Through this work we will amend what appear to be some errors in the previous general studies of Haubold and Mathai \\cite{haubold2}. This paper makes a comprehensive study of the linear perturbation equations for cosmological fluids with gravitational instabilities with application to large scale structure formation. For a historical perspective we note that Lifshitz \\cite{lifshitz} concluded that gravitational instability could not be responsible for the formation of structure in the Universe. The correct conclusion, that gravitational instability suffices, was pointed out by Novikov \\cite{novikov}. As our work details the mathematical structure of the appropriate equations describing cosmological structure formation, we note the nice series of papers by Ratra and Peebles \\cite{ratra} directed to understanding the applications of special functions to the problem of gravitational instability in cosmological models. We further note the nice series by Buchert {\\em et al} \\cite{buchert} concerning analytic results and their relevance to observational cosmology. We will also make a comparison with work done in cosmological plasma physics in an Einstein-deSitter background \\cite{dettmann,plasma,gailis}. This is interesting due to the mathematically very similar form of fluid equations for both type of systems. This similarity is largely due to the similarity of the electromagnetic and gravitational forces. In this paper we will analyze the long wavelength region of the solutions. This corresponds to the Jeans unstable region of parameter space, and requires use of Frobenius methods of expansion of the differential equations. In a follow-up paper \\cite{paper2} we will investigate the short wavelength region of the solutions, which will require a WKB approximation scheme to be developed. In these papers we take explictly the temperature relationship $T \\sim 1/a(t)^2$, where $a(t)$ is the radius of the Universe. We give here the explanation why this is so. Following standard textbook material in Padmanabhan \\cite{padman} (as summarized in equation (3.118)) and Peebles (1993, page 179) we find that the baryons follow this relationship when \\[ 1 + z = 142( \\Omega_b h^2 / 0.024)^{2/5}. \\] Here $\\Omega_b$ has been scaled to the WMAP best fit value. Thus for redshifts below $z \\sim 140$ the baryon temperature drops as $1/a^2$ down to $z \\sim 20$, where the Universe reionizes (probably in a patchy fashion) and the $1/a^2$ scaling no longer holds. This $20 < z < 140$ redshift region is important, as it is where early large scale structure formation is thought to have occurred. The paper is to be organized as follows. In Section~2 we introduce the most general cosmological density perturbation equations in the Newtonian approximation. We review and classify previous work on these equations to put our current work into context, showing what has been achieved, what needs amendment, and where we will seek to expand current knowledge. In Section~3 we revert to the one-component equations, to illustrate some of the basic principles which will be important later in our analysis, and to reveal some apparently new results. This will enable us to begin to tackle the two-component problem in Section~4. In this section the CDM two-component model in an expanding universe will be investigated. This has ties with the previous work cited, and we will demonstrate the limitations of the existing formalism here. We present new results apparently overlooked in the work of Haubold and Mathai \\cite{haubold2}. After this, we are ready to study the most general baryonic and dark matter equations in Section~5, where we will consider the long wavelength approximation, applicable to either HDM or CDM. This is followed by our conclusions in Section~6. ", "conclusions": "A method for determining the small $k$ solutions of a general two-component cosmological density perturbation model has been expounded in this paper. We have only displayed the solutions to first order, but it is possible derive them up to any order by the method in principle. We have explored the mathematical properties and peculiarities of density perturbations influenced by gravitational interaction, particularly contrasting them to plasma modes, and correcting a number of previous misconceptions in the literature. The expanding Universe introduces new features not predictable from simple static spacetime considerations. In particular, totally new structures to the dispersion relations are found, even in the one-component example. We have shown how the mixture Jeans wavenumber enters the solutions, and clarified its role in an expanding universe context. More work is required to investigate the solutions around the critical scale defined by $k_M$. Although the expansions as derived in this paper are applicable to this region, they are not particularly useful, as many terms in the equations need to be retained when the expansion parameter $\\chi/\\chi_c$ is of $O(1)$. It is unclear how an analytical investigation of this region could proceed at present. We have performed some preliminary studies which involved producing a large number of terms in the expansions (\\ref{gB1})--(\\ref{gD4}) using the Mathematica program described, and then substituting in numerical values for the various physical parameters to obtain numerical coefficients with an ascending series in $\\chi$. At present the plots of these expansions over a range of values of $\\chi$ do not yield reliable results---it is possible that many more terms than are practically calculable will be required, and a very high order of numerical precision will have to be maintained. Other methods of analyzing the modes in this interesting region probably need to be investigated. Of ultimate interest is exploring how these type of modes contribute to the power spectrum. More physical effects may need to be introduced, such as a cosmological constant, or the addition of more matter components. To determine the actual density contrast at a given scale $1/k$, the Fourier modes of the density contrast as derived in this paper would also need to be integrated over the whole range $0 < k < 1/k$. It would be of considerable interest to compare the power spectra calculated by such a method with the well-known power spectra of the various cosmological models in existence today. In concluding, we remark that a similar analysis could be carried out in the post-recombination region $140 < z < 1150$, where now the baryons follow the $T \\sim 1/a$ relationship. The differential equations in Section II will now be different, as will be their solutions; but we expect that the ensuing analysis would yield qualitatively similar results but quantitatively different scaling. This would be a useful future study." }, "0607/cond-mat0607396_arXiv.txt": { "abstract": "We study, using numerical simulations, the dynamical evolution of self-gravitating point particles in static euclidean space, starting from a simple class of infinite ``shuffled lattice'' initial conditions. These are obtained by applying independently to each particle on an infinite perfect lattice a small random displacement, and are characterized by a power spectrum (structure factor) of density fluctuations which is quadratic in the wave number $k$, at small $k$. For a specified form of the probability distribution function of the ``shuffling'' applied to each particle, and zero initial velocities, these initial configurations are characterized by a single relevant parameter: the variance $\\delta^2$ of the ``shuffling'' normalized in units of the lattice spacing $\\ell$. The clustering, which develops in time starting from scales around $\\ell$, is qualitatively very similar to that seen in cosmological simulations, which begin from lattices with applied correlated displacements and incorporate an expanding spatial background. From very soon after the formation of the first non-linear structures, a spatio-temporal scaling relation describes well the evolution of the two-point correlations. At larger times the dynamics of these correlations converges to what is termed ``self-similar'' evolution in cosmology, in which the time dependence in the scaling relation is specified entirely by that of the linearized fluid theory. Comparing simulations with different $\\delta$, different resolution, but identical large scale fluctuations, we are able to identify and study features of the dynamics of the system in the transient phase leading to this behavior. In this phase, the discrete nature of the system explicitly plays an essential role. ", "introduction": "The problem of the evolution of self-gravitating classical particles, initially distributed very uniformly in infinite space, is as old as Newton. Modern cosmology poses essentially the same problem as the matter in the universe is now believed to consist predominantly of almost purely self-gravitating particles --- so called dark matter --- which is, at early times, indeed very close to uniformly distributed in the universe, and at densities at which quantum effects are completely negligible. Despite the age of the problem and the impressive advances of modern cosmology in recent years, our understanding of it remains, however, very incomplete. In its essentials, i.e., stripped of the full detail of current cosmological models, it is a simple well posed problem of out of equilibrium statistical mechanics\\footnote{ Strictly speaking it is not actually known whether the problem is well-controlled without a regularization of the singularity in the gravitational force at $r=0$ (see e.g. \\cite{elskens_etal} for a recent discussion and list of references). In practice, in numerical simulation, there is no intrinsic problem in implementing the N-body gravitational dynamics without such a regularisation for typical initial conditions (i.e. in which particles are not placed initially at the same point). In the numerical simulations reported here, as in cosmological simulations, we do, however, use such a regularization. This is done solely for numerical efficiency, and the results analysed are tested numerically for their independence of the associated cut-off (see below).}. In this context, however, it has been relatively neglected, primarily because of the intrinsic difficulties associated with the attractive long-range nature of gravity and its singular behavior at vanishing separation. In recent years there has, however, been renewed interest (see e.g. \\cite{dauxois}) in the physics of systems with long-range interactions, in which context self-gravitating systems are one of the paradigmatic examples (for a review, see e.g. \\cite{Padmanabhan:1989gm}). A considerable amount of work on these systems in this context has focussed on {\\it finite} systems (see e.g. \\cite{ChavanisPRE, ispolatov_cohen, morikawa_virial,morikawa_nongauss}) --- in which, in certain cases, some of the instruments of equilibrium statistical mechanics may be applied \\footnote{We note that in \\cite{devega_sanchez1,devega_sanchez2} a treatment of infinite self-gravitating systems in the framework of equilibrium statistical mechanics is developed by considering a ``dilute'' infinite volume limit, in which $N \\rightarrow \\infty$ and $V \\rightarrow \\infty$ at $N/V^{1/3}={\\rm constant}$, where $N$ is the number of particles and $V$ is the volume (see also \\cite{chavanis_2006_I} for a more recent discussion) . This is not the physically relevant limit for the problem treated here, as we consider the infinite volume limit taken at constant density, i.e., $N \\rightarrow \\infty$ and $V \\rightarrow \\infty$ at $N/V ={\\rm constant}$. In this case, as discussed further below, the system is intrinsically time dependent and never reaches a thermodynamic equilibrium.} --- and on more tractable one-dimensional models (see e.g. \\cite{miller_1dgrav, miller_1dgrav_relaxation, miller_1dgrav_segregation, miller_1dgrav_ergodicity, Tatekawa:2005nj}). In cosmology perturbative approaches to the problem, which treat the very limited range of low to modest amplitude deviations from uniformity, have been developed (see e.g. \\cite{peebles,padmanabhan}), but numerical simulations are essentially the only instrument beyond this regime. While such simulations constitute a very powerful and essential tool, they lack the valuable guidance which a fuller analytic understanding of the problem would provide. The dynamics of infinite self-gravitating systems is thus both a fascinating theoretical problem of out of equilibrium statistical mechanics, directly relevant both in the context of cosmology and, more generally, in the physics of systems with long-range interactions. Approaching the problem in the context of statistical mechanics, as we do here, it is natural to start by reducing as much as possible the complexity of the analagous cosmological problem. We wish to focus on the essential aspects of the problem. Thus we consider clustering without the expansion of the universe, and starting from particularly simple initial conditions. With respect to the motivation from cosmology, there is of course a risk : in simplifying we may loose some essential elements which change the nature of gravitational clustering. Our results suggest that this is not the case. Even it were, it seems unlikely that we will not learn something about the more complex cosmological problem in addressing this slightly different problem. Gravitational clustering in an infinite space --- static or expanding -- starting from quasi-uniform initial conditions, is intrinsically a problem out of equilibrium. By ``quasi-uniform'' initial conditions we mean that the initial state is a particle distribution --- specified, we will assume, by a stochastic point process \\cite{daley} --- which has relative fluctuations at all scales, of small amplitude above the scale characteristic of the particle ``granularity'' and decaying at infinitely large scales \\footnote{It is also implicit in the phrase ``quasi-uniform initial conditions in infinite space'' that, as noted above, the infinite volume limit here is taken at constant particle density, rather than in the ``dilute'' limit studied in \\cite{devega_sanchez1,devega_sanchez2}.}. One of the most basic results (see e.g. \\cite{peebles,padmanabhan} and also the appendices to this paper) about self-gravitating systems, treated in a fluid limit, is that the amplitude of small fluctuations grows monotonically in time, in a way which is independant of the scale. This linearised treatment breaks down at any given scale when the relative fluctuation at the same scale becomes of order unity, signalling the onset of the ``non-linear'' phase of gravitational collapse of the mass in regions of the corresponding size. In an infinite space, in which the initial fluctuations are non-zero and finite at all scales, the collapse of larger and larger scales will continue ad infinitum. The system can therefore never reach a time independent state, and in particular it will never reach a thermodynamic equilibrium \\footnote{This does, of course, not mean that the instruments of equilbrium statistical mechanics are completely irrelevant. Saslaw (see \\cite{saslaw2000} and references therein) notably has developed a treatment of gravitational clustering in an expanding universe which approximates it as a ``quasi-equilbrium'' in which the thermodynamic variables evolve adiabatically with the expansion of the Universe. Another more formal exploration of the usefulness of some standard equilibrium techniques can be found in \\cite{morikawa_cluster}.}. One of the important results from numerical simulations of such systems in the context of cosmology is, however, that the system nevertheless reaches a kind of scaling regime, in which the temporal evolution is equivalent to a rescaling of the spatial variables \\cite{efstathiou_88,smith}. This spatio-temporal scaling relation is referred to as ``self-similarity'' \\footnote{Note that this term is here used in a different sense to that commonly ascribed to it in condensed matter physics. In this context ``self-similarity'' usually implies that the spatial correlations themselves have invariance properties under rescaling (see, e.g., \\cite{huang}). This is not necessarily the case in the present context.}. It is observed, however, only starting from a restricted class of simple initial conditions --- we will describe these in further detail below --- and in the specific Einstein de Sitter (EdS) expanding universe \\cite{peebles}. The range of initial conditions to which it applies has been a point of discussion in the literature, and theoretical explanations of it typically restrict it to quite a narrow range of such initial conditions, and strictly to the EdS expanding universe. To see whether this kind of simple behavior is reproduced in the system we study, is thus a first point of interest. It is in fact the primary focus of this paper. One comment needs to be made about the use of a static (Euclidean) space-time. The problem of bodies interacting by their mutual Newtonian self-gravity in the infinite volume limit, taken at constant mean density, is in fact ill defined: the force on a particle depends on how the limit is taken. In order to remove this ambiguity one adds a negative background to cancel the contribution of the mean density --- the so-called ``Jeans Swindle'' (see e.g. \\cite{binney}). As discussed in \\cite{gabrielli_06}, this is equivalent to taking the limit symmetrically about each particle on which we calculate the total gravitational force\\footnote{See \\cite{kiessling} for a very clear discussion of this issue. It is also shown here that addition of the negative backgound is equivalent to regularizing the problem with a cosmological constant.}. Then only the fluctuations of the density field generate the gravitational force. In the context of cosmological expanding universe solutions, this ``swindle'' is unnecessary as the expansion absorbs the effect of the mean density, and the perturbations to the comoving particle trajectories are indeed sourced only by the fluctuations (see, e.g., \\cite{peebles}). This modification does not necessarily make the gravitational force well defined in general: whether it is well defined depends on the nature of the fluctuations in the density field at large scales. For the case of the shuffled lattice (SL) considered here, we have studied in detail the properties of the gravitational force in \\cite{gabrielli_06}, and shown the force to be well defined in the presence of the canceling background. Previous works in the same spirit as this \\cite{bottaccio3,bottaccio2,Baertschiger:2004tx} have treated primarily the very simplest initial condition one can envisage: Poisson distributed particles with no initial velocity. One of the basic results which has been emphasized in these works is the role of nearest neighbor interactions at early times in forming structures (see also \\cite{saslaw}), giving rise to non-linear density-density correlations which are then observed to be reproduced at larger and larger scales as time evolves. At the same time the effects of amplification at larger scales --- described by the fluid limit in which the granular structure of the matter is irrelevant --- is observed. When trying to address the basic issue of the relative importance of these mechanisms, one runs into the limits imposed by the simple initial conditions: in a Poisson distribution a single parameter --- the particle density, or equivalently mean inter-particle distance --- controls both the amplitude of fluctuations and the ``granularity'' of the mass distribution. This limitation is one of the major motivations for the different class of initial conditions we study in this work, developing further some initial analysis of this case in \\cite{Baertschiger:2004tx}: we consider lattices subjected to small random displacements. In this case there are now two parameters, the inter-particle distance $\\ell$ and the amplitude $\\Delta$ of the ``shuffling''. Given the scale free nature of gravity it is in fact only the dimensionless combination $\\delta = \\Delta/\\ell$ which is physically relevant (while in the case of Poisson initial conditions there is effectively no free adjustable parameter). When the dynamics of the SL is treated in the fluid limit, as we will see, configurations with different $\\delta$ may also be trivially related. In particular we can consider systems with different $\\delta$ which have different discreteness properties which are equivalent in terms of their fluid description. This allows us to understand notably the aspects of the evolution of the system which can be accounted for in a description of the dynamics in a fluid limit, and those which require the discreteness of the system to be {\\it explicitly} taken into account. This is an important point as almost all existing analytic results on infinite self-gravitating systems are derived in this former limit\\footnote{It is also a question which is very relevant in the context of cosmology, as it concerns the understanding of the discreteness effects in simulations of dark matter, which intrinsically limit their precision. These simulations treat the gravitational clustering of point ``macro-particles'', which typically correspond to the order of $10^{70}$ dark matter particles.}. Our initial conditions are similar, but not identical, to those used in cosmological simulations of the formation of structure in the Universe. In this context the initial conditions are usually given by simple cubic lattices, perturbed by {\\it correlated} displacements, with relative displacements between nearest neighbor particles which are small \\cite{efstathiou_init}. The displacements are generated in reciprocal space starting from an input power spectrum (PS), i.e., what is usually called the ``structure factor'' in condensed matter physics, specifying the desired theoretical density fluctuations. In this paper we describe systematically basic results on gravitational dynamics starting from SL initial conditions. Our principal results are the following: \\begin{itemize} \\item Evolution from these initial conditions converges, after a sufficient time, to a ``self-similar'' behavior, in which the two-point correlation function obeys a simple spatio-temporal scaling relation. The time dependence of the scaling ( i.e. the quantity analogous to the dynamical exponent in out of equilibrium statistical mechanics) is in good agreement with that inferred from the linearized fluid approximation. This result is a generalization of what has been observed, for ``redder'' initial PS ($P(k) \\sim k^n$ with $n\\leq 1$), in simulations in an EdS universe \\cite{efstathiou_88, bertschinger, smith}. \\item Between the time at which the first non-linear correlations emerge in a given SL and the convergence to this ``self-similar'' behavior, there is a transient period of significant duration. During this time, the two-point correlation function already approximates well, at the observed non-linear scales, a spatio-temporal scaling relation, but in which the temporal evolution is faster than the asymptotic evolution. This behavior can be understood as an effect of discreteness, which leads to an initial ``lag'' of the temporal evolution at small scales. \\item Simulations with different particle numbers, but the same large scale fluctuations (as characterized by the PS at small $k$), converge after a sufficient time, not only to the same functional form of the correlation function (with the self-similar behavior), but to the same amplitude. This is further evidence that it is indeed the common large scale fluctuations alone which determine the amplitudes of the correlations, which are thus independent of the discreteness scale $\\ell$. At early times, however, we see manifest difference between the systems, typically again characterized as a ``lag'' of simulations with larger $\\ell$ (and smaller $\\delta$). \\item The non-linear correlations when they first develop are very well accounted for solely in terms of two-body correlations. This is naturally explained in terms of the central role of nearest neighbor interaction in the build-up of these first non-linear correlations. \\item This two-body phase extends to the time of onset of the spatio-temporal scaling, and thus the asymptotic form of the correlation function is already established to a good approximation at this time. We briefly discuss the significance of this quite surprising finding. \\end{itemize} The paper is organized as follows. In the next section we briefly define a SL distribution and introduce the main statistical quantities we use in the analysis and their estimators. We discuss the numerical simulations and their analyses in Sec.\\ref{sec3}. Finally in Sec.\\ref{sec4} we summarize our mains results and conclusions, and briefly discuss some of the many open problems which remains for future investigation. ", "conclusions": "\\label{sec4} To conclude we first summarize our conclusions, and then make a few remarks on open questions to be explored in further works. We have studied the evolution under their Newtonian self-gravity, in a static euclidean space, of classical point particles initially distributed in {\\it infinite} space in a quasi-uniform manner. This is a paradigmatic problem of the out of equilibrium statistical mechanics of long range interacting systems, which has received little attention in this context. Specifically we have considered a one relevant parameter class of initial conditions in which the particles are randomly perturbed off a lattice. We have found that our simulations converge aysmptotically (but for times smaller than those at which the size of the finite simulation box becomes relevant) to solutions characterized by a simple spatio-temporal scaling relation in which the temporal dependence of the scaling can be derived from the linearized fluid theory. These results are qualitatively very similar to those observed in numerical studies in the context of cosmology, i.e., for expanding space-times and for more complex initial conditions in which the displacements of the particles off the lattice are correlated in order to produce the PS of fluctuations of cosmological models. More specifically, the observed spatio-temporal scaling is a simple generalization of what is known in the cosmology literature as ``self-similarity'' in an expanding universe to the case of (i) a static universe, and (ii) a PS $P(k) \\propto k^2$. Further we have observed that there is a transient phase to this behavior, in which already, to a good approximation, the same spatio-temporal scaling relation holds for the two-point correlation function $\\xi(r,t)$, but with a more rapid temporal evolution of the scaling factor. We have noted that the lagging of the evolution behind the asymptotic behavior in this regime can be ascribed to effects of discreteness (i.e. non fluid effects) slowing down the evolution of fluctuations at scales comparable to the inter-particle distance which have been quantified in \\cite{joyce_05,marcos_06}. We have seen also that the form of the correlation function emerges already at the very early times when the first non-linear correlations develop due to two-body correlations which develop under the effect of nearest neighbor interactions. The gravitational evolution of a SL in a static universe thus shares the qualitative features of similar, but more complicated models, in cosmology. It thus provides a simplified ``toy model'' in which to study some fundamental problems which remain open concerning the evolution of these systems, which have been studied extensively in numerical simulations but remain poorly understood analytically, notably: \\begin{itemize} \\item The absence of a theory which adequately explains the shape (i.e. functional form) and evolution of the observed non-linear correlations. \\item The absence of a ``theory of discreteness errors''. In cosmology simulations of particles displaced off lattices (or ``glasses'') aim to reproduce the evolution of a self-gravitating fluid. There is currently very little systematic understanding of how well this evolution is actually approximated. We have highlighted in this paper that the SL gives a very well defined, and simplified, framework in which to address this problem. \\end{itemize} Let us remark finally on a few other points: \\begin{itemize} \\item We have worked here with initial velocities set equal to zero. In exploring the analogy with cosmological simulations there is another choice of initial velocities which is natural. This is that corresponding to that given by the Zeldovich approximation discussed above, with $f(t)$ chosen in Eq.~(\\ref{Zeldovich}) to correspond to the purely growing mode of density fluctuations, i.e., $f(t)=e^{t/\\tdyn}$. The initial velocities are then simply the initial displacements divided by $\\tdyn$. This introduces no further new characteristic scales in the initial conditions. Its effect on the evolution will be to make the transient to self-similarity slightly shorter, but it will not significantly change any of our findings or conclusions. \\item We have made a specific choice of PDF for our shuffling, given in Eq.~(\\ref{eq:pu}). We expect different choices again to modify slightly the nature of the transient, but not the self-similarity. This latter, as we have emphasized, depends only on the $k^2$ form of the PS at small $k$, which is in fact the same for any PDF with finite variance. Indeed the coefficient of the $k^2$ is just given by this variance, and the difference between PDFs will manifest themselves in modifications of the fluctuations at small scales (i.e. larger $k$). For example if the two PDF have different fourth moments, this will be reflected in a different coefficient in the $k^4$ correction to the small $k$ PS. Just as in the case of velocities, there is a natural choice if one wishes to maximize the analogy with cosmological simulations: a simple Gaussian PDF which is what is used in this context. In fact this choice is also natural from another point of view, as we will explain in detail in a forthcoming article\\cite{baertschiger_06_prep}: when one considers constructing new particle distributions by a simple ``coarse-graining'' on some scale, the SL with Gaussian PDF, due to the Central Limit Theorem, has the property of being the unique one which is invariant under such a coarse-graining. \\item We have reported in this paper simulations in which the softening $\\varepsilon$ has been kept fixed (in our chosen length units). We have mentioned that we have checked that our results for clustering amplitudes above this scale are robust to the use of significantly smaller values of $\\varepsilon$. A more extensive and systematic study of the role of this parameter would, however, be of interest, specifically with the goal of understanding in detail how the clustering properties are modified by it at small scales. \\end{itemize}" }, "0607/astro-ph0607254_arXiv.txt": { "abstract": "{ We apply the linear filter for the weak-lensing signal of dark-matter halos developed in Maturi et al. (2005) to the cosmic-shear data extracted from the Garching-Bonn-Deep-Survey (GaBoDS).} { We wish to search for dark-matter halos through weak-lensing signatures which are significantly above the random and systematic noise level caused by intervening large-scale structures.} { We employ a linear matched filter which maximises the signal-to-noise ratio by minimising the number of spurious detections caused by the superposition of large-scale structures (LSS). This is achieved by suppressing those spatial frequencies dominated by the LSS contamination.} { We confirm the improved stability and reliability of the detections achieved with our new filter compared to the commonly-used aperture mass (Schneider, 1996; Schneider et al., 1998) and to the aperture mass based on the shear profile expected for NFW haloes (see e.g.~Schirmer et al., 2004; Hennawi \\& Spergel, 2005). Schirmer et al.~(2006) achieved results comparable to our filter, but probably only because of the low average redshift of the background sources in GaBoDS, which keeps the LSS contamination low. For deeper data, the difference will be more important, as shown by Maturi et al. (2005).} { We detect fourteen halos on about eighteen square degrees selected from the survey. Five are known clusters, two are associated with over-densities of galaxies visible in the GaBoDS image, and seven have no known optical or X-ray counterparts. } ", "introduction": "Claims have been made in the literature that dark mass concentrations were significantly detected through their weak-lensing signal \\citep[see e.g.][]{ER00.1,UM00.1,MI02.1,ER03.1}. If confirmed, these detections were extremely exciting because they showed that large, dark mass concentrations could exist which for some reason failed to emit light, either as stellar light from galaxies or X-ray emission from hot intracluster plasma. Since the baryon fraction in clusters should faithfully reflect the cosmic mixture of baryonic and dark mass \\citep{ET03.1,ET06.1}, detections of truly dark cluster-sized halos could shed doubt on our understanding of the formation of non-linear cosmic structures. Obviously, gravitational lensing is the only method able to detect dark mass concentrations. This strength is also its weakness: all density inhomogeneities projected along the line-of-sight cause gravitational light deflection, thus any detection of a lensing signal possibly due to a dark-matter halo is affected to some degree by the large-scale structures projected into the observed field. Superpositions of large-scale structures can create signals mimicking dark-matter halos. Commonly used methods for detecting halos through weak lensing, in particular the aperture mass \\citep{SC96.2,SC98.2}, are optimised for measuring the total projected mass enclosed within an aperture, and for suppressing the cross-correlation of the signal in neighbouring apertures. A clean separation between the signals of halos and large-scale structures is strictly impossible because the large-scale structure can be considered as composed of halos with a broad mass spectrum. Unambiguous noise suppression can thus not be achieved. However, we have argued in \\citep{MAT04.2} that the weak-lensing power spectrum obtained from the \\emph{linearly} evolved dark-matter power spectrum can reasonably be considered as a noise contribution against which the weak-lensing signal of halos can be filtered. The underlying assumption is that those halos which we are searching for do in fact create the non-linear power spectrum. Using simulations, we demonstrated that the linear filter which follows uniquely from this noise model and the assumption that dark-matter halos on average have an NFW density profile does indeed perform as expected \\citep{MAT04.2}. It suppresses the noise from large-scale structures substantially, thus considerably reducing the number of spurious detections, and is much less sensitive than the aperture mass against changes in the filter scale. Removing the LSS contamination is less important for relatively shallow observations, i.e.~for average source redshifts $z_\\mathrm{s}\\lesssim1$, but it is fundamental for deep observations in which the integrated contribution of the matter along the line-of-sight becomes non-negligible. We now apply this filter to the weak-lensing data obtained from the Garching-Bonn-Deep-Survey \\citep[hereafter \\textit{GaBoDS}, see][]{SC03.2}. The purpose of this study is two-fold. First, we wish to compare the halo detections with the new filter and with the aperture mass, and second, we want to assess the reliability of the halo detections and their stability against changes in the filter scale. We emphasise that we do not want to devaluate the aperture mass. It was shown to have many desirable properties for the measurement of cosmic shear. However, its construction for a different purpose than halo detection motivates us to search for an alternative measure. Ultimately, our filter can be seen as a variant of the aperture mass based on a considerably narrower angular weight function, which is characterised by the assumed halo density profile and the noise model. Throughout, we assume that dark matter halos have density profiles of NFW shape and use the angular scale radius of the profile as a free parameter. Moreover, we compute the weak-lensing power spectrum of \\emph{linear} large-scale structures adopting a $\\Lambda$CDM cosmological model, with $\\Omega_0=0.3$, $\\Omega_\\Lambda=0.7$, $\\sigma_8=0.9$ and $h=0.7$. The paper is structured as follows: Sect.~2 briefly summarises the survey characteristics; Sects.~3 and 4 describe the aperture mass and our optimal filter, respectively; and Sect.~5 presents the analysis of the data sample. Our conclusions are presented in Sect.~6. ", "conclusions": "} The application of our filter to the GaBoDS weak-lensing survey confirms the expectations on its performance raised by \\cite{MAT04.2} based on numerical simulations. Our filter is more stable than the aperture mass against changes of the filter size, thus considerably simplifying the interpretation of data. It has better statistical properties compared to the aperture mass given by Eq.~(\\ref{eq:ap_pol}), yielding more reliable results. Although our filter and the optimised aperture mass given by Eq.~(\\ref{eq:ap_tanh}) perform comparably on the GaBoDS data, the LSS suppression characterising our filter will be primarily important for deeper observations, for which the average background-galaxy redshift will be larger than in GaBoDS and the LSS will not be negligible. We measured in our data sample a large contamination from residual $B$-modes, but a statistical analysis of the noise properties allowed us to define a criterion to select a sample of reliable detections. We emphasise that a data reduction procedure which minimises the residual $B$-mode is of fundamental importance for the detection of new cluster-sized dark-matter halos. We rejected all detections below the thresholds even if associated with know clusters, because we wish to search blindly for candidate clusters based exclusively on weak-lensing data. On the $19.6$ square degrees covered by the GaBoDS data, we found $14$ detections with a sufficiently high signal-to-noise ratio, $5$ of which are known clusters, 4 are associated with concentrations of galaxies visible in our data, and 5 detections are not associated with any visible concentration of galaxies. Deep optical and $X$-ray follow-ups of the $9$ unknown detections should be performed to clarify their nature." }, "0607/astro-ph0607581_arXiv.txt": { "abstract": "We describe further observations of QSO J0906+6930, a z=5.48 blazar likely to be detected in $\\gamma$-rays. New radio and X-ray data place significant constraints on any kpc-scale extension of the VLBA-detected jet. Improved optical spectroscopy detects absorption from an intervening galaxy at z=1.849 and raise the possibility that this distant, bright source is lensed. We combine the new data into an improved SED for the blazar core and comment on the Compton keV-GeV flux component. ", "introduction": "Q0906+6930 (=GB6 J0906+6930=SRM J090630.74+693030.8) was discovered in a survey of radio-bright, flat spectrum sources chosen to be like the EGRET blazars \\citep{srm03,set05}. Follow-up observations confirmed a large $z\\approx 5.5$ redshift and found evidence for a very compact pc-scale jet with the VLBA \\citep{ret04}. The source is quite radio-loud with an $R = f_{\\rm 5\\,GHz}/f_{\\rm 440nm}({\\rm rest}) \\approx 500-1000$, supporting the blazar interpretation. These data showed several peculiar properties. For example, the radio spectrum appeared to steepen above 10\\,GHz, while the jet component still appeared to be inverted. Also, the source shows a relatively large $1350$\\AA\\, continuum flux and large kinematic widths for the emission lines. These suggest a large $\\sim 2 \\times 10^9M_\\odot$ black hole mass [note that an error in the $\\lambda F_\\lambda$ luminosity quoted in \\citet{ret04} implied $M\\ge 10^{10}M_\\odot$]. While this source was, at best, a weak background enhancement in the {\\it EGRET} survey, the prospects for detection with GLAST seem quite strong. This is particularly interesting as observations of a cut-off in the blazar spectrum above $\\sim 10$\\,GeV can be used to probe absorption by light produced at the peak of star formation \\citep{crr04}. We report here on further observations of Q0906+6930, which support its identification as a high-z blazar, test the nature of the jet component and probe it's status as a high mass, high luminosity source. ", "conclusions": "The new data generally confirm the SED picture described in \\citet{ret04} and show that at least the X-ray Compton component is present in this blazar. The lack of any resolved kpc-scale jet component suggests that all emission observed is from the core (or a compact pc-scale jet). Also, this does not support the \\citet{sch02} picture where large-scale X-ray jets are produced by up-scatter of the CMB. In this model, the $\\propto (1+z)^4$ increase of the CMB energy density counters the decreasing jet surface brightness; at $z>4$ jet emission is thus expected to exceed core emission. With $\\le 1$\\% of the core X-ray flux in a resolvable jets, our results extend those of \\citet{lopet06}, suggesting core photons, rather than CMB photons dominate the flux up-scattered as Compton X-ray emission. One peculiarity noted in \\citet{ret04} was the relatively large optical and radio flux; we noted the possibility of lensing of the QSO core. This should be re-visited, since the strong intervening metal line absorbers suggest galactic-scale masses along the line-of-sight. The most important is the MgII system at z=1.849, which at $W_0=1.9$\\AA, likely represents a classical $\\sim L_\\ast$ spiral galaxy \\citep{chet05} which should be a Lyman Limit System (LSS) absorber, within an impact parameter $$ R = 38 h^{-1}(L_K/L_K^\\ast)^{0.15} {\\rm kpc}, $$ for an $L_\\ast$ galaxy at the absorber redshift, i.e. within $6.2^{\\prime\\prime}$ for our assumed cosmology. Our limited direct imaging does not detect such a galaxy, but the results are not very constraining. At this $z$, the cosmic age is $\\sim 3.6$\\,Gyr and an $L_\\ast$ Sa formed at z=5 in our cosmology would have $r\\sim 23.9$ for a \\citet{bc93} model, i.e. about $2\\times$ fainter than our imaging $r\\sim 23.3$ limit. Sbc types would be even fainter. Can such a galaxy lens our blazar? The expected image angular scale for a lens mass $M$ is $$ \\theta_0 = \\left [ 4 D_{LS}G\\,M/(D_S D_L c^2)\\right ]^{1/2} $$ where one uses angular diameter distances and $D_{LS} = D_s-(1+z_L)D_L/(1+z_s)$. The geometry is favorable for lensing and the lens scale is then $1.35 \\times 10^{-6}M^{1/2}$arcsec with $M$ in $M_\\odot$. We see no evidence for a double image of the blazar on arcsecond scales, either in the optical or the radio, so this precludes classical macro lensing by an $L_\\ast$ galaxy. Indeed, the core is unresolved in our VLBI maps, where the half-power beam width at 43\\,GHz was 0.55$\\times 0.30$\\,mas, so we can place an upper limit on an effective lens mass of $M\\le 1.4 \\times 10^5 M_\\odot$. The lack of obvious variability on year time-scales suggests a not very constraining lower mass limit of $M\\ge 10^{-3} v_{300}^2 M_\\odot$ for a lens galaxy velocity dispersion of $\\sim 300\\,v_{300}$\\,km/s. We may infer a slightly tighter limit if we note that the large BLR equivalent widths imply that the broad line region must be lensed along with the core. Naively applying the continuum BLR radius correlation of \\citet{kasp05} to our optical continuum luminosity $\\lambda F_\\lambda |_{5100\\AA} \\approx 3 \\times 10^{46} {\\rm erg/s}$ we conclude $$ R_{BLR} \\approx 5.8 \\times 10^{16} \\left ( \\lambda F_\\lambda |_{5100\\AA}/10^{44}{\\rm erg/s} \\right )^{0.69} \\approx 2.9 \\times 10^{18}{\\rm cm}. $$ This in turn implies a minimum lens mass $M\\ge 10^{-2}M_\\odot$. We infer that macro-lensing does not amplify Q0906+6930, but that micro- or milli-lensing are possible. On the other hand, it should be noted that the probability of an unrelated intervening galaxy is not small. Since $dN/dz(W_0>1.9\\AA) \\approx 0.13$ over the redshift range covered by our G3 spectrum \\citep{net05}, we expect 0.06 such absorbers. Detection of one MgII systems is thus not very improbable and so we cannot infer from this that lensing amplification has brightened Q0906+6930. We conclude by re-iterating that our new SED measurements support the picture of Q0906+6930 as a blazar with a bright $\\gamma$-ray Compton component emitting at z=5.48. This Compton radiation must traverse the peak of star formation at $z\\approx 2-3$, where optical/UV emission can attenuate the $>$10\\,GeV photons, just as the IR background can attenuate the TeV emission of blazars at lower $z$, e.g. \\citet{dka05} and references therein. Of course, there is also attenuation from the host photon field and, in principle, from the intervening z=1.89 galaxy. These contributions will make it difficult to extract constraints on the extragalactic background photon field from this (or any one) object. However, statistical studies of high red-shift blazars with GLAST \\citep{crr04}, should still be able to extract global constraints on the extragalactic background light and its evolution." }, "0607/astro-ph0607062_arXiv.txt": { "abstract": "We implement the Elliptical Gauss-Laguerre (EGL) galaxy-shape measurement method proposed by Bernstein \\& Jarvis (2002) [\\citet{BJ02}] and quantify the shear recovery accuracy in weak lensing analysis. This method uses a deconvolution fitting scheme to remove the effects of the point-spread function (PSF). The test simulates $>10^7$ noisy galaxy images convolved with anisotropic PSFs, and attempts to recover an input shear. The tests are designed to be immune to shape noise, selection biases, and crowding. The systematic error in shear recovery is divided into two classes, calibration (multiplicative) and additive, with the latter arising from PSF anisotropy. At S/N $>50$, the deconvolution method measures the galaxy shape and input shear to $\\sim1\\%$ multiplicative accuracy, and suppresses $>99\\%$ of the PSF anisotropy. These systematic errors increase to $\\sim4\\%$ for the worst conditions, with poorly resolved galaxies at S/N $\\simeq20$. The EGL weak lensing analysis has the best demonstrated accuracy to date, sufficient for the next generation of weak lensing surveys. ", "introduction": "Weak gravitational lensing, the shearing of galaxy images by gravitational bending of light, is an effective tool to probe the large-scale matter distribution of the universe. It is also a means to measure the cosmological parameters by comparing observation to numerical simulations of large scale structure growth~\\citep{bartelmann01}. There are many weak lensing (WL) surveys underway to obtain the cosmological parameters to higher precision, and in particular to probe the evolution of the dark energy by observing its effects on the evolution of matter distribution (DLS\\footnote{Deep Lens Survey: {\\tt dls.bell-labs.com/}}, CFHTLS\\footnote{Canada France Hawaii Telescope Legacy Survey: {\\tt www.cfht.hawaii.edu/Science/CFHLS/}}). The WL signal is very subtle, however; it is necessary to measure these small distortions (typical shear $\\gamma\\sim 1\\%$) in the presence of optical distortions and the asymmetric point-spread-function (PSF) of real-life imaging. The level of systematic error in the WL measurement methods are currently above the statistical accuracy expected from future wide and deep WL surveys (Pan-STARRS\\footnote{Panoramic Survey Telescope \\& Rapid Response System: {\\tt pan-starrs.ifa.hawaii.edu/}}, SNAP\\footnote{Supernova / Acceleration Probe: {\\tt snap.lbl.gov/}}, LSST\\footnote{Large Synoptic Survey Telescope: \\tt www.lsst.org/}, SKA\\footnote{Square Kilometre Array: \\tt www.skatelescope.org/}). Because there are no ``standard shear'' lenses on the sky, shear-measurement techniques are tested by applying them to artificial galaxy images and seeing if one can correctly extract a shear applied to the simulation. In most cases, the recovered shear can be written as $\\gamma_{\\rm out} = m \\gamma_{\\rm in} + c$. Departures from the ideal $m=1$ we will term ``calibration'' or ``multiplicative'' errors and quote as percentages. Deviations from the ideal $c=0$ can result from uncorrected asymmetries in the PSF and optics, and will be termed ``additive errors'' or ``incomplete PSF suppression.'' Such tests of the most widely applied analysis method \\citep{KSB}[KSB], find $m=$0.8--0.9, but this coefficient is implementation dependent~\\citep{erben01,bacon01}, and depends upon the characteristics of the simulated galaxies. Hirata \\& Seljak (2003) [\\citet{Hirata03}] demonstrate that various PSF-correction methods can produce shear measurements miscalibrated by a few \\% to 20\\% or more. Heymans et al.\\@ (2005) [Shear TEsting Programme, (\\citet{STEP})] present testing of many existing shear-measurement pipelines using a common ensemble of sheared simulated images. These methods show a median calibration error of 7\\%, although some (the BJ02 rounding kernel method, an implementation of a KSB method, as well as the one described in this paper) show no calibration error, to within the $\\sigma_m\\approx1\\%$ noise level of the first STEP tests. Although the statistical accuracy in past surveys was comparable to the 7\\% systematics, it is expected to be well below 1\\% in future surveys. Hence, understanding and eliminating the WL systematic errors require the most urgent attention today. In this paper, we implement the elliptical Gauss-Laguerre (EGL) deconvolution method as described in BJ02, and subject it to a series of tests designed to be more stringent than any previous test of WL measurements. The deconvolution method is distinct from the \\citet{Jarvis03} method, also described BJ02, in which the anisotropic PSF effects are removed using a ``rounding kernel'' instead. WL testing regimes are of two types: in end-to-end tests ({\\it e.g.} STEP), one produces simulated sky images with a full population of stars and galaxies, analyzes them with the same pipeline as one would real data, then checks the output shear for veracity. We perform here more of a dissection, in which we analyze the performance of the method one galaxy type at a time, and vary the parameters of the galaxy and PSF images to determine which, if any, conditions cause the measurement to fail. While lacking the realism of an end-to-end test, this allows us to isolate and fix weaknesses. If we can demonstrate that the method succeeds under a set of conditions that will circumscribe those found on the real sky, then we can have confidence that our method is reliable, whereas end-to-end testing is reliable only to the extent that the simulated sky reproduces the characteristics of the real sky. We investigate here the performance of our EGL method across the range of noise levels, degree of resolution by the PSF, pixel sampling rates, galaxy ellipticity, and PSF ellipticity, using both highly symmetric and asymmetric galaxy shapes. We test not only the accuracy of shear recovery, but also the accuracy of the shear uncertainty estimates. The EGL method is further elaborated in \\S2, while the implementation, \\glfit, is detailed in \\S3. The shear accuracy test procedure is described in \\S4. The conditions under which the shape measurement succeeds, and the accuracy of its estimates of shear, are presented in \\S5. Previous dissection tests include \\citet{Hirata03} and \\citet{Kuijken}. The former studies the performance of several methodologies on varied galaxy and PSF shapes/sizes in the absence of noise. The latter study verified its ``polar shapelet'' method to better than 1\\% calibration accuracy. In \\S6 and \\S7 we conclude with comparisons to other shape-measurement methodologies and tests, and draw inferences for future surveys. ", "conclusions": "The elliptical Gauss-Laguerre decomposition is one of the most stringently tested methods to characterize shapes of galaxies. With the EGL decomposition, shapes are measured without the need for empirical shear/smear polarizabilities, and PSFs are removed by deconvolution. The shear, obtained from averaging an isotropic ensemble of galaxy shapes, is highly accurate due to the definition of shape as shear. We have demonstrated that the EGL method allows shear recovery of unprecedented accuracy, and quantified its degradation due to PSF, truncation of the EGL decomposition, image noise, and sampling~rate/resolution. However, the current work is limited to the extraction of shapes; further work, including the full pipeline analysis, is required for attaining 0.1\\% accuracy in shear estimation." }, "0607/astro-ph0607548_arXiv.txt": { "abstract": "{} {We present results from deep $\\gamma$-ray observations of the Galactic pulsar wind nebula \\object{HESS\\,J1825--137}\\ performed with the H.E.S.S.\\ array.} { Detailed morphological and spatially resolved spectral studies reveal the very high-energy (VHE) $\\gamma$-ray aspects of this object with unprecedented precision. } { We confirm previous results obtained in a survey of the Galactic Plane in 2004. The $\\gamma$-ray emission extends asymmetrically to the south and south-west of the energetic pulsar \\object{PSR\\,J1826--1334}, that is thought to power the pulsar wind nebula. The differential $\\gamma$-ray spectrum of the whole emission region is measured over more than two orders of magnitude, from 270~GeV to 35~TeV, and shows indications for a deviation from a pure power law. Spectra have also been determined for spatially separated regions of HESS\\,J1825--137. The photon indices from a power-law fit in the different regions show a softening of the spectrum with increasing distance from the pulsar and therefore an energy dependent morphology.} { This is the first time that an energy dependent morphology has been detected in the VHE $\\gamma$-ray regime. The VHE $\\gamma$-ray emission of HESS\\,J1825--137 is phenomenologically discussed in the scenario where the $\\gamma$-rays are produced by VHE electrons via Inverse Compton scattering. The high $\\gamma$-ray luminosity of the source cannot be explained on the basis of constant spin-down power of the pulsar and requires higher injection power in past.} \\authorrunning{F. Aharonian et al.} \\titlerunning{Energy dependent $\\gamma$-ray morphology in HESS\\,J1825--137} ", "introduction": "A growing number of extended objects that seem to be associated with energetic pulsars are detected in the Galactic Plane by their very high-energy (VHE, energy $E_{\\gamma} \\gtrsim 100$~GeV) $\\gamma$-ray emission. Latest results on this class of objects include emission from \\object{MSH--15--5\\emph{2}} (\\object{HESS\\,J1514--591})~\\citep{HESSMSH} and \\object{Vela~X} (\\object{HESS\\,J0835--455})~\\citep{HESSVelaX}, and the two sources in the \\object{Kookaburra} region (\\object{HESS\\,J1420--607} and \\object{HESS\\,J1418--609}) as described in~\\citet{HESSKookaburra}. If these associations are correct, then these objects are pulsar wind nebulae (PWN), objects generally thought to be powered by a relativistic particle outflow (electrons and positrons) from a central source. The central source -- a pulsar -- is a rapidly rotating neutron star generated in a supernova event. The relativistic wind of particles flows freely out until its pressure is balanced by that of the surrounding medium. In that region the wind decelerates and a standing termination shock is formed at which particles are accelerated~\\citep{KennelCoroniti, AhaAtoKif97}. The existence of electrons accelerated to energies $>100$~TeV in such PWN has been established by X-ray observations of synchrotron emission, e.g. in the Crab nebula~\\citep{ChandraCrab}. VHE $\\gamma$-rays can be generated in PWN from the high-energy electrons by non-thermal bremsstrahlung or inverse Compton (IC) scattering on photon target fields, such as the cosmic microwave background (CMBR) or star-light photons. \\begin{figure*} \\centering \\includegraphics[width=0.7\\textwidth]{fig1.eps} \\caption{Acceptance-corrected smoothed excess map (smoothing radius 2.5\\arcmin) of the $2.7\\degr\\ \\times 2.7\\degr$ field of view surrounding HESS\\,J1825--137. The linear colour scale is in units of integrated excess counts within the smoothing radius of 2.5\\arcmin. The excess has been derived from a model of the system acceptance as described in the text. The inset in the bottom left corner shows the PSF of the dataset (smoothed in the same way as the excess map with the black dashed circle denoting the smoothing radius). The dashed black and white contours are linearly spaced and denote the $5 \\sigma$, $10 \\sigma$ and $15 \\sigma$ significance levels (the $5 \\sigma$ contour being the outermost one), determined with a $\\theta = 0.1\\degr$ radius cut. The best fit position of HESS\\,J1825--137 is marked with a black square, the best extension and position angle by a black ellipse (see text). The dotted white contour shows the 95\\% positional confidence contour of the unidentified EGRET source \\object{3EG\\,J1826--1302}. The position of the pulsar PSR\\,J1826--1334 is marked by a white triangle. The bright point-source to the south of HESS\\,J1825--137 is the microquasar \\object{LS\\,5039}\\ (\\object{HESS\\,J1826--148})~\\citep{HESS5039}. The colour scale for this source is truncated in this Figure. The Galactic plane is shown as a white dashed line. Some indication for an additional emission region to the north of the pulsar is seen.} \\label{fig::skymap} \\end{figure*} One such object, HESS\\,J1825--137, has been detected by the High Energy Stereoscopic System (H.E.S.S.) in a survey of the inner Galaxy~\\citep{HESSScan, HESSScanII} and has subsequently been associated with the X-ray PWN G18.0--0.7 surrounding the energetic pulsar PSR\\,J1826--1334~\\citep{HESS1825}. This pulsar PSR\\,J1826--1334 (also known as PSR\\,B1823--13) was detected in the Jodrell Bank 20~cm radio survey~\\citep{RadioDetection} and is among the 20 most energetic pulsars in the current ATNF catalogue (spin down power ${\\dot E} = 3 \\times 10^{36}$ erg/s). The distance of PSR\\,J1826--1334 as measured from the dispersion of the radio pulses is $3.9\\pm0.4$~kpc~\\citep{Cordes_Lazio}. The radio detection further revealed characteristic properties of the system that are similar to those of the well studied Vela pulsar, namely a pulse period of 101\\,ms and a characteristic age of 21.4~kyears (derived by $\\tau = P/2 \\dot{P}$). This age renders PSR\\,J1826--1334 one of the 40 youngest pulsars detected so far~\\citep{ATNF}, and due to this, deep radio observations were performed to find emission associated with the remnant of the Supernova explosion that gave rise to the pulsar. However, deep VLA observations of the 20\\arcmin\\, surrounding the pulsar have failed to detect this Supernova remnant (SNR)~\\citep{VLABraun}. Initial observations of the region in X-rays with ROSAT~\\citep{ROSAT1825} revealed a point source surrounded by an elongated diffuse region of size $\\sim$5\\arcmin. The X-ray emission region was subsequentially observed with the ASCA instrument and the data confirmed the picture of a compact object surrounded by an extended emission region~\\citep{ASCA1825}. While ROSAT data did not provide sufficient statistics, ASCA data lacked the spatial resolution to resolve and interpret the sources in this region. The situation was clarified in an XMM-Newton observation in which high angular resolution observations revealed a compact core of extension 30\\arcsec\\, surrounding PSR\\,J1826--1334, and furthermore an asymmetric diffuse nebula extending at least 5\\arcmin\\, to the south of the pulsar~\\citep{XMM1825}. In this XMM-Newton dataset the signal to noise ratio deteriorates rapidly at offsets larger than 5\\arcmin\\, and for this reason the XMM data cannot place useful constraints on the presence of a faint shell of emission at larger radii as might be produced by an associated SNR. The extended asymmetric structure was attributed to synchrotron emission from the PWN of PSR\\,J1826--1334~\\citep{XMM1825}. The X-ray spectrum in the diffuse emission region follows a power law with photon index $\\Gamma \\sim$2.3 and an X-ray luminosity between 0.5 and 10~keV of $L_x \\sim3 \\times 10^{33}$ erg s$^{-1}$ compared to the X-ray spectrum for the compact core following a power law with $\\Gamma \\sim$1.6 and $L_x \\sim9 \\times 10^{32}$ erg s$^{-1}$ (these luminosities are derived assuming a distance of 4~kpc). \\citet{XMM1825} discussed various scenarios to explain the asymmetry and offset morphology of the PWN G18.0--0.7. The most likely explanation seems to be that a symmetric expansion of the PWN is prevented by dense material to the north of the pulsar which shifts the whole emission to the south. Asymmetric reverse shock interactions of this kind have originally been proposed to explain the offset morphology of the Vela~X PWN based on hydro-dynamical simulations by~\\citet{Blondin}. Indeed recent analyses of CO data show dense material surrounding PSR\\,J1826--1334 (at a distance of 4~kpc) to the north and northeast~\\citep{AnneCO1825}, supporting this picture. It is interesting to note, that H.E.S.S.\\ has now detected offset morphologies from both G18.0--0.7 and Vela~X~\\citep{HESSVelaX}, confirming the existence of a class of at least two offset PWN implied by X-ray observations~\\citep{XMM1825}. Whereas X-rays probe a combination of the thermal and ultrarelativistic components, which could have been mixed at the time when the asymmetric reverse shock interaction took place, the H.E.S.S.\\ results are important in determining the offset morphology of the ultrarelativistic component alone. Based on its proximity and energetics, the pulsar PSR\\,J1826--1334 has been proposed to be associated with the unidentified EGRET source 3EG\\,J1826--1302~\\citep{EGRETCat}. This EGRET source exhibits a hard power law of photon index $2.0 \\pm 0.11$ with no indication of a cut-off. The pulsar lies south of the centre of gravity of the EGRET position and is marginally enclosed in the 95\\% confidence contour (see Fig.~\\ref{fig::skymap}). It has been shown~\\citep{ZhangCheng} that an association between PSR\\,J1826--1334 and 3EG\\,J1826--1302 is plausible based on the pulsar properties (such as pulsar period and magnetic field derived in the frame of an outer gap model), and that the observed $\\gamma$-ray spectrum can be fit to this model. Although an unpulsed excess from EGRET has been reported with a significance of $9 \\sigma$ ~\\citep{NelEGRET}, a significant periodicity could not be established. Additionally an ASCA X-ray source possibly connected to the EGRET data above 1\\,GeV~\\citep{ASCAGeVRoberts} was found in this region. Recently, \\citet{NolanEGRET} reassessed the variability of the EGRET source and found a weak variability, which led the authors to consider the source finally as a PWN candidate in the EGRET high-energy $\\gamma$-ray energy range above 100~MeV. Here we report on re-observations of the VHE $\\gamma$-ray source HESS\\,J1825--137 and the region surrounding PSR\\,J1826--1334 performed with H.E.S.S.\\ in 2005. H.E.S.S.\\ consists of four imaging atmospheric Cherenkov telescopes and detects the faint Cherenkov light from $\\gamma$-ray induced air showers in the atmosphere above an energy threshold of 100~GeV up to several tens of TeV. Each telescope is equipped with a mirror area of $107$~m$^2$~\\citep{HESSOptics} and a 960 photo-multiplier camera for the detection of the faint Cherenkov light. The telescopes are operated in a coincidence mode in which at least two telescopes must have triggered in each event~\\citep{HESSTrigger}. The H.E.S.S.\\ system has a point source sensitivity above 100~GeV of $<2.0\\times\\,10^{-13}$ cm$^{-2}$s$^{-1}$ (1\\% of the flux from the Crab nebula) for a $5 \\sigma$ detection in a 25 hour observation. The system is located in the Khomas Highland of Namibia~\\citep{HESS} and began operation in December 2003. ", "conclusions": "" }, "0607/astro-ph0607312_arXiv.txt": { "abstract": "Cosmic microwave background (CMB) polarization observations will require superb control of systematic errors in order to achieve their full scientific potential, particularly in the case of attempts to detect the $B$ modes that may provide a window on inflation. Interferometry may be a promising way to achieve these goals. This paper presents a formalism for characterizing the effects of a variety of systematic errors on interferometric CMB polarization observations, with particular emphasis on estimates of the $B$-mode power spectrum. The most severe errors are those that couple the temperature anisotropy signal to polarization; such errors include cross-talk within detectors, misalignment of polarizers, and cross-polarization. In a $B$ mode experiment, the next most serious category of errors are those that mix $E$ and $B$ modes, such as gain fluctuations, pointing errors, and beam shape errors. The paper also indicates which sources of error may cause circular polarization (e.g., from foregrounds) to contaminate the cosmologically interesting linear polarization channels, and conversely whether monitoring of the circular polarization channels may yield useful information about the errors themselves. For all the sources of error considered, estimates of the level of control that will be required for both $E$ and $B$ mode experiments are provided. Both experiments that interfere linear polarizations and those that interfere circular polarizations are considered. The fact that circular experiments simultaneously measure both linear polarization Stokes parameters in each baseline mitigates some sources of error. ", "introduction": "Cosmic microwave background (CMB) polarimetry is one of the most exciting frontiers in cosmology. CMB polarization has already been detected \\cite{kovac,wmappol,readhead,leitch,barkats,wmap3pol}, and we may expect future instruments to characterize the polarization signal in much greater detail (e.g., \\cite{korotkov}). In the near future, CMB polarization data are expected to refine estimates of cosmological parameters \\cite{kinney}, probe the ionization history of the Universe \\cite{zal97} and the details of recombination \\cite{peebles}, and measure gravitational lensing due to large-scale structure \\cite{zalsellens}. Most exciting of all, polarization maps may provide a direct probe of an inflationary epoch in the extremely early Universe by detecting the signature of primordial gravitational radiation \\cite{zalsel,selzal,kkslett,kks}. A crucial insight into the analysis of CMB polarization data is the fact that any CMB polarization map can be divided into two components, a scalar component, traditionally denoted $E$, and a pseudoscalar component called $B$. The CMB is weakly polarized, meaning that both of these components are much smaller than the unpolarized (temperature) anisotropy. Furthermore, the $B$ component is expected to be much weaker than $E$, since scalar density perturbations produce only $E$ to linear order \\cite{zalsel,selzal,kkslett,kks}. (See Figure \\ref{fig:spectra}.) Experiments to date have detected only the $E$ component. In the future, the search for the weaker $B$-type polarization will be a high priority, as the $B$ modes may contain the imprint of gravitational waves produced during inflation. \\begin{figure*} \\includegraphics[width=3in]{cls.eps} \\includegraphics[width=3in]{ratios.eps} \\caption{Power spectra for temperature anisotropy (T), $TE$ cross-correlation (X, absolute value plotted), $E$-type polarization, and $B$-type polarization. The best-fit parameters from the three-year WMAP data were used \\cite{wmap3imp} with a tensor-to-scalar ratio $T/S=0.01$. The right panel shows the ratios of the power spectra.} \\label{fig:spectra} \\end{figure*} Characterization of CMB polarization requires both very low noise and exquisite control of systematic errors. In particular, some sources of systematic error may cause the polarization signal to be contaminated by the much larger unpolarized anisotropy, while others mix the $E$ and $B$ components. As efforts to design $B$-mode experiments intensify, it is important to consider carefully the susceptibility of different designs to various kinds of error. Hu \\etal~\\cite{HHZ} have provided a detailed framework for performing such an analysis in the context of an imaging experiment. For interferometric measurements, the issues are somewhat different. The purpose of this paper is to forecast the effects of a variety of systematic errors on interferometric measurements. Interferometric methods have played an important role in measurements of CMB anisotropy and polarization. Pioneering attempts to detect CMB anisotropy with interferometers are described in \\cite{MarPar} and \\cite{Sub}. Several groups have successfully detected primary CMB anisotropies \\citep{CAT1,CAT2,DASIT,CBIT,VSA} and polarization \\citep{readhead,leitch} using interferometers. The formalism for analyzing CMB data from interferometers has been developed by a number of authors \\cite{HobLasJon,HobMag,WCDH,HobMais,Mye,BW} as well as in the experimental papers cited above. In any data set that fails to cover the entire sky, it is impossible to separate the $E$ and $B$ components perfectly \\cite{LCT,bunn,bunnerratum,bunnetal}. The operation of separating a polarization map into $E$ and $B$ components is nonlocal when the map is viewed in real space, but in Fourier space or spherical harmonic space, it can be done locally (mode by mode). Since interferometric data sample the sky in the Fourier domain, $E$-$B$ separation may be cleaner for interferometric data than for maps made with single-dish instruments \\cite{Parketal,ParkNg}. As we will see, a variety of systematic errors in interferometers can be modeled via Jones matrices \\cite{tinbergen,heiles,HHZ} and by deviations of the antenna patterns (including cross-polar contributions) from assumed ideal forms. We will assume that each of these errors can be characterized by small unknown parameters, such as gain fluctuations, cross-talk between detectors, pointing errors, etc. We will first calculate the effect of each error on the measured visibilities. We will then provide a method of quantifying the effects of each of these errors on estimates of the polarization power spectra $C_l^E,C_l^B$ that can be obtained from a hypothetical data set. This paper has the following structure. Section \\ref{sec:formalism} presents the mathematical formalism we will use to describe interferometric visibilities for polarization data. Section \\ref{sec:viserr} presents the effects of various systematic errors on the visibilities extracted from a hypothetical CMB experiment. Section \\ref{sec:spectra} presents a method of forecasting errors on power spectrum estimates from errors on visibilities. Sections \\ref{sec:instresults} and \\ref{sec:beamresults} contain results showing how the error forecasts on both $E$ and $B$ power spectra depend on the parameters that characterize the various systematic errors. Section \\ref{sec:discuss} presents a discussion of the implications of these results, and a brief appendix contains a useful mathematical result. Sections \\ref{sec:spectra} through \\ref{sec:beamresults} contain quite a bit of technical detail. The particularly busy or impatient reader should note that the key ideas of Section \\ref{sec:spectra} are summarized at the beginning, and the final results of Sections \\ref{sec:instresults} and \\ref{sec:beamresults} are summarized in Section \\ref{sec:discuss} and Table \\ref{table}. ", "conclusions": "\\label{sec:discuss} This paper has presented a method of quantifying the effects of a variety of systematic errors on estimates of the CMB polarization power spectra and have applied the method to a variety of possible errors. Let us begin by summarizing these results in a more compact form. To illustrate the relative magnitudes of the various sources of error, let us consider a fiducial set of experimental parameters. Let us assume that the true power spectra in the range of multipoles probed by our experiment are in the ratio \\beq C^T:C^E:C^B = 300^2:300:1, \\eeq roughly typical for subdegree-scale experiments. Furthermore, let us assume a fiducial value of \\beq \\ssq=0.02, \\eeq which corresponds roughly to a baseline formed by a pair of antennas separated by three times the antenna diameter. Having chosen these fiducial values we can work out the effect of any particular error source. For instance, consider the effect of gain errors on a linear experiment aiming to measure $B$ polarization. The leading contribution to the error is the one that couples $EE$ to $B$, with \\beq \\kappasub{B,EE,\\gamma_2}=\\tfrac{1}{2}(\\ssq)^{1/2}=0.071. \\eeq The effect on the measurement of $C^B$ is \\beq \\frac{\\delta\\hat C^B_{\\rm rms}}{C^B}=\\kappasub{B,EE,\\gamma_2}\\gamma_2 \\frac{C^E}{C^B}=21\\gamma_2. \\eeq Say for instance that we wish systematic errors to have at most a 10\\% effect. Then $21\\gamma_2<0.1$ or $\\gamma_2<5\\times 10^{-3}$. Of course $\\gamma_2$ here represents the r.m.s.\\ value of an unknown gain fluctuation, so this should be interpreted as an estimate of the level to which gain fluctuations must be understood. Table \\ref{table} summarizes the results of such calculations for the various errors considered in this paper. A horizontal line separates instrument from sky errors. In each case, the dominant term listed is the one that involves the largest input power spectra. In cases where $\\ssq$ is small, an error term that is lower in the hierarchy may be of comparable significance to the dominant term. The table therefore lists a second contribution to each error where appropriate. This second contribution has $\\kappa$ more weakly dependent on $\\ssq$ than the dominant contribution, so for large antenna separation it may be the more important term (although for the fiducial parameters adopted here it never is). In the cases of coupling errors and cross-polarization in an $E$-mode measurement, the dominant term is independent of $\\ssq$, so there is no need to consider a second term. In all entries in the table, the coefficients are averaged over $\\alpha$ and calculated with the leading-order term in an expansion in $\\ssq$, As figures \\ref{fig:glkappa2}-\\ref{fig:xp} indicate, the latter approximation is excellent. In all cases, the error parameters should be taken as r.m.s.\\ residuals after known errors have been removed. For instance, as we noted in the previous section, sky curvature can induce cross-polarization characterized by $\\mu=\\frac{1}{2}\\sigma^2$. Presumably that effect would be known and accounted for; the parameter $\\mu$ in Table \\ref{table} represents an unknown and hence unmodeled additional component. Not surprisingly, the coupling parameters $\\varepsilon$ and cross-polarization $\\mu$ are of the greatest concern, since they couple the temperature power spectrum to polarization measurements. in particular, if we want $\\delta\\hat{C}^B/C^B$ to be, say at most 10\\%, then these parameters must be $\\varepsilon,\\mu\\lesssim 10^{-4}$. Recall that for a linear experiment the coupling parameters can be used to describe errors in the alignment of the polarizers, so a $B$ mode experiment would require alignment with a precision $\\sim 10^{-4}$ radians or $\\sim 0.3'$. For the $E$ power spectrum, on the other hand, the required tolerance is about $0.3^\\circ$. For pointing and beam shape errors, circular experiments have an advantage over linear experiments, because errors that differ between measurement of $V_Q$ and $V_U$ (parameterized by $\\delta_-,\\zeta_-$) are absent. Gain errors, on the other hand, are worse in a circular experiment. All of the errors in Table \\ref{table} are expressed as couplings between band powers. In the case of instrument errors, we have seen that the visibility errors can be expressed as linear combinations of the visibilities themselves. In other words, the Fourier-space window functions associated with the errors have exactly the same shape as the visibilities themselves. In the case of beam errors, this is not strictly true: $\\delta V_Q$, for instance, has a different window function from $V_Q$. However, for all of the errors considered in this paper, differences in Fourier space sensitivity introduced by the errors are relatively small: in all cases, the errors sample regions of Fourier space centered near ${\\bf k}=2\\pi{\\bf u}$ with widths $\\Delta u \\sim \\sigma^{-1}$, just as the visibilities themselves do. In short, the errors do not couple greatly different angular scales to each other. This contrasts with single-dish imaging experiments, in which scale-scale coupling induced by systematic errors is an important consideration \\cite{HHZ}. The results above were calculated using a simple and relatively conservative model for propagating errors from visibilities to power spectrum estimates. In a real data set, each resolution element in the Fourier plane would be sampled by multiple visibilities rather than just one pair. If we can assume that the errors in all of these visibilities are independent of each other and have ``nice'' probability distributions (particularly that the errors are centered on zero), then the estimates should be reduced by a factor of $\\sqrt N$ where $N$ is the number of independent visibility pairs $(V_Q,V_U)$ per resolution element. However, since systematic errors often do not have nice statistical properties, a more conservative approach may be warranted. Even if errors do not need to be {\\it removed} to the levels indicated here, it seems safe to say that their properties need to be {\\it studied} down to this level of precision in order to have confidence in the results. Although there is expected to be no cosmological circular polarization in the CMB, it is worthwhile to consider the effects of circular polarization in the context of systematic errors. On the one hand, various errors can couple any intrinsic circular polarization that does exist (e.g., from foregrounds) into the linear polarization channels, resulting in spurious $E$ and $B$ signals. On a more positive note, assuming that there is no intrinsic circular polarization, monitoring the circular polarization visibilities $V_V$ may provide a way to assess systematic errors. In particular, in a linear experiment coupling errors (including polarizer misalignments) lead to a contribution to $V_V$ that is correlated with the temperature anisotropy. Considering the level of control of these errors that is required in a $B$ mode experiment, such a diagnostic may prove quite useful." }, "0607/astro-ph0607130_arXiv.txt": { "abstract": "This paper presents the cosmological applications of the quasispherical Szekeres model. The quasispherical Szekeres model is an exact solution of the Einstein field equations, which represents a time--dependent mass dipole superposed on a monopole and therefore is suitable for modelling double structures such as voids and adjourning galaxy superclusters. Moreover, as the Szekeres model is an exact solution of the Einstein equations it enables tracing light and estimation of the impact of cosmic structures on light propagation. This paper presents the evolution of a void and adjourning supercluster and also reports on how the Szekeres model might be employed either for the estimation of mass of galaxies clusters or for the estimation of the luminosity distance. ", "introduction": "When performing astronomical observations one has to keep in mind that light from observed objects propagates through the Universe which is a complicated and evolving system. Although the evolution of the Universe can be neglected on small scales, it cannot be done so if distant objects, such as high--redshift galaxies, quasars or very remote supernovae are observed. Additionally, on large scales the impact of the cosmic structures on light propagation must be considered. Since in the Newtonian mechanics matter does not affect light propagation, the general relativity must be employed. This paper shows that such problems of high--redshift astronomy can be solved by employing the quasispherical Szekeres model (Szekeres 1975a). The quasispherical Szekeres model is an exact solution of the Einstein field equations, which represents a time--dependent mass dipole superposed on a monopole and therefore it is suitable for modelling double structures such as voids and adjourning galaxy superclusters. Moreover, as the Szekeres model is an exact solution of the Einstein equations it enables tracing light and estimation of the impact of cosmic structures on light propagation. The structure of this paper is as follows: Sec. 2 presents the astronomical observations of the Universe; Sec. 3 presents the theoretical approach to these data; Sec. 4 presents the Szekeres model; in Sec. 6 the evolution of a void and an adjourning cluster within the quasispherical Szekeres model is studied; the algorithm which was employed for these calculations is presented in Sec. 5; Sec. 7 presents how the quasispherical Szekeres model can be adopted to analysis of the astronomical observations. ", "conclusions": "This paper presents the cosmological application of the quasispherical Szekeres model. This model is an exact solution of the Einstein field equations, which represents a time--dependent mass dipole superposed on a monopole. Therefore, the Szekeres model is suitable for modelling double structures such as voids and adjourning galaxy superclusters. The models based on the Szekeres solution have also one more advantage --- they can be employed in solving problems of light propagation, which is impossible within the N--body simulations. The Szekeres model has a great, and so far unused, potential for applications in cosmology. It is not only suitable for studying the interactions between cosmic structures, but can also be used for estimation of the impact of matter inhomogeneities on light propagation. The Szekeres model is suitable for the investigation of following issues: --- the mass estimation based of the dynamics of galaxies, --- the luminosity of distant objects, such as high--redshift supernovae, --- the age of the Universe." }, "0607/astro-ph0607306_arXiv.txt": { "abstract": "We measure the ages, stellar masses, and star formation histories of $z\\sim 6$ galaxies, observed within 1\\,Gyr of the Big Bang. We use imaging from the {\\em Hubble Space Telescope (HST)} and the {\\em Spitzer Space Telescope} from the public ``Great Observatories Origins Deep Survey'' (GOODS), coupled with ground-based near-infrared imaging, to measure their spectral energy distributions (SEDs) from $0.8-5\\,\\mu$m, spanning the rest-frame UV and optical. From our sample of $\\approx 50$ `$i'$-drop' Lyman-break star-forming galaxies in GOODS-South with $z'_{AB}<27$, we focus on $\\approx 30$ with reliable photometric or spectroscopic redshifts. Half of these are confused with foreground sources at {\\em Spitzer} resolution, but from the 16 with clean photometry we find that a surprisingly large fraction (40\\%) have evidence for substantial Balmer/4000\\,\\AA\\ spectral breaks. This indicates the presence of old underlying stellar populations that dominate the stellar masses. For these objects, we find ages of $\\sim 200 - 700$\\,Myr, implying formation redshifts of $7\\le z_f \\le 18$, and large stellar masses in the range $\\sim 1 - 3\\times 10^{10}\\,M_{\\odot}$. Analysis of seven $i'$-drops that are undetected at 3.6\\,$\\mu$m indicates that these are younger, considerably less massive systems. We calculate that emission line contamination should not severely affect our photometry or derived results. Using SED fits out to $8\\,\\mu$m, we find little evidence for substantial intrinsic dust reddening in our sources. We use our individual galaxy results to obtain an estimate of the global stellar mass density at $z\\sim 6$. Correcting for incompleteness in our sample, we find the $z\\sim 6$ comoving stellar mass density to be $2.5\\times 10^{6}M_{\\odot}\\,{\\rm Mpc}^{-3}$. This is a lower limit, as post-starburst and dust-obscured objects, and also galaxies below our selection thresholds, are not accounted for. From our results, we are able to explore the star formation histories of our selected galaxies, and we suggest that the past global star formation rate may have been much higher than that observed at the $z\\sim 6$ epoch. The associated UV flux we infer at $z>7$ could have played a major role in reionizing the universe. ", "introduction": "\\label{sec:INTRO} Studying the stellar populations in the most distant objects known could provide a key insight into galaxy formation, potentially revealing the star formation history at even earlier epochs. The current frontier for spectroscopically-confirmed galaxies is $z\\sim 6$, with unconfirmed candidates based on photometric redshifts at perhaps even higher redshifts. The $i'$-drop technique is based on the Lyman-break technique (Steidel et al.\\ 1996) and robustly selects star-forming galaxies at $z\\sim 6$ (Stanway, Bunker \\& McMahon 2003; Bunker et al.\\ 2004; Bouwens et al.\\ 2004a; Yan \\& Windhorst 2004; Giavalisco et al.\\ 2004), when the universe was less than 1\\,Gyr old. If some of these $i'$-drop galaxies can be shown to harbour stellar populations with ages of a few hundred Myr, then this pushes their formation epoch to $z\\sim 10$. Measurements of the stellar masses of individual $z\\sim 6$ galaxies can also constrain structure formation paradigms; in a simple hierarchical model, massive galaxies assemble at later times through merging, so it might be expected that in this scenario the number density of massive evolved galaxies in the first Gyr would be low. In Eyles et al.\\ (2005), we presented {\\em Spitzer}/IRAC imaging of $z\\sim 6$ $i'$-drop galaxies with known spectroscopic redshifts, sampling the rest-frame optical at $3.6-8\\,\\mu$m. Our previous work with optical {\\em HST}/ACS data and near-infrared imaging from VLT/ISAAC and {\\em HST}/NICMOS (Stanway, McMahon \\& Bunker 2005) explored the rest-frame ultraviolet (UV) in these galaxies, which is dominated by recent or ongoing star formation. The addition of {\\em Spitzer}/IRAC imaging allowed us to fit spectral energy distributions from the multi-wavelength broad-band photometry to stellar population synthesis models. We were able to constrain the stellar masses and ages, and hence explore the preceding star formation history and the formation epochs. We concentrated primarily on two bright, well-detected objects (SBM03\\#1\\,\\&\\,\\#3 from Stanway, Bunker \\& McMahon 2003) with spectroscopic redshifts of $z\\approx 5.8$ (Stanway et al.\\ 2004a,b; Bunker et al.\\ 2003). For these two sources, we found evidence for the presence of Balmer/4000\\,\\AA\\ spectral breaks, indicating significant old stellar populations, with ages of the order of a few hundred million years (see Section~\\ref{sec:DETECTIONS}). From this we inferred formation redshifts of $z_{f}\\sim 7.5 - 13.5$, and that even more vigorous star formation had taken place prior to the time of observation. Our work was confirmed independently by Yan et al.\\ (2005) who, in their selection of galaxies, studied one of the same objects, SBM03\\#1 (see also Finlator, Dav\\'e \\& Oppenheimer 2006). The $z\\sim 6$ epoch marks a pivotal point in the history of the Universe -- the end of the reionisation era (Becker et al.\\ 2001; Kogut et al.\\ 2003; Spergel et al.\\ 2006). Hence we suggested that if objects such as these were commonplace at $z\\ge 6$, the UV flux from their initial vigorous starbursts may have played a key role in the reionisation of the Universe, supporting earlier work of Bunker et al.\\ (2004) and Egami et al.\\ (2005). The evidence found for the presence of significant Balmer/4000\\,\\AA\\ breaks in our two well-detected $i'$-drops raises the question: are these breaks rare in $z\\sim 6$ objects, or are they commonplace? In this paper, we now look to expand on our case studies of a few individual $z\\sim 6$ sources, by considering the observed optical-infrared properties of a larger population of $i'$-drop galaxies. If Balmer/4000\\,\\AA\\ breaks are found to be rare occurrences a possible scenario is that most star-forming $i'$-drops could be young ``protogalaxies'' undergoing their first bout of star formation at $z\\sim 6$. On the other hand, if these breaks are commonplace at $z\\sim 6$, it could be inferred that there is a significant population of well-established objects in place 1\\,Gyr after the Big Bang. In these systems, vigorous star formation would be required at $z\\gg 6$ in order to assemble the bulk of the stellar mass. Thus our primary goal is to obtain a robust estimate of the stellar mass density at $z\\sim 6$ from our sample of $i'$-drop galaxies; coupled with age estimates (also derived from our photometry), we may be able to uncover the preceding star formation history. This is immensely important when considering galaxy assembly scenarios and also the reionisation of the Universe. The contribution of star-forming galaxies at $z\\ge 6$ to the UV ionizing background, and hence the reionisation of the Universe, is still debated. Bunker et al.\\ (2004) discovered $\\sim 50$ $i'$-drop galaxies in the {\\em Hubble} Ultra Deep Field (HUDF) with $z'_{AB}<28.5$ ($10\\,\\sigma$). Star formation rates of these galaxies extend down to $1\\,M_{\\odot}\\,{\\rm yr}^{-1}$, derived from the rest-frame UV continuum assuming a standard Salpeter initial mass function (IMF). Bunker et al.\\ concluded that the star formation density from these observed sources would be insufficient to reionize the Universe at $z\\sim 6$, even with large escape fractions for the Lyman continuum photons. Yan \\& Windhorst (2004) independently confirmed the Bunker et al.\\ $i'$-drops in the HUDF (see Stanway \\& Bunker 2004 for a comparison), but suggested that unobserved galaxies below the detection limit could contribute much of the flux if the faint-end slope of the rest-UV luminosity function was much steeper than $\\alpha\\sim -1.6$ seen for Lyman break galaxies at lower redshifts ($z\\sim 3-4$; Steidel et al.\\ 1996). Stiavelli, Fall \\& Panagia (2004) suggested that a warmer intergalactic medium (IGM), a top-heavy IMF and substantially lower metallicity (e.g., Population III) at $z\\sim 6$ might just provide sufficient ionizing flux if the escape fraction was $f_{esc}\\sim 0.5$, much higher than observed at $z=0-4$ ($f_{esc}=0.01-0.1$; Inoue, Iwata \\& Deharveng 2006). Whilst the slumping star formation rate density from the observed $i'$-drops (Bunker et al.\\ 2004, Bouwens et al.\\ 2004a) might be insufficient to account for reionisation at $z\\sim 6$, it is possible that earlier more intense star formation played a significant role in achieving reionisation at higher redshifts. Using public imaging taken as part of the ``Great Observatories Origins Deep Survey'' (GOODS; Dickinson \\& Giavalisco 2003; Dickinson et al.\\ {\\em in prep}), our group has explored the stellar mass density and ages of $v$-drop galaxies at $z\\sim 5$ (Stark et al.\\ 2006). Comparing the inferred previous star formation histories of these $v$-drops with observations of the star formation rate density at higher redshifts, Stark et al.\\ concluded that perhaps as much as half of the star formation occurring at $z > 5$ goes unobserved. Potential reasons for this include a high contribution from low-luminosity sources, dust obscuration and/or a yet to be observed phase of star formation at very high redshift ($z\\gg 6$). Recently, Yan et al.\\ (2006) have provided an estimate to the stellar mass density at $z\\sim 6$, and concluded that the bulk of reionising photons must have been provided by other sources, perhaps by objects that are below current detection limits. It should be noted that such studies of Lyman break galaxies (LBGs) place {\\em lower} limits on the stellar mass density at the epoch of observation. The LBG selection technique is reliant upon an objects' detection in the rest-frame UV (e.g., the $z'$-band for $i'$-drops at $z\\sim 6$), so there must be at least {\\em some} ongoing star formation at the epoch of observation for the galaxy to be selected. Attempts to find post-starburst (Balmer-break) objects at $z\\sim 6$ are extremely difficult and uncertain due to a large population of lower redshift interloping galaxies. For example, Mobasher et al.\\ (2005) have recently suggested the presence of a massive post-starburst galaxy, HUDF-JD2, with a photometric redshift of $z_{phot}\\approx 6.5$ (high-redshift solutions with $z>5$ being preferred 85 per cent of the time). This comes from the re-analysis of an IRAC-detected Extremely Red Object (IERO) identified in the HUDF by Yan et al.\\ (2004), who originally derived a photometric redshift of $z_{phot}\\sim 3.4$ (see also Chen \\& Marzke 2004). If the higher redshift from Mobasher et al.\\ (2005) is correct, then the SED suggests a remarkably large stellar mass of $6\\times 10^{11}\\,M_{\\odot}$ at $z\\approx 6.5$ -- a factor of ten greater than the masses we presented in Eyles et al.\\ (2005). However, deep spectroscopy has yet to yield a redshift for this source, and the high photometric redshift estimate of Mobasher et al.\\ has been disputed by Dunlop et al.\\ (2006), who suggest that $z_{phot}=2.2$ is more plausible. Rather than study all galaxies in the GOODS fields with potential photometric redshifts at $z\\sim 6$, in this paper we restrict ourselves to the $i'$-drop selection, which has proven to be a reliable technique in isolating $z\\approx 6$ star-forming galaxies (see Bunker et al.\\ 2005 for a review) and so should minimize low redshift contaminant sources. Hence by investigating the stellar masses of our $z\\sim 6$ Lyman-break galaxies, we should be able to derive lower limits on the stellar mass density at this epoch. In this paper we examine the SEDs of $i'$-drop galaxies in the GOODS-South field, using observations from {\\em HST}/ACS in the optical, VLT/ISAAC in the near-IR, and {\\em Spitzer}/IRAC to span the rest-frame UV/optical. The work presented in this paper is an independent analysis of the $i'$-drop population, and differs from the recent work of Yan et al.\\ (2006) in several ways. In addition to photometric data gathered in {\\em HST}/ACS \\& {\\em Spitzer}/IRAC wavebands, we also use ground-based near-IR imaging of the GOODS-South field to better constrain the SEDs and the stellar population fitting (see Sections~\\ref{sec:PHOT}\\,\\&\\,\\ref{sec:SEDS}). Rather than assign a common redshift of $z=6$ to all galaxies in our sample, we choose to use $i'$-drops with either spectroscopically confirmed or robust photometric redshifts, and provide photometry for each source and details of the best-fit stellar population models for individual galaxies (Sections~\\ref{sec:REDUCT}\\,\\&\\,\\ref{sec:ANALYSIS}). For those objects which suffer from confusion with neighbouring sources, we attempt to subtract the contaminating objects in order to obtain accurate aperture photometry, rather than simply discarding these $i'$-drops from our sample (see Section~\\ref{sec:GALFIT}). We look to build on our previous work in Eyles et al.\\ (2005) by now considering a full sample of $i'$-drop candidates, and exploiting new, improved imaging datasets provided by the GOODS team (see Section~\\ref{sec:IMAGING}), including both epochs of the {\\em Spitzer}/IRAC imaging. An outline of this paper is as follows: Section~\\ref{sec:OBSERVATIONS} provides a summary of the imaging datasets used in this study and our selection of $i'$-drop galaxies. In Section~\\ref{sec:REDUCT} we describe the photometry and the removal of contaminating sources, and also the fitting of stellar population synthesis models to the observed spectral energy distributions. We discuss our results and their implications in Section~\\ref{sec:ANALYSIS}, and our conclusions are presented in Section~\\ref{sec:CONCLUSIONS}. Throughout we adopt the standard ``concordance'' cosmology of $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$, and use $H_0\\,=\\,70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$, which is within $2\\,\\sigma$ of the latest WMAP determination (Spergel et al.\\ 2006) -- in our adopted cosmology, the Universe today is 13.67\\,Gyr old, and at $z=6$ its age was 914\\,Myr. All quoted magnitudes are on the $AB$ system (Oke \\& Gunn 1983). ", "conclusions": "\\label{sec:CONCLUSIONS} We have used multi-waveband imaging of the GOODS-South field ({\\em HST}/ACS, VLT/ISAAC \\& {\\em Spitzer}/IRAC) to measure the stellar mass density of the $z\\sim 6$ $i'$-drop galaxy population. From an original catalog of 52 $i'$-drop candidates, we concentrated on 31 objects for which either spectroscopic or robust photometric redshifts from GOODS-MUSIC were available. A further 14 of these were eliminated from our analysis due to IRAC-confusion with neighbouring sources, which could not be satisfactorily removed. One further source, singled out as having peculiar photometry, had its SED matched very well to the spectra of a T-dwarf object -- an interloper in our sample. Fitting of SEDs to B\\&C spectral synthesis models was conducted on the remaining 16 $z\\sim 6$ galaxies, and properties including ages, stellar masses and star formation histories were constrained. Of these, nine were detected at IRAC wavelengths, and six of these showed evidence for significant Balmer/4000\\,\\AA\\ spectral breaks, brightening across the break by up to a factor of $2-3$ in $f_{\\nu}$. These indicate the presence of old stellar populations that dominate the stellar masses of these galaxies, with inferred ages of $\\sim 200 - 700$\\,Myr, and stellar masses of $\\sim 1.0 - 3.0\\times 10^{10}\\,M_{\\odot}$. During the SED modelling process, we considered the possibility of intrinsic dust reddening, and also the effects of differing metallicities and IMF models. We do not find evidence of substantial dust reddening in our $i'$-drop galaxies, and using differing metallicities did not have any significant effect on our derived properties. Use of a Chabrier, rather than Salpeter IMF had little effect on our inferred galaxy ages, but did reduce our stellar mass values by $\\approx 30$\\% due to a mass re-scaling. The results of the SED fitting of the three other fainter IRAC-detected sources were inconclusive. For the seven objects undetected in the IRAC wavebands, their SED fitting inferred much younger, less massive systems than their detected counterparts. Using the constrained properties of our $i'$-drop sample, we were able to calculate a value for the $z\\sim 6$ stellar mass density of $2.5\\times 10^{6}M_{\\odot}\\,{\\rm Mpc}^{-3}$, correcting for those objects eliminated from our analysis due to their un-treatable IRAC-confusion and those lacking GOODS-MUSIC photometric redshifts. Using a somewhat uncertain correction in order to account for the stellar mass in objects below our $z'$-band magnitude selection limit, this value could perhaps be $5 - 8\\times 10^{6}M_{\\odot}\\,{\\rm Mpc}^{-3}$. Any post-starburst and dust-obscured $z\\sim 6$ sources would not be found using the $i'$-drop selection technique, and hence our $z\\sim 6$ stellar mass density value is necessarily a lower limit, and is consistent with the estimates of Yan et al.\\ (2006). Exploring the previous star formation histories of our $i'$-drops, as inferred from their SED fitting, we suggest that the global star formation of these sources may have been substantially higher prior to the epoch of observation, and the resultant UV flux at $z>7$ may have played an important role in reionizing the Universe. \\subsection*" }, "0607/astro-ph0607595_arXiv.txt": { "abstract": "% We study the influence of X-ray radiation on the wind parameters of O stars. For this purpose we use our own NLTE wind code. The X-ray emission (assumed to be generated in wind shocks) is treated as an input quantity. We study its influence on the mass-loss rate, terminal velocity and ionization state of the stellar wind of Galactic O stars. ", "introduction": "From the observation of hot stars it is known that they emit X-ray radiation \\citet[e.g. Bergh\\\"ofer et al.,] [hereafter BSC] {rosat}. O star X-ray luminosity $L_\\text{X}$ scales with stellar luminosity $L$ roughly as $L_\\text{X}\\approx10^{-7}L$. The X-rays in O~stars are generated most likely by shocks that emerge in the supersonic stellar wind either due to the wind instability \\citep{felpulpal} or due to the stellar magnetic field \\citep{udo}. ", "conclusions": "" }, "0607/astro-ph0607240_arXiv.txt": { "abstract": "We present a simultaneous analysis of 10 galaxy lenses having time-delay measurements. For each lens we derive a detailed free-form mass map, with uncertainties, and with the additional requirement of a shared value of the Hubble parameter across all the lenses. We test the prior involved in the lens reconstruction against a galaxy-formation simulation. Assuming a concordance cosmology, we obtain $\\Ht=13.5^{+2.5}_{-1.3}\\;\\hbox{Gyr}$. ", "introduction": "If an object at cosmological distance is lensed into multiple images, the light travel time for individual images differs. For variable sources, the differences are observable as time delays. The delays are of order \\begin{equation} \\Delta t \\sim \\frac{GM}{c^3} \\sim (\\Delta\\theta)^2 \\, \\Ht \\label{delay-order} \\end{equation} where $M$ is the lens mass and $\\Delta\\theta$ is the image separation (in radians). As \\cite{refsdal64} first pointed out, the effect provides an independent way of measuring $\\Ht$. Time-delay measurements have made much progress over the past decade and now at least 15 are available (details below). While Eq.~(\\ref{delay-order}) provides the order of magnitude, to determine the precise factor relating time delays and $\\Ht$ one has to model the mass distribution. An observed set of image positions, rings, magnification ratios, and time delays is generically reproducible by many different mass models. This results in a large model-dependent uncertainty on the inferred Hubble parameter, even with perfect lensing data. To appreciate how serious this model-dependence is, compare the models of B0957+561 by \\cite{kundic97} and \\cite{bernstein99}: the results are $H_0=64\\pm13$ and $77^{+29}_{-24}\\;\\locunit$ respectively, both at $95\\%$ confidence; the more general models in the latter paper yield {\\em larger\\/} error-bars. Alternatively, consider the nice summary in Fig.~12 of \\cite{courbin03} of published $H_0$ estimates and uncertainties from individual lenses. Among the lenses shown, B1608+656 has all three of its independent time delays measured, B1115+080 has two delays measured, whereas the others have one each. One would expect these two best-measured lenses to be the best constrained. Yet B1608+656 has the largest error-bars on $H_0$ and B1115+080 the second-largest. This suggests that in the less-constrained lenses the real uncertainties are much larger, but have been underestimated because the fewness of constraints did not force sufficient exploration of model-dependence. A general strategy for dealing with the non-uniqueness problem is to search through a large ensemble of models that can all reproduce the observations \\citep{ws00,oguri04,jakobsson05}. In this paper we will follow such a strategy, simultaneously modeling 10 time-delay lenses coupled by a shared Hubble parameter. The basic method is the same as in \\cite{sw04} and the accompanying {\\em PixeLens\\/} code, but a number of refinements have been made. ", "conclusions": "We have expressed our main result (Fig.~\\ref{pg}) preferentially in terms of $\\Ht$ rather than $H_0$ because the former appears more naturally in lensing theory. But it is interesting to continue with $\\Ht$ in comparing with other techniques, because $\\Ht$ has a simple interpretation quite generally: it is $a/\\dot a$ or the doubling-time for metric distances at the current expansion rate. Coincidentally,\\footnote{Though a spoof paper by D.~Scott ({\\tt astro-ph/0604011}) develops a conspiracy theory for this.} in the concordance cosmology ($K=0,\\Omega_m\\simeq\\frac14,w=-1$) $\\Ht$ also equals the expansion age of the universe, within uncertainties. In particular, $\\Ht$ estimates can be immediately compared with globular-cluster ages, such as in \\cite{krauss03}. The well-known recent measurements of $\\Ht$, expressed in Gyr are: \\begin{enumerate} \\item $13.6\\pm1.5$ from \\cite{freedman01}, who combine several different indicators calibrated using Cepheids; \\item $15.7\\pm0.3\\;\\hbox{(statistical)}\\pm1.2\\;\\hbox{(systematic)}$ from \\cite{sandage06}, using SN\\thinspace Ia distances calibrated using Cepheids; \\item $13.6\\pm0.6$ from \\cite{spergel06}, using the CMB fluctuation spectrum. \\end{enumerate} Our result is consistent with any of these. It is worth noting, however, that the Hubble parameter appears in very different guises in different techniques. The distance-ladder methods measure the local cosmological expansion rate, independent of the global geometry. By contrast, in the CMB, $\\Ht$ is one parameter in a global cosmological model. Lensing is different again: here one assumes a global geometry and then measures a single scale parameter. The same is true of Sunyaev-Zel'dovich and X-ray clusters. The latter technique has made significant progress recently \\citep{jones05} but thus far still relies on strong assumptions: spherical symmetry of the cluster potential and hydrostatic equilibrium of the gas. In principle, lensing time delays can determine the global geometry as well \\citep{refsdal66} but the amount of data needed is not observationally viable yet. Whether lensing time delays can get the uncertainties in the Hubble parameter down to the 5\\% level is an open question. Maybe galaxy-lens models can be constrained enough to determine $\\Ht$ to better than 5\\%, thus making lensing the preferred method \\citep{schech04}; or maybe the approach is best used in reverse, inputting $\\Ht$ to constrain galaxy structure \\citep{kochanek06}. Fortunately, either outcome is worthwhile, and the basic technique will be the same. So whether the optimists or the pessimists are right, the usual cliches of ``more data!'' [time-delay measurements] and ``more theory!'' [lens models] are both apt. \\newpage" }, "0607/astro-ph0607526_arXiv.txt": { "abstract": "Amplified Tau-airshower at horizons may well open a novel powerful windows, at PeV-EeV energy, to Neutrino Astronomy. Neutrino induced air-showering astronomy rise because of neutrino masses, their mixing and the consequent replenishment of tau flavor during neutrino flight into spaces; Tau-Air-Showers escaping the Earth are the main traces amplified by its millions muon, billions gamma and thousand billions photon secondaries. Earth edges and its sharp shadows is the huge beam-dump detector for UHE neutrino and the almost noise-free screen for tau air-showers (as well as for PeVs anti-neutrino electron scattering on air electrons by resonant interactions). Crown array detectors for horizontal Cherenkov signals on mountains, on balloons and satellites widening the solid angle view are being elaborated; deep and wide valleys are considered. Tau Air-Showers neutrinos at EeV energies might rise in AUGER, facing the Ande shadows and eventually linking twin fluorescence telescopes to better test horizontal-inclined Cerenkov blazing photons. Tau air-showers may be revealed in ARGO being located within a deep valley testing inclined showers from the mountain sides; MILAGRO (and MILAGRITO) on a mountain top might already hide records of horizontal up-going muon bundles due to far UHECR and less far (but rarer) tau air-showers at EeV. MAGIC (or Veritas and Shalon) Telescopes pointing downward to terrestrial grounds acts, for EeV Tau neutrino air-showers astronomy, as a massive tens of $km^3$ water equivalent detector, making (in a given direction) it at present the most powerful dedicated neutrino telescope. MAGIC facing the sea edges must also reveal mirrored downward UHECR Air-showers (Cherenkov) flashes. Magic-crown systems may lead to largest neutrino detectors in near future. They maybe located on top mountains, on planes or balloons or in satellite arrays. They may be screened in deep valleys. Finally within cosmological relic light neutrino mass bounds ( suggesting $\\Sigma m_i\\simeq 0.18 eV$) a nearly degenerated neutrino mass $m_i \\simeq 0.06 eV$ rises offering future possible Z-Showering signals originated by $E_{\\nu} \\simeq 60 ZeV $ and UHECR secondaries above GZK cut-off, up to $E_p \\simeq ZeV $. ", "introduction": "Since Galileo we enjoyed of an optical view first of the planets, stars, and later galaxy maps while, since last century we enlarged the astronomical electromagnetic windows in radio, infrared, UV, X, $\\gamma$ with great success. Now a more compelling UHE $\\nu$ astronomy at EeVs energy is waiting at the corner. It is somehow linked to a very expected new particle astronomy: the UHECR at GZK energy $\\geq 4\\cdot10^{19}$eV: it must be a limited and nearby one (tens Mpc) because of cosmic BBR opacity. There have been since now two successfully neutrino astronomy at opposite low energy windows: the solar and the supernova ones. The solar ones has been explained by Davis,Gallex, SK, SNO experiment in last four decades opening the $\\nu$ physics to a solar neutrino mass splitting and a clear probe to its mixing behavior. The supernova SN 1987A was an unique event that anyway had a particular expected signatures at tens MeV. On going experiment on cosmic supernova background in S.K. are at the threshold edges, possibly ready to a discover of this cosmic background. However there is a more exciting and energetic $\\nu$ astronomy at PeV and EeV energy associate to the evidence of charged UHECR spectra at EeV and tens hundred of EeV band. Indeed any EeV CR originated nearby an AGN or GRB or BL Lac jet will be partially screened by the same source lights leading to a consequent photo-pion production, associated with PeV secondary neutrinos. In a much simpler and guaranteed way, at energy about $4\\cdot10^{19}$eV, UHECR should propagate in cosmic photon black body, being partially arrested by photopion productions, (GZK cut-off), leading to EeV neutrinos all along the Universe confines. These UHE $\\nu$ components, consequence of the GZK cut-off, are called cosmogenic or GZK neutrinos. Their flux may be estimated by general arguments and there is quite a wide consensus on such neutrino GZK flux at EeV energies. These \\textit{guaranteed} neutrinos may be complementary to possible \\textit{expected} higher energy neutrinos (at ZeV energies) whose role might explain UHECR isotropy and homogeneity being originated at cosmic distances. In this model UHECR born as nucleons via $\\nu+\\bar{\\nu_R}\\rightarrow Z\\rightarrow X+N$ (Z-burst or Z-shower model) \\cite{Fargion-Mele-Salis99},\\cite{Weiler97},\\cite{Yoshida1998} are overcoming present (AGASA,HIRES,AUGER) un-observed local (VIRGO,PERSEUS) source distribution, as would be prescribed by naive GZK cut-off. However GZK neutrinos and Z-Burst neutrinos at EeV are making comparable flux predictions and we shall restrict to the simplest GZK flux assumption. In conclusion we remind the role of neutrino masses in calibrating the Z-Burst showering and the influence in UHECR spectra. ", "conclusions": "Because muon tracks are mostly of downward atmospheric nature the underground neutrino telescope are tracing rarer upward ones mostly of atmospheric origin; higher energy up-going astrophysical neutrino signals are partially suppressed by Earth opacity, and are unique tracks. On the contrary $\\tau$ air shower at horizons is spreading its signal in a wider area leading to populated (millions-billions) muon and gamma (as well as electron pairs) bundles in their showering secondary mode. This amplified signal may be observed and disentangled from farer and air filtered UHECR, in different ways and places: mountains, balloons, satellites with different detector array area and thresholds. The advantage to be in high quota is to be extending the visible target terrestrial area and solid angle, as well as to let a longer tau flight distance (and energy), and to enlarge the air shower area; a too high altitude, however , looses solid angle. To make an intuitive estimate the Tau air-shower size area, at tens PeVs-EeVs ,( detectable at horizons within a lateral distance as large as $3$ km. from the main shower axis by a telescope like Magic),it is nearly $30 km^2$; at EeV energy the equivalent detection depth crossed by the tau lepton before the exit from the Earth reaches $10-20 km$ distances; the corresponding detection Neutrino volume (inside the narrow, conic $10^{-3} sr.$, shower beam) is within $30-60 km^3$, in any given direction , see \\cite{Fargion2005}. A few events of GRBs a year may be located within these horizons, as well as AGN and BL Lac in their flare activity. In such occasions Magic, Veritas and Hess array are the most sensitive neutrino telescope at PeVs-EeV energy. Even on average, for a present $2\\cdot2^o$ view of Magic, at present energy thresholds, such telescopes (for Neutrino at Glashow PeVs energy windows), are testing a total mass-solid angle a comparable or larger than to $10^{-2} km^3 sr$, an order of magnitude comparable with the present AMANDA detector. Moreover the light neutrino mass that seem to converge (by recent cosmic constrains) toward a light degenerated neutrino mass value $m_i \\simeq 0.6 eV$ seem to suggest a UHE primary neutrino at $E_{\\nu} = 60 ZeV$ and a UHECR secondary bump at $E=3.3 ZeV $. In this view UHECR modulation at GZK energies may reflect lightest neutrino masses. In conclusion a maximal alert for the Neutrino air-showering within the Earth shadows is needed: in AUGER, Milagro, Argo, as well as in ASHRA, CRTNT, Shalon Telescopes the signal is beyond the corner. In particular the Magic (and Veritas) arrays telescopes facing from the mountains the Horizons edges may soon test our proposal leading to such crown arrays. In a sentence we believe that the UHE Neutrino Astronomy is beyond the corner, Tau is its courier and its sky lay just beneath our own sky: the Earth." }, "0607/astro-ph0607460_arXiv.txt": { "abstract": "Over the last years a new generation of model atmosphere codes, which include the effects of metal line-blanketing of millions of spectral lines in NLTE, has been used to re-determine the properties of massive stars through quantitative spectral analysis methods applied to optical, IR and UV spectra. This has resulted in a significant change of the effective temperature scale of early type stars and a revision of mass-loss rates. Observed mass-loss rates and effective temperatures depend strongly on metallicity, both in agreement with theoretical predictions. The new model atmospheres in conjunction with the new generation of 10m-class telescopes equipped with efficient multi-object spectrographs have made it possible to study blue supergiants in galaxies far beyond the Local Group in spectroscopic detail to determine accurate chemical composition, extinction and distances. A new distance determination method, the flux weighted gravity - luminosity relationship, is discussed as a very promising complement to existing stellar distance indicators. Observationally, there are still fundamental uncertainties in the determination of stellar mass-loss rates, which are caused by the fact that there is evidence that the winds are inhomogeneous and clumped. This may lead to major revisions of the observed rates of mass-loss. ", "introduction": "Hot massive stars are cosmic engines of fundamental importance not only in the local but also in the early universe. A first generation of very massive stars has very likely influenced the formation and evolution of the first building blocks of galaxies. The spectral appearance of Lyman break galaxies and Ly$\\alpha$-emitters at high redshift is dominated by an intrinsic population of hot massive stars. Gamma-ray bursters are very likely the result of terminal collapses of very massive stars and may allow to trace the star formation history of the universe to extreme redshifts. It is obvious that the understanding of important processes of star and galaxy formation in the early universe is intimately linked to our understanding of the physics of massive stars. The observational constraints of the latter are provided by quantitative spectroscopic diagnostics of the population of hot massive stars in the local universe. It is the goal of this contribution to provide an overview about the dramatic progress, which has been made in this field over the last five years. There are two factors which have contributed to this progress, new observational facilities such as the optical/infrared telescopes of the 10m-class on the ground and observatories in space allowing for spectroscopy in the UV (HST, FUSE, GALLEX), IR (ISO, Spitzer) and at X-ray wavelengths (XMM, Chandra) and the enormous advancement of model atmosphere and radiative transfer techniques. As for the latter, it is important to realize that modelling the atmospheres of hot stars is a tremendous challenge. Their physics are complex and very different from standard stellar atmosphere models. They are dominated by a radiation field with energy densities larger than or of the same order as the energy density of the atmospheric matter. This has two important consequences. First, severe departures from Local Thermodynamic Equilibrium of the level populations in the entire atmosphere are induced, because radiative transitions between ionic energy levels become much more important than inelastic collisions. Second, supersonic hydrodynamic outflow of atmospheric matter is initiated by line absorption of photons transferring outwardly directed momentum to the atmospheric plasma. This latter effect is responsible for the existence of the strong stellar winds observed and requires the use of NLTE model atmospheres, which include the hydrodynamic effects of stellar winds (see \\cite{kud98} for a detailed description of the physics of hot star atmospheres). The winds of hot massive stars are fundamentally important. Their energy and momentum input into the ISM is significant creating circumstellar shells, stellar wind bubbles and initiating further star formation. They affect the stellar evolution by modifying evolutionary timescales, chemical profiles, surface abundances and stellar luminosities. They also have substantial effects on the structure of the stellar atmospheres. They dominate the density stratification and the radiative transfer through their transonic velocity fields and they modify the amount of the emergent ionizing radiation significantly (Gabler et al., 1989, 1991, 1992, Najarro et al., 1996). While the basic concepts for the hydrodynamic atmospheres of hot stars and the spectroscopic diagnostics of their parameters and stellar winds have been developed in the eighties and nineties (see reviews by \\cite{kuhu90}, \\cite{kud98} and \\cite{kupu00}), the development of the most recent generation of model atmospheres, which for the first time accounts selfconsistently for the effects of NLTE metal line-blanketing (see section 2), has lead to a dramatic change of the diagnostic results. A new effective temperature scale has been obtained for the O-star spectral types, which are apparently significantly cooler than originally assumed and, thus, on average have a lower luminosity and mass and also provide less ionizing photons, than previously thought. We will discuss these effects in detail in section 2, 3, and 4. Note that we will focus our review on ``normal'' hot massive stars in a mass-range between 15 and 100 $M_{\\odot}$ in well-established evolutionary stages such as dwarfs, giants, and supergiants (Fig.\\,\\ref{hrd}). We will not discuss objects with extreme winds such as Wolf-Rayet stars, luminous blue variables in outbursts, etc. For those, we refer the reader to the contributions by Paul Crowther and Nathan Smith in this volume. \\begin{figure} \\centerline{\\hbox{ \\psfig{figure=kudfig1.ps,width=7cm} }} \\caption[]{ Evolutionary tracks of massive stars in the HRD with (solid) and without (dashed) stellar rotation, \\cite{meynet00}) and the domains of spectral types discussed in this review. \\label{hrd}} \\end{figure} The new model atmospheres in conjunction with the new generation of 10m-class telescopes equipped with efficient multi-object spectrographs allow to study blue supergiants in galaxies far beyond the Local Group in great spectroscopic detail to determine accurate chemical composition, extinction and distances. We will report recent results in this new field of ``extragalactic stellar astronomy'' in section 5. In section 6, we will discuss a new distance determination method based on stellar photospheric spectroscopy, the flux weighted gravity - luminosity relationship. The diagnostics of stellar winds, in particular of mass-loss rates, are also affected by the proper accounting for line-blanketing affects. While the general scaling relations of stellar wind parameters with stellar luminosity, mass, radius, and metallicity as described in \\cite{kupu00} remain qualitatively unchanged, important quantitative changes have been found. There are still puzzling uncertainties with regard to the observed rates of mass-loss related to the inhomogeneous structure of the stellar wind outflows. We will present and discuss the most recent results in section 7. As indicated above, the new generation of model atmosphere codes has important applications on the interpretation of spectra of star-bursting galaxies in the early universe. Fig.\\,\\ref{starburst} is a nice example. This work by \\cite{rix04} has been used to constrain star formation rates and metallicities in high redshift Lyman-break galaxies. Another example is the work by \\cite{barton04}, which uses stellar atmosphere model predictions by \\cite{bromm01} and \\cite{schaerer03} to explore the possibility to detect the first generation of very massive stars at redshifts around ten with present day 8m-class telescopes and with future (diffraction limited) optical/IR telescopes of 30m aperture such as the GSMT (see also report by the GSMT Science Working Group, \\cite{kud03a}). As it turns out, not only would the GSMT be able to detect such objects it would also allow to constrain the IMF of massive stars in the early universe from the relative comparison of L$_{\\alpha}$ and HeII1640 recombination lines caused by the population of the first stars. \\begin{figure} \\centerline{\\hbox{ \\psfig{figure=kudfig2.ps,width=9cm} }} \\caption[]{ Model UV spectra of an integrated population of massive stars in a starburst galaxy displayed as a function of metallicity. The model atmosphere code used was WMBasic (\\cite{pauldrach01}) and the population synthesis was done with Starburst99. (From \\cite{rix04}). \\label{starburst}} \\end{figure} ", "conclusions": "" }, "0607/astro-ph0607656_arXiv.txt": { "abstract": "The main subject of my talk is the question: in what kind of astrophysical systems magnetic reconnection is interesting and/or important? To address this question, I first put forward three general criteria for selecting the relevant astrophysical environments. Namely, reconnection should be: fast; energetically important; and observable. From this, I deduce that the gas density should be low, so that the plasma is: collisionless; force-free; and optically thin. Thus, for example, the requirement that reconnection is fast implies that Petschek's reconnection mechanism must be operating, which is possible, apparently, only in the collisionless regime. Next, I argue that the force-free condition implies that the magnetic field be produced in, and anchored by, a nearby dense massive object, e.g., a star or a disk, strongly stratified by gravity. I then stress the importance of field-line opening (e.g., by differential rotation) as a means to form a reconnecting current sheet. Correspondingly, I suggest the Y-point helmet streamer as a generic prototypical magnetic configuration relevant to large-scale reconnection in astrophysics. Finally, I discuss several specific astrophysical systems where the above criteria are met: stellar coronae, magnetically-interacting star--disk systems, and magnetized coronae above turbulent accretion disks. In the Appendix I apply the ideas put forward in this talk to the solar coronal heating problem. ", "introduction": "\\label{sec-fast} First, what is {\\it fast reconnection}? Usually in the magnetic reconnection literature, a reconnection mechanism is called ``fast'' if the reconnection rate is independent (or scales only logarithmically with) the classical resistivity. In my talk, I will use a somewhat broader definition. I will call a reconnection process fast when the reconnection rate (defined as the ratio of reconnection velocity to the Alfv\\'en speed) is independent of (or depends only relatively weakly on) the global size~$L$ of the system. In other words, in this definition reconnection is fast when its dimensionless rate is determined (almost) entirely by the local physical parameters near the center of the reconnection layer. In this sense, for example, the classical Sweet--Parker reconnection is slow, because its rate scales as~$L^{1/2}$, wheres the maximum Petschek reconnection rate is (almost) fast, since it scales only logarithmically with~$L$. \\subsection{Main Mechanism of Fast Reconnection in Astrophysics} \\label{subsec-handwaves} In Astrophysics, {\\it reconnection} is a magic word (a ``tooth fairy'', as it would be called in Peyton Hall),% \\footnote {Another well-known astrophysical ``tooth fairy'' is {\\it magnetic field} itself.} in the sense that it has become customary to invoke reconnection when it is needed to solve one's problems and to assume that it always works when called upon. It has to be noted that the main reconnection mechanism in Astrophysics is NOT Petschek reconnection, nor is it Hall reconnection, nor anomalous-resistivity reconnection. No, the most important reconnection mechanism in Astrophysics invokes waves, a certain type of waves, in fact. Called {\\it handwaves}% \\footnote {I acknowledge first hearing a satyrical mentioning of handwaves as a playful ``real'' physical mechanism from Henk Spruit (2005, private communication).} (See Fig.~1). The mechanism works like this: {\\it Well, we know that fast reconnection happens in the Solar corona, and in the Earth magnetosphere. So it should also happen in OUR astrophysical system.} Following this well-established and respected astrophysical tradition, I will also make extensive use of hand-waving arguments throughout my talk :) \\begin{figure} [h] \\centerline{\\psfig{file=handwaves-color.eps,width=4 in}} \\figcaption{Main Reconnection Mechanism in Astrophysics. \\label{fig-handwaves}} \\end{figure} \\subsection{Fast Reconnection: Petschek's Legacy} \\label{subsec-Petschek} Those who are not satisfied with the mechanism described in the previous subsection, have to rely on actual thought. An excellent example of a brilliant thinker, who contributed a lot to our understanding of reconnection, is Harry Petschek. Why are Petschek's ideas on reconnection important? People now associate fast reconnection with Petschek's (1964) reconnection mechanism.% \\footnote {By the way, I will assume the audience to be familiar with both the Sweet-Parker and Petschek models of reconnection.} A great non-trivial idea due to Petschek is that the main bottleneck stifling the reconnection process in the classical Sweet--Parker (Sweet~1958; Parker~1957) model can be circumvented if one can set up a special magnetic field and flow configuration --- the Petschek configuration. The bottleneck arises because of the necessity to have a reconnection layer that is simultaneously thin enough for the resistivity to be important and thick enough for the plasma to be able to flow out of the layer. Petschek's key idea was that the thin current layer and the thick outflow channel do not have to be the same if the reconnection region is not a simple rectangular box, as it was in the Sweet--Parker theory, but has a somewhat more complicated structure, with four standing slow shocks attached to a central diffusion region. As a result, there is an additional geometric factor that can lead to faster reconnection. This geometrical enhancement is especially important in astrophysics, for the following reason. What distinguishes astrophysical and space systems from Earth-bound laboratory experiments is a huge contrast in length-scales. Astronomical systems are astronomically large: the system size~$L$ is usually much greater than the microscopic physical scales, e.g., the ion gyro-radius~$\\rho_i$, the ion collisionless skin-depth~$d_i$, and the Sweet--Parker reconnection layer thickness~$\\delta_{\\rm SP}$. Hence, the reconnection rate problem is especially severe in astrophysical systems. Hence, we need a clever idea. For example, the idea of geometric enhancement due to Harry Petschek. What this means is that, unless one has a special mechanism like this, no microphysics (e.g., the Hall effect or anomalous resistivity) can give reconnection rates that are rapid enough to be of interest to observations. For example, when the reconnection layer's thickness becomes comparable with the collisionless ion skin depth~$d_i$, the layer enters the Hall regime. Then, in a simple Sweet--Parker-like analysis, the mass conservation condition would result in a reconnection velocity that is by a factor of~$d_i/L$ smaller than the Alfv\\'en speed~$V_A$. Since $d_i$ is usually much smaller than~$L$, this rate would be much too slow to be of practical interest. \\subsection{Fast Reconnection Means Collisionless Reconnection} \\label{subsec-collisionless} Thus, we see that Petschek's mechanism, or a variation thereof, is absolutely indispensable for astrophysical reconnection. Unfortunately, however, several numerical and analytical studies (e.g., Biskamp~1986; Scholer~1989; Uzdensky \\& Kulsrud~2000; Erkaev~et~al. 2001; Kulsrud~2001; Malyshkin et al.~2005) have shown that in resistive MHD with uniform resistivity (and, by inference, with resistivity that is a smooth function of plasma parameters, e.g., Spitzer) Petschek's mechanism fails and Sweet--Parker scaling applies instead. The same conclusion was achieved in laboratory studies by the Magnetic Reconnection Experiment (MRX), lead by Masaaki Yamada at Princeton Plasma Physics Laboratory, in the high-collisionality regime (Ji~et~al.~1998; Trintchouk~et~al.~2003). What all this means is that, whenever classical resistive MHD applies, one does not get fast reconnection. This implies that fast reconnection can happen {\\it only} when the plasma is relatively collisionless so that resistive MHD doesn't apply. This condition of fast reconnection can be formulated roughly% \\footnote {If the guide component of the magnetic field is not zero, this condition may be somewhat different, although similar in concept. For simplicity, however, in this paper we shall consider the fast reconnection condition only as given by equation~(\\ref{eq-1}).} as (e.g., Yamada~et~al. 2006) \\beq \\delta_{\\rm SP}\\ll d_i \\equiv {c\\over{\\omega_{pi}}} \\, . \\label{eq-1} \\eeq What this condition means is the following. As a reconnection layer is forming, its thickness $\\delta$ is getting smaller and smaller. If condition~(\\ref{eq-1}) is not satisfied, then this thinning saturates at $\\delta=\\delta_{\\rm SP}$, and reconnection then proceeds in the slow Sweet--Parker regime. However, if condition~(\\ref{eq-1}) is satisfied, then various two-fluid and/or kinetic effects kick in as soon as~$\\delta$ drops down to about~$d_i$ or so, well before the collisional resistive effects become important. Then, the reconnection procesess necessarily involves collisionless, non-classical-resistive-MHD physics. Thus, in the collisional regime, when classical resistive MHD applies, fast Petschek reconnection does not appear to be possible. Does going to the collisionless regime help? There is a growing consensus that the answer to this question is YES. In Space/Solar physics, of course, there has long been a very serious evidence for fast collisionless reconnection; it has been further significantly strengthened by recent laboratory measurements in the MRX (Ji~et~al. 1998; Yamada~et~al. 2006). These measurements, however, have not been able to elucidate the special role of the Petschek mechanism in accelerating reconnection. On the other hand, over the past decade or so, several theoretical and numerical studies have indicated that fast reconnection enhanced by the Petschek mechanism (or a variation thereof) does indeed take place in the collisionless regime. It appears that there may be two regimes of collisionless reconnection. Physically, these two possibilities are very different from each other; nevertheless, they both appear to lead to the establishment of a Petschek-like configuration, which enhances the reconnection rate. The two regimes in question are: \\begin{itemize} \\item{{\\it Hall-MHD reconnection}, involving two-fluid effects in a {\\it laminar} flow configuration (e.g., Shay~et~al. 1998; Birn~et~al.~2001; Bhattacharjee~et~al. 2001). [See, however, recent particle simulations by Daughton et~al. (2006) and by Fujimoto (2006) that cast doubt on Hall reconnection as a possible fast reconnection mechanism.]} \\item{Spatially-localized {\\it anomalous resistivity} due to micro-turbulence; this seems to lead to a Petschek configuration with the inner diffusion region having a width of the order of the resistivity localization scale (e.g., Ugai \\& Tsuda 1977; Sato \\& Hayashi 1979; Scholer~1989; Biskamp \\& Schwarz 2001; Erkaev~et~al. 2001; Kulsrud~2001; Malyshkin~et~al. 2005).} \\end{itemize} At present, it is still not clear which one of these two mechanisms works in a given physical situation (if at all). Also not known is whether these two mechanisms can coexist and perhaps even enhance each other. Recent experimental evidence from the MRX experiment suggests that both regimes do exist in reality and that they may operate simultaneously in a given system (Yamada~2006, private communication). In any case, it seems that one does get a Petschek-enhanced fast reconnection process if the plasma is collisionless [in the sense of equation~(\\ref{eq-3})]. To sum up, in order for astrophysical reconnection to be fast, it needs Petschek's mechanism to operate and that in turn requires the reconnection layer to be collisionless. Thus, for the purposes of this talk, I will put an equal sign between collisionless reconnection and fast Petschek's reconnection. In fact, whenever we observe violent and rapid energetic phenomena that we interpret as reconnection, it is always in relatively tenuous plasmas. Please correct me if this is not so. I would be very interested in learning about counter-examples. Is there any evidence for fast large-scale reconnection events in collisional astrophysical environments? \\subsection{Fast Reconnection: Range of Densities in Astrophysics} \\label{subsec-densities} Note that the statement of collisionality is scale-dependent. This is because~$d_i$ is a microscopic scale, independent of~$L$, whereas $\\delta_{\\rm SP}\\sim\\sqrt{L}$. In astrophysics $L\\gg d_i$ is very large, and so one might expect $d_i\\ll \\delta_{\\rm SP}$ for large enough systems. However, in practice, this doesn't always have to be so. Indeed, notice that $\\delta_{\\rm SP}\\sim 1/\\sqrt{S} \\sim V_A^{-1/2} \\sim \\rho^{-1/4}$, whereas $d_i\\sim\\rho^{-1/2}$. Thus, $$ {{\\delta_{\\rm SP}}\\over{d_i}} \\sim \\rho^{1/4} \\, . $$ so it depends on density. This dependence is weak but, in astrophysics, one has to deal not only with a huge dynamic range in~$L$ but also with an enormous dynamic range in~$\\rho$. For example, \\begin{itemize} \\item{$\\rho\\sim 10^{14}\\,{\\rm g/cm}^3$ inside a neutron star;} \\item{$\\rho\\sim 1\\,{\\rm g/cm}^3$ average solar density;} \\item{$n_e\\sim 10^{10}\\,{\\rm cm}^{-3}$ in the solar corona;} \\item{$n_e\\sim 1\\,{\\rm cm}^{-3}$ in the ISM.} \\end{itemize} --- 38 orders of magnitude variation in density! Thus, one can readily find many astrophysical systems that are not only large but also rarefied and collisionless. \\subsection{The Fast Reconnection Condition} \\label{subsec-condition} How can one quantify condition~(\\ref{eq-1}) of collisionless reconnection? It is pretty straight-forward to show (Yamada~et~al. 2006) that \\beq {{\\delta_{\\rm SP}}\\over{d_i}} \\sim \\biggl({L\\over{\\lambda_{e,\\rm mfp}}}\\biggr)^{1/2}\\, \\biggl(\\beta_e\\,{m_e\\over{m_i}}\\biggr)^{1/4} \\, , \\label{eq-2} \\eeq where I have neglected numerical factors of order~1. Here, $\\beta_e$ is the ratio of the plasma pressure inside the layer to the pressure of the reconnecting magnetic field component ($B_0^2/8\\pi$) outside the layer; $\\lambda_{e,\\rm mfp}$ is the classical electron mean free path due to Coulomb collisions. Thus, using equation~(\\ref{eq-1}), we see that reconnection is collisionless when \\beq \\lambda_{e,\\rm mfp} > L\\sqrt{\\beta m_e/m_i} \\simeq L\\beta^{1/2} /40 \\, . \\label{eq-3} \\eeq [The condition suggested by Yamada~et~al. (2006) differs from equation~(\\ref{eq-3}) by a factor of~2. Since the discussion here is very qualitative, I regard this difference as unessential. We are not going to quibble about factors of~2, are we?] We can go a little bit further. The mean free path $\\lambda_{e,\\rm mfp}$ can be written as \\beq \\lambda_{e,\\rm mfp} \\simeq 7\\cdot 10^{7}{\\rm cm}\\, n_{10}^{-1}\\, T_7^2 \\, , \\label{eq-lambda-1} \\eeq where we have taken the Coulomb logarithm equal to~20 and where~$n_{10}$ and~$T_7$ are the electron density~$n_e$ and temperature~$T_e$ given in units of $10^{10}\\, {\\rm cm}^{-3}$ and~$10^7$~K, respectively. These parameters are to be taken at the center of the reconnection layer. Combining equations~(\\ref{eq-3}) and~(\\ref{eq-lambda-1}), the criterion for fast collisionless reconnection can now be formulated as a condition on the layer's length~$L$ in terms of the central values of~$n_e$ and~$T_e$: \\beq L < L_c \\equiv 40\\, \\beta^{-1/2}\\, \\lambda_{e,\\rm mfp} \\simeq 3\\cdot 10^{9}{\\rm cm}\\ \\beta^{-1/2}\\, n_{10}^{-1}\\, T_7^2 \\, . \\label{eq-L_c-1} \\eeq Now, what about the plasma-$\\beta$ parameter? For definiteness, let us focus on the extreme case of a reconnecting configuration with no guide field. Then the condition of pressure balance across the layer dictates that \\beq \\beta \\equiv {{8\\pi n_e k_B (T_e+T_i)}\\over{B_0^2}} = 1 \\, , \\label{eq-pressure-balance} \\eeq where we have neglected the outside thermal pressure. Furthermore, assuming for simplicity that $T_e=T_i$, we can then express the central electron temperature in terms of~$B_0$ and~$n_e$ as \\beq T_e = {{B_0^2/8\\pi}\\over{2k_B n_e}} \\simeq 1.4 \\cdot 10^7\\, {\\rm K} \\, B_{1.5}^2 \\, n_{10}^{-1} \\, , \\label{eq-T_e} \\eeq where $B_{1.5}$ is the outside magnetic field $B_0$ expressed in units of~30~G. Upon substituting this estimate into equation~(\\ref{eq-lambda-1}), we get \\beq \\lambda_{e,\\rm mfp} \\simeq 1.5\\cdot 10^8{\\rm cm}\\, n_{10}^{-3}\\, B_{1.5}^4 \\, , \\label{eq-lambda-2} \\eeq and correspondingly, \\beq L_c(n,B_0) \\simeq 6\\cdot 10^{9}{\\rm cm}\\, n_{10}^{-3}\\, B_{1.5}^4 \\, . \\label{eq-L_c-2} \\eeq We see that the condition $L< L_c$ is easily satisfied for solar flares, for example. Also, this result has interesting implications for the coronal heating problem, as I will discuss in the Appendix. A reservation: this was just a simple example, presented in order to illustrate the basic idea. In general, the physics is more complicated and less certain. In particular, in the above example I have assumed the plasma pressure at the center of the current layer to be equal to the outside magnetic pressure outside (the background gas pressure in the corona outside the layer is negligible since the corona is almost force-free), i.e., that $\\beta\\simeq 1$. However, if there is a guide magnetic field, then the cross-layer pressure balance is modified; in particular, $\\beta$ can be much less then~1, determined by thermal transport processes along the layer, e.g., the electron thermal conduction (radiative losses are small on the timescale of transit through the layer). Just as important, since collisions are rare, and since the collisional electron-ion energy-equilibration rate is suppressed due to the large mass ratio, the electron and ion temperatures in the layer need not be equal. For example, ions may be much hotter than the electrons and may provide the bulk of the pressure support against the outside magnetic field. The electron temperature in this case would be far below the equipartition value (about $10^7$ K). Correspondingly, the electron mean-free path would be much lower than that given by equation~(\\ref {eq-lambda-2}). \\subsection{Fast Reconnection: Caveats and Alternatives} \\label{subsec-caveats} With all this said, there is still room for caveats and alternative ideas. I will mention just some of them here: \\begin{itemize} \\item{Recent numerical work by Cassak~et~al. (2005) shows an intriguing evidence for {\\bf bistable reconnection}: once fast Hall reconnection has begun, it is hard to switch it back to the slow Sweet--Parker mode, even if the resistivity is raised to the level that violates (\\ref{eq-1}).} \\item{{\\bf Turbulent Reconnection:} Lazarian \\& Vishniac (1999) suggested that fast reconnection may happen in pure resistive MHD, although only in 3D, in the presence of externally imposed MHD turbulence (see also Bhattacharjee \\& Hameiri 1986; Strauss~1988; Kim \\& Diamond~2001). We don't know really whether this mechanism works, but certainly it is an interesting idea. And a testable one! Numerical tests should now be possible. It is worth pursuing and may be a good topic for a PhD thesis project! } \\item{{\\bf Bursty, impulsive reconnection:} e.g., Bhattacharjee (2004)} \\item{{\\bf 3D reconnection}: e.g., Longcope (1996)} \\item{{\\bf Additional Physics:} e.g., reconnection in partially-ionized plasmas in the context of molecular clouds (e.g., Zweibel~1989).} \\end{itemize} ", "conclusions": "" }, "0607/hep-th0607001_arXiv.txt": { "abstract": "\\addtolength{\\baselineskip}{1.2mm} The power spectrum of M-theory cascade inflation is derived. It possesses three distinctive signatures: a decisive power suppression at small scales, oscillations around the scales that cross the horizon when the inflaton potential jumps and stepwise decrease in the scalar spectral index. All three properties result from features in the inflaton potential. Cascade inflation realizes assisted inflation in heterotic M-theory and is driven by non-perturbative interactions of $N$ M5-branes. The features in the inflaton potential are generated whenever two M5-branes collide with the boundaries. The derived small-scale power suppression serves as a possible explanation for the dearth of observed dwarf galaxies in the Milky Way halo. The oscillations, furthermore, allow to directly probe M-theory by measurements of the spectral index and to distinguish cascade inflation observationally from other string inflation models. ", "introduction": "For a long time it seemed difficult to connect inflation to string-theory. In its low-energy approximation string-theory is described by supergravity. Inflation based on the F-term potentials of 4-dimensional ${\\cal N}=1$ supergravities resulting from string/M-compactifications suffers from a large slow-roll parameter $\\eta$. The origin of this problem traces back to the appearance of the K\\\"ahler-factor $\\exp(K)$ in the F-term potential. New possibilities to address this problem arose with the advent of D-branes \\cite{Polchinski:1995mt}. They allowed to identify the inflaton with open string modes such as the geometrical distance between two D-branes \\cite{Dvali:1998pa}. An inflaton requires a very shallow potential. Hence, a priori, moduli serve as natural candidates. To provide them with a non-trivial potential supersymmetry needs to be broken which can be done in various ways. One might add anti D-branes to the open string sector \\cite{Kachru:2003sx} or supersymmetry breaking fluxes to the closed string sector \\cite{Dine:1985rz}. Also the inclusion of non-perturbative instanton effects leads to spontaneous supersymmetry breaking in the low-energy supergravity \\cite{Curio:2001qi}. Assuming just a single inflaton, the task for deriving inflation from string-theory becomes then finding a way of breaking supersymmetry which leaves the inflaton with a sufficiently flat potential while endowing all other moduli with steep stabilizing potentials. All standard methods of breaking supersymmetry generate, however, steep potentials, not flat ones. One way to generate a flat inflaton potential nevertheless is to study brane-antibrane inflation in warped backgrounds with the inflaton being identified with the brane-antibrane distance \\cite{Kachru:2003sx}. Warped geometries arise in the presence of branes and fluxes. The eventual stabilization of the volume modulus, however, modifies the inflaton potential and renders it too steep for inflation unless finetuning is applied \\cite{McAllister:2005mq}. Here, we focus on an alternative mechanism to generate inflation in M/string-theory, the multi brane inflation proposal \\cite{Becker:2005sg}, \\cite{ta} (see also \\cite{Cline:2005ty}). One starts with a multi inflaton scenario associating one inflaton with each brane position. The presence of several branes is indeed generically enforced by tadpole cancellation conditions. The interesting advantage of this mechanism lies in the fact that the potentials for the individual inflatons need no longer be flat. The reason is that the Hubble friction experienced by every inflaton becomes large -- simply by increasing the number of inflatons -- regardless of the steepness of the potentials. This had first been pointed out in \\cite{Liddle:1998jc} in the context of 4-dimensional Friedmann-Robertson-Walker (FRW) cosmologies based on exponential potentials which generate power-law inflation. The premise under which this mechanism operates is the suppression of strong cross-couplings among the inflatons. This suppression is given in multi brane inflation models since interactions between non-neighboring branes which could generate cross-couplings are suppressed by longer distances. In M-theory cascade inflation \\cite{ta} there is an exponential suppression of such cross-couplings since interactions between the relevant M5-branes arise from non-perturbative open M2-instantons. In this work, after highlighting the needed ingredients of M-theory cascade inflation, we focus on the determination of its power spectrum and the resulting observable signatures. Beyond demonstrating the compatibility of the power spectrum with present cosmological constraints, we find that it exhibits {\\em three distinctive signatures -- power suppression at small distances, stepwise decrease in the spectral index and oscillations in the spectrum.} The power suppression which follows in cascade inflation from M-theory dynamics might serve as an explanation for the scarceness of observed dwarf galaxies in the Milky Way halo, as suggested in \\cite{Kamionkowski:1999vp}, \\cite{Sigurdson:2003vy}. This is not explained by standard cosmology which overpredicts their abundance by an order of magnitude. The oscillations and stepwise decreases, on the other hand, provide a unique signature which allow to probe M-theory observationally by measuring the spectral index. It furthermore clearly distinguishes M-theory cascade inflation observationally from other string inflation models. ", "conclusions": "" }, "0607/astro-ph0607436_arXiv.txt": { "abstract": "We present a series of cosmological \\nbody\\ simulations which make use of the hydrodynamic approach to the evolution of structures (Dom{\\'\\i}nguez 2000). This approach addresses explicitly the existence of a finite spatial resolution and the dynamical effect of subresolution degrees of freedom. We adapt this method to cosmological simulations of the standard \\LCDM\\ structure formation scenario and study the effects induced at redshift $z=0$ by this novel approach on the large--scale clustering patterns as well as (individual) dark matter halos. Comparing these simulations to usual \\nbody\\ simulations, we find that (i) the new (hydrodynamic) model entails a proliferation of low--mass halos, and (ii) dark matter halos have a higher degree of rotational support. These results agree with the theoretical expectation about the qualitative behaviour of the \"correction terms\" introduced by the hydrodynamic approach: these terms act as a drain of inflow kinetic energy and a source of vorticity by the small--scale tidal torques and shear stresses. ", "introduction": "Gravitational instability is commonly accepted as the basic mechanism for structure formation on large scales. Combined with the CDM model it leads to the picture of hierarchical clustering with wide support from deep galaxy and cluster observations. During the recent phase of cosmic evolution groups and clusters of galaxies condense from large scale density enhancements, and they grow by accretion and merger processes of the environmental cosmic matter. But despite the fact that the currently favoured \\LCDM\\ model has proven to be remarkably successful on large scales (cf. WMAP results, Spergel et al. 2003, Spergel et al. 2006), recent high--resolution \\nbody\\ simulations still seem to be in contradiction with observation on sub--galactic scales. There is, for instance, the problem with the steep central densities of galactic halos as the highest resolution simulations favor a cusp with a logarithmic inner slope for the density profile of approximately $-1.2$ (Diemand, Moore \\& Stadel 2005; Fukushige, Kawai \\& Makino 2004; Power~\\ea 2003), whereas high resolution observations of low surface brightness galaxies are best fit by halos with a core of constant density (Simon~\\ea 2005; de Block \\& Bosma 2002; Swaters~\\ea 2003). A further problem relates to the overabundance of small--sized (satellite) halos; there are many more subhaloes predicted by cosmological simulations than actually observed in nearby galaxies (e.g., Moore et al.\\ 1998, Klypin et al.\\ 1999, Gottl\\o ber et al.\\ 2003). The lack of observational evidence for these satellites has led to the suggestion that they are completely (or almost completely) dark, with strongly suppressed star formation due to the removal of gas from the small protogalaxies by the ionising radiation from the first stars and quasars (Bullock et~al. 2000; Tully et~al. 2002; Somerville 2002; Hoeft~\\ea 2005). Other authors suggest that perhaps low mass satellites never formed in the predicted numbers in the first place, indicating problems with the \\LCDM\\ model in general, which is replaced with Warm Dark Matter instead (Knebe~\\ea 2002; Bode, Ostriker~\\& Turok 2001; Col{\\'\\i}n~\\ea 2000). Suggested solutions also include the introduction of self-interactions into collisionless \\nbody\\ simulations (e.g. Spergel~\\& Steinhardt 2000; Bento~\\ea 2000), and non-standard modifications to an otherwise unperturbed CDM power spectrum (e.g.\\ bumpy power spectra: Little, Knebe~\\& Islam 2003; tilted CDM: Bullock 2001c). Recent results from (strong) lensing statistics though suggest that the predicted excess of substructure is in fact required to reconcile some observations with theory (Dahle~\\ea 2003, Dalal~\\& Kochanek 2002), but this conclusion has not been universally accepted (Sand~\\ea 2004; Schechter~\\& Wambsganss 2002; Evans~\\& Witt 2003). The discovery of the mismatch between observations and simulations is a result of the increase in the resolution of \\nbody\\ simulations over the last years. This has emphasized the importance of the influence of subresolution scales on the simulated dynamics, at least when it comes to halo properties. The purpose of this work is to study the hydrodynamic--like formulation of the formation of cosmologial structure proposed recently by Dom{\\'\\i}nguez~(2000), dubbed SSE (small--size expansion). The SSE addresses explicitly the existence of a finite spatial resolution and the dynamical effect of subresolution degrees of freedom. Although developed independently, the SSE approach is close in spirit to the Large--Eddy Simulation of turbulent flow (see, e.g., Pope 2000 and refs.\\ therein). This is a method devised to simulate only the relevant, large--scale degrees of freedom according to the Navier--Stokes equations describing flow in the high--Reynolds number (i.e., turbulent) regime: physically meaningful approximations are made in order to model the coupling to the neglected, small--scale degrees of freedom. The SSE starts from the microscopic equations of motion for a set of $N$ particles under their mutual gravity and provides a set of hydrodynamic-like equations for the (coarse--grained) mass density and velocity fields. These new equations now contain \"correction terms\" which describe the effects of the coarse--graining procedure and correct for them, respectively. It therefore only appears natural to implement these correction terms into an (adaptive) mesh \\nbody\\ code where the density is treated in a coarse--grained fashion, too: mesh--based Poisson solvers frequently used for cosmological simulations smooth the particle distribution onto a grid and hence deal with a coarse--grained density field when solving for the potential via Poisson's equation. For this purpose we will adapt the open source \\nbody\\ code \\mlapm\\ (Multi-Level Adaptive Particle--Mesh)\\footnote{Available at \\texttt{http://www.aip.de/People/AKnebe/MLAPM}}. The particles of the \\nbody\\ simulations presented in this study are effectively hydrodynamical Lagrangian particles which move under the action not only of the mesh--computed gravitational force, but also of the additional, correction terms modeling sub--grid degrees of freedom in the context of the SSE. The rest of the work is structured as follows: in Sec.\\ \\ref{HAPPI} we describe the theoretical background of the SSE and provide the link to mesh--based \\nbody\\ codes. In Sec.\\ \\ref{Nbody} we present the simulation of several models (standard \\LCDM\\ model, \\LWDM\\ model, and two runs incorporating the SSE corrections). In Sec.\\ \\ref{Analysis} we perform a comparative analysis of the four runs at redshift $z=0$ from two complementary points of view: properties of the mass density and velocity fields, and properties of halos. Finally, Sec.\\ \\ref{Conclusions} includes a discussion of the results and the conclusions. ", "conclusions": "\\label{Conclusions} We have presented a series of cosmological \\nbody\\ simulations which made use of the hydrodynamic approach to the evolution of structures (Dom{\\'\\i}nguez 2000). This approach is novel in that it deals with the mass density and velocity fields with explicit account of the coarse-grained nature intrinsic to any approach of solving, for instance, Poisson's equation via Monte Carlo sampling of phase-space. This \\nbody\\ approach unavoidably introduces finite resolution effects and there have been systematic studies of the consequences in the context of cosmological structure formation (Kuhlman, Melott \\& Shandarin 1996; Splinter \\ea 1998; Moore \\ea 1998; Knebe \\ea 2000; Power \\ea 2003). \\nbody\\ simulations invariably neglect the dynamical effect of sub--resolution degrees of freedom altogether. For the first time we have run simulations including a physical model of the coupling to the neglected scales. \\nbody\\ codes are usually viewed as integrators of the Vlasov--Poisson system of equations. However, we have argued how grid--based \\nbody\\ codes such as \\mlapm\\ can be reinterpreted to integrate hydrodynamic--like equations for the mass density and velocity fields. The additional, correction term introduced in the hydrodynamic approach is proportional to a \"coupling constant\" $B$ which depends on the smoothing window used to calculate the coarse-grained fields. It is found to be $B=1/4$ for the triangular--shaped--cloud window used throughout the \\nbody\\ code \\mlapm. In order to get a better understanding of the effects of the correction term onto the evolution of cosmic structures we also performed a simulation with a higher value $B=1$ --- this later model is not physically motivated but rather serves as an \"academic toy model\" for comparison. The standard \\LCDM\\ simulation can be understood as another HAPPI run with the value $B=0$. In order to allow for a better comparison with other feasible alternatives to the concordance \\LCDM\\ model as well as to better gauge the influence of the correction term, we also simulated the evolution of the same structures in a \\LWDM\\ universe. In this work we concentrated on the comparison of the four simulations at redshift $z=0$. We analyzed the resulting structures in two complementary manners: global properties of the mass density and velocity fields, on the one hand, and specific properties of DM halos, on the other hand. We find appreciable differences between the $B\\neq 0$ runs and the reference ($B=0$) \\LCDM\\ run, even though the force due to the correction terms are for most particles one or even two orders of magnitude smaller than the total force (cf.\\ \\Fig{HAPPIdens}). Most remarkably, the correction term favors the proliferation of low--mass halos, giving the mass distribution a more \"grainy\" aspect, as well as the gain of angular momentum specially by low--mass halos, which also shows up in a velocity field with a larger vorticity. These effects are quantitatively larger as the value of $B$ increases; for $B=1/4$ the differences lie at the $(10-20)\\%$ level (and even higher for the specific angular momentum at low masses). A feature in which the $B=1/4$ and $B=1$ runs exhibit an opposite trend with respect to the $B=0$ run is the concentration of high--mass halos: the $B=1$ run results in an overabundance of high--mass halos with a lower concentration % This is paralleled by a smaller circular velocity of these halos, and by the relative lack of power in the spectrum of density fluctuations at sufficiently small scales, so that the maximum density reached in the $B=1$ run is much smaller than in the other runs. The $B=1/4$ run, however, shows precisely the opposite tendency with respect to the reference run. One can conjecture that this discrepancy between the $B=1/4$ and $B=1$ runs lies in a difference in the rate of shear and vorticity generation and of kinetic energy drainage by the correction term. A comparative study of the structures at different redshifts is required in order to obtain more precise conclusions about this issue. The relatively small quantitative differences between the $B=1/4$ and the $B=0$ runs evidenced in the properties that we have measured suggest that the $B=1/4$ correction term could be considered a small perturbation to the $B=0$ evolution. By contrast, the results of the $B=1$ run indicate that the correction term should not be treated as a perturbation in this case. Our results agree with the theoretical expectation for the qualitative behavior of the correction term, which models the effect of small--scale tidal torques and shear stresses (Dom{\\'\\i}nguez 2000, 2002; Buchert \\& Dom{\\'\\i}nguez 2005). We observed that the correction term is dominant preferentially in walls at high redshifts, and later on in filaments, regions of mass accretion onto halos, halo centers as well as in regions of particular dynamical activity (e.g., mergers), that is, regions of large gradients in the fields, in concordance with the form of the correction term~(\\ref{correction}). The term is expected to act as a drain of kinetic energy in collapsing regions: this can explain the formation of small clusters of particles, which would otherwise fly by each other --- instead, they can be gravitationally confined by a potential well that is lower than in the $B=0$ model. % This could explain the low mass halo proliferation for $B = 1/4 $ or $B = 1$, as well as the observed tendency of halos to attain a slightly more concentrated configuration in the case of $B=1/4$, when the correction term can be considered a small perturbation. If $B=1$, on the other hand, the loss of kinetic energy is apparently so important that, in some cases of halos in regions of high dynamical activity, dynamical relaxation and coalescence are slowed down noticeably, leading to a multiple-core structure. These not completely relaxed halos would then have a lower concentration and a lower mass than their LCDM counterparts, similarly to the simulation results. We further confirmed explicitly that the correction term acts as a source of vorticity. This relates directly to the gain in angular momentum of halos, which tends to increase with growing value of $B$, especially at the low--mass end of the halo distribution. Finally, we remark that our findings agree qualitatively with conclusions following from a comparative study of identical initial conditions evolved at different resolutions. We have run a series of test simulations where we either switched on the HAPPI correction term or increased the actual force resolution; both methods lead to comparable results that are in qualitative agreement with the conclusions presented here. For a more quantitative analysis we though refer the reader to a future paper in preparation where we will investigate the relationship between HAPPI simulations and higher-resolution ones in more detail. The proliferation of small halos with increasing resolution has also been reported by other authors in and around (massive) halos (Klypin et al. 1999, Moore \\ea 1998) as well as in voids (Gottl\\o ber \\ea 2003). Concerning the generation of angular momentum though, the relevance for the formation of realistic disk galaxies has yet to be determined but there are clear indications that this task requires good mass and force resolution % (Governato~\\ea 2004). In conclusion, the HAPPI implementation seems indeed to be qualitatively consistent with what one expects from higher resolution simulations and hence may provide a framework for a better understanding of resolution effects in \\nbody\\ simulations. However, further work is required to substantiate this possibility." }, "0607/astro-ph0607600_arXiv.txt": { "abstract": "At present, the heliosphere is embedded in a warm low density interstellar cloud that belongs to a cloud system flowing through the local standard of rest with a velocity near $\\sim$18 \\kms. The velocity structure of the nearest interstellar material (ISM), combined with theoretical models of the local interstellar cloud (LIC), suggest that the Sun passes through cloudlets on timescales of $\\le 10^3$--10$^4$ yr, so the heliosphere has been, and will be, exposed to different interstellar environments over time. By means of a multi-fluid model that treats plasma and neutral hydrogen self-consistently, the interaction of the solar wind with a variety of partially ionized ISM is investigated, with the focus on low density cloudlets such as are currently near the Sun. Under the assumption that the basic solar wind parameters remain/were as they are today, a range of ISM parameters (from cold neutral to hot ionized, with various densities and velocities) is considered. In response to different interstellar boundary conditions, the heliospheric size and structure change, as does the abundance of interstellar and secondary neutrals in the inner heliosphere, and the cosmic ray level in the vicinity of Earth. Some empirical relations between interstellar parameters and heliospheric boundary locations, as well as neutral densities, are extracted from the models. ", "introduction": "The heliosphere is a low density cavity that is carved out from the local interstellar medium (LISM) by the solar wind. The size and particle content of the heliosphere are determined by the solar wind -- LISM interaction, and they vary in response to the Galactic environment of the Sun as the Sun and interstellar clouds move through space. The path of the Sun has taken us through the Local Bubble void \\citep[galactic longitudes 180\\deeg$\\ltsim$\\glong$\\ltsim$270\\deeg,][]{FrischYork:1986}, and we have recently ($\\ltsim$ 10$^3 - 10^5$ yr ago, depending on cloud shapes and densities) entered a clumpy flow of low density interstellar material \\citep[][]{Frisch94}. This clumpy flow, the ``cluster of local interstellar cloudlets'' (CLIC), is flowing away from the Sco-Cen association and extends 10--30 pc into the Galactic center hemisphere and $\\ltsim$3 pc for many directions in the anticenter hemisphere. Inhomogeneities in the CLIC create temporal variations in the dynamic interstellar pressure at the solar location, which may produce significant variations in heliosphere properties over geologically short timescales \\citep{Frisch:1993a,Frisch:1997,ZankFrisch99,Florinskietal03a,Frischetal:2002,Frisch:2004igpp}. The heliosphere itself is a dynamically changing object which is highly sensitive to interstellar pressure \\citep[e.g.][]{Holzer:1989,Zank99}. The interaction of the ISM with the fully ionized solar wind gives rise to the heliospheric morphology which includes the heliopause (HP), a tangential discontinuity separating solar wind and LISM, and the termination shock (TS) where the solar wind becomes subsonic and is diverted downstream to form a heliotail. Depending on the pressure of the surrounding interstellar material, an interstellar bow shock (BS) may form upwind of the heliopause. These general boundaries are created by the plasma interaction, yet the presence of neutral H and its coupling to the plasma protons via charge exchange greatly influences the details of the heliospheric morphology and the location of its boundaries (see \\citet{Zank99} for a review). The sensitivity of the heliosphere to variations in the physical characteristics of the interstellar cloud surrounding the solar system is poorly understood, and in this paper we focus on heliosphere variations due to encounters with a range of low density clouds such as expected in the immediate past and future solar history. Different interstellar environments may produce noticeable changes in the interplanetary environment of the inner heliosphere, as indicated by the amount of neutral H, anomalous, and galactic cosmic rays (GCR) at 1 AU. There is some evidence that lunar soils contain an archive of isotopic abundances that are different from the particle environment of the present era \\citep{rfw00}, and antarctic ice cores show signatures that may be interpreted as cosmic ray background variations at Earth \\citep{Be10,SonettJokipii:1987,FlorinskiZank05}. These possibilities have motivated our study of the behavior of the global heliosphere under variable boundary conditions resulting from passage through interstellar clouds. Given the inhomogeneity of the local solar neighborhood and the galactic environment in general, we test the heliosphere response to a range of local interstellar boundary parameters using about two dozen specific parameter sets. Our choices are justified in \\S \\ref{ss:nearbyism}. Four highlights of the corresponding heliospheric models are detailed in \\S \\ref{ss:exmodels}. The results of all the heliospheric models calculated for this study suggest relationships of the heliospheric boundary locations and the neutral particle densities with the interstellar parameters, discussed in \\S \\ref{ss:results}. The synopsis of all the individual model results through these relations is a quantitative expression of the sensitivity of the heliosphere to changing interstellar boundary conditions. Furthermore, the relations allow for a prediction of boundary locations and particle content for heliospheres with yet different boundary parameters, without actually engaging in a complex, non-linear global heliosphere simulation. This predictive power can also be used in the emerging field of astrospheres, which are the analogues of heliospheres around solar-like cool stars. We discuss the response of the global heliosphere to variable interstellar properties, and speculate on aspects of the implications of these variations for the 1 AU location of the Earth in \\S \\ref{ss:discuss}. ", "conclusions": "On its path through the galaxy, the Sun has encountered (and will encounter) different interstellar environments. This motivates a parameter study to investigate the response of the heliosphere to these changing conditions under the assumption of a constant solar wind. For conditions that are not too far from the contemporary LISM environment, the following findings emerge from analyzing \\nummodels\\ self-consistent multi-fluid models. \\begin{enumerate} \\item Allowing generous assumptions about the LIC morphology, the LIC column density towards nearby stars indicates the Sun first encountered the LIC gas within the past 40,000 yr, and the CLIC within the past $\\sim$60,000/$\\tilde{f}$ yr (where $\\tilde{f}$ represents the fraction of space filled by the CLIC). The Sun is expected to exit the LISM gas cloudlet, which is characterized by the common LIC velocity, sometime within the next $\\sim$0--4000 yr. In general, passage through interstellar clouds will lead to variations in the heliosphere boundary conditions over timescales possibly as short as 10$^3$ yr. Nearby ISM generally resembles low column density ISM observed elsewhere. \\item The size of the heliosphere is determined by the balance of solar wind and interstellar pressure. For the investigated parameter range, in which the LISM is mostly ram-pressure dominated, the upwind termination shock distance can be estimated by equation (\\ref{eq-ts2corr}), using equations (\\ref{eq-pb}) and (\\ref{eq-TS2}). This relation is derived from a pressure balance argument modified by an empirical factor expressing the efficiency of the neutral pressure contribution to the overall interstellar pressure. \\item Heliocentric distances of interest such as the heliopause, the bow shock, or the upwind and downwind termination shock scale linearly with each other (e.g.\\ $ r_{\\rm HP} = 1.39 \\,\\, r_{\\rm TS}$, $r_{\\rm BS} = 1.95 \\,\\, r_{\\rm HP}$). Therefore, when the upwind termination shock distance is predicted in an absolute way as above, the other distances can be predicted as well. However, the scalability and predictability of the heliosphere size with these relations are only applicable to parameter sets in which the LISM flow is supersonic. The subsonic cases, when the Sun is surrounded by hot plasma or alternately when the Sun and the surrounding interstellar cloud are comoving in space, obey a different set of correlations, and are generally more difficult to model numerically. \\item For low interstellar velocities, the heliosheath and heliotail plasma are subsonic throughout, and the ratio of downwind to upwind termination shock distance (TS asymmetry) is 2.1. For higher velocities, the heliosphere assumes a rocket shape, with a modified pressure balance in the downwind directions. \\item Neutral hydrogen results such as the filtration ratio, the peak hydrogen wall density, or the density at 5 AU upwind of the Sun, correlate with each other. Their absolute value is weakly correlated to the interstellar velocity with which the neutrals arrive at the respective heliosphere, as the charge exchange mean free path depends on this velocity, and higher velocities shorten the heliocentric distances to the heliospheric boundaries. \\item For encounters with a high density interstellar cloud ($\\sim$15 \\cc, about 50 times the contemporary value), the particle fluxes arriving at Earth orbit, including interstellar neutrals, neutral solar wind, and cosmic rays will increase markedly. These changes potentially affect Earth's atmosphere and its climate. The changes in particle fluxes just due to a higher interstellar velocity are less pronounced. \\item For the period when the Sun was embedded in the Local Bubble, particle fluxes were reduced substantially. Secondary particles like anomalous cosmic rays and neutral solar wind were entirely absent, and the galactic cosmic ray flux arriving at Earth was comparable to the contemporary flux, or even reduced, depending on the modulation model. \\end{enumerate}" }, "0607/astro-ph0607166_arXiv.txt": { "abstract": "We present calculations of the electronic structure of various atoms and molecules in strong magnetic fields ranging from $B=10^{12}$~G to $2\\times10^{15}$~G, appropriate for radio pulsars and magnetars. For these field strengths, the magnetic forces on the electrons dominate over the Coulomb forces, and to a good approximation the electrons are confined to the ground Landau level. Our calculations are based on the density functional theory, and use a local magnetic exchange-correlation function which is tested to be reliable in the strong field regime. Numerical results of the ground-state energies are given for H$_N$ (up to $N=10$), He$_N$ (up to $N=8$), C$_N$ (up to $N=5$), and Fe$_N$ (up to $N=3$), as well as for various ionized atoms. Fitting formulae for the $B$-dependence of the energies are also given. In general, as $N$ increases, the binding energy per atom in a molecule, $|E_N|/N$, increases and approaches a constant value. For all the field strengths considered in this paper, hydrogen, helium, and carbon molecules are found to be bound relative to individual atoms (although for $B$ less than a few $\\times 10^{12}$~G, carbon molecules are very weakly bound relative to individual atoms). Iron molecules are not bound at $B\\alt 10^{13}$~G, but become energetically more favorable than individual atoms at larger field strengths. ", "introduction": "Neutron stars (NSs) are endowed with magnetic fields far beyond the reach of terrestrial laboratories \\citep{meszaros92,reisenegger05,harding06}. Most of the $\\sim 1600$ known radio pulsars have surface magnetic fields in the range of $10^{11}-10^{13}$~G, as inferred from their measured spin periods and period derivatives and the assumption that the spindown is due to magnetic dipole radiation. A smaller population of older, millisecond pulsars have $B\\sim 10^8-10^9$~G\\@. For about a dozen accreting neutron stars in binary systems, electron cyclotron features have been detected, implying surface fields of $B\\sim 10^{12}-10^{13}$~G\\@. An important development in astrophysics in the last decade centered on the so-called anomalous x-ray pulsars and soft gamma repeaters \\citep{woods05}: there has been mounting observational evidence in recent years that supports the idea that these are magnetars, neutron stars whose radiations are powered by superstrong magnetic fields of order $10^{14}-10^{15}$~G or higher \\citep{duncan92,thompson95,thompson96}. By contrast, the highest static magnetic field currently produced in a terrestrial laboratory is $5\\times 10^5$~G; transient fields approaching $10^9$~G have recently been generated during high-intensity laser interactions with dense plasmas \\citep{wagner04}. It is well-known that the properties of matter can be drastically modified by strong magnetic fields found on neutron star surfaces. The natural atomic unit for the magnetic field strength, $B_0$, is set by equating the electron cyclotron energy $\\hbar\\omega_{Be}=\\hbar (eB/m_ec)=11.577\\,B_{12}$~keV, where $B_{12}=B/(10^{12}~{\\rm G})$, to the characteristic atomic energy $e^2/a_0=2\\times 13.6$~eV (where $a_0$ is the Bohr radius): \\be B_0={m_e^2e^3c\\over\\hbar^3}=2.3505\\times 10^9\\, {\\rm G}. \\label{eqb0} \\ee For $b=B/B_0\\agt 1$, the usual perturbative treatment of the magnetic effects on matter (e.g., Zeeman splitting of atomic energy levels) does not apply. Instead, in the transverse direction (perpendicular to the field) the Coulomb forces act as a perturbation to the magnetic forces, and the electrons in an atom settle into the ground Landau level. Because of the extreme confinement of the electrons in the transverse direction, the Coulomb force becomes much more effective in binding the electrons along the magnetic field direction. The atom attains a cylindrical structure. Moreover, it is possible for these elongated atoms to form molecular chains by covalent bonding along the field direction. Interactions between the linear chains can then lead to the formation of three-dimensional condensed matter \\citep{ruderman74,ruder94,lai01}. Our main motivation for studying matter in such strong magnetic fields arises from the importance of understanding neutron star surface layers, which play a key role in many neutron star processes and observed phenomena. Theoretical models of pulsar and magnetar magnetospheres depend on the cohesive properties of the surface matter in strong magnetic fields \\citep{ruderman75,arons79,usov96,harding98,beloborodov06}. For example, depending on the cohesive energy of the surface matter, an acceleration zone (``polar gap'') above the polar cap of a pulsar may or may not form. More importantly, the surface layer directly mediates the thermal radiations from neutron stars. The advent of x-ray telescopes in recent years has made detailed study of neutron star surface emission a reality. Such study can potentially provide invaluable information on the physical properties and evolution of NSs: equation of state at supernuclear densities, superfluidity, cooling history, magnetic field, surface composition, different NS populations, etc. (see, e.g., Ref.~\\citep{yakovlev04}). More than two dozen isolated neutron stars (including radio pulsars, radio-quiet neutron stars and magnetars) have clearly detected thermal surface emission \\citep{kaspi05,haberl05,harding06}. While some neutron stars show featureless spectra, absorption lines or features have been detected in half a dozen or so systems \\citep{haberl05}. Indeed, many of the observed neutron stars have sufficiently low surface temperatures and/or high magnetic fields, such that bound atoms or molecules are expected to be present in their atmospheres \\citep{lai97,potekhin99,ho03,potekhin04}. It is even possible that the atmosphere is condensed into a solid or liquid form from which radiation directly emerges \\citep{lai97,vanadelsberg05,lai01}. Thus, in order to properly interpret various observations of neutron stars, it is crucial to have a detailed understanding of the properties of atoms, molecules and condensed matter in strong magnetic fields ($B\\sim 10^{11}$-$10^{16}$~G). \\subsection{Previous works} H and He atoms at almost all field strengths have been well studied \\citep{ruder94,jones99,alhujaj03}, including the nontrivial effect associated with the center-of-mass motion of a H atom \\citep{potekhin98}. \\citet{neuhauser87} presented numerical results for several atoms up to $Z=26$ (Fe) at $B\\sim 10^{12}$~G based on calculations using a one-dimensional Hartree-Fock method (see also Ref.~\\citep{mori02} for $Z$ up to 10). Some results [based on a two-dimensional (2D) mesh Hartree-Fock method] for atoms (up to $Z=10$) at the field strengths $B/B_0=0.5-10^4$ are also available \\citep{ivanov00,alhujaj04a, alhujaj04b}. The Hartree-Fock method is approximate because electron correlations are neglected. Due to their mutual repulsion, any pair of electrons tend to be more distant from each other than the Hartree-Fock wave function would indicate. In zero-field, this correlation effect is especially pronounced for the spin-singlet states of electrons for which the spatial wave function is symmetrical. In strong magnetic fields ($B\\gg B_0$), the electron spins (in the ground state) are all aligned antiparallel to the magnetic field, and the multielectron spatial wave function is antisymmetric with respect to the interchange of two electrons. Thus the error in the Hartree-Fock approach is expected to be less than the $1\\%$ accuracy characteristic of zero-field Hartree-Fock calculations \\citep{neuhauser87,schmelcher99}. Other calculations of heavy atoms in strong magnetic fields include Thomas-Fermi type statistical models \\citep{fushiki92,lieb94a,lieb94b} and density functional theory \\citep{jones85,jones86, kossl88,relovsky96}. The Thomas-Fermi type models are useful in establishing asymptotic scaling relations, but are not adequate for obtaining accurate binding and excitation energies. The density functional theory can potentially give results as accurate as the Hartree-Fock method after proper calibration is made \\citep{vignale87,vignale88}. Quantitative results for the energies of hydrogen molecules H$_N$ with $N=2,3,4,5$ in a wide range of field strengths ($B\\gg B_0$) were obtained (based on the Hartree-Fock method) by Lai {\\it et al.}~\\citep{lai92,lai01} and molecular excitations were studied in Refs.~\\citep{lai96,schmelcher01} (more complete references can be found in Ref.~\\citep{lai01}). Quantum Monte Carlo calculations of H$_2$ in strong magnetic fields have been performed \\citep{ortiz95}. Some numerical results of He$_2$ for various field strengths are also available \\citep{lai01}. Hartree-Fock results of diatomic molecules (from H$_2$ up to C$_2$) and several larger molecules (up to H$_5$ and He$_4$) at $B/B_0=1000$ are given in Ref.~\\citep{demeur94}. \\subsection{Plan of this paper} In this paper and its companion paper \\citep{medin06a}, we develop a density-functional-theory calculation of the ground-state energy of matter for a wide range of magnetic field strengths, from $10^{12}$~G (typical of radio pulsars) to $2\\times 10^{15}$~G (magnetar fields). We consider H, He, C, and Fe, which represent the most likely composition of the outermost layer of neutron stars (e.g., Ref.~\\citep{harding06}). The present paper focuses on atoms (and related ions) and small molecules. Because of additional complications related to the treatment of band structure, calculations of infinite molecular chains and condensed matter are presented in Ref.~\\citep{medin06a}. Our calculations are based on density functional theory \\citep{hohen64,kohn65,jones89}. As mentioned above, the Hartree-Fock method is expected to be highly accurate, particularly in the strong field regime where the electron spins are aligned with each other. In this regime the density functional method is not as accurate, due to the lack of an exact correlation function for electrons in strong magnetic fields. However, in dealing with systems with many electrons, the Hartree-Fock method becomes increasingly impractical as the magnetic field increases, since more and more Landau orbitals (even though electrons remain in the ground Landau level) are occupied and keeping track of the direct and exchange interactions between electrons in various orbitals becomes computationally rather tedious. Our density-functional calculations allow us to obtain the energies of atoms and small molecules and the energy of condensed matter using the same method, thus providing reliable cohesive energy of condensed surface of magnetic neutron stars, a main goal of our study. Compared to previous density-functional-theory calculations \\citep{jones85,jones86, kossl88,relovsky96}, we use an improved exchange-correlation function for highly magnetized electron gases, we calibrate our density functional code with previous results (when available) based on other methods, and (for calculations of condensed matter) adopt a more accurate treatment of the band structure. Moreover, our calculations extend to the magnetar-like field regime ($B\\sim 10^{15}$~G). Note that in this paper we neglect the motions of the nuclei, due to electron-nucleus interactions or finite temperatures. The center-of-mass motions of the atoms and molecules induce the motional Stark effect, which can change the internal structure of the bound states (see, e.g., Refs.~\\citep{lai01,potekhin98}). Such issues are beyond the scope of this paper. After summarizing the approximate scaling relations for atoms and molecules in strong magnetic fields in Sec.~II, we describe our method in Sec.~III and present numerical results in Sec.~IV\\@. Some technical details are given in the Appendix. ", "conclusions": "We have presented density-functional-theory calculations of the ground-state energies of various atoms and molecular chains (H$_N$ up to $N=10$, He$_N$ up to $N=8$, C$_N$ up to $N=5$, and Fe$_N$ up to $N=3$) in strong magnetic fields ranging from $B=10^{12}$~G to $2\\times10^{15}$~G\\@. These atoms and molecules may be present in the surface layers of magnetized neutron stars, such as radio pulsars and magnetars. While previous results (based on Hartree-Fock or density-functional-theory calculations) are available for some small molecules at selected field strengths (e.g., Refs.~\\citep{lai92,lai01,demeur94}) many other systems (e.g., larger C molecules and Fe molecules) are also computed in this paper. We have made an effort to present our numerical results systematically, including fitting formulae for the $B$-dependence of the energies. Comparison with previous results (when available) show that our density-functional calculations tend to overestimate the binding energy $|E_N|$ by about $10\\%$. Since it is advantageous to use the density functional theory to study systems containing large number of electrons (e.g., condensed matter; see Ref.~\\citep{medin06a}), it would be useful to find ways to improve upon this accuracy. At $B_{12} \\ge 1$, hydrogen, helium, and carbon molecules are all more energetically favorable than their atomic counterparts (although for carbon, the relative binding between the atom and molecule is rather small), but iron is not. Iron molecules start to become bound at $B_{12}\\agt 10$, and are not decidedly more favorable than isolated atoms until about $B_{12}=100$. For the bound molecules considered here, the ground-state energy per atom approaches an asymptotic value as $N$ gets large. The molecule then essentially becomes a one-dimensional infinite chain. We will study such condensed matter in our companion paper \\citep{medin06a}." }, "0607/astro-ph0607350_arXiv.txt": { "abstract": "The idea of a unified model for all astrophysical jets has been considered for quite some time. We present here a few scaling laws relevant to all type of astrophysical jets, analogous to those of \\citet{sams96} which are widely used for astrophysical black holes. We use Buckingham's $ \\Pi $ theorem of dimensional analysis to obtain a family of dimensional relations among the physical quantities associated to astrophysical jets. ", "introduction": "Although the first report of an astrophysical jet was made by \\citet{curtis}, these objects were extensively studied much later with radio astronomy techniques \\citep{reber}. Quasars, and radiogalaxies were discovered and later gathered in a unified model which proposed a dusty torus around the nucleus of the source \\citep{antonucci85}. Years later, some galactic sources showed similar features to the ones presented by quasars and radiogalaxies, i.e. relativistic fluxes, a central engine, symmetrical collimated jets, radiating lobes, and apparent superluminal motions \\citep[cf.][]{sunyaev91}. These objects are usually identified as \\( \\mu \\)--quasars. Optical and X-ray observations showed other similar non--relativistic sources in the galaxy associated to H--H objects \\citep[cf.][]{pino04}. Lately, the strong explosions found in long Gamma Ray Bursts (GRB) have been modelled as collapsars, in which a jet of a very short lifetime is associated to the observed phenomena \\citep[cf.][]{kulkarni99,castro99}. The similarities between all astrophysical jets, mainly those between quasars and micro--quasars, and the scaling laws for black holes proposed by \\citet{sams96} and \\citet{rees98} made us search for the possible existence of some scaling laws that may occur to astrophysical jets in terms of very simple physical parameters, such as the magnetic field associated to the accretion disc, accretion rate and mass of the central object. These sort of relations have been studied in a very different way by \\citet{heinz03,heinz05} giving scalings between the flux \\( F_\\nu \\) at a frequency \\( \\nu \\) and the mass of the central object. The present work presents a few mathematical relations that appear naturally as a consequence of dimensional analysis and Buckingham's \\( \\Pi \\) theorem. We begin by considering some of the most natural physical dimensional quantities that have to be included to describe some of the physical phenomena related to all classes of jets. We then calculate the dimensional relations associated to these quantities. Finally, we briefly discuss these relations and their physical relevance to astrophysical jets. ", "conclusions": "According to the previous analysis astrophysical jets exist due to a precise combination of electromagnetic, mechanical and gravitational processes independently of the physical mechanisms behind the central engine. A trivial dimensionless parameter that is obtained using Buckingham's \\( \\Pi \\) theorem of dimensional analysis applied to equation~\\eqref{ec4} is the ratio \\( v / c \\). With this, it is possible to form another dimensionless parameter given by \\( \\Pi_2 \\Pi_4^2 \\Pi_5 / \\left( v / c \\right)^2 \\), which leads to the dimensionless quantity \\( \\Lambda := \\rho G l^2 / v^2 \\) used by \\citet{mendoza05} in order to obtain a maximum length for an astrophysical jet. As explained by \\citeauthor{mendoza05}, this maximum size is most probably determined by the interaction of the jet and its cocoon with their surrounding environment, leading to the generation of Kelvin--Helmholtz instabilities. Applying the results of equations \\eqref{r_j} and \\eqref{ls'} to GRB jets with a canonical magnetic field \\( B \\sim 10^{16} \\, \\textrm{G} \\), leads to wrong output kinetic luminosities and typical sizes of GRB jets. To correct these, proportionality factors of \\( \\sim 10^9 \\) in the jet's length \\( r_j \\) and of \\( 10^{-5} \\) in the kinetic power \\( L \\) in equations \\eqref{r_j} and \\eqref{ls'} have to be used respectively. The reason for this might be due to the fact that these jets have very short life times and so, they hardly resemble a traditional steady jet. Also, there might be some particular physical mechanisms that make a dimensionless combination changing the proportionality factors in equations \\eqref{r_j} and \\eqref{ls'} in such a way that they give the correct value needed for these class of jets. There are well known kinetic luminosities that appear in the literature related to the ejection of jets from different sources. As an example, in the \\citet{blandfordpayne} model the luminosity takes the following form \\begin{equation} L = B_\\text{p}^2 R^3 \\Omega, \\label{bland-payne} \\end{equation} \\noindent where \\( \\Omega \\) is the angular velocity of the poloidal component of the magnetic field \\( B_\\text{p} \\) and \\( R \\) is the size of the rotating region. It is possible to obtain equation~\\eqref{bland-payne} with the model presented in this article if we proceed as follows. Let us include in the functional relation~\\eqref{ec4} an important parameter of the problem, namely an angular velocity \\( \\Omega \\), so that \\begin{equation} L = L(\\dot{M},\\, M,\\, c,\\, G,\\,B,\\, l,\\, v,\\, \\rho,\\, \\Omega). \\label{lum-gen} \\end{equation} \\noindent It is then possible to build another dimensionless parameter \\( \\Pi_7 \\) given by \\begin{equation} \\Pi_7 := \\frac{ \\Omega M }{ \\dot{M} }. \\label{pi7} \\end{equation} The \\citet{blandfordpayne} kinetic luminosity is obtained using equations \\eqref{pis} and \\eqref{pi7} with the introduction of a new dimensionless parameter \\( \\Pi_8 := \\Pi_3^2 \\, \\Pi_4^3 \\, \\Pi_7 \\, / \\, \\Pi_1 \\). In the same manner, let us define a new dimensionless parameter \\( \\Pi_9 := \\Pi_3^2 \\, \\Pi_4^4 \\, \\Pi_7^2 \\, / \\, \\Pi_1 \\) which reproduces the general dimensional shape of the kinetic luminosity of the \\citet{blandfordznajek} model, giving \\citep{meier02} \\begin{equation} L \\propto \\frac{ 1 }{ c } B^2 R^4 \\Omega^2. \\label{blandford-znajek} \\end{equation} \\noindent The constant of proportionality in this equation has a value of \\( \\sim \\, 0.1 \\, \\)--\\( \\, 0.03 \\) depending on the geometry of the problem. Of all our results, it is striking the fact that the jet power is inversely proportional to the accretion rate associated with it. This is probably due to the following. For a fixed value of the mass of the central object (in any case, for the time that accretion takes place, the mass of the central object does not increase too much) when the accretion mass rate increases, the magnetic field lines anchored to the plasma tend to pack up and thus, the field's intensity increases in such a way as to get the correct result given by equation~\\eqref{ls'}. All the results presented in this article are in agreement with powerful extragalactic jets and jets associated with \\( \\mu \\)--quasars and Herbig--Haro objects. They are also in agreement with jets associated to long gamma--ray bursts if proportionality factors are adjusted for the corresponding physical quantities. Our main result, that the kinetic luminosity of the jet as a function of the magnetic field, the accretion rate and the mass of the central object is of general validity, independently of whether the same physical mechanism produces jets from galactic sources (i.e. Herbig--Haro, \\( \\mu \\)-quasar and possibly jets associated to Gamma Ray Bursts) or extragalactic ones. Indeed, if the physics behind the accretion--ejection mechanism that occurs in jets works differently for each class of jets, then the luminosity output will be given by different mathematical expressions. For example, the \\citet{blandfordpayne} and \\citet{blandfordznajek} models, are given by equations \\eqref{bland-payne} and \\eqref{blandford-znajek} respectively. However, for these two particular cases, we have proven above that if the Luminosity is a general function given by equation \\eqref{lum-gen}, then both cases lead to the luminosity relation \\eqref{ls}. On the other hand, if the physical mechanism that generates jets at all scales is the same, relation \\eqref{ls} is also of general validity if the luminosity is a function described by \\eqref{ec4}, or equivalently by \\eqref{lum-gen}." }, "0607/astro-ph0607399_arXiv.txt": { "abstract": "Following the recent outburst of the recurrent nova RS~Oph on 2006 Feb 12, we measured its near-infrared size using the IOTA, Keck, and PTI Interferometers at multiple epochs. The characteristic size of $\\sim$3~milliarcseconds hardly changed over the first 60 days of the outburst, ruling out currently-popular models whereby the near-infrared emission arises from hot gas in the expanding shock. The emission was also found to be significantly asymmetric, evidenced by non-zero closure phases detected by IOTA. The physical interpretation of these data depend strongly on the adopted distance to RS~Oph. Our data can be interpreted as the first direct detection of the underlying RS~Oph binary, lending support to the recent ``reborn red giant'' models of Hachisu \\& Kato. However, this result hinges on an RS~Oph distance of $\\simle$540pc, in strong disagreement with the widely-adopted distance of $\\sim$1.6~kpc. At the farther distance, our observations imply instead the existence of a non-expanding, dense and ionized circumbinary gaseous disk or reservoir responsible for the bulk of the near-infrared emission. Longer-baseline infrared interferometry is uniquely suited to distinguish between these models and to ultimately determine the distance, binary orbit, and component masses for RS~Oph, one of the closest-known (candidate) SNIa progenitor systems. ", "introduction": "Most astronomers are familiar with Classical Novae, exploding stars in which an accreting white dwarf (WD) in an interacting binary system accumulates enough material for it to become unstable to hydrogen burning. The expanding blastwave from one such event Nova Aql 2005 was recently observed by infrared (IR) interferometry \\citep{lane2005aas} and the geometric distance was estimated based on velocities from spectral line observations. This result is consistent with the ``optically-thick fireball'' model which has been successfully used for twenty years to explain the time-evolution of the spectral energy distribution of classical novae \\citep{gehrz1988}. While classical novae are expected to recur, very few actually have in recorded history. RS~Oph is one of the handful of so-called ``recurrent novae'' with (now) 6~outbursts since 1898 \\citep{warner1976}. The most recent outburst occurred on February 12, 2006 \\citep{narumi2006}, and this unusual event motivated intense monitoring by the IR interferometry community. The special nature of RS~Oph is thought to stem from two causes. First, the WD is likely extremely close to the Chandrasekhar limit, since the amount of hydrogen needed to trigger an outburst decreases dramatically as the WD mass increases. Indeed, detailed models indicate the WD mass is within 1\\% of exploding as a type Ia supernova \\citep[e.g.,][]{hachisu2001}. Second, the mass-losing companion for RS~Oph is a red giant (RG) with a wind, providing a high density medium for accretion onto the WD as well as for the exploding blastwave to interact with. \\citet{bode1985} have produced the most successful model for recurrent novae, drawing a clear analogy to extragalactic supernovae and explaining the radio and X-ray light curves in this context. \\citet{evans1988} were the first to study in detail the IR time evolution of a recurrent nova. They monitored closely the 1-3.5 micron flux of RS~Oph for about 3 years after the 1985 eruption. They found the light curve had a characteristic (2-mag) decay time scale of about 30~days, and compared their observations to the generic predictions of the \\citet{bode1985} model. They concluded that their observations could come from the hot, post-shock gas -- this model would predict that the IR source should be seen linearly expanding at a rate of about $\\sim$1~milliarcseconds per day at a distance of 1.6kpc. This distance estimate is based (most securely) on the expanding size of the radio emission observed in 1985 by \\citet{hjellming1986} and \\citet{taylor1989}, assuming association with the forward shock \\citep[new radio data re-confirm the 1985 observations;][]{rupen2006,obrien2006}. As will become clear, the distance to RS~Oph is key to the interpretation of the IR interferometry data presented here. Challenging this interpretation, \\citet[][see also Kato 1991]{hachisu2001} have recently produced a comprehensive model for recurrent novae meant to explain a wide range of the known novae properties, and makes a specific prediction for the origin of the near-infrared (NIR) continuum which is very different from \\citet{evans1988}. Following onset of the thermonuclear runaway of the hydrogen shell around the WD, the shell expands to AU~size, in effect turning the WD back into a red giant. The shell stably burns hydrogen for a few weeks, shrinking back to the size of white dwarf. According to this model, hot post-shock gas plays no role in forming the IR continuum. Furthermore, \\citet{hachisu2001} prefer a much closer distance of 600~pc implying a binary separation of 2.9 milliarcseconds, easily detectable with current interferometers. In this Letter, we report first-ever size measurements for RS~Oph in the NIR using long-baseline interferometry. Our results are surprising, ruling out the favored expanding fireball model, raising doubts about the established distance to RS~Oph, and motivating a new model for the NIR emission. ", "conclusions": "\\label{discussion} Because the expanding fireball model fails to explain the nearly static size scale of the IR emission, we now seek suitable alternative emission mechanisms for the time-variable IR emission from the recurrent nova RS~Oph. We have pursued the reborn red giant (thermonuclear runaway) model of \\citet{hachisu2001} and found that indeed our 3-interferometer combined dataset can be explained by a simple binary model with separation of $\\sim$3.2~milliarcsecond, PA ~30$\\arcdeg$ E of N, and a brightness ratio varying from 2.5:1 to 5:1. Next we subject the binary hypothesis to further scrutiny. \\subsection{Binary Interpretation of Near-Interferometry Data} \\label{binary2} Based on single-line radial velocity data, \\citet{fekel2000} finds the RS~Oph binary orbit to be roughly circular with a period of 455.72$\\pm$0.83 days and mass function $f=0.221\\pm0.038$~\\msun. RGs in symbiotic systems are typically 1-3\\msun \\citep{dobrzycka1994} and we expect recurrent novae to contain a Chandrasekhar mass WD (1.4\\msun); these facts combined with the known mass function rule out RG masses greater than 2\\msun. Assuming the RS~Oph system mass to be 2.4-3.4$\\msun$, we find the component separation to be 1.55-1.74~AU, or (unprojected) 2.59-2.90~mas at the 600pc distance preferred by \\citet{hachisu2001} -- only slightly smaller than our observed separation of 3.2~mas. Since the RS~Oph outburst took place only 1 month before maximum redshifted velocity \\citep{fekel2000}, our measured binary parameters represent the true orbital semi-major axis and orbital $\\Omega$ for RS~Oph with only weak $sin(i)$ dependencies. Thus, a small reduction in the distance estimate (540~pc) brings the interferometer binary model in agreement with expectations from Kepler's laws. The binary model fits (Table~\\ref{table_model}) show evidence for a change in the brightness ratio over time. While the \\citet{hachisu2001} theory predicts a time-changing brightness ratio, it is beyond the scope of this paper to test the compatibility with the observed IR light curves due to complications from the role of the irradiated RG photosphere and the presence of a post-outburst WD accretion disk. \\subsection{Circumbinary Reservoir of Hot Gas} The distance estimate of $\\simle$540~pc derived in the last section stands in strong contrast to estimates more commonly adopted in the literature. The most significant constraints on distance are set by resolved radio observations of the previous and current burst \\citep{hjellming1986,taylor1989,rupen2006,obrien2006}. By assuming that observed radio proper motions (on the sky) can be ascribed to the fast-moving ejecta or forward shock, workers consistently derive a distance of $\\sim$1600~pc. Similarly large distances were found considering interstellar UV absorption lines \\citep{snijders1987} and HI absorption measurements \\citep{hjellming1986}. Given the strength of the evidence, we now consider the implications of the $d=1600$~pc distance. This distance would rule out the binary interpretation of the near-IR interferometry data laid out in \\S\\ref{binary2}, given existing binary constraints. Instead, we hypothesize that the IR emission arises from a quasi-stationary\\footnote{ It is possible to fit our data with a an expanding wind or jet component, but this requires fine-tuning the relative proportions of multiple components and/or a very asymmetric jetlike emission oriented perpendicular to our long (northeast) baselines. These possibilties will be investigated more thoroughly in future work with an extended dataset.} hot gas reservoir that contributes a combination of emission lines and free-free/bound-free emission in the NIR bands. The observed FWHM of $\\sim$3 mas is $\\sim$5~AU at 1600~pc, about 3$\\times$ the expected RS~Oph binary separation. This size is reasonable for a circumbinary disk or reservoir of hot gas, perhaps kept ionized by the outward moving blastwave or soft X-ray luminosity from the WD itself following outburst. This gas reservoir might be analogous to the ``fallback disk'' inferred to form after some supernovae \\citep[e.g.,][]{wang2006}. Clearly, the hypothesized gas reservoir must be elongated and somewhat off-center with respect to the central source in order to fit the combined IOTA, KI, and PTI interferometry data, especially the non-zero closure phases. It is beyond the scope of this Letter to investigate the details here, and we defer development of this model to a future paper. \\subsection{Future Work} While we have ruled out the important class of expanding fireball models for explaining the IR emission from the recurrent nova RS~Oph, more work lies ahead to test the other emission mechanisms discussed in this Letter. A future study will attempt to synthesize a {\\em self-consistent} model that can explain the time evolution of the IR spectrum, NIR and mid-IR interferometer data, and multi-wavelength light curves at the same time\\footnote{Note added: future modeling should address the asymmetric radio nebula and jet observed by \\citet{obrien2006}.}. If the close distance $d\\simle540$~pc is confirmed, we have a spectacular opportunity to study in detail a likely SNIa progenitor and to learn about unexpected shock physics controlling the non-thermal radio emission. Alternatively, the further distance $d\\sim1600$~pc suggests we have discovered a significant and new component to the RS Oph Nova remnant and future work will characterize the hot circumbinary gas reservoir for the first time." }, "0607/astro-ph0607020_arXiv.txt": { "abstract": "{This paper is part of the series presenting the final results obtained by the ESO Imaging Survey (EIS) project. It presents new \\j\\ and \\k\\ data obtained from observations conducted at the ESO 3.5m New Technology Telescope (NTT) using the SOFI camera. These data were taken as part of the Deep Public Survey (DPS) carried out by the ESO Imaging Survey program, significantly extending the earlier optical/infrared EIS-DEEP survey presented in a previous paper of this series. The DPS-IR survey comprises two observing strategies: shallow \\k\\ observations providing nearly full coverage of pointings with complementary multi-band (in general $UBVRI$) optical data obtained using ESO's wide-field imager (WFI) and deeper \\j\\ and \\k\\ observations of the central parts of these fields. Currently, the DPS-IR survey provides a coverage of roughly 2.1~square~degrees ($\\sim300$~SOFI pointings) in \\k\\ with 0.63~square~degrees to fainter magnitudes and also covered in \\j, over three independent regions of the sky. The goal of the present paper is to briefly describe the observations, the data reduction procedures, and to present the final survey products which include fully calibrated pixel-maps and catalogs extracted from them. The astrometric solution with an estimated accuracy of $\\lesssim0\\farcs15$ is based on the USNO catalog and limited only by the accuracy of the reference catalog. The final stacked images presented here number 89 and 272, in \\j\\ and \\k, respectively, the latter reflecting the larger surveyed area. The \\j\\ and \\k\\ images were taken with a median seeing of $0\\farcs77$ and $0\\farcs8$. The images reach a median $5\\sigma$ limiting magnitude of $J_{AB}\\sim23.06$ as measured within an aperture of 2\\arcsec, while the corresponding limiting magnitude in \\k$_{AB}$ is $\\sim21.41$ and $\\sim$22.16~mag for the shallow and deep strategies. Although some spatial variation due to varying observing conditions is observed, overall the observed limiting magnitudes are consistent with those originally proposed. The quality of the data has been assessed by comparing the measured magnitude of sources at the bright end directly with those reported by the 2MASS survey and at the faint end by comparing the counts of galaxies and stars with those of other surveys to comparable depth and to model predictions. The final science-grade catalogs together with the astrometrically and photometrically calibrated co-added images are available at CDS\\thanks{Available at CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/}. ", "introduction": "\\label{sec:intro} Deep multi-wavelength observations of selected regions of the sky from space and ground-based facilities, combined with spectroscopic observations from large-aperture telescopes, offer the most promising means to probe the distant universe and study in a comprehensive way the evolution of galaxies and large-scale structures over a broad interval of look-back time. Foreseeing the need to complement observations being carried out with the HST, Chandra, XMM, Spitzer and GALEX telescopes with ground-based multi-passband data in optical and infrared, the Working Group for public surveys (SWG) at ESO recommended the ESO Imaging Survey (EIS, Renzini \\& da Costa, 1997) project to undertake deep, optical/infrared observations. This paper is part of a series describing recent data releases from the ESO Imaging Survey using the newly implemented EIS Data Reduction System \\citep{dacosta04}. The reduction of optical data and details of the released products are described in \\citet[ hereafter Paper~I]{dietrich05} in the context of the optical follow-up of the XMM-Newton Serendipitous Sky Survey. In \\citet[ hereafter Paper~II]{olsen06a} the reduction of infrared data is described based on the infrared part of the EIS-DEEP survey. The last survey discussed in the present series is the EIS Deep Public Survey which combines optical \\citep[][ hereafter Paper~III]{mignano06} and infrared observations, with the latter being the topic of the present paper. The original infrared observations, conducted in 1998, targeted the HDF-S and CDF-S regions and fully calibrated images, source catalogs and high-redshift galaxy candidates were immediately made public prior to the Science Verification of the first unit of the VLT \\citep[see preprints by ][]{dacosta98, rengelink98}. In spite of some shortcomings in the original reduction of the infrared images (see Paper~II for a more detailed discussion) these were used in combination with the optical data, taken with the SUSI2 detector, to produce color catalogs and to color-select high-redshift galaxy candidates. These candidates were later used in spectroscopic observations, conducted during commissioning of the VLT and the FORS spectrograph \\citep[e.g., ][]{Cristiani00}, which, in general, confirmed their estimated photometric redshifts. However, the total area of the survey was small and, in particular, only a small fraction of the field of view of Chandra was covered by the available data. In order to expand the original infrared coverage of EIS, complementing other ongoing optical surveys, and provide full coverage of the CDF-S $\\lesssim350$~square~arcmin field and of the flanking regions, the SWG designed the following strategy for the Deep Public Survey (DPS). The EIS DPS survey consists of two parts: first, a deep, optical multi-passband ($UBVRI$) survey reaching limiting magnitudes $m_{AB}\\sim25$~mag and covering three distinct stripes (DEEP1, DEEP2 and DEEP3) of the sky; second, the contiguous coverage in infrared of the same regions, the focus of the present paper. The infrared coverage also has two parts: regions covered in \\k\\ down to a proposed 5$\\sigma$ limiting magnitude of 21.3 (in the AB system) as measured in a 2\\arcsec aperture, and a smaller area covered down to $J_{AB}=$ 23.4 and \\k$_{AB}=$ 22.7~mag. As explained in Paper~III the three regions were selected for the following reasons: 1) DEEP1 because it overlaps with a deep radio survey; 2) DEEP2 because it includes the CDF-S field (DEEP2c); and 3) DEEP3 due to its convenient characteristics and location in the northern galactic cap. The optical observations were carried out using the wide-field imager (WFI), mounted on the ESO/MPG 2.2m telescope at La~Silla covering about $0.25$ square degrees per pointing. Each stripe consists of four adjacent pointings (denoted by a-d in decreasing order of right ascension), yielding an area of one square degree. Further details about DPS can be found in the EIS web pages\\footnote{http://www.eso.org/science/eis}. The primary goal of the DPS survey is to produce a data set from which statistical samples of galaxies can be drawn to study the large scale structures at high redshift. However, these data are also valuable for many other areas of investigation, in particular for the cross-identification with X-ray sources detected by Chandra and XMM, as a photometric reference for the ISAAC mosaic built by the GOODS survey \\citep{dickinson03,giavalisco04}, and combined with the optical data described in Paper~III to improve photometric redshifts and to search for rare high-redshift quasars and galaxies \\citep[e.g., ][]{kong05}. The purpose of the present paper is to describe and present the data from new infrared observations carried out as part of DPS; the optical data are presented in a separate paper of this series (Paper~III). The shallow \\k\\ survey is one of the largest currently available, and the deep \\j\\k\\ is the largest to similar depth \\citep[see][for a review]{elston05}. The paper is organized as follows: Sect.~\\ref{sec:observations} reviews the overall observing strategy and describes the observations. Sect.~\\ref{sec:reductions} briefly describes the techniques employed to process the images, and to astrometrically and photometrically calibrate them. This section also includes a discussion of the assessment made of the quality of the images and of the photometric calibration. Sect.~\\ref{sec:products} presents the final products which include the stacked images and the single-passband catalogs extracted from them. Even though the observations are planned as mosaics, the creation of image mosaics and associated catalogs are beyond the scope of the present paper. Also left for future papers are the production and analysis of optical/infrared color catalogs. While the goal of this paper is not to interpret the data, Sect.~\\ref{sec:discussion} presents the results of comparisons between the present observations and those of other authors. This is done to assess the quality of the astrometry and photometry of the data set. A brief summary is presented in Sect.~\\ref{sec:summary}. ", "conclusions": "\\label{sec:discussion} \\subsection{Internal photometric comparison} \\begin{table} \\caption{Mean magnitude offsets and standard deviations for each mosaic} \\label{tab:int_phot_check} \\begin{tabular}{lcccc} \\hline\\hline Field & Passband & \\# overlaps & Mag. offset & Std.dev.\\\\ \\hline Deep1a & \\j & 24 & -0.02 & 0.11\\\\ Deep1b & \\j & 15 & 0.05 & 0.11\\\\ Deep2b & \\j & 24 & 0.03 & 0.05\\\\ Deep2c & \\j & 12 & -0.03 & 0.08\\\\ Deep3a & \\j & 24 & -0.02 & 0.06\\\\ Deep3b & \\j & 24 & 0.02 & 0.04\\\\ \\hline Shallow\\\\ \\hline Deep1a & \\k & 36 & 0.07 & 0.13\\\\ Deep1b & \\k & 21 & 0.02 & 0.18\\\\ Deep2a & \\k & 39 & 0.08 & 0.12\\\\ Deep2b & \\k & 8 & 0.06 & 0.22\\\\ Deep2c & \\k & 34 & 0.05 & 0.13\\\\ Deep3a & \\k & 15 & 0.06 & 0.08\\\\ Deep3b & \\k & 14 & 0.01 & 0.10\\\\ \\hline Deep\\\\ \\hline Deep1a & \\k & 14 & 0.00 & 0.10\\\\ Deep1b & \\k & 16 & 0.03 & 0.11\\\\ Deep2b & \\k & 24 & 0.03 & 0.09\\\\ Deep2c & \\k & 12 & 0.00 & 0.08\\\\ Deep3a & \\k & 19 & 0.03 & 0.07\\\\ Deep3b & \\k & 24 & 0.05 & 0.08\\\\ \\hline \\end{tabular} \\end{table} The consistency of the photometric calibration of the individual pointings, and therefore of the mosaic, was assessed by comparing the magnitudes of objects in common between adjacent frames. For each mosaic the mean and standard deviation of the magnitude difference of all pairs of sources in common in adjacent frames were computed, considering only objects with a magnitude error $<0.2$ ($\\sim5\\sigma$). Ideally, one would use only stars for this, but because the number of objects in the overlaps is small, it proved necessary to include galaxies as well. In Table~\\ref{tab:int_phot_check} these mean offsets and standard deviations are listed for each mosaic. The table lists in Col.~1 the field name, in Col.~2 the passband, in Col.~3 the number of overlaps in the mosaic, in Cols.~4 and 5 the mean and standard deviation of the magnitude offsets. From the table one finds that on average the offsets are small, typically $\\lesssim0.05$~mag. The scatter is typically less than 0.1~mag. In all cases the mean offset is smaller than the standard deviation. Combined these results suggest a homogeneous photometric calibration across the mosaics with fairly small field-to-field variations. In an attempt to improve the match among different zero points obtained on different nights, correct for data taken under non-photometric conditions, and hence minimize field-to-field variations, the method employed by \\cite{maddox90} was utilized. Taking into account the computed corrections, one finds that, in general, the mean offsets do not change significantly and the scatter decreases only marginally. This result may be partly due to the large uncertainty in each individual offset caused by the small number of sources in the overlap regions and the necessity to include galaxies. However, it may also indicate a good match among the original zero points. \\subsection{External comparison} \\begin{figure*} \\begin{center} \\resizebox{0.4\\textwidth}{!}{\\includegraphics{5019f2a.ps}} \\resizebox{0.4\\textwidth}{!}{\\includegraphics{5019f2b.ps}} \\resizebox{0.4\\textwidth}{!}{\\includegraphics{5019f2c.ps}} \\resizebox{0.4\\textwidth}{!}{\\includegraphics{5019f2d.ps}} \\resizebox{0.4\\textwidth}{!}{\\includegraphics{5019f2e.ps}} \\resizebox{0.4\\textwidth}{!}{\\includegraphics{5019f2f.ps}} \\end{center} \\caption{Astrometric (left column) and photometric (right column) comparison between objects in common between the EIS and 2MASS surveys. The offsets are computed as 2MASS~$-$~EIS. The rows correspond to each adopted strategy as indicated. Dotted lines indicate the zero offset location, while the dashed lines in the right panels depict the mean offset computed for the bright objects as discussed in the text. In this plot magnitudes are in the Vega system.} \\label{fig:comp-2mass} \\end{figure*} In order to externally verify the survey products, the source lists extracted from the final images as described in the previous section were cross-correlated with the point source catalog available from the 2MASS survey \\citep{skrutskie06}. This was carried out adopting a search radius of 1\\arcsec. A total of 1056, 966, 1539 objects were found in common for the \\j, and the deep and shallow \\k~-band final images, respectively, over all regions covered by the survey. The differences in position and magnitude for objects in common and considered ``good'' in both catalogs (5$\\sigma$ detections and in the case of 2MASS not contaminated by neighboring objects) were computed and their distribution as a function of right ascension and declination (left panels) and magnitude (right panels) are shown in Fig.~\\ref{fig:comp-2mass}. The mean offset in position is negligible and amounts to at most $\\sim0\\farcs07$ for the shallow \\k~strategy. The scatter is of $0\\farcs23$ and $0\\farcs27$ in right ascension and declination, respectively, for all strategies. This is consistent with an astrometric accuracy of $\\sim0\\farcs19$ for the EIS images, limited by the typical internal accuracy of the reference catalog used. Considering objects brighter than $J=16.0$ and $K_s=14.5$ (Vega system) the measured magnitudes of the two data sets are in good agreement. At fainter magnitudes the scatter increases and the effects of the Malmquist bias introduced by the depth of the 2MASS data are seen. Objects with magnitudes brighter than 12~mag are saturated in the EIS data and therefore discarded from the plots. For the deep strategies the systematic offset in magnitude is (discarding the outlying objects with offsets in excess of 0.5~mag) $\\sim -0.04$~mag with a scatter of 0.07 in \\k\\ and $\\sim0.0$ with a scatter of 0.08 in \\j , consistent with the estimated uncertainty of the overall zero-point (Sect.~\\ref{sec:reductions}). For the shallow \\k\\ -survey one finds a mean offset of $\\sim0.08$~mag and a scatter of $\\sim0.13$~mag. The amplitude of the scatter is remarkably small, considering the values obtained in the previous section, and demonstrates the uniformity of the photometric calibration of the present data across the sky. \\subsection{Comparison of galaxy and star counts} \\begin{figure*} \\center \\resizebox{0.8\\textwidth}{!}{\\includegraphics{5019f3.ps}} \\caption{Galaxy counts for each individual pointing grouped by mosaic (lines) compared with those of \\citet[ solid circles]{iovino05}.} \\label{fig:galcounts_all} \\end{figure*} An alternative way to evaluate the data is to compare the counts of galaxies and stars with those computed by other authors and/or predicted by models. For galaxies this is illustrated in Fig.~\\ref{fig:galcounts_all}. The figure shows for each passband and strategy (columns) the galaxy counts obtained for each pointing (thin lines) grouped by mosaic (rows), as indicated in the panels of the leftmost panel of each row. These are compared with the counts recently reported by \\citet[ solid circles]{iovino05}. Here galaxies are 3$\\sigma$ detections with CLASS\\_STAR$<0.9$. From the figure one can easily: 1) identify individual cases that depart significantly from the mean; 2) have a clear measure of the variance introduced from field-to-field variations due to the relatively small field of view of SOFI; and 3) identify the variation in depth of a few pointings, such as in the case of Deep1a \\k-shallow among others. Moreover, the relatively small scatter at intermediate magnitudes among the subfields forming a mosaic indicates that the relative photometry between them is reasonable, reinforcing the evidence mentioned above. More importantly, on average, there is an excellent agreement with the counts of \\cite{iovino05}, which reach a comparable depth to the deep \\j\\k\\ strategy. This can be seen in Fig.~\\ref{fig:counts_avrg} which shows the average counts obtained taking into account all fields and regions for the $J$- and $K_s$-band data. In all cases each field contributes only at magnitudes brighter than their respective 80\\% completeness limit, and no completeness corrections have been attempted. These mean counts show an excellent agreement with those recently obtained by different authors \\citep{vaisanen00,martini01, iovino05}, being further strong evidence of the overall quality of the data. The mean counts for the \\j- and \\k-bands are given in Tables~\\ref{tab:counts_j} and \\ref{tab:counts_k} listing in Col.~1 the center of the magnitude bin in the Vega system, in Col.~2 the number of galaxies per 0.5~mag and per square degree, and in Col.~3 the standard deviation among the contributing fields. \\begin{figure} \\begin{center} \\resizebox{\\columnwidth}{!}{\\includegraphics{5019f4a.ps}} \\resizebox{\\columnwidth}{!}{\\includegraphics{5019f4b.ps}} \\end{center} \\caption{Comparison of galaxy number counts in the $J$- (top) and $K_s$-band (bottom) between the present work (solid circles with error bars); diamonds denote \\cite{vaisanen00}, triangles \\cite{martini01} and squares \\cite{iovino05}. } \\label{fig:counts_avrg} \\end{figure} \\begin{table} \\caption{The derived raw galaxy counts for the \\j -band. The counts are given per mag per square degree. Each field only contributes to the bins brighter than their 80\\% completeness limit. No completeness correction has been attempted.} \\label{tab:counts_j} \\begin{center} \\begin{tabular}{rrr} \\hline\\hline Mag & N & $\\sigma_N$ \\\\ \\hline 10.25 & 0.0 & 0.0\\\\ 10.75 & 0.0 & 0.0\\\\ 11.25 & 0.0 & 0.0\\\\ 11.75 & 0.0 & 0.0\\\\ 12.25 & 2.9 & 26.9\\\\ 12.75 & 2.8 & 26.0\\\\ 13.25 & 2.8 & 26.0\\\\ 13.75 & 10.9 & 50.8\\\\ 14.25 & 11.2 & 52.2\\\\ 14.75 & 47.4 & 110.8\\\\ 15.25 & 73.7 & 135.3\\\\ 15.75 & 93.7 & 156.4\\\\ 16.25 & 283.3 & 255.6\\\\ 16.75 & 444.5 & 381.0\\\\ 17.25 & 766.0 & 463.1\\\\ 17.75 & 1350.6 & 606.4\\\\ 18.25 & 2059.3 & 746.3\\\\ 18.75 & 3189.6 & 1083.7\\\\ 19.25 & 4988.9 & 1731.7\\\\ 19.75 & 7804.7 & 2398.9\\\\ 20.25 & 11671.5 & 3354.3\\\\ 20.75 & 16064.9 & 3381.6\\\\ 21.25 & 18383.9 & 5991.3\\\\ 21.75 & 12933.6 & 8897.6\\\\ 22.25 & 10949.9 & 10678.0\\\\ \\hline \\end{tabular} \\end{center} \\end{table} \\begin{table} \\caption{The derived raw galaxy counts for the \\k -band. The counts are given per mag per square degree. Each field only contributes to the bins brighter than their 80\\% completeness limit. No completeness correction has been attempted.} \\label{tab:counts_k} \\begin{center} \\begin{tabular}{rrr} \\hline\\hline Mag & N & $\\sigma_N$ \\\\ \\hline 11.25 & 0.9 & 14.6\\\\ 11.75 & 4.1 & 33.6\\\\ 12.25 & 2.1 & 24.2\\\\ 12.75 & 11.9 & 54.9\\\\ 13.25 & 20.2 & 77.1\\\\ 13.75 & 51.8 & 118.7\\\\ 14.25 & 121.2 & 209.2\\\\ 14.75 & 224.4 & 254.5\\\\ 15.25 & 471.6 & 421.6\\\\ 15.75 & 747.9 & 509.5\\\\ 16.25 & 1282.1 & 664.6\\\\ 16.75 & 2144.8 & 932.2\\\\ 17.25 & 3598.5 & 1366.0\\\\ 17.75 & 5721.6 & 1832.1\\\\ 18.25 & 8872.0 & 2904.3\\\\ 18.75 & 11606.3 & 3772.5\\\\ 19.25 & 10333.8 & 5998.6\\\\ 19.75 & 8330.2 & 5731.2\\\\ 20.25 & 5295.8 & 5289.0\\\\ \\hline \\end{tabular} \\end{center} \\end{table} Finally, the corresponding stellar counts, sources with CLASS\\_STAR$\\geq0.9$, are shown in Fig.~\\ref{fig:starcounts}. The counts (solid circles connected by a solid line) are compared to predictions adopting the model proposed by \\citet[thin lines]{girardi05}. Despite the large scatter due to the small field-of-view, the agreement is remarkable, considering the different regions and that the model was parameterized using optical data. \\begin{figure*} \\center \\resizebox{0.8\\textwidth}{!}{\\includegraphics{5019f5.ps}} \\caption{Comparison of star counts (filled circles connected by solid line) for each mosaic (rows) and strategy (columns) compared with model predictions by \\citet[ thin line]{girardi05}. } \\label{fig:starcounts} \\end{figure*} This paper presents the infrared data accumulated by the Deep Public Survey conducted by the EIS program. The paper presents the observations, reductions, and some of the main products generated and administrated by the EIS Data Reduction System which allows the unsupervised reduction of large amounts of data. For instance, the data set described in the present paper was reduced overnight using eight dual-processors computers running Linux. The highthroughput of the system also makes it ideal for handling the new generation of infrared surveys using wide-field cameras. In general, the depth of the survey is relatively homogeneous and the coverage rather uniform with exception of the shallow coverage of Deep1. However, considerable improvement in coverage and photometric accuracy could have been reached if the survey was carried out in service rather than visitor mode allowing a better optimization depending on the actual weather conditions, in particular ensuring sufficient photometric data for reliable calibration of all the fields. The near-infrared dataset being released is one of largest currently available and its quality has been assessed qualitatively by visual inspection of the individual images and quantitatively by direct comparison with measurements made by the 2MASS survey limited to the bright end and by the comparison of galaxy and star counts with those of other authors and with model predictions. The depth and areal coverage in the near-infrared combined with the deep multi-wavelength optical data reported in Paper~III make this data set extremely useful to further explore the nature of extremely red objects \\citep[e.g. ][]{kong05} and to search for distant galaxy clusters. With this paper, all the infrared data accumulated by the EIS project up to late 2004 are now in the public domain\\footnote{The science grade images and catalogs are available at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/}." }, "0607/astro-ph0607216_arXiv.txt": { "abstract": " ", "introduction": " ", "conclusions": "\\subsection{What's next~?} The theory of {\\it steady} jet production from Keplerian accretion discs is now completed. The physical conditions required to thermo-magnetically drive jets are known, all relevant physical processes have been included in the framework of mean field dynamics. Of course, there are still many unsolved questions: \\\\ {\\it (i)} Can a sustained MHD turbulence maintain $\\alpha_m \\sim 1$? This is a huge constraint that deserves a thorough investigation.\\\\ {\\it (ii)} Observations of T~Tauri jets favor solutions with large ejection efficiencies ($\\xi \\sim 0.1$, \\citealt{pese04}) requiring additional heating at the disc surface. A theoretical assessment of this heating must be undertaken.\\\\ {\\it (iii)} What is the stability of MAES? As will be seen below, there was some claims that MAES were unstable but they were proven to be wrong. On the other hand, jets do show time dependent features and one must clearly go beyond steady state models. On that respect, numerical simulations will be very helpful.\\\\ {\\it (iv)} Disc driven winds do not treat the star-disc interaction. Understanding the whole process of star formation requires now to address this crucial issue as it pinpoints the problem of the stellar angular momentum removal. This is further discussed below. \\subsection{Biases of self-similarity} Self-similarity allows to take into account all dynamical terms in the equations and, as such, is the best means to solve in a self-consistent way the steady-state accretion-ejection problem. However, there is a price to pay...\\index{self-similarity, bias} {\\bf (i) The asymptotic behaviour}\\index{asymptotic equilibrium} is obviously biased since, for instance, neither inner nor outer pressures can be taken into account. In fact, no realistic \"radial\" boundary condition whatsoever can be dealt with. When modeling an astrophysical jet, this implies for instance to truncate the solution at one inner and outer radius. But there is another aspect, less known and more subtle. \\citet{cont94a} and \\citet{ostr97} obtained jet solutions within the same self-similar framework but with different asymptotic behaviors. The reason stems from the fact that they played around with $\\beta$ (flux function $a(r)\\propto r^\\beta$) as if it were a free function whereas the mathematical matching with a Keplerian disc imposes its value. On the other hand, \\citet{pell92} obtained also recollimating\\index{recollimation} {\\it non self-similar} solutions, which indicates that recollimation can indeed be physical and not entirely due to self-similarity. In fact, it can be shown that recollimation of a jet launched from a Keplerian accretion disc is possible whenever the radial profile of the ejection efficiency $\\xi$ is smooth enough \\citep{ferr97}. {\\bf (ii) The regularity conditions} are to be imposed at the modified points\\index{critical points, modified} and not at the usual ones (see Tsinganos, this volume). However, these locations coincide for both the slow (SM) and the Alfv\\'en points so that one can be confident that there is no bias there. However, this is not so for the fast magnetosonic point. Self-similar trans-FM solutions require an Alfv\\'en surface very close to the disc \\citep{ferr04}, which can only be done by the action of a large pressure in the sub-Alfv\\'enic region. This is obviously a strong bias since it is not clear whether such a pressure is indeed provided in astrophysical objects. Note however that crossing this modified FM point is more a theoretician satisfaction than anything else: it gives no additional physical insight on the disc physics. {\\bf (iii) The local disc physical conditions} as obtained with self-similar solutions are not biased. The physical processes are well identified and understood and can be sometimes even obtained in a pure analytical manner. They have been also confirmed by numerical experiments of \\citet{cass02, cass04, zann04} (although one might object that numerical experiments were actually tested with the help of semi-analytical solutions). \\subsection{Is accretion-ejection unstable?} There has been some claims in the literature that the accretion-ejection process itself was unstable\\index{instabilities, accretion-ejection} \\citep{lubo94b,cao02}. The idea was the following. Start from a steady picture where the accretion velocity $u_r$ at the disc midplane is due to the jet torque. It leads to a bending of the poloidal field lines described by an angle $\\theta$ with the vertical. Now imagine a small perturbation $\\delta u_r$ enhancing the accretion velocity. Then, according to these authors, the field lines would be more bent ($\\theta$ increases) which would lead to lower the altitude of the sonic point. Because the sonic point would be located deeper in the disc atmosphere, where the density is higher, more mass would be henceforth ejected which would then increase the total angular momentum carried away by the jet. This means that the torque due to the jet is enhanced and will, in turn, act to increase the accretion velocity. Thus, the accretion-ejection process is inherently unstable. The whole idea of this instability is based upon a crude approximation of the disc vertical equilibrium. In fact, the magnetic field produces a strong vertical compression so that, as $\\theta$ is increased, {\\it less} mass is being ejected, not more. This has been pointed out by \\citet{koni96} and \\citet{koni04} and is indeed verified in the full MAES calculations reported here. \\subsection{Magnetic fields in accretion discs} The necessary condition for launching a self-collimated jet from a Keplerian accretion disc is the presence of a large scale vertical magnetic field close to equipartition\\index{magnetic field, equipartition} \\citep{ferr95}, namely \\begin{equation} B_z \\simeq 0.2\\ \\left ( \\frac{M}{M_\\odot}\\right )^{1/4} \\left ( \\frac{\\dot M_a}{10^{-7} M_\\odot/yr} \\right )^{1/2} \\left ( \\frac{r_o}{1\\mbox{ AU}}\\right )^{-5/4 + \\xi/2} \\mbox{ G,} \\end{equation} This value is far smaller than the one estimated from the interstellar magnetic field (see M\\'enard's contribution), assuming either ideal MHD or $B \\propto n^{1/2}$ \\citep{heil93,basu94,basu95b}. This implies some decoupling between the infalling/accreting material and the magnetic field in order to get rid off this field. This issue is still under debate. The question is therefore whether accretion discs can build up their own large scale magnetic field (dynamo)\\index{magnetic field, dynamo} or if they can drag in and amplify the interstellar magnetic field? Although no large scale fields have been provided by a self-consistent disc dynamo, this scenario cannot be excluded. But the latter scenario (advection) seems a bit more natural. Let us assume that the disc material is always ionized enough to allow for some coupling with the magnetic field (and use MHD). The outer parts of the accretion disc will probably take the form of a SAD with no jets and almost straight (${\\cal R}_m\\sim 1$) field lines \\citep{lubo94a}. In that case, the steady-state solution of the induction equation for the poloidal field is $B_z \\propto r^{-{\\cal R}_m}$ \\citep{ferr06a}. Hence, as a result of both advection and (turbulent) diffusion, the magnetic field in a SAD will be a power-law of the radius. Can a SAD\\index{disc, standard} transport $B_z$ and allow for a transition to an inner JED\\index{magnetic field, advection}? This will be so if there is some transition radius (the outer JED radius $r_e$) where $\\mu= B_z^2/\\mu_oP$ becomes of order unity. In a SAD the total pressure writes $P = \\frac{\\dot M_a \\Omega_K^2 h}{6\\pi \\nu_v} \\propto r^{-3/2 - \\delta}$ with $h(r)\\propto r^\\delta$. Since $\\delta$ is always close to unity in circumstellar discs, one gets $\\mu \\propto r^{-\\epsilon}$ with $\\epsilon \\sim 1$. Thus, it can be readily seen that it is indeed reasonable to expect such a transition (computing it is another matter), at least in some objects. The recent Zeeman observation of a magnetic field in the accretion disc of FU Or supports this conclusion \\citep{dona05}. \\begin{figure}[t] \\centering \\includegraphics[width=0.9\\textwidth]{fig8.eps} \\caption{Two classes of stationary accretion powered disc winds. {\\bf (a)} \"extended disc winds\", when the magnetic flux threading the disc is large enough so that a large radial extension of the whole accretion disc drives jets ($r_e \\gg r_i$). The Alfv\\'en surface $S_{\\mbox A}$ is expected to adopt a rather conical shape. {\\bf (b)} \"X-winds\", when the magnetic flux is small and only a tiny disc region is driving jets. The Alfv\\'en surface can be either convex or concave, although the latter is probably more physical (since less material can be ejected at the two extremes and the Alfv\\'en point is rejected to infinity). Adapted from \\citet{ferr06b}.} \\label{fig:winds} \\end{figure} \\subsection{X-winds and disc winds} The X-wind\\index{wind models, disc wind, X-wind} model \\citep{shu94a,naji94,shu95,shan98, shan02} is a rich and complex model but, contrary to common belief, it is an accretion-powered wind launched from the accretion disc. In practice, if the amount of magnetic flux threading the disc is large so that $r_e\\gg r_i$, then one gets an \"extended disc wind\", whereas if the magnetic flux is tiny with $r_e \\geq r_i$, one gets an \"X-wind\" (Fig.~\\ref{fig:winds}). The dynamics and asymptotic behaviour of jets will differ strongly between an extended disc wind and an X-wind and can thereby be tested against observations \\citep{ferr06b}. But this difference arises mainly because of the restricted range in radii in the X-wind case, not because the underlying disc physics is different. The basic phenomena described in Section~2 apply as well for the portion of the disc launching the X-wind. Thus, equipartition fields are required, the \"viscous\" torque is negligible with respect to the jet torque and the angular momentum carried away by the X-wind is {\\it exactly} the same amount lost by the accreting material. As a consequence, X-winds cannot take away any angular momentum from the central star. Published material on the dynamics of X-winds contains: (i) a scenario for the origin of $B_z$ (stellar) and the star-disc interaction (leading to the opening of some magnetic field lines); (ii) the calculation of the sub-Alfv\\'enic ideal MHD jet (elliptic domain defined by prescribed boundary conditions); (iii) a somewhat mysterious \"interpolation\" to a simple jet asymptotic solution. The following questions remain therefore to be addressed:\\\\ {\\bf (1)} Can the disc afford the imposed mass flux and field geometry? Indeed, the assumed ejection to accretion mass flux ratio of 1/3 from such a tiny region is huge and would require a fantastic ejection efficiency ($\\xi$ of order unity or larger). The calculations of JEDs showed that this is unfeasible in a steady way. However, the huge magnetic field gradients required in the X-wind launching region provide a significantly different situation. This has never been analyzed.\\\\ {\\bf (2)} How good is the transfield equilibrium satisfied? There is no mathematical procedure to find a solution of mixed type (elliptic-hyperbolic) PDEs when the singular surfaces are unknown. The trick used for X-winds provides an incomplete solution, but there is maybe some means to fulfill the transverse equilibrium by using an iterative scheme. In any case, this important point is missing in the current published material. \\subsection{Magnetic star-disc interactions} Nowadays it seems accepted that a lot if not all young stars have a magnetospheric interaction with their circumstellar accretion disc (see Alencar's contribution, this volume). If one assumes that the disc is threaded by a large scale magnetic field, then the question of how this field is connected to the stellar field arises. First ideas are always simple and so is the stellar magnetic field\\index{magnetic field, stellar dipole}, assumed up to now to be dipolar and axisymmetric (see Mohanty et al. 2006). We define here the magnetopause\\index{magnetospheric accretion, truncation radius} as the radius $r_m$ below which all field lines threading the disc are tied to the star whereas beyond $r_m$, they are disconnected from the star. The case envisioned within the X-wind scenario assumes a stellar magnetic moment anti-parallel to the disc magnetic field \\citep{shu94a}. As a consequence, a neutral surface (where $B=0$) appears above each disc surface, illustrated by a limiting poloidal field with a Y shape (Fig.~\\ref{fig:magn}). The other case, a stellar magnetic moment parallel to the disc magnetic field, has been proposed by \\citet{ferr00}. The two fields then cancel each other at the disc midplane, defining a neutral line\\index{magnetic, neutral line} at a radius $r_X$ where reconnection takes place. This configuration gives rise to \"Reconnection X-winds\" (hereafter ReX-winds) specifically above this zone. \\begin{figure}[t] \\centering \\includegraphics[width=0.9\\textwidth]{fig9.eps} \\caption{Two simple axisymmetric star-disc magnetospheric interactions. {\\bf (a)} \"Y-type\" interaction obtained when the stellar magnetic moment is anti-parallel to the disc magnetic field. A current sheet is formed at the interface between the open stellar field and the disc field. Such a configuration cannot produce per se a wind. {\\bf (b)} \"X-type\" configuration obtained when the stellar magnetic moment is parallel to the disc magnetic field. A magnetic X-point is generated at the disc midplane where the two fields cancel each other. Unsteady ejection (\"ReX-winds\") can be launched above this reconnection site. Adapted from \\citet{ferr06b}.} \\label{fig:magn} \\end{figure} \\subsubsection{Accretion curtains} The first question here is can these simple topologies allow for accretion below\\footnote{Accretion is realized beyond $r_m$ by e.g. the jet torque within the JED and, farther away in the SAD, by the turbulent \"viscous\" torque.} $r_m$? Disc material\\index{magnetospheric accretion, accretion columns} will accrete only if it looses angular momentum and this depends on both turbulence and the magnetic torque due to the magnetosphere. The magnetosphere will try to make the disc corotate with the protostar so the sign of the torque depends on their relative angular velocity. The corotation radius\\index{magnetospheric accretion, corotation radius}, $r_{co}= (GM/\\Omega_*^2)^{1/3}$, is defined as the radius where the stellar angular velocity $\\Omega_*$ is equal to the Keplerian one. This gives an estimate of the real angular velocity of the disc (since the disc magnetic field introduces already a deviation). Roughly speaking, if $r_m> r_{co}$ the star rotates faster than the disc and deposits its angular momentum, whereas if $r_m< r_{co}$, the star rotates slower and thus spins down the disc. Note that $r_m$ denotes roughly the radius where the stellar magnetic field becomes dynamically dominant, namely $\\mu >1$. Thus, unless a very efficient turbulent mechanism\\footnote{Note that it should be operating when $\\mu >1$, while the magneto-rotational instability\\index{instabilities, magnetorotational} is already quenched at $\\mu \\sim 1$ \\citep{balb03}.} is operating and transports radially the stellar angular momentum, no accretion is possible when $r_m> r_{co}$ (although such a \"propeller\" regime is favorable for ejection). As a consequence, both X-type and Y-type interactions allow for a magnetospheric accretion as long as $r_m< r_{co}$. It is interesting to note that both configurations require an equatorial reconnection zone\\index{reconnection} (interesting for sudden energy dissipation and chondrules, \\citealt{shu01,goun06}). In the case of a Y-type interaction, it arises because of the requirement that the magnetospheric field makes an angle with the vertical large enough in order to allow the disc material to flow inwards. This assumption implies a magnetic neutral \"belt\"\\index{magnetic, neutral line} at the disc midplane (see Fig.~1 in \\citealt{ostr95}), but whose origin and dynamics were not discussed and remain therefore major unsolved issues. In the case of an X-type interaction, the presence of the magnetic neutral line is due to and maintained by the cancellation of the two fields (see fig.\\ref{fig:magn}). The accreting disc material can cross the resistive MHD region and is lifted vertically by the strong Lorentz force above the reconnection site. The transition from an accretion disc to accretion curtains can be quite smooth in that case. \\subsubsection{Stellar spin down} The second question is the issue of the stellar angular momentum removal by winds\\index{angular momentum transport, stellar spin down} (see \\citet{matt05} for more details and the necessity of winds). As explained earlier, X-winds carry away the angular momentum of the accreting disc material. Thus, such a configuration cannot brake down the protostar (as initially claimed). On the contrary, the X-type configuration provides a very efficient means to do it \\citep{ferr00}. The reason is the possibility to launch disc material above the reconnection site. The scenario is the following (Fig.~\\ref{fig:rex}). A stationary extended JED is settled in the innermost regions of the accretion disc and provides open magnetic flux to the star. This magnetic field reconnects at $r_X$ with closed stellar field lines: the disc field contributes thereby to transform closed magnetospheric flux into open flux. At the reconnection site, the disc material is lifted vertically and loaded onto these newly opened field lines, tied to the rotating star. Whenever $r_X> r_{co}$, the star is rotating faster than the loaded material and it undergoes a strong magneto-centrifugal acceleration. This gives rise to the so-called ReX-wind\\index{wind models, ReX-wind}, whose energy and angular momentum are those of the star. Using a toy-model for the magnetic interaction \\citet{ferr00} showed that such winds could brake down a {\\it contracting} protostar\\index{protostar} on time scales that are comparable to the duration of the embedded phase (Class 0 and I sources)\\index{protostar, Class 0}. The protostar was assumed to rotate initially at breakup speed and, after some $10^5$ to $10^6$ yrs, it has been spun down to 10\\% of it despite its contraction and mass accretion. \\begin{figure}[t] \\centering \\includegraphics[height=4cm]{fig10.eps} \\caption{The ReX-wind configuration \\citep{ferr00}. A MAES is established around a protostar\\index{protostar} whose magnetic moment is parallel to the disc magnetic field. This is a natural situation if both fields (disc and stellar) have the same origin. Left: black solid lines are streamlines, white dashed lines are contours of equal total velocity (mainly rotation inside the disc) and the background color scale shows the density stratification. The ReX-wind (arrows) would be confined and channeled by the outer disc wind. Right: sketch of the magnetic configuration leading to Rex-winds and accretion curtains around the magnetic neutral line at $r_X$. Arrows show the expected time-dependent plasma motion.} \\label{fig:rex} \\end{figure} ReX-winds seem therefore to offer a serious possibility to brake down protostars (to my knowledge, there is no other model in the literature). Note that ReX-winds are probably intermittent by nature because of the unavoidable radial drift of the reconnection site (there is no ejection whenever $r_X< r_{co}$). Dynamically speaking, such an unsteady \"wind\" should be better described as bullets flowing inside the hollow disc wind. Remarkably the basic features of X-type configurations remain if the stellar dipole is inclined: one would observe in that case precessing bullets channeled by the outer disc wind. Heavy numerical simulations will be required to test and analyze this scenario." }, "0607/astro-ph0607202_arXiv.txt": { "abstract": "{The $B$--Mode of the Cosmic Microwave Background Polarization (CMBP) promises to detect the gravitational wave background left by Inflation and explore this very early period of the Universe. In spite of its importance, however, the cosmic signal is tiny and can be severely limited by astrophysical foregrounds. In this contribution we discuss about one of the main contaminant, the diffuse synchrotron emission of the Galaxy. We briefly report about recent deep observations at high Galactic latitudes, the most interesting for CMB purposes because of the low emission, and discuss the contraints in CMBP investigations. The contamination competes with CMB models with $T/S = 10^{-2}$--$10^{-3}$, close to the intrinsic limit for a 15\\% portion of the sky (which is $T/S \\sim 10^{-3}$). If confirmed by future surveys with larger sky coverage, this gives interesting perpectives for experiments, that, targeting selected low emission regions, could reach this theoretical limit. } ", "introduction": "The polarization of the Cosmic Microwave Background radiation (CMBP) is a powerful tool to investigate the early Universe. The tensor pertubations of the primordial Gravitational Wave (GW) background left by Inflation, and the reionization history of the Universe can be effectively studied by the two components the CMBP can be expanded in: the $E$- and $B$-mode (Zaldarriaga \\& Seljak 1997). The $B$--mode component is faint, but can give us the first tool to investigate the physics of the Inflation. In fact, the $B$-mode level on degree scales is related to the amount of the primordial GW background emitted by Inflation, which is usually measured through the tensor-to-scalar perturbation power ratio $T/S$ (Figure~\\ref{cbFig} and, e.g., \\cite{boyle06}~2006, \\cite{kinney06}~2006). The $B$-mode spectrum normalization has a linear dependence on it (see Figure~\\ref{cbFig}), while the other CMB components are almost insensitive to these tensor perturbations\\footnote{The temperature spectrum $C^T$ has some sensitivity for values $T/S > 0.1$.}. \\begin{figure} \\includegraphics[angle=0, width=0.9\\hsize]{carretti_Oxford2006_Fig1.ps} \\caption{Angular power spectra of temperature anisotropy ($C^T$), $E$-mode ($C^E$) and $B$-mode ($C^B$) of the CMB. Cosmological parameteres of the so-called {\\it concordance} model after the WMAP data are assumed (\\cite{spergel06}~2006). In particular an optical depth of the reionized medium of $\\tau = 0.1$ is used. Spectra for some values of the unknown $T/S$ parameter are plotted. \\label{cbFig} } \\end{figure} In combination with parameters measured by the CMB Temperature spectrum (the scalar pertubation spectral index $n$ and its running $dn / d\\ln k$), it distinguishes among the several Inflation models. However, $T/S$ is unknown so far and only upper limits exist ($T/S < 0.22$, 95\\% C.L., \\cite{seljak06}~2006). The many Inflation models predict it to range within orders of magnitude (approximately $10^{-4} < T/S < 10^{-1}$), which corresponds to signals as low as 3--100~nK. The detection of the $B$-mode is thus fundamental to distinguish among different models and to measure their relevant parameters, like the energy density of the Universe when Inflation itself occurred. Still undetected, the measurement of this CMBP component is a hot topic in cosmology. ", "conclusions": "The results reported in Sect.~\\ref{synchSec} depict a situation with significant contamination even at high Galactic latitudes, but with better conditions in the low emission regions. Such {\\it low lands} represent about 15\\% of the sky and could be the right place where to conduct deep CMBP observations to look for the $B$-mode. To limit observations in small regions, however, imposes intrinsic limitation on the minimum detectable $T/S$, mainly because of the leakage from $E$- into the weaker $B$-mode. In fact, \\cite{amarie05}~(2005) find that an all-sky survey would allow a detection of $T/S$ with a theoretical sensitivity limit\\footnote{i.e. in the ideal case of negligible instrument noise.} of $\\Delta T/S = 1.5 \\times 10^{-5}$, that becomes $\\Delta T/S = 3.2 \\times 10^{-5}$ when 70\\% of the sky is available, $\\Delta T/S = 10^{-3}$ for 15\\%, and $\\Delta T/S = 10^{-2}$ in 1\\% (3-$\\sigma$ C.L.). It is worth noting that the class of Inflation models with minimal fine-tuning have $T/S$ values ranging within $10^{-3}$ and $10^{-1}$ (\\cite{boyle06}~2006), for which a 15\\% sky portion would be large enough for the first detection of the tensor CMBP component. \\begin{figure} \\centering \\includegraphics[angle=0, width=0.7\\hsize]{carretti_Oxford2006_Fig4.ps} \\caption{The PGMS survey coverage in Galactic coordinates. \\label{stripFig} } \\end{figure} \\begin{table} \\centering \\caption{Main features of the PGMS survey.} \\begin{tabular}{@{}lr@{}} \\hline Central frequency & 2.332-GHz \\\\ Bandwidth & $128$-MHz \\\\ FWHM & $8.8'$ \\\\ Galactic longitude & $l \\sim 255^\\circ$ \\\\ Area size & $5.0^{\\circ}\\times 90.0^{\\circ}$ \\\\ $Q$, $U$ sensitivity in a beam pixel & 0.5-mK \\\\ Telescope & Parkes \\\\ \\hline \\end{tabular} \\label{pgmsTab} \\end{table} Despite of these limitation, the minimum $T/S$ value detectable in 15\\% of the sky almost matches the limit imposed by the foregrounds in low emission regions. The search for the $B$-mode in a sky portion of such a size could thus be a good trade-off between instrinsic and foregrounds limits. In addition, the weakness of the $B$-mode signal makes already a challenge to detect the signal for $T/S=0.1$ with the present technology (\\cite{cortiglioni06}~2006). It is likely that significant technological improvements will be necessary before scientists can face an all-sky mapping mission with a sensitivity able to match the intrinsic limit of $\\Delta(T/S)\\sim 10^{-5}$. Therefore, an experiment aiming at detecting the $B$-mode in a smaller region (10-15\\% of the sky) can be an interesting intermediate step that would allow us to probe Inflation models down to $T/S = 10^{-3}$. The observations conducted in the {\\it low lands} cover small areas. Although taken in three independent samples, surveys of large portions of the sky are necessary to understand whether the areas observed so far are peculiar {\\it lucky} lowest emission cases or are representative of the conditions of the {\\it low lands}. New information are expected to come from the recently completed all-sky mapping at 1.4-GHz, even if, as mentioned in Sect.~\\ref{foregSec}, these are likely modified by the FR action up to $|b|\\sim50^\\circ$. First analyses using these data have been reported by \\cite{laporta06}~(2006). However, they regard half a sky and contain both large local high emission regions and the Galactic disc that is strongly modified by FR. Deeper and higher frequency data will come soon from the Parkes Galactic Meridian Survey (PGMS, \\cite{carretti05c}~2005c) aimed at surveying a Galactic meridian at 2.3-GHz (Figure~\\ref{stripFig}). The survey is in progress at the Parkes radiotelescope and main features are reported in Table~\\ref{pgmsTab}. It will explore the polarized Galactic diffuse emission along the meridian $l=255^\\circ$ from Galactic plane to south Pole. The frequency is high enough to avoid significant FR effects and main goals are to reveal the intrinsic Galactic emission and study its behaviour with the Galactic latitude. This meridian goes through one of the low emission regions visible in the WMAP data and more extended information about the Galactic emission level in the {\\it low lands} are expected." }, "0607/astro-ph0607034_arXiv.txt": { "abstract": "It is argued that there is a linear correlation between star formation rate (SFR) and accretion rate for normal bright active galactic nuclei (AGNs). However, it is still unclear whether this correlation holds for LINERs, of which the accretion rates are relatively lower than those of normal bright AGNs. The radiatively inefficient accretion flows (RIAFs) are believed to be present in these LINERs. In this work, we derive accretion rates for a sample of LINERs from their hard X-ray luminosities based on spectral calculations for RIAFs. We find that LINERs follow the same correlation between star formation rate and accretion rate defined by normal bright AGNs, when reasonable parameters are adopted for RIAFs. It means that the gases feed the black hole and star formation in these low-luminosity LINERs may follow the same way as that in normal bright AGNs, which is roughly consistent with recent numerical simulations on quasar evolution. ", "introduction": "Black hole (BH) accretion is thought to power AGNs, and the UV/optical bump observed in bright quasars is naturally interpreted as blackbody emission from standard thin accretion disks \\citep*[e.g.,][]{sm89}. The AGN activity may be switched off when the gas near the black hole is exhausted \\citep[see][ for a recent review and references therein]{narayan02}. Most nearby galactic nuclei are much less active and show very different properties with bright quasars, such as low ionization, lack ``UV bump'' etc \\citep*[see][for a review]{ho05}. The accretion mode in these nearby galaxies may be different from that in powerful AGNs. \\citet{ny94} proposed that the standard accretion disk should transit to a RIAF when the accretion rate $\\dot{m}$(=$\\dot{M}/\\dot{M}_{\\rm Edd}$) declines below a critical value $\\dot{m}_{\\rm crit}$ within a certain transition radius \\citep*[e.g.,][]{emn97,yn04,llg04}. RIAFs are optically thin, geometrically thick and very hot, which are supposed to be present in many low-luminosity AGNs \\citep*[e.g.,][]{l96,qe99,gnb99,ymfb02,yn04,nt05} and our Galactic center Sgr $\\rm A^{*}$ \\citep*[e.g.,][]{ny95}. A very low nuclear-luminosity class of low-ionization nuclear emission-line region (LINER) galaxies were identified by \\citet*[][]{heckman80}, in which are found approximately 2/3 low luminosity AGNs \\citep{hfs97a}. An important fraction of LINERs are clearly weak manifestations of quasar-like phenomena, as demonstrated by the presence of broad H$\\alpha$ emission, which are almost certainly accretion-powered \\citep*[][]{ho97b}. These LINERs can be identified with quiescent BH remnants from the quasar era. In the present epoch, the supply of the gas available for powering central engines is much curtailed. The low mass accretion rates inferred for many LINERs may suggest that these sources represent the ``missing link'' between powerful AGNs and normal galaxies as our own. The recent discovery of the correlation between BH mass and stellar velocity dispersion in nearby galaxies \\citep*[e.g.,][]{ge00,fm00} demonstrates a fundamental link between the growth of supermassive BH and bulge formation. The connection between SFR and accretion rate has been explored by some authors \\citep*[e.g.,][]{hec04,hao05,sa05,du05}. Using the SDSS (Sloan Digital Sky Survey) observations of 123,000 low-redshift galaxies, \\citet{hec04} found that the global volume-averaged SFR/$\\dot{M}$ ratio is approximately 1000 in bulge-dominated systems, which is in agreement with the ratio of bulge to BH mass implied by the $M_{\\rm bh}-\\sigma$ relation \\citep{mh03}. \\citet{hao05} also yielded similar results for several ten AGNs (including infrared-selected QSOs, optically-selected QSOs, NLS1s). It was found that the ratio SFR/$\\dot{M}$ of LINERs is apparently different from that of normal bright AGNs, of which the mass accretion rate $\\dot{M}$ is estimated by assuming a constant radiative efficiency for all sources. The mass accretion rates $\\dot{M}$ derived for LINERs may probably be incorrect \\citep{sa05}, because RIAFs are believed to be in these sources and their radiative efficiencies are lower than these for standard thin disks. In this paper, we explore whether the LINERs follow the same correlation between accretion rate and SFR defined by normal bright AGNs. The accretion rates are re-estimated for these LINERs based on our spectral calculations for RIAFs. ", "conclusions": "The bolometric luminosities of AGNs can be estimated by integrating their spectral energy distribution or using an empirical relation \\citep*[e.g.,][]{wu02,el94,ho99}, and then their mass accretion rates are derived if the radiative efficiency is known. A linear correlation between SFR and $\\dot{M}$ was found by \\citet{sa05} for normal bright AGNs, while the LINERs in their sample obviously deviate from this correlation, if the same constant radiative efficiency 0.1 is adopted for the whole sample. It was found that the accretion rates for LINERs are significantly lower than those of normal bright AGNs for given SFRs \\citep*[see Fig. 10 in][]{sa05}. RIAFs are believed present in these low-luminosity LINERs, and their radiative efficiency $\\epsilon_{r}$ is usually much lower than that of standard disks \\citep{ny94}. This implies the accretion rates of LINERs are underestimated by \\citet{sa05}. In this paper, we derive the mass accretion rates of these LINERs from their X-ray luminosities based on RIAF spectral calculations, which should be more reliable. The standard thin disk is believed present in Seyferts and QSOs and its radiative efficiency $\\epsilon_{r}^{0}\\sim0.1$. The slim disk may be present in the NLS1s and its radiative efficiency should be also lower than that of standard disk due to the photon-trapping effects \\citep*[e.g.,][]{ab88,om02,cw04}. However, all the Eddington ratios of NLS1s in our sample are less than 6 \\citep*[$L_{\\rm Bol}/L_{\\rm Edd}$, \\ see table 2\\ in][]{sa05} and the radiative efficiencies of slim disks with such Eddington ratios may not deviate much from that of the standard disk \\citep*[see, dashed line and solid line of Fig. 1 in ][]{om02}. So, the derived mass accretion rates should be accurate to a factor of 2 for NLS1s in these sample, which will not alter our main conclusion. The SFRs is calculated from the FIR luminosity for the whole sample since the FIR luminosity is widely used as a tracer of SFR in galaxies \\citep*[e.g.,][]{kenn98,kew02,sa05}, though we can not rule out some contribution of dust heating by the AGNs and old stars in elliptical galaxies. However, there are lack of correlations between the Mid-IR and FIR for well-studied PG quasars, which suggests that they are dominated by different heating sources\\citep*[e.g.,][]{haas99}. If AGN heating dominates the cooler FIR-emitting dust, then there should be a correlation between quasar OUV and FIR luminosity, while such a correlation has not been found\\citep*[e.g.,][]{mc99,isaak02,priddey03}. In addition, QSOs and Seyferts roughly follow the same universal correlation between the FIR and radio emission deduced from ``normal\" galaxies, which strongly suggests that FIR emission is still powered by star formation rather by AGNs \\citep*[e.g.,][]{nar88,lw05}. Although the FIR emission contributed by old stars in some elliptical galaxies could be important (e.g., very few sources in \\citet{bell03} sample have old star contribution, about 4 of 249, higher 80 percent, see Fig. 6 in that paper), the average contribution of the old stars is around 30 percent. In our present SFR calculations, we use the calibrated formula, in which the old star contribution has been subtracted \\citep{bell03}. So, in statistical sense, the SFRs estimated in our paper are reliable, though we cannot rule out the SFRs of a very few sources in this sample have been overestimated 3-4 times. Even in this case, we believe this has not affected our main statistical results. We plot our re-estimated accretion rate $\\dot{M}$ in Figure 3, and find that all LINERs follow the same correlation defined by the normal bright AGNs (e.g., Seyfert, Quasar and NLS1). The linear correlation between SFR and $\\dot{M}$ indicates that black hole accretion evolves in the same way as star formation, which are both regulated by the interstellar gas in the host galaxies. Our results show that the gases feed the black hole and star formation in low luminosity LINERs follow the similar way as luminous normal AGNs. Recent numerical simulations on the quasar activity triggered by the galaxy merger indeed show that the accretion rate $\\dot{M}$ and SFR decreases simultaneously with time over several orders of magnitude after the quasar shines at its Eddington limit \\citep{sdh04,msh05}. In their simulations, the gas in the galaxies is driven away by the bright quasar radiation, and then both the star formation and BH accretion are quenched simultaneously. The SFR is nearly linearly varying with $\\dot{M}$ in nearly three orders (e.g., $\\dot{M}\\thicksim5\\times10^{-4}-5\\times10^{-1}\\ \\rm M_{\\odot}yr^{-1}$, SFR$\\thicksim5\\times10^{-3}-5 \\ \\rm M_{\\odot}yr^{-1}$ corresponding to $\\sim1.7-2.0\\ \\rm Gyr$) for the model with galaxies of virial velocity $\\rm V_{\\rm vir}=160km/s$ \\citep*[middle red lines in Fig. 2 of][]{msh05}. Our results are roughly consistent with their simulations, though the accretion mode transition has not been considered in their simulations. There are 8 radio galaxies (Eddington ratio less than $10^{-2}$) in 14 radio galaxies, of which the accretion rates are derived from their X-ray emission. It is still debating whether the X-ray emission are dominated by the jet emission for some radio galaxies \\citep{fkm04,yw05}. If the jet emission in X-ray bands is important, the accretion rates derived in this paper are only the upper limits, and those 8 triangles in Figure 3 should be shifted towards left direction. As discussed in Sect. 4.2, we have not found these 8 sources deviating from the correlation between SFR and $\\dot{M}$, which may imply X-ray emission from the jets is unimportant at least for these 8 sources. We find that our results are insensitive to any model parameters except $\\delta$, which describes the fraction of the viscously dissipated energy directly heating the electrons in the accretion flow. In our calculations, $\\delta=0.1$ is adopted, which is a typical value successfully used to model the observed spectra for some low-luminosity AGNs \\citep*[e.g.,][]{qn99}. Although winds may be present in RIAFs, the detailed physics is still unclear and such winds are only described by an artificial power-law parameter $p_{\\rm w}$ \\citep*[e.g.,][]{bb99}. In RIAF spectral calculations for X-ray wavebands, the value $\\delta$ is somewhat degenerate with $p_{\\rm w}$ \\citep*[e.g.,][]{qn99}. Our calculations show that similar conclusion can be obtained if $\\delta=0.3$ and $p_{\\rm w}=0.9$ are adopted for RIAF with winds. The present sample is a mixture of AGNs, of which both BH masses and bolometric luminosities are estimated. It is difficult to evaluate to what extent the sample is affected by the selection effects. A more robust sample is desired for testing this relation between $\\dot{M}$ and SFR." }, "0607/astro-ph0607344_arXiv.txt": { "abstract": "The evidence for particle acceleration in supernova shells comes from electrons whose synchrotron emission is observed in radio and X-rays. Recent observations by the HESS instrument reveal that supernova remnants also emit TeV \\grays; long awaited experimental evidence that supernova remnants can accelerate cosmic rays up to the ``knee'' energies. Still, uncertainty exists whether these \\grays\\ are produced by electrons via inverse Compton scattering or by protons via $\\pi^0$-decay. The multi-wavelength spectra of supernova remnants can be fitted with both mechanisms, although a preference is often given to $\\pi^0$-decay due to the spectral shape at very high energies. A recent study of the interstellar radiation field indicates that its energy density, especially in the inner Galaxy, is higher than previously thought. In this paper we evaluate the effect of the interstellar radiation field on the inverse Compton emission of electrons accelerated in a supernova remnant located at different distances from the Galactic Centre. We show that contribution of optical and infra-red photons to the inverse Compton emission may exceed the contribution of cosmic microwave background and in some cases broaden the resulted \\gray\\ spectrum. Additionally, we show that if a supernova remnant is located close to the Galactic Centre its \\gray\\ spectrum will exhibit a ``universal'' cutoff at very high energies due to the Klein-Nishina effect and not due to the cut-off of the electron spectrum. As an example, we apply our calculations to the supernova remnants RX J1713.7-3946 and G0.9+0.1 recently observed by HESS. ", "introduction": "A new calculation \\citep{PS05} of the Galactic interstellar radiation field (ISRF) consistent with multi-wavelength observations by DIRBE and FIRAS indicates that the energy density of the ISRF is higher, particularly in the inner Galaxy, than previously thought. This has implications for the inverse Compton (IC) scattering of electrons and positrons and other electromagnetic processes in the interstellar medium \\citep{Moskalenko2000,Moskalenko2006}. Another place where the enhanced ISRF may play a role is the IC emission off very high energy (VHE) electrons accelerated in a supernova remnant (SNR) shock. SNRs are believed to be the primary sources of cosmic rays in the Galaxy. Observations of X-ray \\citep{Koyama1995} and \\gray\\ emission \\citep{Aharonian2005b,Aharonian2006} from SNR shocks reveal the presence of energetic particles, thus testifying to efficient acceleration processes. Acceleration of particles in collisionless shocks is a matter of intensive research in conjunction with the problem of cosmic ray origin \\citep{Drury1983,Blandford1987,Jones1991}. Current models include nonlinear effects \\citep[e.g.,][]{Berezhko2006} and treat particle acceleration using hydrodynamic codes \\citep[e.g.,][]{Ellison2005}. The predicted spectrum of accelerated particles has a power-law form in rigidity with index which may slightly vary around $-2.0$. Conventional estimates put the maximum reachable particle energy (protons) at or just below the knee energy; for the case of electrons the maximum energy is lower due to the synchrotron and IC energy losses \\citep{Ellison2005}. Young SNRs may be capable of particle acceleration up to $10^{17}$ eV due to the effect of magnetic field amplification \\citep{Bell2001} around the shock and assuming Bohm diffusion and a Sedov expansion law. The VHE \\gray\\ emission from shell-type SNRs has been modelled using leptonic (IC) and hadronic ($\\pi^0$-decay) scenarios \\citep[e.g.,][]{Baring2005}. The leptonic scenario fits the broad-band spectrum of a SNR assuming a pool of accelerated electrons scattered off the cosmic microwave background (CMB) producing VHE \\grays\\ while the magnetic field and electron spectrum cut-off are tuned to fit the radio and X-ray data \\citep[e.g.,][]{Lazendic2004}. The hadronic model fits the VHE \\gray\\ spectrum assuming a beam of accelerated protons hits a target, such as a nearby molecular cloud \\citep{Aharonian2002}. The latter, if definitively proven, would be the first experimental evidence of proton acceleration in SNRs. While there is no clear distinction between different models of VHE emission from SNRs, some authors tend to prefer the hadronic scenario since it fits better the observed spectral shape \\citep{Aharonian2006,Berezhko2006}. Such a preference, however, is made based on a simplified ``one-zone'' model which typically includes CMB photons only. In this paper we evaluate the effect of the ISRF on the IC emission of electrons accelerated in SNRs located at different distances from the Galactic Centre (GC). As examples, we apply our calculations to the shell-type SNR RX J1713.7-3946 and composite SNR G0.9+0.1 recently observed by HESS \\citep{Aharonian2005a, Aharonian2006}. ", "conclusions": "We have evaluated the effect of the ISRF on the IC emission of electrons accelerated in SNRs. Our calculation shows that for a SNR located in the inner Galaxy the contribution of optical and IR photons to the IC emission exceeds the contribution of the CMB and in some cases broadens the resultant \\gray\\ spectrum. SNRs located close to the GC exhibit a ``universal'' cut-off at VHEs due to the KN effect and not due to a cut-off in the electron spectrum. We have made a calculation of the IC and synchrotron emission for simple power-law electron spectra including the contribution by the ISRF using the shell-type SNR RXJ1713.7-3946 and composite SNR G0.9+0.1 as examples. Within the confines of a simple one-zone leptonic model, it is possible to fit the observed flux spectrum with a reasonable combination of model parameters. Observations in the GeV to sub-TeV range by the GLAST experiment will be critical in distinguishing between leptonic or hadronic scenarios of \\gray\\ production in SNRs as the predictions for the spectral shape in this energy range are distinctly different. Finally, we point out at least two cases for the leptonic model where the contribution of the ISRF will be critical. First will be when fitting the synchrotron peak dictates a relatively {\\it low-energy} cut-off in the electron spectrum. In this case, the VHE \\grays\\ will be produced by lower-energy electrons scattered off optical and IR photons. Second will be when fitting the synchrotron peak dictates a {\\it high-energy} cut-off in the electron spectrum. An equal or exceeding contribution of optical and IR protons, but with lower energy cut-off in the spectrum of upscattered \\grays\\ due to the KN effect, will effectively broaden the IC peak. Observations of the TeV emission from SNRs in the inner Galaxy may thus serve as a probe of the ISRF near their location. In turn, observations of shell-type SNRs in the outer Galaxy where the CMB photons will provide the majority of the IC emission may be used to evaluate the spectrum of accelerated electrons in the shell and for studies of shock acceleration." }, "0607/astro-ph0607172_arXiv.txt": { "abstract": "We present results of fitting the 50-day time series of photometry of $\\alpha$ Cen A taken by the WIRE satellite in 1999. Both power spectrum and autocovariance function (ACF) fitting techniques were used in an attempt to determine mode frequencies, rotational splittings, lifetimes and amplitudes of low-$\\ell$ p-modes. In all, using both techniques, we managed to fit 18 modes (seven $\\ell$ = 0, eight $\\ell$ = 1 and three $\\ell$ = 2) with frequencies determined to within 1 - 2 $\\mu$Hz. These estimates are shown to be 0.6 $\\pm$ 0.3 $\\mu$Hz lower, on average, than the frequencies determined from two other more recent studies (\\citeauthor{Bouchy2002} \\citeyear{Bouchy2002}; \\citeauthor{Bedding2004} \\citeyear{Bedding2004}), which used data gathered about 19 months after the WIRE observations. This could be indicative of an activity cycle, although due to the large uncertainty, more data would be needed to confirm this. Over a range of 1700 to 2650 $\\mu$Hz we were also able to use the ACF fitting to determine an average lifetime of 3.9 $\\pm$ 1.4 days, and an average rotational splitting of 0.54 $\\pm$ 0.22 $\\mu$Hz, which is the first ever reliable estimate of this parameter. In contrast to the ACF, the power spectrum fitting was shown to return significantly biased results for these parameters. ", "introduction": "The past ten years have seen a number of increasingly successful attempts to detect and measure solar-like oscillations in other stars. Due to its proximity and similarity to the Sun, many of these studies have been focused on the star $\\alpha$ Cen A. The first clear detection of p-mode oscillations on this star was made by \\cite{Schou2001} from photometry using the Wide-Field Infrared Explorer (WIRE) satellite taken over a 50-day period. \\cite{Schou2001} correctly determined the large frequency separation but, unfortunately, wrong $\\ell$ identifications were made and hence an incorrect value for the small separation was determined. Further detections and the first correct mode identifications were made by \\cite{Bouchy2002} using a 13-day run of velocity measurements taken by the CORALIE spectrograph. More recently, \\cite{Bedding2004} determined the frequencies for over 40 individual modes from observations by the UVES and UCLES instruments taken over a period of 5 nights. The main driving force behind each subsequent study has been to improve the signal-to-noise ratio (SNR) in order to initially detect as many modes as possible and then to better constrain the limits placed on the determined frequencies. It is of course also important to improve resolution, but practical constraints have meant all these studies were limited in the length of observations that could be made. This has meant accurate determination of mode parameters such as power, rotational splitting and lifetime has been difficult. Here, we apply two sophisticated fitting procedures to the WIRE $\\alpha$ Cen A data collected in 1999 in order to improve the parameter determinations. Although this data set has the poorest SNR of the three studies mentioned above, it does have the longest time series. Hence, we would expect to extract more reliable estimates of the average lifetime and rotational splitting of the $\\alpha$ Cen A modes. The first fitting procedure we applied was a traditional power spectrum fitting method. This involved taking the Fourier transform of the time series and then fitting a Lorentzian-like model to the various mode peaks in the resulting power spectrum. The second procedure used is a new technique based on fitting the autocovariance function (ACF) of the time series (i.e., the unnormalized autocorrelation function). Since the modes seen in the ACF are all superimposed, one must first filter the time series in order to isolate the modes one is hoping to fit. The ACF can then be computed and a model based on an exponentially decaying, periodic function fitted to the result. This technique was first introduced as a possible method for mode determination by \\cite{Gabriel1998} and developed more fully by \\cite{Fletcher2004} in an attempt to better constrain mode parameters of long solar p-mode data sets. In the initial investigation of the WIRE data by \\cite{Schou2001} only a handful of modes were identified due to the poor SNR in the data set. However, a distinct advantage of revisiting this data comes in having a large number of robust mode identifications from the aforementioned two later studies. This provides additional a-priori information that we can use as initial `guess' values for our fitting procedures. In order to fully test our fitting of the WIRE data we also generated a set of artificial time series. These data were created specifically to mimic the WIRE data allowing us to explicitly determine the precision and robustness of our fitted parameters. We detail the creation of this simulated data in Section~\\ref{Data_Sec}. In Section~\\ref{ModelFit_Sec} we describe in detail the procedure involved in applying the two fitting techniques to the data. Finally, in Section~\\ref{Results_Sec}, we go on to present and analyse the results of our fitting for the mode frequencies, amplitudes, lifetimes and rotational splitting parameters. ", "conclusions": "\\label{Summary_Sec} The 50-day time series of photometry observations taken in 1999 by the WIRE spacecraft has been reanalysed using power spectrum and autocovariance fitting methods. With the help of a-priori information, regarding the location in frequency of modes from other recent $\\alpha$ Cen A studies, we have managed to fit 18 different modes in the power spectrum and autocovariance function (ACF), 16 of which were fitted in both. The values of the fitted frequencies are slightly lower than those determined by \\citeauthor{Bedding2004} \\citeyear{Bedding2004} and \\citeauthor{Bouchy2002} \\citeyear{Bouchy2002}, although without better data we cannot say with any confidence whether this is indicative of an activity cycle for $\\alpha$ Cen A. In addition to the frequencies we have also been able to estimate an average rotational splitting across the fitted modes of 0.54 $\\pm$ 0.22 $\\mu$Hz using the ACF. An average lifetime was also estimated by fitting the ACF and was found to be 3.9 $\\pm$ 1.4 days. Although the actual fitted value is larger than the lifetime estimated by \\cite{Kjeldsen2005}, the error bars do overlap. Estimates of the amplitude were also obtained, however they were rather poorly constrained, especially for the weaker modes. Simulated time series made to mimic the WIRE data were created in order to test the accuracy and precision of the fitting methods using a Monte Carlo approach. We found that for the most part, the fitted parameter estimates averaged over a number of realizations agreed with the input values used to create the data. However, we did find that fits to the power spectrum tended to underestimate the linewidths and overestimate the splittings. The bias on both of these parameter estimates were reduced when fitting the ACF. There still may be opportunities to refine this work further. For example, in this analysis, a fairly basic approach to dealing with the window function was employed, simply allowing for the subsequent sidebands in the models. A more sophisticated approach for fitting the power spectrum would be to convolve the spectral window with the model and fit that to the data. Also, because of the Weiner-Khinchine relation, that states that the ACF is actually the Fourier transform of the power spectrum, this technique can probably be applied to the ACF fitting approach as well. Additionally, there is now a new set of WIRE $\\alpha$-Cen observations that was taken in January 2004 and lasted for around 30 days. If modes can be identified and fitted from this data as well, it will give an excellent comparison with the 1999 time series and should allow for a better investigation into a possible activity cycle." }, "0607/astro-ph0607491_arXiv.txt": { "abstract": "The Northern HIPASS catalogue (NHICAT) is the northern extension of the HIPASS catalogue, HICAT \\citep{meyer04}. This extension adds the sky area between the declination range of $+2^{\\circ} < \\delta < +25^{\\circ}30\\arcmin$ to HICAT's declination range of $-90^{\\circ} < \\delta < +2^{\\circ}$. HIPASS is a blind \\HI\\ survey using the Parkes Radio Telescope covering $71\\%$ of the sky (including this northern extension) and a heliocentric velocity range of -1,280 \\kms\\ to 12,700 \\kms\\ . The entire Virgo Cluster region has been observed in the Northern HIPASS. The galaxy catalogue, NHICAT, contains {\\em{1002}} sources with $v_{\\rm{hel}} > 300$ \\kms\\ . Sources with $-300$ \\kms\\ $< v_{\\rm{hel}} < 300$ \\kms\\ were excluded to avoid contamination by Galactic emission. In total, the entire HIPASS survey has found {\\em{5317}} galaxies identified purely by their HI content. The full galaxy catalogue is publicly-available at $\\langle${\\bf{\\tt http://hipass.aus-vo.org}}$\\rangle$. ", "introduction": "The \\HI\\ Parkes All-Sky Survey (HIPASS) survey is a blind \\HI\\ survey using the Parkes Radio Telescope\\footnote{The Parkes telescope is part of the Australia Telescope which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO.}, and the Northern extension increases this survey by a further 37 percent in sky area. The primary objective of extending Southern HIPASS to the north is to complement the southern census of gas-rich galaxies in the local Universe. The \\HI\\ mass function, HIMF, \\citep{zwaan03} and galaxy two-point correlation function \\citep{meyer05} based on Southern HIPASS showed that the statistical measures of the galaxy population from HIPASS are limited by cosmic variance. Recently, \\citet{zwaan05} used HICAT to investigate the effects of the local galaxy density on the HIMF. Using the $n$-th nearest neighbour statistic, they found tentative evidence that the low-mass end of the HIMF becomes steeper in higher density regions. These authors were able to examine the trend in the slope of the HIMF for different values of $n$ in the statistic. Larger values of $n$ correspond to sampling the density on larger scales. For each value up to $n=10$, the slope became systematically steeper as the density increased. Thus it appears that the \\HI\\ properties of galaxies might be influenced by environmental effects on quite large scales (where a typical separation of the $n=10$ nearest neighbours is $\\sim$ 5 Mpc), in addition to the well-known local effects, such as tidal interactions between neighbouring galaxies. Previously, \\citet{rosenberg02} found $\\alpha \\approx -1.2$ and $\\alpha \\approx -1.5$ for the slope of the HIMF in the immediate and field regions of the Virgo Cluster, respectively. Also, \\citet{spring05}, found flatter slopes to the low mass end of the HIMF in higher density regions. However their galaxy sample was selected optically. Since Northern HIPASS covers the entire Virgo Cluster region, the Northern HIPASS catalogue (NHICAT) can be used in conjunction with the Southern HIPASS catalogue (HICAT) to explore these trends and investigate the effects of cosmic variance on HIPASS galaxy catalogue statistics. It is worth noting that the Northern HIPASS also provides the first blind \\HI\\ survey of the entire region in and around the Virgo Cluster. Assuming a Virgo distance of 16 Mpc and an integrated flux limit of 15 Jy \\kms\\, this corresponds to a mass sensitivity of $9 \\times 10^8$ \\msun. Thus the survey will detect any \\HI\\ clouds above this mass limit in the vicinity of the Virgo Cluster, regardless of stellar content. The Virgo Cluster provides a nearby example of processes that are more common at higher redshifts, such as galaxy-galaxy and galaxy-intracluster medium interactions. Northern HIPASS will be used to investigate the role of \\HI\\ in a cluster environment in individual galaxies as well as statistically across the whole cluster. Understanding the role of \\HI\\ is vital for galaxy evolution models. \\citet{kenney04} found 6 galaxies in the Virgo Cluster showing distorted \\HI\\ morphology. Using N-body simulations, \\citet{vollmer01} investigated the effect of ram pressure stripping in the Virgo Cluster and found that \\HI\\ deficiency is dependent on galaxy orbits within the cluster. They concluded that all the galaxies showing some form of distorted \\HI\\ distribution have already passed through the centre of the cluster and are not infalling for the first time. The catalogue of extragalactic \\HI\\ sources from HIPASS was named HICAT and was presented in \\citet{meyer04} (hereafter known as MZ04), while the completeness and reliability of HICAT was assessed by \\citet{zwaan04}. Here we present a catalogue of extragalactic \\HI\\ sources from Northern HIPASS, named NHICAT. The basic parameters of HIPASS, Northern HIPASS, HICAT and NHICAT are given in Table~\\ref{params}. Apart from the declination coverage, the main difference between the two surveys and catalogues is the higher noise level in Northern HIPASS. For a full summary of parameters of existing blind \\HI\\ surveys, including subsets of HIPASS (HIPASS Bright Galaxy Catalogue, \\citet{korib04}; the South Celestial Cap catalogue, \\citet{kilborn02}) see Table 1 of MZ04. The Northern HIPASS Zone of Avoidance (NHIZOA) survey by \\citet{donley05} covers northern declinations of the Galaxy - a subset of the Northern HIPASS area - at a higher sensitivity (RMS = 6 mJy beam$^{-1}$). Optical identification of NHICAT sources will use similar techniques to HICAT \\citep{doyle05} and will be presented in a later paper. With a total spatial coverage of 29,343 square degrees and 5317 sources, the combined HICAT and NHICAT catalogue is the largest purely \\HI-selected galaxy catalogue to date. The Arecibo L-Band Feed Array (ALFALFA) surveys will eventually cover the same region of sky as Northern HIPASS and will extend up to a declination of $+36^{\\circ}$. More information about the progress of ALFALFA can be found online at $\\langle${\\bf{\\tt http://egg.astro.cornell.edu/alfalfa}}$\\rangle$ \\citep{giov05}. In this paper we present NHICAT, together with the completeness and reliability analysis of the catalogue. Section 2 reviews Northern HIPASS and its properties. The source identification and the generation of the catalogue is described in Section~\\ref{datasection}. Section~\\ref{noisesection} discusses the noise statistics of Northern HIPASS and the completeness of NHICAT is analysed in Section 3.1. The narrowband follow-up observations and the reliability of NHICAT will be described in Section 3.2. \\begin{table*} \\begin{center} \\caption{Survey and catalogue parameters.} \\label{params} \\begin{tabular}{lccccc} \\hline Survey name & Survey range & RMS & Catalogue Name & Catalogue range & Sources \\\\ & ($deg$, \\kms\\ ) & (mJy beam$^{-1}$) & & ($deg$, \\kms\\ ) & \\\\ \\hline HIPASS & $\\delta < +2^{\\circ}$, & 13 & HICAT & $\\delta < +2^{\\circ}$, & 4315\\\\ (MZ04) & $-1280$ 3 s). A comparison with the few high magnetic field radio pulsars (having magnetic field strengths above the quantum critical field line; Figure \\ref{fig1}) is not statistically significant, having only four objects in this class. However we note that all their luminosities, as well as the radio luminosity of \\xte \\citep{crh+06}, are greater than the limits obtained for the AXPs. \\begin{figure} \\begin{center} \\includegraphics[width=8truecm]{burgayf6.ps} \\caption{Distribution in luminosity at 1400 MHz of the observed pulsar population of the galactic field (solid line) and of the intrinsic population (dashed line; \\citealt{lfl+06}). The shaded region indicates the luminosity limits reached in the present search.} \\label{fig6} \\end{center} \\end{figure} Note however that, if we compare our results with the {\\it{intrinsic}} luminosity distribution of the radio pulsars (dashed line in figure \\ref{fig6}; \\citealt{lfl+06}), we obtain a probability of $\\sim 76$\\% that our observations are not deep enough to detect a radio pulsed signal from our targets. In considering the causes of the non detection of a radio pulsed signal from our four targets, besides the luminosity bias, we must take into account the possibility that, although the X-ray beam is pointing toward us, the radio beam, usually narrower, is not. Assuming a pulse duty cycle of $\\sim 5\\%$ (hence a beam semi-aperture $\\geq 9^{\\circ}$), typical of long period pulsars and similar to that of the radio pulsed signal detected by \\citet{crh+06} in the transient AXP \\xte ($\\sim 4\\%$ at 1.4 GHz), we can calculate, following e.g. \\citet{bbp+03}, that the probability that such a narrow radio beam misses the earth is $\\leq 77\\%$. The composite probability that the beams of all four AXPs are not pointing toward us is hence $\\leq 34\\%$. The non detection of RRAT-like bursts from any of these AXPs, despite our long exposures, seems to weaken the hypothesis that RRAT bursts might be related to the short bursts observed from the magnetars leaving us with other plausible conjectures of a relation with other classes of neutron stars such as middle aged radio pulsars displaying giant pulses \\citep{rbg+06,wsrw06} or with X-ray Dim Isolated Neutron Stars \\citep{ptp06}. The only case of a detection of radio pulsations from an AXP concerns the only confirmed transient magnetar \\xte\\, \\citep{crh+06}. Radio emission from this source is strongly related with the occurrence of an outburst of its X-ray emission \\citep{hgb+05}, as well as an IR enhancement \\citep{rti+04}. Furthermore, whereas the X-ray flux is decaying exponentially with timescale of a few hundreds days \\citep{gh05}, \\xte\\, radio emission is still on more than 3 years after the X-ray outburst. Interestingly the sole other possible transient AXP is the candidate \\axj, one of our targets. Our radio observations of this source were performed more than six years after its possible X-ray outburst occurred in 1993, hence unfortunately nothing can be safely concluded from our upper limits, in favor or against the possible radio and X-ray correlation during the outbursts of this source. However, assuming that \\axj experienced, after the X-ray outburst, a phase of radio emission similar to that of \\xte, our null detection implies that the fading of the radio emission has a time scale of the order of few years: in particular, if \\axj at the onset of its putative radio emission phase had a similar luminosity as \\xte, this would imply a decrease in $L_{1400}$ of a factor of $\\sim 20$ over six years." }, "0607/astro-ph0607422_arXiv.txt": { "abstract": "The sequence of massive star supernova types IIP (plateau light curve), IIL (linear light curve), IIb, IIn (narrow line), Ib, and Ic roughly represents a sequence of increasing mass loss during the stellar evolution. The mass loss affects the velocity distribution of the ejecta composition; in particular, only the IIP's typically end up with H moving at low velocity. Radio and X-ray observations of extragalactic supernovae show varying mass loss properties that are in line with expectations for the progenitor stars. For young supernova remnants, pulsar wind nebulae and circumstellar interaction provide probes of the inner ejecta and higher velocity ejecta, respectively. Among the young remnants, there is evidence for supernovae over a range of types, including those that exploded with much of the H envelope present (Crab Nebula, 3C 58, 0540--69) and those that exploded after having lost most of their H envelope (Cas A, G292.0+1.8). ", "introduction": "Core collapse supernovae show considerable diversity among their properties. A basic observational division is into the SNe II (Type II supernovae), which have hydrogen in their spectra, and SNe Ib/c, which do not (or have weak hydrogen lines). The reason for the difference is that the progenitors of the SN Ib/c have lost their H envelopes, and perhaps more, during their evolution leading up to the supernova. The mass loss can occur either through the winds from a single star or can be aided by interaction with a binary companion. The SNe II show strong diversity themselves. Their observational classification is based on a variety of factors, but it is clear that presupernova mass loss plays a significant role in determining the type. Two types are distinguished by their light curves: IIP (plateau) and IIL (linear). Models of Type IIP light curves have long showed that the likely progenitors of the SNe IIP are the red supergiants that end their lives with most of their H envelopes retained (\\cite[Grasberg et al. 1971]{GIN71}, \\cite[Chevalier 1976]{Chev76}). The plateau phase of the light curve is due to the internal energy deposited by the initial explosion. This progenitor hypothesis has been directly confirmed by observations of the progenitors of a number of SNe IIP (\\cite[Hendry et al. 2006]{Hen06} and references therein). While the SNe IIP might explode with a hydrogen envelope of $\\sim10\\Msun$, the more rapid decline of the SNe IIL imply that they explode with an envelope of $\\sim1\\Msun$ (\\cite[Blinnikov \\& Bartunov 1993]{BB93}). Because of higher rates of mass loss for more luminous stars, the reduced H envelope is expected to occur for single stars with initial masses of $\\gsim20\\Msun$. Alternatively, mass loss in a binary system could play a role in the reduced envelope mass. The prototype of the SNe IIb was SN 1993J, which made a transition from a Type II at early times to a Type Ib/c at late times, based on spectroscopic observations. The H envelope mass required for SN 1993J was $\\sim0.2\\Msun$ (\\cite[Woosley et al. 1994]{Woo94}). For this to occur in a single star requires special timing, so a binary origin is preferred. A likely binary companion for SN 1993J has been directly observed (\\cite[Maund et al. 2005]{Mau05}). SNe IIn have the spectroscopic feature of narrow emission lines (\\cite[Schlegel 1990]{Sch90}), typically H$\\alpha$, which indicates that circumstellar interaction plays a role in the emission from early times. A supernova can be a Type IIn and another type; e.g., SN 1998S was both a IIn and IIL. Because of the strong circumstellar interaction, it can be difficult to determine the nature of the photospheric emission in a SN IIn. The H emission from circumstellar interaction implies strong mass loss before the supernova in a SN IIn, so the H envelope is likely to be depleted at the time of the supernova. The most noteworthy peculiar SN II is the nearby SN 1987A, which was relatively compact at the time of the explosion although it had a massive H envelope. The best explanation for the explosion as a blue supergiant star and the axisymmetric ring features around it is probably that it was in a binary system (\\cite[Podsiadlowski 1992]{Pod92}). The Type Ib/c supernovae are believed to be H-poor Wolf-Rayet stars at the time of their explosion. There is some observational evidence that SNe Ib, which have He lines, can have a small amount of high velocity H at the time of the explosion (\\cite[Elmhamdi et al. 2006]{Elm06}). Although the presence of a H envelope in a massive star typically leads to the formation of a red supergiant in the late evolutionary stages, a small amount of H mass ($\\lsim0.01\\Msun$) is not expected to support an extended envelope. These considerations show that a major factor in the determination of supernova type is the amount of H left in the envelope at the time of the supernova. If the H envelope mass is greater than the core mass, then the core is effectively decelerated by the envelope and there is mixing between them by Rayleigh-Taylor instabilities. There is not only the outward mixing of heavy elements, but also the inward mixing of H to low velocities. This can be directly observed in the late spectra of SNe IIP; e.g., late spectra of SN 1999em showed H moving at several $100\\kms$ (\\cite[Elmhamdi et al. 2003]{Elm03}). The relative numbers of the different kinds of supernovae is uncertain from an observational point of view. If all the supernovae came from single stars, the stellar mass function was the Salpeter function ($n(M)\\propto M^{-2.35}$), and Type IIP supernovae came from $8-20\\Msun$ stars, Type IIL $20-25\\Msun$, and Type Ib/c $>25\\Msun$, the relative fractions of IIP:IIL:Ib/c would be 0.71:0.08:0.21. Binary evolution could increase the relative number of IIL and Ib/c events (\\cite[Nomoto et al. 1996]{Nom96}). However, SNe IIP are likely to be an important component of the core collapse supernovae. After the explosion of the progenitor star, the event has an afterlife in two ways: through its interaction with the surrounding medium and through the possible activity of a central compact remnant (neutron star or black hole). I will discuss how the expectations for the afterlife phases depend on the supernova type and what can be learned about these phases from observations. The early circumstellar interaction observed in extragalactic supernovae is discussed in Section 2, the pulsar wind nebula expansion inside the supernova in Section 3, and circumstellar interaction in young remnants in Section 4. More details on the material in Sections 3 and 4 can be found in \\cite{Chev05}. Section 5 contains a discussion of future prospects. ", "conclusions": "\\label{sec:dc} There are excellent future prospects for developing a more complete picture of the massive star evolution leading up to a supernova and the subsequent expansion of the supernova into the circumstellar medium. The increasing number of {\\it Hubble Space Telescope} images of galaxies has improved the prospects for identifying the progenitor stars of nearby supernovae. Follow up observations at radio and X-ray wavelengths can then reveal the mass loss environment for that particular progenitor. There is a growing number of young remnants that have observed pulsar wind nebulae and/or circumstellar interaction. Many of these have been well observed at X-ray wavelengths (owing to {\\it Chandra} and {\\it XMM}), but are less well observed at optical and infrared wavelengths. Infrared observations seem especially important because a number of the objects have high extinction. As the amount of information increases, there is the possibility of looking for correlations between the nature of the compact object in a remnant and the nature of the surrounding supernova. An initial examination of this point (\\cite[Chevalier 2005]{Chev05}) did not reveal any correlations. Along with these endeavors, hydrodynamic modeling of the variety of supernova events, along with their interaction with mass loss, is needed. The result will be a better understanding of the final evolution of massive stars and the variety of possible outcomes." }, "0607/astro-ph0607278_arXiv.txt": { "abstract": "From examination of only 4 deg$^2$ of sky in the NOAO Deep Wide-Field Survey (NDWFS) region, we have identified the first radio-loud quasar at a redshift $z>6$. The object, FIRST J1427385+331241, was discovered by matching the FLAMEX IR survey to FIRST survey radio sources with NDWFS counterparts. One candidate $z>6$ quasar was found, and spectroscopy with the Keck II telescope confirmed its identification, yielding a redshift $z=6.12$. The object is a Broad Absorption Line (BAL) quasar with an optical luminosity of $M_B\\sim-26.9$ and a radio-to-optical flux ratio $\\sim 60$. Two \\ion{Mg}{2} absorptions systems are present at redshifts of $z=2.18$ and $z=2.20$. We briefly discuss the implications of this discovery for the high-redshift quasar population. ", "introduction": "High-redshift quasars provide both interesting constraints on the growth of the first supermassive black holes, and light sources with which to probe the ionization history of the Universe. The Sloan Digital Sky Survey (SDSS) broke the $z=6$ barrier after covering its first 1550 deg$^2$ \\citep{Fan01} and has subsequently identified a total of nine objects with $z \\geq 6$ drawn from sky coverage of 6550 deg$^2$ \\citep{Fan01,Fan03,Fan04,Fan06}. Owing to the relatively bright limiting $z'$-band magnitude of the SDSS survey, all of these quasars are luminous, with $M_B<-26.5$. \\citet{Fan04} derive estimates for the space density ($6\\pm 2 \\times 10^{-10}$ Mpc$^{-3}$ for $M_{1450}<-26.7$) and luminosity function using nine SDSS-discovered quasars with $z>5.7$, concluding $\\Psi (L) \\sim L^{-3.2\\pm0.7}$. This steep luminosity function suggests deeper surveys could locate large numbers of high-redshift objects. All of the $z>6$ quasars discovered to date are radio-quiet, with ratios of radio to optical flux $<10$. One source has been detected at 20~cm with a flux density of $55\\mu$Jy, while a second has a measured value of $26\\pm 12~\\mu$Jy \\citep{carilli04}. The most distant quasar known with a flux density level of $>1$~mJy is SDSS J083643.85+005453.3 at $z=5.82$. However, based on semi-analytic models of dark halo formation, \\citet{hqb04} predict high surface densities of high-redshift radio-loud quasars; in particular, they suggest roughly four $z>6$ quasars per square degree should be found in the catalog constructed from Faint Images of the Radio Sky at Twenty-cm (FIRST -- \\citealt{FIRST}). Motivated by this suggestion, we have matched the FIRST catalog to the NOAO Deep Wide-Field Survey (NDWFS -- \\citealt{ndwfs}) and the recently released FLAMINGOS EXtragalactic Survey (FLAMEX -- \\citealt{flamex}), yielding a net survey area of 4.1 deg$^2$. The result is the detection of the first $z>6$ radio-loud quasar. In Section 2, we briefly describe the catalogs used for this project. We then go on to present our search procedure which produced one good candidate object (\\S 3). Section 4 describes the spectroscopic confirmation of this quasar, \\S 5 its continuum properties from rest UV to IR, and \\S 6 its absorption line systems. In \\S 7 we discuss the implication of this discovery for the high-redshift quasar population as well as for future searches for high-redshift objects. Magnitudes are reported in the Vega system unless otherwise noted. Photometry in the optical to near-infrared was obtained from SExtractor \\citep{1996A&AS..117..393B} MAG\\_AUTO magnitudes. Throughout the paper we use the standard cosmological model, with parameters $\\Omega_{\\lambda} = 0.7,\\ \\Omega_M = 0.3,\\ {\\rm and}\\ H_0 = 70$~km s$^{-1}$ Mpc$^{-1}$ (Spergel et al. 2003). ", "conclusions": "Approximately 10\\% of the normal SDSS quasars are detected in the FIRST survey (White et al. 2006). The fact that SDSS has found one $z>6$ quasar (all radio quiet) per $\\sim 730$ deg$^2$ searched and we have found a radio-loud object in a 4 deg$^2$ survey requires examination,\\footnote{One possible explanation is that we are, as previously documented, just lucky (Becker, Helfand, and White 1992; White, Kinney, and Becker 1993; Stern et al. 2005).} as a naive estimate would predict our chances of success to be 4.1~deg$^2/730$~deg$^2 \\times 0.1 = 6 \\times 10^{-4}$. Allowing for the possibility that the quasar is being gravitationally magnified, we might have reached $\\sim 1$ mag deeper into the quasar luminosity function than Fan et al. Even with a steep luminosity function $\\Psi (L) \\sim L^{-3.2}$, this would only increase the number density by a factor of $\\sim 20$; the number density of $6\\times10^{-10}\\ {\\rm Mpc}^{-3}$ from \\citet{Fan04} would predict $\\sim 0.1$ quasars in 4 deg$^2$ in the redshift range $65.8$ quasars (see Figure~\\ref{maglimits}); it is the object's red optical colors that prevented it from being detected by SDSS. As at low-redshift, radio selection has located quasars that are redder than those found by optical criteria \\citep{white03}. The improbable discovery of this quasar suggests that dust obscuration may be important at high redshift, and that current estimates for the quasar number density at $z\\sim6$ are too low. Radio searches such as this one, and mid- to far-IR searches such as that of \\citet{cool} are likely to find a new population of sources at high redshift. The issue of the redshift dependence of quasar radio emission has been debated in the literature for decades. The consensus ten years ago appeared to be that ``the fraction of radio-loud quasars decreases with increasing redshift'' (Schmidt et al. 1995 and references therein). From observations of forty optically selected quasars with $3.33.9$) Schmidt et al. detected three objects at a 20~cm flux density threshold of 0.2~mJy, far below their expectation of 9-18 detections. However, all such studies relied on rather small and heterogenous samples of objects with different selection criteria. Recently, we have studied the radio properties of the 41,295 quasars from the SDSS DR3 catalog that fall within the FIRST survey area (White et al. 2006). An important feature of our analysis is that we include {\\it all} quasars -- especially the $\\sim 90\\%$ which fall below the radio detection threshold -- by stacking the radio images. \\citet{white06} find a very tight correlation of radio luminosity with absolute magnitude at all redshifts, but, importantly, the slope of the relation is not unity: $L_R \\sim L_{opt}^{0.72}$. Thus, the higher luminosity optical objects -- inevitably those found in magnitude-limited high-$z$ surveys -- are underluminous in the radio when compared with the mean of the lower-redshift population. When normalized for this dependence on optical luminosity, we find a remarkably flat distribution of the radio-to-optical flux ratio $R$ for quasars with redshift: the mean value of $R$ changes by less than 0.1 from $z=0$ to $z=5$ (White et al. 2006, figure 11; see also Cirasuolo et al. 2006 and \\citealt{petric}). The cumulative fraction of quasars above the Schmidt et al. (1995) limit of 0.2~mJy at 6~cm is $\\sim 0.12$ in our SDSS sample (assuming a radio spectral slope of $\\alpha = -0.5$ between 20 and 6~cm). Correcting for the (high) mean optical luminosity of the Schmidt et al. sample ($M_B = -26.3$) reduces the expected fraction by 20\\% to 0.10. Thus, we would expect four detections in their sample of forty; three objects were detected. We conclude that the expectation of discovering a radio-loud quasar at high $z$ is not reduced significantly by a decline in radio emission with redshift; indeed, the fact that two of nineteen quasars now known above $z=5.7$ lie above the FIRST survey threshold is not unexpected. Our discovery of a $z>6$ radio-loud quasar in a 4-deg$^2$ survey suggests that tractable wider-area surveys with deep $K$-magnitude (and deeper radio) limits would be highly productive. We have not yet exhausted the possibilities in the FLAMEX region, as the $\\sim200$ radio sources showing extended emission have yet to be matched. The detection of more, higher-$z$ objects could provide useful probes of the epoch of reionization through redshifted 21~cm absorption measurements." }, "0607/hep-ph0607267_arXiv.txt": { "abstract": "We consider the decoupling of neutrinos in the early Universe in presence of non-standard neutral current neutrino-electron interactions (NSI). We first discuss a semi-analytical approach to solve the relevant kinetic equations and then present the results of fully numerical and momentum-dependent calculations, including flavor neutrino oscillations. We present our results in terms of both the effective number of neutrino species ($N_{\\rm eff}$) and the impact on the abundance of $^4$He produced during Big Bang Nucleosynthesis. We find that the presence of neutrino-electron NSI may enhance the entropy transfer from electron-positron pairs into neutrinos instead of photons, up to a value of $N_{\\rm eff}\\simeq 3.12$ for NSI parameters within the ranges allowed by present laboratory data, which is almost three times the effect that appears for standard weak interactions. Thus non-standard neutrino-electron interactions do not essentially modify the density of relic neutrinos nor the bounds on neutrino properties from cosmological observables, such as their mass. ", "introduction": "\\label{sec:introduction} In the early Universe, neutrinos were kept in thermal contact with the electromagnetic primordial plasma by rapid weak interactions with electrons and positrons. When the temperature dropped below a few MeV, these weak processes became ineffective and the process of neutrino decoupling took place, while shortly after the $e^\\pm$ pairs began to annihilate almost entirely into photons thus producing a difference between the temperatures of the relic photons and neutrinos. This difference can be easily calculated if we assume that neutrinos were completely decoupled when the $e^\\pm$ pairs transferred their entropy to photons, leading to the well-known temperature ratio $T_\\gamma/T_\\nu=(11/4)^{1/3}\\simeq 1.40102$. Indeed, this simplified picture should be improved since some relic interactions between $e^\\pm$ and neutrinos exist all along the $e^\\pm$ annihilation stage, leading to a slightly smaller increase of the comoving photon temperature and to small distortions (at the percent level) of the neutrino momentum distributions. Presently, there exist compelling evidences for flavor neutrino oscillations from a variety of experimental data on solar, atmospheric, reactor and accelerator neutrinos (see e.g.\\ \\cite{Maltoni:2004ei,Fogli:2005cq}). These results are well understood by assuming that neutrinos have masses and mix, which in turn seems to point out the necessity of some new physics beyond the Standard Model (SM) of fundamental interactions. Interestingly, non-zero neutrino masses usually come with non-standard interactions (NSI) that might violate leptonic flavor and/or break weak universality. Recent analyses \\cite{Berezhiani:2001rs,Berezhiani:2001rt,Davidson:2003ha,Barranco:2005ps} have considered the neutral current NSI in a phenomenological way, showing that they can be bound using measurements of neutrino-electron scattering, as well as data from LEP and from related charged lepton processes. The aim of this paper is to study the neutrino decoupling process in presence of additional interactions between neutrinos and electrons, a possibility already noted in \\cite{Berezhiani:2001rs,Davidson:2003ha}. In this case, neutrinos could be kept in longer contact with $e^\\pm$ and thus share a larger amount of the total entropy transfer than in the SM. Actually, if the non-standard neutrino-electron interactions were large enough, the neutrino momentum distribution would be significantly different from the standard case. In turn, this would modify the final yield of light nuclei during the epoch of Big Bang Nucleosynthesis (BBN), as well as the radiation content of the Universe, affecting the anisotropies of the Cosmic Microwave Background (CMB) and the power spectrum of Large Scale Structures (LSS). Our goal is to calculate how the decoupling is modified taking into account NSI with couplings which are still allowed by present laboratory data, and to discuss the possibility that cosmological observations can be used as a complementary way to bound these exotic scenarios. The paper is organized as follows. We begin in Sec.\\ \\ref{NSIsummary} by describing the formalism adopted for the non-standard electron-neutrino interactions and summarize the current bounds from a variety of experimental data. We then consider the process of relic neutrino decoupling in the presence of non-standard electron-neutrino interactions, giving first an estimate by using a semi-analy\\-tical approach in Sec.\\ \\ref{sec:decoupling}. Finally, in Sec.\\ \\ref{numerical} we report the results of the full momentum-dependent numerical calculations for the neutrino spectra and the effect on the primordial $^4$He yield and other cosmological observables. We present our conclusions in Sec.~\\ref{conclus}. ", "conclusions": "The process of relic neutrino decoupling in the early Universe is sensitive to the strength of the interactions between neutrinos and the plasma formed by electrons and positrons. If neutrinos were kept in longer thermal contact with them than in the standard case, they would share a larger fraction of the entropy release from $e^\\pm$ annihilations. This would affect the predicted characteristics of the cosmic background of neutrinos (C$\\nu$B), which in turn could modify the late evolution of the Universe and the bounds on neutrino properties from the analysis of cosmological observables. In this paper we have considered how the decoupling of relic neutrinos is modified in presence of non-standard neutral current neutrino-electron interactions. First, we have provided a rough estimate of the size of NSI couplings needed to affect in a relevant way the properties of the C$\\nu$B and of related observables. We find that needed couplings are by far larger than the existing laboratory bounds, as also confirmed by a semi-analytical solution of the relevant kinetic equations. Thus, NSI can only play a minor role in shaping the C$\\nu$B. However, NSI might still contribute with comparable or larger distortions in the neutrino spectra than predicted in the SM. In order to quantify these deviations, we have performed fully numerical and momentum-dependent calculations of the density-matrix equations relevant for neutrino evolution in the early Universe, including the effects of flavor neutrino oscillations. Typically, we find that the enhancement in observables like $\\neff$, which measures the change in the radiation content, or the $^4$He mass yield in the BBN can be up to three times larger than the ones found in the standard case \\cite{Mangano:2005cc}, for values of all NSI parameters close to the limits placed by laboratory experiments. Nevertheless, even the variation of up to $0.2\\%$ that we found in the $^4$He abundance is yet too small (about one order of magnitude) when compared to the observational error on $Y_p$, which is unfortunately dominated by systematics, and as such difficult to pin down in the near future. Instead, a value of $\\neff\\simeq 3.12$, which is still possible within the present parameter space for NSI couplings, might even be barely detectable. Indeed, it has been shown that the CMB satellite PLANCK will soon provide temperature and polarization data that will probably measure the value of $\\neff$ with an uncertainty of order $\\sigma(\\neff)\\simeq 0.2$ \\cite{Bowen:2001in,Bashinsky:2003tk} (for a previous, more optimistic forecast see \\cite{Lopez:1998aq}), or better when combined with data from a large galaxy redshift survey such as SDSS \\cite{Lesgourgues:2004ps}. Future CMB missions may reach the sensitivity of $0.04-0.05$ needed to test the standard scenario \\cite{Bashinsky:2003tk}, and might achieve a 2$\\,\\sigma$ hint for $\\neff\\simeq 3.12$ and eventually test the effect of large neutrino-electron NSI. However, barring some extreme cases, we conclude that the prediction of $\\neff\\simeq 3$ within a few \\% is quite robust even when taking into account neutrino-electron NSI of the four-fermion type. Thus, their existence can not modify in a significant way the bounds on neutrino properties from cosmological observables, in particular on their masses, as recently reviewed in \\cite{Lesgourgues:2006nd}. Since it is likely that future experiments may narrow further the allowed range of NSI couplings, their role in a cosmological context will be even smaller, at most a sub-leading correction to the standard prediction of the C$\\nu$B properties. Turning the argument around, however, one can conclude that a significant deviation of $\\neff$ from 3 may require major revisions of the cosmological model, like the introduction of new relativistic relics and/or of an exotic thermal history, or both (see e.g.\\ \\cite{Serpico:2004nm,Cuoco:2005qr})." }, "0607/astro-ph0607015_arXiv.txt": { "abstract": "We report on detailed spectroscopic studies performed for the secondary star in the black hole binary (micro-quasar) V4641 Sgr in order to examine its surface chemical composition and to see if its surface shows any signature of pollution by ejecta from a supernova explosion. High-resolution spectra of V4641 Sgr observed in the quiescent state in the blue-visual region are compared with those of the two bright well-studied B9 stars (14 Cyg and $\\nu$ Cap) observed with the same instrument. The effective temperature of V4641 Sgr (10500 $\\pm$ 200 K) is estimated from the strengths of He~{\\sc i} lines, while its rotational velocity, $\\it v$ sin $\\it i$ (95 $\\pm$ 10 km s${}^{-1}$), is estimated from the profile of the Mg~{\\sc ii} line at 4481 \\AA. We obtain abundances of 10 elements and find definite over--abundances of N (by 0.8 dex or more) and Na (by 0.8 dex) in V4641 Sgr. From line-by-line comparisons of eight other elements (C, O, Mg, Al, Si, Ti, Cr, and Fe) between V4641 Sgr and the two reference stars, we conclude that there is no apparent difference in the abundances of these elements between V4641 Sgr and the two normal late B-type stars, which have been reported to have solar abundances. An evolutionary model of a massive close binary system has been constructed to explain the abundances observed in V4641 Sgr. The model suggests that the progenitor of the black hole forming supernova was as massive as $\\sim 35 M_\\odot$ on the main-sequence and, after becoming a $\\sim 10 M_\\odot$ He star, underwent \"dark\" explosion which ejected only N and Na-rich outer layer of the He star without radioactive $^{56}$Ni. ", "introduction": "We estimate the rotational velocity, $\\it v$ sin $\\it i$, of V4641 Sgr from the observed profile of the Mg~{\\sc ii} line at 4481 \\AA~, as shown in figure 2. First, we tried to reproduce the profiles of the line in 14 Cyg and $\\nu$ Cap by adopting the mean values of published data of $\\it v$ sin $\\it i$. They are 30 and 25 km s${}^{-1}$ for 14 Cyg (\\cite{abt95b}; \\cite{royer02}) and $\\nu$ Cap (\\cite{abt95a}; \\cite{dwo00}; \\cite{fekel03}; \\cite{royer02}), respectively. We find that acceptable reproductions of the observed line profile can be obtained using adopted values of $\\it v$ sin $\\it i$ for the two stars, as shown in figure 8. Next, we searched for a suitable value of $\\it v$ sin $\\it i$ in V4641 Sgr by exploring values between 60 and 150 km s${}^{-1}$; we found that the best fit is achieved at 95 $\\pm$ 10 km s${}^{-1}$. If we use $\\it v$ sin $\\it i$ = 123 km s${}^{-1}$, as obtained in \\citet{orosz01}, the simulated profile becomes too shallow and too wide (figure 2). We conclude that their result of $\\it v$ sin $\\it i$ is too large, most probably resulting from the low resolution of their spectral data. We tried to estimate the abundance of C in V4641 Sgr using the C~{\\sc ii} line at 4267 \\AA~ (figure 3). The line consists of two major components at 4267.001 \\AA~ and at 4267.261 \\AA~ (table 2). The C~{\\sc ii} feature in $\\nu$ Cap can be reproduced with a solar abundance of C, while the observed feature in 14 Cyg appears to be too weak. This may suggest either a slight under--abundance of C or an error in the adopted atmospheric parameters for 14 Cyg. We find that the broard feature at 4267 \\AA~ in V4641 Sgr can be reproduced by assuming a solar abundance of C. We find a shallow and broard absorption feature at 6483 \\AA~ in V4641 Sgr in both the May 20 and June 19 data. Comparisons with the two reference stars (figure 4) shows that the feature in V4641 Sgr appears to be too deep. The feature coincides with the positions of the five components of N~{\\sc i} lines (Multiplet No. 21). We simulated spectra of both reference stars, assuming solar abundances, and find that observations can be reasonably reproduced when excluding the sharp components due to atmospheric absorption. The only observable feature is the Fe~{\\sc ii} line at 6482.204 \\AA. On the other hand, the observed feature of V4641 Sgr at 6483 \\AA~ cannot be reproduced when we assume solar abundances of N and Fe. A reasonable reproduction can be obtained when a significant over--abundance of N (by around 1.0 dex) is assumed, as shown in figure 4. Allowing for a slight uncertainty in the continuum level because of the low S/N ratio, we conclude that N is indeed over--abundant in V4641 Sgr by at least 0.8 dex. This conclusion is in accordance with the result noted in \\citet{orosz01}, who obtained an over--abundance of N by 1.0 dex from an analysis of the same spectral feature. Analyses of weak O~{\\sc i} lines near 6455 \\AA~ are shown in figure 5. We conclude that the O~{\\sc i} feature can be reasonably fit by using a solar O abundance. This is in contrast with the result given in \\citet{orosz01}, who obtained an over--abundance of O (by a factor of three), when analysing the same spectral feature. Next, we analyse the Na~{\\sc i} D lines. The D lines are usually contaminated by interstellar absorption superposed on the stellar components, as shown for 14 Cyg and $\\nu$ Cap (figure 6). Fortunately, spectral lines of V4641 Sgr observed on July 19 are significantly redshifted (by +210 km s${}^{-1}$), and we can analyse both of the D1 and D2 components, which are unaffected by interstellar absorptions. The broad D lines in V4641 Sgr are too strong to be accounted for by the solar Na abundance. As is shown in figure 6, we need to enhance the abundance of Na by 0.7 dex in order to fit the D2 line. Further enhancement (by 0.2 dex) of the Na abundance is needed to fit the D1 line (although the D1 line is heavily affected by atmospheric absorption). Upon averaging the results obtained from the D1 and D2 lines, we concluded that Na is over--abundant by at least 0.8 dex in V4641 Sgr. The abundances of six elements (Mg, Si, Al, Ti, Cr, and Fe) in V4641 Sgr were estimated by comparing the major absorption features of each element with those seen in the two reference stars. figures 7 to 11 display the results. In these figures, the solid lines show simulated spectra for each star, assuming solar abundances (except for figure 8, where an enhanced abundance of Mg is assumed), taken from \\citet{gre98}. Neutral and singly ionized lines of Mg are shown in figures 7 and 8, respectively. We find that the Mg~{\\sc i} triplet lines can be reproduced with a solar abundance of Mg, while a slight enhancement (0.20 dex) of Mg is suggested from the observed profile of the Mg~{\\sc ii} line at 4481 \\AA. On the other hand, the Mg~{\\sc ii} feature at 4390.5 \\AA~ in V4641 Sgr can be reproduced by assuming a solar Mg abundance. Thus, we conclude that the abundance of Mg in V4641 Sgr coincides with that of the Sun, which differs from the results reported by \\citet{orosz01}, who concluded that this star exhibits an enhanced (by a factor of seven) abundance of Mg. We find that an Al~{\\sc ii} line at 4663 05 \\AA~ is clearly present in both of the two reference stars, and in V4641 Sgr (figure 9). Our spectral analyses of this line suggests that Al may by slightly over--abundant (by +0.2 dex) in V4641 Sgr, although the relatively poor S/N ratio in the blue region does not allow for a reliable analysis of the line. A pair of Si~{\\sc ii} lines near 5050 \\AA~ are compared in figure 10. We conclude that the abundance of Si in V4641 Sgr is close to the solar abundance. In figure 11, we compare Ti~{\\sc ii}, Cr~{\\sc ii}, and Fe~{\\sc ii}~ lines, where all of the absorption features can be reproduced by assuming the same (solar) abundances in the three stars. We analyse several other strong Ti~{\\sc ii}, Cr~{\\sc ii}, and Fe~{\\sc ii}~ lines (noted in table 1) and found that all of these features in V4641 Sgr can be explained using solar abundances of Ti, Cr, and Fe. Our result for Ti is again in contrast with the result given in \\citet{orosz01}, who obtained an over--abundance of Ti (by a factor of 10). They used four Ti~{\\sc ii} lines (at 5129.2 \\AA, 5185.9 \\AA, 5188.7 \\AA, and at 5226.5 \\AA) to obtain the Ti abundance, and pointed out a high abundance of Ti from the two lines (at 5129.2 \\AA~ and 5226.5 \\AA) using data of spectral resolution of about 4 \\AA. We examined all of these lines on our high-resolution data and find that all these lines are very weak in the reference stars, and also in V4641 Sgr, when compared to those Ti~{\\sc ii} lines listed in table 1. We concluded that the four lines used in \\citet{orosz01} are inadequate to be used for the abundance analysis of Ti. Our final derived abundances for 10 observed elements (11 ions) are summarized in table 2, together with their expected errors. We estimate uncertainties in the abundances of each ion introduced by errors in the adopted parameters: 200 K in {\\it $T_{\\rm eff}$}, 0.5 in log $\\it g$, and 0.5 km s${}^{-1}$ in $\\xi$${}_{\\rm t}$. When these errors are combined, we conclude that the our abundance analysis results are reliable within 0.25 dex (table 2). We examined the effect of a difference in spectral resolution on the resulting abundances by a simple test. The original data of both 14 Cyg and V4641 Sgr were degraded to around $R = 8000$, the highest resolution used in \\citet{orosz01}, by convolving with an appropriate Gaussian function. We then repeated abundance analyses using the degraded data for several spectral features such as the Mg~{\\sc ii} line at 4481 \\AA~, the Mg~{\\sc i} triplet lines, and the three Ti~{\\sc ii} lines listed in table 1. Fairly good agreements (within 0.05 dex) were found for 14 Cyg from both high and low resolution data. On the other hand, differences as large as 0.15 dex were found between abundance results obtained from weak and noisy spectral features in the case of V4641 Sgr. We infer that these differences are mainly resulted from the relatively poor SN ratio in the V4641 Sgr data, but not from the difference in the spectral resolution. When the limited S/N ratio of our observation and the high rotational velocity of V4641 Sgr are taken into account, the expected error in the abundances should be increased to around 0.3 dex. ", "conclusions": "We obtained abundances of 10 elements in V4641 Sgr, and found definite over--abundances of only two elements (N and Na). The abundances of the eight other elements in V4641 Sgr have been shown to be the same as those in the two reference stars (solar abundances), except for a possible enhancement of Mg suggested from the Mg~{\\sc ii}~line at 4481 \\AA~ and that of Al. However, when averaged with the result obtained from the Mg~{\\sc i} triplet lines and the Mg~{\\sc ii} line at 4390.5 \\AA, the abundance of Mg is coincident with that in the reference stars within the expected error ([Mg/H] = +0.10 $\\pm$ 0.30). The above conclusions are in contrast to the results noted in \\citet{orosz01} except for N, where they concluded over--abundances of N (1.0 dex), O (0.48 dex), Mg (0.85 dex), and Ti (1.0 dex) in V4641 Sgr when compared to the Sun. We suggest that the primary reason for obtaining discordant abundances for O, Mg, and Ti is the difference in the spectral resolution of the data. \\citet{orosz01} used a much lower resolving power ($\\it R$ ranging from 1200 to 7700) than obtained in the present study ($\\it R$ $\\sim$ 40000). Our results for the abundances of the light elements in the secondary star of V4641 Sgr [definite over--abundances of N and Na, normal (solar) abundances of O, and the $\\alpha$-elements Mg, Si, and Ti] are unique when compared with the results obtained for other X-ray binaries. Abundances obtained in four secondary stars in X-ray binaries are compared in figure 12 [V4641 Sgr (this study), GRO J1655-40 \\citep{isra99}, A0620-00 \\citep{gonz04}, and Cen X4 \\citep{gonz05}]. We note that all four stars show distinct abundance patterns. The $\\alpha$-elements (O, Mg, Si, S, and Ti) are definitely over--abundant in GRO J1655-40, while they appear to be normal in V4641 Sgr. Fe is over--abundant only in Cen X-4. N and Na are over--abundant only in V4641 Sgr. The difference in the abundances of O between V4641 Sgr and GRO J1655-40 is impressive. The observed abundance pattern in V4641 Sgr (enhanced N and Na, and normal $\\alpha$-elements) seems to be different from those of the usual supernova models that predict the enhancement of $\\alpha$-elements (\\cite{podsi}, \\cite{gonz04}, and \\cite{gonz05}). However, these variations of abundance patters can be explained with the variations of the abundances of supernova explosions that are associated with black hole formation \\citep{umeda03}. In order to explain the abundances in V4641 Sgr quantitatively, we calculated the evolution of the star with the initial mass of 40 $M_\\odot$ and the solar metallicity from the main-sequence to collapse as in \\citet{umeda05} and constructed the following evolutionary models for V4641 Sgr. In the close binary system, this 40 $M_\\odot$ primary star underwent a common envelope phase and lost most of its H-rich envelope until it became a He star of mass 15.14 $M_\\odot$. The system also lost its angular momentum and became compact with the orbital period as short as that observed. Figure 13 shows the abundance distribution near the surface of the He star at the onset of collapse. In the He-rich layer, $^{14}$N and $^{23}$Na were enhanced by the CNO-cycle and Ne-Na cycle (proton captures on $^{21}$Ne and $^{22}$Ne) during H-burning. In the deeper He layer, the $^{14}$N abundance was decreased by successive $\\alpha$-capture to produce $^{22}$Ne during weak He shell burning. The 15 $M_\\odot$ He star is massive enough to eventually formed a black hole (BH). We assume that the collapse induced a relatively weak explosion. Generally, an explosion with a smaller energy leads to a larger amount of fall back materials and thus a smaller amount of ejecta (e.g., \\cite{iwamoto05}). In order to reproduce the observations of V4641 Sgr, we assume that the kinetic energy of explosion was as small as $E=6\\times 10^{49}$ ergs. In such a weak explosion, only 0.5 $M_\\odot$ materials above $M_r=14.66M_\\odot$ were ejected. The abundance distribution in the He layer in figure 13 does not change in the explosion. A part of the ejected materials must be captured by the secondary star. The captured (accreted) materials were then mixed with the materials of solar metallicity in the atmosphere of the secondary star. Since the ejecta is relatively N- and Na-rich without any enhancement of $\\alpha$-elements, this could explain the observed abundance pattern of V4641 Sgr. If the accreted material is mixed with 40-times larger amount of secondary star materials, the final abundance pattern would be consistent with the observed abundance pattern of the secondary star (figure 14). Here, most of the heavy elements above Na originated from the materials of V1641 Sgr. Such a partial mixing (e.g., slow rotational mixing) may be realized because the surface temperature of V4641 Sgr is too high for deep convective mixing to occur. We should note that an alternative scenario is possible. If the stellar wind of the He star of $\\sim$ 15 $M_\\odot$ blows at a high enough rate, a part of the N- and Na-rich materials in the He layer would have been blown off and captured by the secondary star. If the energy of supernova explosion was even smaller than $\\sim 6 \\times 10^{49}$, no mass ejection occurred. These processes could lead to the observed abundance pattern of V4641 Sgr. The above 40 $M_\\odot$ model formed a BH of 14 $M_\\odot$. The initially 30 $M_\\odot$ model produces a similar abundance pattern by forming a 7.2 $M_\\odot$ BH. Since the observed BH mass of V4641 Sgr is $\\sim 9.6 M_\\odot$, the progenitor of the BH in V4641 is likely a $\\sim 35 M_\\odot$ star. It is highly uncertain in the current supernova models under what condition the BH formation can induce a supernova explosion and how much explosion energy can be released; it may depend on the rotation of the BH and the progenitor. The case of V4641 Sgr suggests the BH- forming supernova was really {\\sl dark}, because no radioactive $^{56}$Ni was ejected. Such a {\\sl dark} supernova corresponds to the extreme end of the {\\sl faint} supernova branch \\citep{nomoto05}. Another possible scenario is the contamination by rotationally induced mixing in the secondary star, itself. However, the observed rotational velocity ($v \\sin i \\sim 100$~km s$^{-1}$) and estimated mass may be lower than those predicted for the simultaneous enhancements of N and Na, although they strongly depend on the uncertain inclination. \\vskip 5mm We thank Drs. K. Matsumoto and M. Uemura for comments and suggestions. This research was partly supported by Grants-in-Aid from the Ministry of Education, Culture, Sports, Science and Technology (No. 15540236 KS, No. 17540218 MTH, No. 17740163 NI, and No. 17033002 KN). T.C.B. ackowledges partial support from grant AST 04-06784, as well as from grant PHY 02-16783, Physics Frontier Center/Joint Institute for Nuclear Astrophysics (JINA), awarded by the US National Science Foundation." }, "0607/astro-ph0607223_arXiv.txt": { "abstract": "Starting with the first detection of an afterglow from a short-duration hard-spectrum $\\gamma$-ray burst (SHB) by Swift last year, a growing body of evidence has suggested that SHBs are associated with an older and lower-redshift galactic population than long-soft GRBs and, in a few cases, with large ($\\age 10$ kpc) projected offsets from the centers of their putative host galaxies. Here we present observations of the field of GRB\\, 060502B, a SHB detected by Swift and localized by the X-ray Telescope (XRT). We find a massive red galaxy at a redshift of $z=0.287$ at an angular distance of 17.1\\arcsec\\ from our revised XRT position. Using associative and probabilistic arguments we suggest that this galaxy hosted the progenitor of GRB\\, 060502B. If true, this offset would correspond to a physical displacement of $73 \\pm 19$ kpc in projection, about twice the largest offset inferred for any SHB to date and almost an order of magnitude larger than a typical long-soft burst offset. Spectra and modeling of the star-formation history of this possible host show it to have undergone a large ancient starburst. If the progenitor of GRB\\, 060502B was formed in this starburst episode, the time of the GRB explosion since birth is $\\tau \\approx 1.3 \\pm 0.2$ Gyr and the minimum kick velocity of the SHB progenitor is $v_{\\rm kick, min} = 55 \\pm 15$ km s$^{-1}$. ", "introduction": "Since the seminal work of \\citet{kmf+93}, a consensus view has emerged that short-duration hard-spectrum GRBs (SHBs) arise from a separate physical population than long-duration soft-spectrum GRBs (LSBs). The populations are distinguished phenomenologically by an observed bimodality in the GRB duration distribution \\citep{mgi+81,ncdt84} and an apparent corresponding bimodality in spectral hardness. While most LSB progenitors are now believed to be due to the death of massive stars, without a successful detection of an afterglow or a host galaxy the nature of the SHBs remained a mystery until recently. In May 2005, the Swift satellite detected and localized SHB\\, 050509B and, for the first time, found a fading X-ray afterglow \\citep{gso+05}; this was the first SHB localized quickly ($\\ale 10$ s) and accurately ($< 100$ arcsec$^{2}$). Ground-based followup observations led to the discovery of an early-type galaxy at a redshift of $z=0.258$ approximately 10\\arcsec\\ from the X-ray afterglow position \\citep{bpp+06}. A chance association with such a galaxy was deemed unlikely even under conservative assumptions ($P <$ few percent) and stood in stark contrast with the lines-of-sight of LSBs, with which no association of with an early-type was ever made. Both the nature of the burst itself (lacking any supernova signature; \\citealt{hsg+05}) and the location (in the halo of a red galaxy with very little star formation) suggested a progenitor of a very different nature from the purported progenitors of LSBs. In particular, these observations were in close agreement with predictions \\citep{bp95,bsp99,fwh99} for the nature of the environment -- particularly, the offset from host galaxy and the type of the host -- associated with the merger of a degenerate binary (e.g.~\\citealt{npp92}). Further Swift and {\\it HETE-2} detections of SHBs have continued to support this hypothesis, though SHBs are not universally at large offsets and are not always associated with early-type galaxies (see \\citealt{bp06} for a review). SHB\\, 050724 \\citep{bpc+05,pbc+06,gcg+06} and 050813 \\citep{pbc+06}, like 050509B, were found to be in close association with old, red galaxies (see also \\citealt{ltf+06}). SHB\\, 050724 had optical and radio afterglow emission that pinpointed its location to be within its red host, making the association completely unambiguous, though the association of 050813 with any single host remains somewhat tentative. Not all hosts lack active star formation; SHB\\, 050709 \\citep{vlr+05,hwf+05,ffp+05,cmi+06} and 051221A \\citep{sbk+06} both had optical afterglows and were associated with galaxies with evidence for current star formation. However, despite the availability of both X-ray and optical afterglow locations, no nearby host has successfully been identified for either SHB\\, 060121 or SHB\\, 060313 (although see \\citealt{hjo+06}). In this article we examine the field of Swift SHB\\,060502B \\citep{tbb+06} and, in \\S \\ref{sec:0502b}, we present imaging and spectroscopy of a bright red galaxy near the X-ray afterglow position. In \\S \\ref{sec:host} we present evidence that supports the notion that the progenitor of SHB\\,060502B was born in that galaxy. Accepting this connection we discuss the implications of the nature of the host and offset for the progenitors of SHBs. Though the association of this galaxy with the GRB is the most tenuous of SHB--host associations thus far proposed, we conclude in \\S \\ref{sec:conc} that there are both observational and theoretical motivations to accept this association for this and (similarly configured) future SHBs. Some of our work on this GRB was given preliminarily in \\citet{bpk+06}; our results presented herein are consistent with, but supersede that reference. Throughout this {\\it Letter} we assume $H_0 = 71~h_{71}$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m = 0.3$, $\\Omega_\\Lambda = 0.7$. ", "conclusions": "\\label{sec:conc} The large offset from what we have argued is a plausible host, if true, holds important ramifications for both the sort of viable progenitors and where they are born. First, the large offset would seem to be at odds with the hypothesis of a degenerate binary origin where systematic kicks are small (such as in globular clusters [GCs]; \\citealt{gpm06}). While the expected number density of GCs at 75 kpc is exceedingly small \\citepeg{bbb+05}, there certainly could be a GC at $z = 0.287$ in the XRT error circle (it would appear as faint red point source with magnitude $R \\approx 29$, in principle observable with HST imaging). Alternatively, $G^*$ could have undergone a major merger leaving behind a progenitor system at the XRT position. Second, if the progenitor was created during what appears to be the last starburst in the putative host, then the time since zero age main sequence would be $\\tau \\approx 1.3 \\pm 0.2$ Gyr (90\\% confidence). At the inferred offset, this would imply a minimum systemic kick velocity of $v_{\\rm kick, min} = r/\\tau \\approx 55 \\pm 14$ km s$^{-1}$. Such a kick velocity is comparable to the models for degenerate binaries \\citep{fwh99} and observations of Galactic double NS systems \\citep{dpp05}. The kick could have been significantly larger, implying that the progenitor orbited about the host before the GRB event. Indeed with the inferred stellar mass $7 \\times 10^{11} M_{\\sun}\\, h_{71}^{-2}$ of the putative host, unless the progenitor was born on the outskirts of the host gravitational potential, the true $v_{\\rm kick}$ would have to have been comparable to or greater than dispersion velocity of the host. If the progenitor remains gravitationally bound to $G^*$ then the systemic orbital velocity of progenitor spends most time near zero velocity, with its initial kinetic energy stored as gravitational potential. That is, we nominally expect an orbiting progenitor to produce a burst near the maximal distance from its host. Indeed if all the energy is stored as potential, then for SHB 060502B, the gravitational potential of the progenitor system is $$ \\epsilon_{\\rm pot} = \\frac{G\\, M_{G^*}}{d} \\approx 6 \\times 10^{14} ~~{\\rm erg~gm}^{-1} \\left(\\frac{M_{G^*}}{10^{12} M_\\odot}\\right) \\left(\\frac{d}{73~{\\rm kpc}}\\right)^{-1}. $$ Upon birth, the kinetic energy per unit mass imparted to the progenitor must have been: $$ \\epsilon_{\\rm kin} = \\frac{1}{2} v_{\\rm kick}^2 \\approx 1 \\times 10^{14} ~~{\\rm erg~gm}^{-1} \\left(\\frac{v_{\\rm kick}}{160\\, {\\rm km~s}^{-1}}\\right)^2. $$ Here we have taken the nominal velocity of the kick as the geometric mean of the dispersion velocity ($\\approx 460$ km s$^{-1}$) and $v_{\\rm kick, min} $; that is, we assume $v_{\\rm kick} = 160$ km s$^{-1}$. That $\\epsilon_{\\rm kin}$ is even within an order of magnitude of $\\epsilon_{\\rm pot}$ is either a remarkable coincidence\\footnotemark\\footnotetext{If the progenitor is a double neutron star with $M_{\\rm prog} = 2.8 M_\\odot$, then the total respective energies are $E_{\\rm pot} = M_{\\rm prog} \\epsilon_{\\rm pot} = 3 \\times 10^{48}$ erg and $E_{\\rm kin} = M_{\\rm prog} \\epsilon_{\\rm kin} = 7 \\times 10^{47}$ erg. We can think of no progenitor model to explain why $E_{\\rm kin}$ and $E_{\\rm pot}$ is also comparable to $E_{\\rm iso, \\gamma}$.} or, we suggest, indicative of support on dynamical grounds for the ejection hypothesis. We end by acknowledging the difficulty of confirming, beyond reasonable doubt, our hypothesis that $G^*$ hosted the birth of the progenitor of SHB\\,060205b. The progenitors of most LSBs, owing to their connection with massive stars, allowed for unambiguous associations with putative hosts --- most with probability of chance alignment $P \\ale 10^{-3}$ \\citep{bkd02}. With SHB\\,060502b we have estimated under mildly conservative assumptions (ie.\\ without regard to host type) that the chance of a spurious assignment with $G^*$ is $P \\ale$ 10\\%. The $\\epsilon_{\\rm kin} \\approx \\epsilon_{\\rm pot}$ argument and the similarity with GRB\\, 050509b likely strengthen this particular association. Yet with SHBs, especially if the majority of progenitors are long-lived high-velocity degenerate mergers, the community must accept that an appreciable fraction of host assignments relative to LSBs will be spurious \\citep{btw97}. Of course absorption line redshifts of SHB afterglows, one of the remaining observational goals of the field, will help to significantly cull the number density of viable hosts on the sky." }, "0607/astro-ph0607479_arXiv.txt": { "abstract": "We show that inverse Compton scattering of cosmic-microwave-background and starlight photons by cosmic-ray electrons in the interstellar and intergalactic space explains well the spectrum and intensity of the diffuse gamma-ray background radiation (GBR), which was measured by EGRET aboard the Compton Gamma Ray Observatory (CGRO) in directions away from the Galactic disk and centre. The Gamma Ray Large Area Space Telescope (GLAST) will be able to separate the Galactic foreground from the extragalactic gamma-rays, and to provide stringent tests of the theory. ", "introduction": " ", "conclusions": "" }, "0607/astro-ph0607153_arXiv.txt": { "abstract": "We examine statistical isotropy of large scale anisotropies of the Internal Linear Combination (ILC) map, based on three year WMAP data. Our study reveals no significant deviation from statistical isotropy on large angular scales of 3-year ILC map. Comparing statistical isotropy of 3-year ILC map and 1-year ILC map, we find a significant improvement in 3-year ILC map which can be due to the gain model, improved ILC map processing and foreground minimization. ", "introduction": "The Cosmic Microwave Background (CMB) anisotropy has been shown to be a very powerful observational probe of cosmology. Detailed measurements of the anisotropies in the CMB can provide a wealth of information about the global properties, constituents and history of the Universe. In standard cosmology, the CMB anisotropy is expected to be statistically isotropic, {\\it i.e.}, statistical expectation values of the temperature fluctuations (and in particular the angular correlation function) are preserved under rotations of the sky. This property of CMB anisotropy has been under scrutiny after the release of the first year of WMAP data \\cite{erik04a, Copi:2003kt, Schwarz:2004gk, Hansen:2004vq, angelwmap, Land:2004bs,Land:2005ad, Land:2005dq, Land:2005jq,Land:2005cg, Bielewicz:2004en, Bielewicz:2005zu,Copi:2005ff, Copi:2006tu, Naselsky:2004gm, Prunet:2004zy,Gluck:2005td, Stannard:2004yp, Bernui:2005pz, Bernui:2006ft, SadeghMovahed:2006em, Freeman:2005nx, us_apj, Chen:2005ev}. We use a method based on bipolar expansion of the two point correlation function which is an improved and enhanced follow up on our previous work on first year WMAP data \\cite{us_apj, us_bigpaper}. This method is shown to be sensitive to structures and patterns in the underlying two-point correlation function. We apply our method to the improved Internal Linear Combination (ILC) map \\cite{Hinshaw:2006ia}, based on three year WMAP data\\footnote{This map is available on LAMBDA as part of the three-year data release \\cite{Hinshaw:2006ia,Jarosik:2006ib,Page:2006hz,Spergel:2006hy}.}. We choose the ILC map for testing statistical isotropy (SI) of the CMB anisotropy for the following reasons \\begin{enumerate} \\item{The ILC is a full-sky map and hence is easier to work with. Masking the sky results in violation of statistical isotropy. An originally SI CMB anisotropy map deviates from SI after masking \\cite{us_bigpaper}. } \\item{Residuals from Galactic removal errors in the three-year ILC map are estimated to be less than 5 $\\mu K$ on angular scales greater than $\\sim10\\deg$ \\cite{Hinshaw:2006ia}. Hence at low-$l$, multipoles are not significantly affected by foregrounds. In addition, it is interesting to examine the above statement by testing the statistical isotropy of the ILC map.} \\item{On large scales, the three-year ILC map is believed to provide a reliable estimate of the CMB signal, with negligible instrument noise, over the full sky \\cite{Hinshaw:2006ia}. } \\end{enumerate} These properties of the ILC map allow us to study the cosmological signal on large scales. In addition, there are theoretical motivations for hunting for SI violation on large scales of CMB anisotropy. Topologically compact spaces~\\cite{ell71, Lachieze-Rey:1995kj, lev02, linde} and anisotropic cosmological models~\\cite{ellis&maccallum, collins&hawking1, collins&hawking2, doroshke, barrow, Jaffe:2006fh, us_bianchi} are examples of this. Both observational artifacts and the above theoretical models cause a departure from statistical isotropy and it has been shown that our method is a useful tool to find out these deviations (see {\\it e.g.} \\cite{us_bianchi, us_foregrounds, us_prl}). The rest of this paper is organized as follows: Section \\ref{characterization} is a brief introduction to temperature anisotropy of CMB. Section \\ref{SIsection} describes the formulation of statistical isotropy in general. Section \\ref{estimators} is a description of estimators we use to look for deviations from statistical isotropy. We present the application of our method on the WMAP data in Section \\ref{data}. Section \\ref{discussion} contains discussion on the cosmological implications of our null detection of deviations from statistical isotropy on large angular scales in 3-year ILC map of WMAP data, and in Section \\ref{summary} we summarize our results. ", "conclusions": "\\label{discussion} The null results of search for departure from statistical isotropy has implications for the observation and data analysis techniques used to create the CMB anisotropy maps. Observational artifacts such as non-circular beam, inhomogeneous noise correlation, residual striping patterns, and residuals from foregrounds are potential sources of SI breakdown. Our null results confirm that these artifacts do not significantly contribute to large scale anisotropies of 3-year ILC map. We have also quantified the differences between 1-year and 3-year ILC maps. It is shown that 3-year ILC map is ``cleaner'' than 1-year ILC map at $l\\le 25$. This can be due to the gain model and improved ILC map processing and foreground minimization. It has also been observed that at large $l$ deviations from statistical isotropy occur which we think is because of residuals from foregrounds. However we limit ourselves to the low-$l$ limit because in addition to observational artifacts, there are theoretical motivations for hunting for SI violation on large scales of CMB anisotropy. Topologically compact spaces~\\cite{ell71, Lachieze-Rey:1995kj, lev02, linde} and anisotropic cosmological models~\\cite{ellis&maccallum, collins&hawking1, collins&hawking2, doroshke, barrow, Jaffe:2006fh, us_bianchi} are examples of this. Each of these models will cause departures from statistical isotropy in CMB anisotropy maps. And a null detection of departure from statistical isotropy at low $l$ in the WMAP data can be used to put constraints on these models. Our measure is sensitive to axial asymmetries in the two point correlation of the temperature anisotropy \\cite{us_prl}. And this is even more significant now because the new measure of reduced bipolar coefficients does retain directional information. Our analysis doesn't show a significant detection of an ``axis of evil'' in the WMAP data. We have redone our analysis on ILC map filtered with a low-pass filter that only keeps $l=2,3,4$ to search for a preferred direction at low multipoles. We have not been able to detect any significant deviation from statistical isotropy using various filters. We could not test the effect of alignment of low multipoles on statistical isotropy because we had no theory or model to explain them. Validity of statistical isotropy at large angular scales can put tight constraints on anisotropic mechanisms that are candidates of explaining the low quadrupole of the WMAP and COBE data. It is worth noticing that our method can be extended to polarization maps of CMB anisotropy. Analysis of statistical isotropy of full-sky polarization maps of WMAP are currently under progress and will be reported in a separate publication." }, "0607/astro-ph0607365_arXiv.txt": { "abstract": "We report deep K-band (18-27\\,GHz) observations with the 100-m Green Bank Telescope of HCN$(1-0)$ line emission towards the two submillimeter-selected galaxies (SMGs) SMM\\,J02399$-$0136 ($z=2.81$) and SMM\\,J16359$+$6612 ($z=2.52$). For both sources we have obtained spectra with channel-to-channel rms noise of $\\sigma \\le 0.5$\\,mJy, resulting in velocity-integrated line fluxes better than $\\ls 0.1$\\,Jy\\,km\\,s$^{-1}$, although we do not detect either source. Such sensitive observations -- aided by gravitational lensing of the sources -- permit us to put upper limits of $L'_{\\mbox{\\tiny{HCN}}}\\ls 2\\times 10^{10}\\,$K\\,km\\,s$^{-1}\\,$pc$^2$ on the intrinsic HCN$(1-0)$ line luminosities of the two SMGs. The far-infrared (FIR) luminosities for all three SMGs with sensitive HCN$(1-0)$ observations to date are found to be consistent with the tight FIR-HCN luminosity correlation observed in Galactic molecular clouds, quiescent spirals and (ultra) luminous infrared galaxies in the local Universe. Thus, the observed HCN luminosities remain in accordance with what is expected from the universal star formation efficiency per {\\it dense} molecular gas mass implied by the aforementioned correlation, and more sensitive observations with today's large aperture radio telescopes hold the promise of detecting HCN$(1-0)$ emission in similar objects in the distant Universe. ", "introduction": "The search for molecular lines at high redshifts ($z\\gs 1$) offers one of the richest and most exciting avenues to follow in the study of galaxy formation and evolution (Solomon \\& Vanden Bout 2005). The detections of rotational transitions of CO in extremely distant quasars (Carilli et al.\\ 2002; Walter et al.\\ 2003), high-z radio galaxies (Papadopoulos et al.\\ 2000; De Breuck et al.\\ 2005) and submillimeter (submm) selected galaxies (Frayer et al.\\ 1998, 1999; Neri et al.\\ 2003), have established that the most luminous and extreme objects in the early Universe harbour vast amounts of molecular gas ($\\sim 10^{10-11}\\,\\Msolar$ -- Neri et al.\\ 2003; Greve et al.\\ 2005) and have large dynamical masses ($\\sim 10^{11}\\,\\Msolar$ -- Genzel et al.\\ 2003; Tacconi et al.\\ 2006). Modelling of the bulk physical conditions of the molecular gas in high-$z$ QSOs have even been possible in cases where several CO lines have been detected (Wei\\ss\\ et al.\\ 2005). The success of CO observations is primarily due to its much larger abundance compared to other species as well as its easily excitable low-$J$ (CO $J+1\\longrightarrow J, J<3$) rotational lines, which makes it an excellent tracer of the total amount of {\\it diffuse} ($n(\\mbox{H}_2) \\sim 10^{2-3}$\\,cm$^{-3}$) molecular gas. However, the diffuse gas only serves as a reservoir of {\\it potential} fuel available for star formation, and is not actively involved in forming stars. In our own Galaxy, the sites of active star formation coincide with the dense cores of giant molecular clouds (GMCs). This dense gas phase ($n(\\mbox{H}_2) \\gs 10^{4}$\\,cm$^{-3}$) is best traced using molecules with high critical densities such as HCN. This basic picture also appears to hold in nearby galaxies, as suggested by HCN$(1-0)$ surveys of local ($z\\ls 0.1$) Luminous Infrared Galaxies (LIRGs -- $L_{\\mbox{\\tiny{IR}}}\\sim 10^{11}\\,\\Lsolar$) and Ultra Luminous Infrared Galaxies (ULIRGs -- $L_{\\mbox{\\tiny{IR}}}\\sim 10^{12}\\,\\Lsolar$), which found a remarkably tight correlation between IR luminosity and HCN line luminosity (Solomon et al.\\ 1992; Gao \\& Solomon 2004a,b). Recently, this relation was shown to extend all the way down to single GMCs, thus extending over 7-8 orders of magnitude in IR luminosity (Wu et al.\\ 2005). This was interpreted as suggesting that the true star formation efficiency (i.e.\\ the star formation rate per unit mass of {\\it dense} gas) is the same in these widely different systems. Attempts have been made at extending the correlation between IR and HCN luminosity to primordial galaxies at high redshifts (Barvainis et al.\\ 1997; Solomon et al.\\ 2003; Isaak et al.\\ 2004; Carilli et al.\\ 2005), resulting in four objects detected in HCN out of 8 objects observed. So far the results have been consistent with the local correlation, however, the small number of sources precludes a statistically robust conclusion as to whether the relation also holds at high redshift. Furthermore, the majority of high-$z$ objects observed have been extremely luminous QSOs and HzRGs dominated by Active Galactic Nuclei (AGN). The detection of HCN$(1-0)$ in the $z=2.286$ QSO F\\,10214$+$4724 (Vanden Bout, Solomon \\& Maddalena 2004), which marked the first high-$z$ detection of HCN using the Green Bank Telescope (GBT), motivated us to test the capability of the GBT to probe the dense gas in a population of high-$z$ galaxies less extreme, i.e.\\ less AGN-dominated than F\\,10214$+$4724, but more representative of starburst galaxies in the early Universe. The goal of our study was to observe two of the most prominent examples of submillimeter-selected galaxies (see Blain et al.\\ 2002 for a review) -- a population likely to constitute the progenitors of today's massive spheroids, and thus fundamental to our understanding of galaxy formation and evolution. Throughout this paper we adopt a flat cosmology, with $\\Omega_m=0.27$, $\\Omega_\\Lambda=0.73$ and $H_0=71$\\,km\\,${\\mbox{s}^{-1}}$\\,Mpc$^{-1}$ (Spergel et al.\\ 2003). ", "conclusions": "\\label{section:results} The final HCN$(1-0)$ spectra of J02399 and J16359, reduced using the two independent methods described in the previous section, are shown in Figure \\ref{figure:hcn-spectra}. The spectra have been smoothed to $50$\\,km\\,s$^{-1}$ bins. Both methods fail to get completely rid of residual baseline wiggles in the final spectra. However, the spectra reduced with GETNOD appear to be slightly less noisy than the spectra reduced with the second reduction method. The channel-to-channel rms noise of the spectra reduced using the GETNOD routine are 0.1 and 0.3\\,mJy, respectively. This is somewhat higher than the theoretical noise estimates of 0.07 and 0.12\\,mJy, calculated for a typical system temperature $T_{sys}=38\\,$K and integration times corresponding to the ones of J02399 and J16359, respectively. Thus, it would seem that although the noise does integrate down as $t^{-1/2}$, the residual baseline wiggles prevents us from reaching the theoretical noise limit. No emission is detected significantly above the noise at $V_{LSR}=0$\\,km\\,s$^{-1}$ (or any other $V_{LSR}$), which corresponds to the CO redshift and thus the expected position of the HCN$(1-0)$ line. However, the sensitivity of the observations allow us to put upper limits on the HCN$(1-0)$ line luminosity of these two SMGs. In doing so we use the upper line flux limits derived from the spectra produced by GETNOD, as it results in the lowest noise spectra. The 3-$\\sigma$ upper limits are calculated following Seaquist, Ivison \\& Hall (1995) \\begin{equation} S_{\\mbox{\\tiny{HCN(1-0)}}}\\Delta V \\le 3 \\sigma (\\delta v \\Delta v_{\\mbox{\\tiny{fwhm}}} ) ^{1/2}, \\label{equation:upper-limits} \\end{equation} where $\\sigma$ is the channel-to-channel rms noise, $\\delta v$ the velocity resolution and $\\Delta v_{\\mbox{\\tiny{fwhm}}}$ the line width. The spectra were binned to a velocity resolution of 50\\,km\\,s$^{-1}$. We set the HCN$(1-0)$ line widths equal to that of the CO lines -- see Greve et al.\\ (2005). This is probably a conservative estimate since in local ULIRGs, the HCN line widths are rarely larger than those of CO (Solomon et al.\\ 1992; Gao \\& Solomon 2004a). The resulting upper line flux limits are 0.08 and 0.13\\,Jy\\,km\\,s$^{-1}$ for J02399 and J16359, respectively. From these flux limits we derive upper limits on the apparent HCN$(1-0)$ line luminosities of $L'_{\\mbox{\\tiny{HCN(1-0)}}}\\le 5.0\\times 10^{10}$ and $\\le 6.6\\times 10^{10}$\\,K\\,km\\,s$^{-1}$pc$^2$ for J02399 and J16359, respectively. Correcting for gravitational lensing we find intrinsic line luminosities of $L'_{\\mbox{\\tiny{HCN(1-0)}}}\\le 2.0\\times 10^{10}$ and $\\le 0.3 \\times 10^{10}$\\,K\\,km\\,s$^{-1}$pc$^2$. These upper limits on the HCN line luminosity are similar to those achieved towards high-$z$ QSOs using the Very Large Array (e.g.\\ Isaak et al.\\ 2004; Carilli et al.\\ 2005). Table \\ref{table:results} lists our findings along with all high-$z$ HCN observations published in the literature at the time of writing. It is seen that of 10 sources observed to date only 4 have been detected -- the remainder being upper limits. In order to determine the star formation efficiency per dense gas mass, and the dense gas fraction, as gauged by the $L_{\\mbox{\\tiny{FIR}}}/L'_{\\mbox{\\tiny{HCN}}}$ and $L'_{\\mbox{\\tiny{HCN}}}/L'_{\\mbox{\\tiny{CO}}}$ ratio, respectively (Gao \\& Solomon 2004a,b), we need accurate estimates of the FIR and CO(1-0) luminosities.\\\\ \\indent The FIR luminosities of SMGs are generally difficult to estimate accurately, largely due to the poor sampling of their FIR/submm/radio spectral energy distributions. Furthermore, in our case the situation is complicated by the fact that most, if not all, of the sources in Table \\ref{table:results} contain AGN, which may be at least partly responsible for heating the dust and thus powering the FIR emission. Unfortunately, FIR/submm data from the literature do not allow for a detailed modelling of a hot ($\\gs 100$\\,K) dust component (heated by the AGN), and we are therefore unable to correct for any AGN contamination. The one exception is the $z=3.9$ BAL quasar, APM\\,08279$+$5255, where the separate AGN vs.\\ starburst contributions to the FIR luminosity have been determined by Rowan-Robinson (2000) who find an apparent FIR luminosity of $L_{\\mbox{\\tiny{FIR}}}\\simeq 1.0\\times 10^{14}\\,\\Lsolar$ for the starburst (corrected to the cosmology adopted here). In the case of SMGs, extremely deep X-ray observations strongly suggest that while every SMG probably harbours an AGN, it is the starburst which in almost all cases powers the bulk (70-90 percent) of the FIR luminosity (Alexander et al.\\ 2005), and any AGN contamination would therefore not dramatically affect our conclusions. While optical spectroscopy of J16359 shows no evidence to suggest the presence of strong nuclear activity in this source (Kneib et al.\\ 2004), this is not the case in J02399, which appears to be a type-2 QSO judging from its X-ray emission and optical spectrum properties (Ivison et al.\\ 1998; Vernet \\& Cimatti 2001). Thus, J02399 could be a rare case where the AGN contribution is substantial. However, given our inability to correct its FIR luminosity for AGN contamination, we estimate the FIR luminosities of both sources simply by fitting a modified black body with a fixed $\\beta =1.5$ to their rest-frame SEDs and integrating it over the wavelength range $40-120\\mu$m, see Table \\ref{table:results}. This is consistent with the way in which the FIR-luminosities of the objects observed by Carilli et al.\\ (2005) were calculated. The FIR luminosities for the remaining objects in Table \\ref{table:results} were taken from Carilli et al.\\ (2005) and converted to the cosmology adopted in this paper.\\\\ \\indent The CO luminosities were taken from the compilation of high-$z$ CO detections in Greve et al.\\ (2005). These are mostly high-$J$ CO detections ($J=3-2$ or $4-3$), and in order to derive the CO$(1-0)$ line luminosity, we assume optically thick, thermalised CO line ratios, i.e.\\ $(3-2)/(1-0)$ and $(4-3)/(3-2) \\sim 1$, see Table \\ref{table:results}. It should be noted, however, that studies on the ISM in local starburst nuclei yield CO $(3-2)/(1-0) \\sim 0.64$ (Devereux et al.\\ 1994), while first results on ULIRGs yield a CO $(6-5)/(4-3)\\sim0.6$ ratio but with CO transitions of $J=4-3$ and higher tracing a different H$_2$ gas phase than the lower three (Papadopoulos, Isaak \\& van der Werf 2006). These results along with known examples of very low global (high$-J$)/(low$-J$) CO ratios in high-$z$ starbursts (e.g.\\ Papadopoulos \\& Ivison 2002; Greve et al.\\ 2003; Hainline et al. 2006) suggest that we may underestimate the SMG CO$(1-0)$ luminosities when we assume line ratios of unity, especially if observed CO $J+1\\longrightarrow J$, $J+1\\ge 4$ lines are used.\\\\ \\indent In Figure \\ref{figure:lhcn-lfir} a) -- c) we plot $L_{\\mbox{\\tiny{FIR}}}$ against $L'_{\\mbox{\\tiny{HCN}}}$, $L'_{\\mbox{\\tiny{HCN}}}$ against $L'_{\\mbox{\\tiny{CO}}}$, and $L_{\\mbox{\\tiny{FIR}}}$ against $L'_{\\mbox{\\tiny{HCN}}}/L'_{\\mbox{\\tiny{CO}}}$, respectively, for our two sources along with all other existing high-$z$ HCN observations to date. For comparison we have also plotted the sample of local (U)LIRGs (open circles) observed by Gao \\& Solomon (2004a,b), which defines the local $L_{\\mbox{\\tiny{IR}}}$-$L'_{\\mbox{\\tiny{HCN}}}$ relation: $L_{\\mbox{\\tiny{IR}}}/L'_{\\mbox{\\tiny{HCN}}} = 900\\,\\Lsolar\\,$(K\\,km\\,s$^{-1}$\\,pc$^2$)$^{-1}$. Since we are interested in the correlations with FIR-luminosity, we carefully fitted modified black-body SEDs to the Gao \\& Solomon sample and derived FIR-luminosities by integrating the SEDs from $40-120\\mu$m. To this end we used all available FIR/submm data in the literature for each source. The resulting local FIR-HCN correlation is given by $\\log L_{\\mbox{\\tiny{FIR}}} = (0.97\\pm 0.07) \\log L'_{\\mbox{\\tiny{HCN}}} + (2.9\\pm 0.5)$, and is shown as the solid line in Figure \\ref{figure:lhcn-lfir} a). The slope of this correlation is consistent with unity, suggesting that the local FIR-HCN correlation is linear (Carilli et al.\\ 2005). In the following we shall use this correlation, along with its 1-$\\sigma$ limits allowed by the fitting errors (dotted lines in Figure \\ref{figure:lhcn-lfir} a)), to discuss whether the SMGs and other high-$z$ sources follow the local FIR-HCN correlation. Secondly, motivated by what is seen in local (U)LIRGs, we shall impose a minimum dense gas fraction on the high-$z$ population and from that derive a lower limit on their HCN-luminosity. Finally, we assume the SMGs follow the local FIR-HCN correlation (within 1-$\\sigma$) and explore what lower limits can be put on their dense gas fraction. With only four HCN detections to date and little more than a handful of upper limits, including the ones presented here, it is difficult to determine whether the local $L_{\\mbox{\\tiny{FIR}}} - L'_{\\mbox{\\tiny{HCN}}}$ relation extends to higher redshifts and luminosities. It is interesting to note in Figure \\ref{figure:lhcn-lfir} a), however, that while the three SMGs with HCN observations (our two sources plus SMM\\,J14011$+$0252) all are consistent within $1$-$\\sigma$ of the local relation extrapolated to higher luminosities, 5/7 of the QSOs observed in HCN fall above the $1$-$\\sigma$ envelope. This may suggest that QSOs have higher FIR luminosities for a fixed HCN luminosity than the SMGs. Although, this is not statistically significant, it is consistent with the expected picture in which the AGN in QSOs contribute a higher fraction to the FIR luminosity than in SMGs. Correcting for the AGN contribution in QSOs would lower them from where they are currently plotted, bringing them more in line with the local $L_{\\mbox{\\tiny{FIR}}} - L'_{\\mbox{\\tiny{HCN}}}$ correlation.\\\\ \\indent Alternatively, the data could in principle also be interpreted as a steepening of the slope of the correlation at higher luminosities, meaning that the AGN contribution in both QSOs and SMGs is higher than in local (U)LIRGs -- a scenario which would be in line with the suggested increase in AGN dominance with FIR luminosity in local ULIRGs (Veilleux et al.\\ 1999; Tran et al.\\ 2001). Although, X-ray constraints suggest that AGN contribute $\\ls 30$ percent of to the FIR luminosity of SMGs, they do not allows us to completely rule out the above scenario since it is possible many SMGs harbour obscured, Compton-thick AGN (Alexander et al.\\ 2005). However, the fact that the star formation rates for SMGs derived from reddening-corrected H$\\alpha$ luminosities agree with those derive from their FIR luminosities, suggest that AGN are unlikely to contribute signficantly to the FIR. We see from Figure \\ref{figure:lhcn-lfir} b) that all the high-$z$ objects lie above the $L'_{\\mbox{\\tiny{HCN}}} - L'_{\\mbox{\\tiny{CO}}}$ correlation fitted to local galaxies with FIR luminosities less than $10^{11}\\,\\Lsolar$ and with a fixed slope of unity. This is consistent with the the steepening of the $L'_{\\mbox{\\tiny{HCN}}} - L'_{\\mbox{\\tiny{CO}}}$ correlation at higher FIR luminosities seen within the local (U)LIRGs sample itself as reported by Gao \\& Solomon (2004b). They argued that the dense molecular gas fraction, as gauged by the $L'_{\\mbox{\\tiny{HCN}}}/L'_{\\mbox{\\tiny{CO}}}$ ratio, is a powerful indicator of starburst dominated systems. Also, locally they found that all galaxies with $L'_{\\mbox{\\tiny{HCN}}}/L'_{\\mbox{\\tiny{CO}}} \\ge 0.06$ have $L_{\\mbox{\\tiny{IR}}} \\gs 10^{11}\\,\\Lsolar$. In Figure \\ref{figure:lhcn-lfir} c) we have plotted our estimates of $L_{\\mbox{\\tiny{FIR}}}$ for the local (U)LIRG sample against their $L'_{\\mbox{\\tiny{HCN}}}/L'_{\\mbox{\\tiny{CO}}}$ ratios, and find that apart from two sources, the same criterion can be applied to the FIR luminosity, i.e.\\ sources with $L'_{\\mbox{\\tiny{HCN}}}/L'_{\\mbox{\\tiny{CO}}} \\ge 0.06$ have $L_{\\mbox{\\tiny{FIR}}} \\gs 10^{11}\\,\\Lsolar$. It is also seen that all the high-$z$ sources observed so far -- all of which have been very luminous -- obey the same criterium. Given their large FIR luminosities, it therefore seems reasonable to assume that the SMGs too have high dense gas fractions, i.e.\\ $L'_{\\mbox{\\tiny{HCN}}}/L'_{\\mbox{\\tiny{CO}}}\\ge 0.06$. In that case, one can put a {\\it lower} limit on their HCN luminosity and investigate where they would lie in the FIR-HCN diagram. The range of allowed HCN luminosities derived in this way for each SMG (and other high-$z$ objects) is shown as horizontal lines in Figure \\ref{figure:lhcn-lfir}a. It is seen that while J02399 and J16359, and the other high-$z$ objects, could lie above the upper dotted line and still have a high dense gas fraction ($L'_{\\mbox{\\tiny{HCN}}}/L'_{\\mbox{\\tiny{CO}}}\\ge 0.06$), this is not the case for J14011. Thus, if our above assumption is correct, we can conclude that at least one of the SMGs (J14011), and possibly more, is consistent with the local FIR-HCN correlation, and that this SMG should be detectable in HCN if the sensitivity is improved by a factor $\\times 2$. APM\\,08279$+$5255 has the highest measured dense gas fraction, based on detections, in the high-$z$ sample, even though it falls above the local $L_{\\mbox{\\tiny{FIR}}} - L'_{\\mbox{\\tiny{HCN}}}$ relation in Figure \\ref{figure:lhcn-lfir} a). Naively, this would indicate that the ISM in this system is dominated by dense gas, and that star formation contributes substantially to its FIR-luminosity. However, caution should be taken since we have calculated the HCN/CO ratio using the nuclear CO line luminosity of APM\\,08279 (Papadopoulos et al.\\ 2001). Adopting the global CO line luminosity, which includes emission extended on 10-kpc scales (Papadopoulos et al.\\ 2001), i.e.\\ scales comparable to those probed by the HCN observations ({\\sc fwhm}~$\\simeq 7\\arcsecs$ -- Wagg et al.\\ 2005), results in a HCN/CO ratio almost an order of magnitude lower. Similarly, we see that the upper limits on the dense molecular gas fraction in J02399 and J16359 are rather high, yet both sources lie close to the $L_{\\mbox{\\tiny{FIR}}} - L'_{\\mbox{\\tiny{HCN}}}$ correlation in Figure \\ref{figure:lhcn-lfir} a). What would the dense gas fraction of the SMGs be if they were required to lie between the upper $1$-$\\sigma$ line of the local FIR-HCN correlation and their observed HCN limits (for unchanged FIR luminosities)? The answer is given in Figure \\ref{figure:lhcn-lfir} c), where the permitted range in HCN/CO luminosity ratios are marked as horizontal lines. It is seen that J14011 has a dense gas fraction well within the range of most ULIRGs, and could have $L'_{\\mbox{\\tiny{HCN}}}/L'_{\\mbox{\\tiny{CO}}} \\ls 0.06$ and still be consistent with the FIR-HCN correlation. For the other two SMGs we see that even in the most conservative case, where the SMGs are just barely consistent with the $1$-$\\sigma$ line, they would still have some of the highest dense molecular gas fractions observed, comparable to the most extreme starburst dominated ULIRGs seen locally. Thus we conclude that if the SMGs do in fact follow the FIR-HCN correlation (or are within $1$-$\\sigma$ of it), then J02399 and J16359 are likely to have higher dense gas fractions than J14011. Note, however, that if an AGN contributes significantly to the FIR luminosity, the above argument would yield a higher HCN luminosity, and thus higher HCN/CO ratios than its true value, but we emphasize that the energy output at FIR/submm wavelengths from AGN in SMGs is found to be negligble compared to the starburst (Alexander et al.\\ 2005)." }, "0607/astro-ph0607403_arXiv.txt": { "abstract": "We solve the problem of coherent Mikheyev-Smirnov-Wolfenstein (MSW) resonant active-to-sterile neutrino flavor conversion driven by an initial lepton number in the early universe. We find incomplete destruction of lepton number in this process and a sterile neutrino energy distribution with a distinctive cusp and high energy tail. These features imply alteration of the non-zero lepton number primordial nucleosynthesis paradigm when there exist sterile neutrinos with rest masses $m_{\\rm s} \\sim 1\\,{\\rm eV}$. This could result in better light element probes of (constraints on) these particles. ", "introduction": " ", "conclusions": "" }, "0607/astro-ph0607129_arXiv.txt": { "abstract": "Diffuse interstellar bands (DIBs) still await an explanation. One expects that some progress in this field will be possible when all the known DIBs are divided into families in such a way that only one carrier is responsible for all bands belonging to the given family. Analyzing high resolution optical spectra of reddened stars we try to find out spectroscopic families for two prominent DIBs, at 5780 and 5797 angstroms. Among the DIBs, observed in the spectral range from 5590 to 6830 angstroms, we have found 8 candidates to belong to 5780 spectroscopic family and the other 12 DIBs candidating to family of 5797 structure. ", "introduction": "The DIBs are absorption features which are generated in the interstellar medium by still unidentified set of carriers. They are found in the visual and near infrared spectra, between 4000 and 13500 angstroms. The discovery of the first DIBs in stellar spectra dates back to the pioneering years of stellar spectroscopy. The original report on the discovery of two spectral features, centered near 5780 and 5797 angstroms, in spectra of some spectroscopic binaries was published in 1922 by Heger. The extended review paper presenting DIB problematics was published by Herbig (1995). To date more than 300 DIBs are detected and the number is still increasing (see e.g. Galazutdinov at all., 2000). None of them has been identified! The identification of the carriers of DIBs is one of the most difficult challenges for spectroscopists. The solution of the mystery of the carriers of DIBs is expected from interdisciplinary spectroscopic collaboration between molecular physicists, molecular chemists and astronomers. ", "conclusions": "" }, "0607/astro-ph0607635_arXiv.txt": { "abstract": "Cosmic dust extinction alters the flux of type Ia supernovae. Inhomogeneities in the dust distribution induce correlated fluctuations of the SN fluxes. We find that such correlation can be up to $60\\%$ of the signal caused by gravitational lensing magnification, with an opposite sign. Therefore if not corrected, cosmic dust extinction is the dominant source of systematic uncertainty for future SNe Ia lensing measurement limiting the overall S/N to be $\\la 10$. On the other hand, SN flux correlation measurements can be used in combination with other lensing data to infer the level of dust extinction. This will provide a viable method to eliminate gray dust contamination from the SN Ia Hubble diagram. ", "introduction": "Gravitational lensing causes several observable effects such as distortion of galaxy shape ({\\em cosmic shear}), variation of galaxy number density ({\\em cosmic magnification}) and mode-coupling in cosmic backgrounds. Over the upcoming years measurements of these effects will provide an accurate mapping of the matter distribution in the universe (for reviews see \\cite{Bartelmann01,Refregier03}). Recently, several other lensing reconstruction methods have been proposed. One possibility is to measure the spatial correlation of lensing induced supernova (SN) flux fluctuations. In fact due to lensing magnification\\footnote{{ Throughout this paper, the term lensing magnification refers to both the cases of magnification ($\\mu>1$) and de-magnification ($\\mu<1$). To be more specific, the spatial correlation functions and the corresponding power spectra ($C_{\\kappa}$ and $C_{\\kappa\\delta \\tau}$) investigated hereafter are averaged over the full distribution of $\\mu$.}}, the SN flux is altered such that $F\\rightarrow F\\mu\\simeq F(1+2\\kappa)$, where $F$ is the intrinsic SN flux, $\\mu$ is the lensing magnification and $\\kappa$ is the lensing convergence. Intrinsic fluctuations of the SN flux are random (analogous to intrinsic galaxy ellipticities in cosmic shear measurement). In contrast those induced by lensing magnification (see e.g. \\citet{Kantowski95,Frieman96,Holz98,Dalal03}) are correlated with the overall matter distribution (analogous to the shear signal). Therefore the lensing signature can be inferred either from spatial correlation measurements of SN fluxes \\citep{Cooray06} or from the root-mean-square of flux fluctuations of high redshift SNe for which the lensing signal is dominant \\citep{Dodelson06}. Gravitational lensing also induces scatter in the galaxy fundamental plane through magnification of the effective radius, $R_e\\rightarrow R_e\\mu^{1/2}\\simeq R_e(1+\\kappa)$. Since intrinsic scatters in the fundamental plane are random, spatial correlation measurements can be used to infer the lensing signal \\citep{Bertin06}. A similar analysis can be applied to the Tully-Fisher relation as well. Astrophysical effects may limit the accuracy of these methods. For instance extinction by cosmic gray dust can be an important source of systematic uncertainty. This is because dust absorption changes the apparent SN flux and may induce correlation of the flux fluctuations. It also induces scatters in the fundamental plane by dimming the galaxy surface brightness and affects the Tully-Fisher relation through dimming the galaxy flux. These effects potentially cause non-negligible systematics in the corresponding lensing measurements. Although the existence of gray dust in the intergalactic medium (IGM) remains untested, this scenario could account for the metal enrichment of the IGM (\\citet{Bianchi} and reference therein). Testing the gray dust hypothesis is also relevant for cosmological parameter inference from SN Ia luminosity distance measurements. Recently \\citet{Corasaniti06} has pointed out that gray dust models which pass current astrophysical constraints can induce a $\\sim\\,20\\%$ bias in the estimate of the dark energy equation of state $w$ using the Hubble diagram of future SN Ia experiments. {In this paper, we study the impact of cosmic gray dust on SN lensing measurements, under the optimistic assumption that contaminations of reddening dust can be perfectly corrected}. The effects on lensing reconstruction based on the fundamental plane and the Tully-Fisher relation can be estimated similarly. For supernova, the key point is that extinction caused by dust inhomogeneities along the line of sight causes flux fluctuations which are anti-correlated with the lensing magnification signal and thus wash-out its imprint. In particular we find that dust induced correlation can bias SN lensing measurements by $10-60\\%$. Therefore this effect is likely to be the dominant source of systematics for future SN surveys characterized by large sky coverage and sufficiently high surface number density. If not corrected, the dust induced correlation would limit the signal-to-noise to ${\\rm S/N}\\la 10$. This is low compared to the S/N achieved by current cosmic shear measurements (e.g. \\citet{Jarvis05,VanWaerbeke05,Hoekstra05}) and that of proposed methods such as CMB lensing \\citep{Seljak99,Zaldarriaga99,Hu01,Hu02}, 21cm background lensing \\citep{Cooray04,Pen04,Zahn05,Mandel05} and cosmic magnification of 21cm emitting galaxies \\citep{Zhang05,Zhang06}. Nevertheless we suggest that measurements of the SN flux correlation still carry valuable information. In fact in combination with other lensing data they will provide a viable method to detect and eliminate cosmic gray dust contamination from future SN Ia luminosity distance measurements. ", "conclusions": "Several new methods have been proposed for inferring the lensing magnification signal from a variety of correlation measurements. These involve SN Ia flux, the fundamental plane and Tully-Fisher relation of optical galaxies. In this paper we have shown that contamination of cosmic dust extinction may severely degrade such measurements. As an example inhomogeneities in the cosmic dust distribution may limit the S/N of SN lensing measurements to $\\la 10$ level. Billions of galaxies can be detected/resolved by the Square Kilometer Array\\footnote{SKA: http://www.skatelescope.org/} through the 21cm hyperfine transition line emission which is not affected by dust extinction. In such a case the only scatters other than intrinsic ones in the Tully-Fisher relation ($L\\propto v_c^4$) are induced by lensing magnification, $L\\rightarrow L(1+2\\kappa)$. Therefore lensing reconstruction using these galaxies is an attractive possibility, since it is free of some systematics associated with cosmic shear such as shape distortion induced by point spread function. On the other hand, spatial correlation measurements of SN fluxes or galaxy-quasar will constrain the amount of cosmic gray dust and its clustering properties to high accuracy. This will not only provide a better understanding of IGM dust physics, but also a valuable handing of dust contamination in the SN Ia Hubble diagram. {\\it Acknowledgments}---. We thank Brice Menard and Ryan Scranton for helpful discussions on dust contamination in SDSS samples. We are also thankful to Yipeng Jing, Peter Nugent and Alexandre Refregier for useful discussions. PJZ is supported by the One-Hundred-Talents Program of Chinese Academy of Science and the NSFC grants (No. 10543004, 10533030)." }, "0607/astro-ph0607545_arXiv.txt": { "abstract": "To understand magnetic diffusion, momentum transport, and mixing in the interior of the sun, we consider an idealized model of the tachocline, namely magnetohydrodynamics (MHD) turbulence on a $\\beta$ plane subject to a large scale shear (provided by the latitudinal differential rotation). This model enables us to self-consistently derive the influence of shear, Rossby and Alfv\\'en waves on the transport properties of turbulence. In the strong magnetic field regime, we find that the turbulent viscosity and diffusivity are reduced by magnetic fields only, similarly to the two-dimensional MHD case (without Rossby waves). In the weak magnetic field regime, we find a crossover scale ($L_R$) from a Alfv\\'en dominated regime (on small scales) to a Rossby dominated regime (on large scales). For parameter values typical of the tachocline, $L_R$ is larger that the solar radius so that Rossby waves are unlikely to play an important role in the transport of magnetic field and angular momentum. This is mainly due to the enhancement of magnetic back-reaction by shearing which efficiently generates small scales, thus strong currents. ", "introduction": "Data from global helioseismology \\citep{Thompson03} have shed some light on the internal rotation of the sun. Throughout the convective envelope, the rotation rate decreases monotonically toward the poles. Near the base of the convection zone, there is a sharp transition between differential rotation in the convective envelope and nearly uniform rotation in the radiative interior. This transition region has become known as the solar tachocline \\citep{Spiegel72}, and its thickness has been constrained to be only a few percent of the solar radius by helioseismic data \\citep{Charbonneau99}. The rotation rate of the radiative interior is intermediate between the equatorial and polar regions of the convection zone. Thus, the radial angular velocity gradient across the tachocline is positive at low latitudes and negative at high latitudes. The tachocline is most probably located below the convection zone \\citep{Charbonneau99}, but part of it might be included in the overshooting region where turbulent plumes coming from the convection zone can induce turbulent motions. This overshooting layer contains a prominent latitudinal differential rotation. It is thus of primary importance to understand the physics pertinent to such a turbulent layer in presence of strong shear. In particular, as the tachocline links two regions of very different transport properties, theoretical predictions of transport coefficients are essential to understand its long term dynamics (e.g. its thickness). For instance, different mechanisms have been invoked to stabilize this transition layer against the radiative broadening during the solar evolution. First, since the radial mixing is ineffective due to the stable stratification in the tachocline, \\citet{Spiegel92} suggested that the efficient turbulent transport of angular momentum (eddy-viscosity) in the horizontal direction could be responsible for the confinement of the tachocline. Other models rely on the existence of a magnetic field fully contained in the radiative zone \\citep{Rudiger96,Gough98,MacGregor99}. Specifically, \\citet{Rudiger96} have shown that this configuration could explain the dynamics of the tachocline if the effective viscosity is larger by at least 4 orders of magnitude than its molecular value. \\citet{Gough98} and \\citet{MacGregor99} have estimated the thickness of the tachocline by balancing the shearing of the poloidal field with the resistive dissipation of toroidal magnetic field. However, in this estimate, the turbulent diffusivity and viscosity might be more relevant than the molecular values. The tachocline is also important for the generation of magnetic field inside the sun (by means of a dynamo process), whose operation crucially depends on the values of effective (eddy) viscosity and magnetic diffusivity \\citep[e.g.][]{Parker93}. The purpose of this paper is to provide a consistent theory of momentum transport and magnetic field diffusion in the tachocline. In view of the tendency of two-dimensional (2D) turbulence in a stably stratified layer, we focus on the dynamics in the (local) horizontal plane orthogonal to density gradient, by incorporating the effects of rotation and toroidal magnetic fields. Specifically, the latitudinal differential rotation is represented by a large-scale shear flow while the latitudinal variation in the Coriolis force is captured by using a local $\\beta$ plane. This model permits us to study the influence on turbulent transport of shear flow, Rossby and Alfv\\'en waves. The shear flow efficiently generates small scales by shearing, thereby enhancing the back-reaction of fluctuating magnetic fields. We show that in the limit of strong magnetic field, the turbulent viscosity and magnetic diffusivity are reduced by magnetic fields only, with the same results as in the 2D magnetohydrodynamics (MHD) case \\citep{Kim01}. In the weak magnetic field regime, we find a crossover from a Alfv\\'en dominated regime (on small scales) to a Rossby dominated regime (on large scales). Using parameter values typical of the tachocline, we identify these two regimes and show that Rossby waves are unlikely to play an important role in the tachoclinic turbulent transport. This work complements our previous study of 3D hydrodynamical (HD) turbulence in the tachocline \\citep{Kim05,2Shears}, within which the effects of shear flows (due to radial and latitudinal differential rotations) and the average rotation were investigated. The remainder of the paper is organised as follows: we provide our model of the tachocline and the governing equations in section \\ref{Model}, the results for the hydrodynamical (HD) and magnetohydrodynamical (MHD) cases in section \\ref{Hydrodynamics} and \\ref{MHD} respectively. In section \\ref{Conclusion}, we discuss the implication of our findings for the tachocline. ", "conclusions": "\\label{Conclusion} In order to elucidate the turbulent transport in the tachocline, we considered a 2D model of turbulence on a $\\beta$ plane, in the presence of a large-scale latitudinal shear (with shearing rate $\\A$) and a uniform toroidal magnetic field (intensity $B_0$). By using the quasi-linear approximation, we have computed the turbulent viscosity and diffusivity (of magnetic fields or particle) in the two limiting cases of weak and strong magnetic field. The intensity of turbulence is not altered much by the presence of waves whereas the turbulent transport can be severely affected due to the effect of these waves on the phase of the field. In the case of a weak magnetic field, we have shown that there is a crossover scale $L_R = 2 \\pi (\\A B_0^3 / \\beta^3 \\eta)^{1/4}$ from an Alfv\\'en dominated (on small scales) to a Rossby dominated turbulence (on scale large enough). In the Rossby dominated regime, we found a negative turbulent viscosity, suppressed only by the shear: $\\nu_T \\propto - \\A^{-2}$. In the Alv\\'en dominated regime, the turbulent viscosity is positive with the following scaling: $\\nu_T \\propto B_0^2 (\\beta \\A)^{-2}$. We also have shown that the turbulent diffusivity is not very much affected by a weak magnetic field and that for all scales, the turbulent diffusivity is suppressed only by Rossby waves: $\\eta_T \\propto \\beta^{-2}$ (the same as in the HD case). In the case of a strong magnetic field, we found that the turbulent viscosity and diffusivity are positive and the same as in the 2D MHD case (except for an additional factor in the turbulent viscosity which is very small in the tachocline context) and consequently are suppressed only by magnetic fields: $\\eta_T \\propto \\nu_T \\propto B_0^{-2}$. These findings have interesting implications for the dynamics of the tachocline. Using typical solar values, $B_0 \\sim 1 \\; \\mbox{T}$, $\\A \\sim 6 \\times 10^{-8} \\; \\mbox{s}^{-1}$, $\\beta \\sim 2.7 \\times 10^{-15} \\; \\mbox{m}^{-1} \\mbox{s}^{-1}$ , $\\nu \\sim 10^{-2} \\mbox{m}^2 \\mbox{s}^{-1}$ and $\\eta \\sim 1 \\mbox{m}^2 \\mbox{s}^{-1}$, we can calculate the boundary between the weak and the strong magnetic regime, $L_B = 2 \\pi B_0 / \\A \\sim 10^8 \\; \\mbox{m}$ and the scale at which the transition from Alfv\\'en to Rossby turbulence occurs, $L_R = 2 \\pi (\\A B_0^3 / \\beta^3 \\eta)^{1/4} \\sim 10^{10} \\; \\mbox{m}$. Let us recall that this scale was calculated in the weak magnetic field case ($L > L_B$) which is consistent in the tachoclinic case as $L_R > L_B$. Figure \\ref{TachoScale} shows the relative positions of these two scales compared to the radius of the Sun. We also show on this figure the minimal scale $L_s = 2 \\pi (\\nu / \\A)^{1/2} \\sim 10^3 \\; \\mbox{m}$ for which our strong-shear approximation holds. \\begin{figure}[h] \\epsscale{.80} \\plotone{f1.eps} \\caption{\\label{TachoScale} Sketch of the different relevant scales evidenced in this paper for tachoclinic values of the parameters. $L_R$ is the scale which separates the strong magnetic field regime from the weak one and $L_B$ (obtained in the weak field approximation) is the scale which separates a Alfv\\'en dominated regime from a Rossby dominated one. We also included the smallest scale $L_S$ for which the large shear limit is valid.} \\end{figure} As relevant motion in the tachocline are probably on scales smaller than the solar radius, the dynamics of the tachocline is likely to be insensitive to the effect of Rossby waves. In the worst case, the dynamics would be that of the weak magnetic field with Alfv\\'en dominated turbulence (meaning a positive turbulent viscosity). These results are mainly due to the enhancement of the back-reaction of fluctuating magnetic fields by shearing. The reduction of the turbulent diffusion and transport induced by Alfv\\'en waves may be problematic in the solar context, invalidating some of the mechanisms believed to be important for the tachocline dynamics \\citep[see also][]{Diamond06}, e.g. the turbulent horizontal viscosity model of \\citet{Spiegel92}. However, the important point in this particular scenario is that the horizontal viscosity can be much larger than the vertical one. We have shown \\citep{2Shears} that it is the case for a turbulence sheared more strongly in the radial direction than in the latitudinal one (as is the case in the tachocline where the radial shear is at least one order of magnitude greater than the latitudinal one). If the effect of the magnetic field is the same in the two directions (reduction by a factor $B_0^2$), this result should not be altered by the inclusion of a magnetic field. However, a detail analysis of this situation requires the study of a three-dimensional model and is outside the scope of this paper. Finally, we note that in this paper, we assumed the toroidal magnetic field to be uniform on a local Cartesian ($\\beta$) plane. However, on a sphere, there is a possibility of instability in presence of a band of toroidal magnetic field and latitudinal differential rotation \\citep{Gilman97}. The turbulence arising from this instability can then be considered as a source of part of the forcing in our model." }, "0607/astro-ph0607290_arXiv.txt": { "abstract": "We discuss recent evidence that currently accepted mass-loss rates may need to be revised downwards, as a consequence of previously neglected ``clumping'' of the wind. New results on the radial stratification of the corresponding clumping factors are summarized. We investigate the influence of clumping on the ionization equilibrium of phosphorus, which is of major relevance when deriving constraints on the clumping factors from an analysis of the FUV P{\\sc v} resonance lines. ", "introduction": "While our understanding of the outflows from luminous OB-stars was thought to be well established, recent evidence indicates that currently accepted mass-loss rates {\\it may} need to be revised downwards {\\it by as much as a factor of ten}. This is a consequence of previously neglected ``clumping'' of the wind, which affects mostly those diagnostics which are sensitive to the {\\it square} of the density, $\\rho$ (such as recombination lines or free-free continua). Considering that numerous stellar-evolution calculations have demonstrated that changing the mass-loss rates of massive stars by even a factor of two has a dramatic effect on their evolution (e.g., \\citealt{Meynet94}), it is evident that such revisions would have enormous implications, not only regarding evolution, but also regarding the feed-back from massive stars. In this article, we will summarize the knowledge which has been accumulated lately and consider the question concerning the REAL mass-loss rates from massive star winds. ", "conclusions": "" }, "0607/hep-th0607034_arXiv.txt": { "abstract": "We have studied the wave dynamics and the Hawking radiation for the scalar field as well as the brane-localized gravitational field in the background of the braneworld black hole with tidal charge containing information of the extra dimension. Comparing with the four-dimensional black holes, we have observed the signature of the tidal charge which presents the signals of the extra dimension both in the wave dynamics and the Hawking radiation. ", "introduction": "In the past years there has been growing interest on studying models with extra dimensions in which the standard model fields are confined on a 3-brane playing the role of our 4-dimensional world, while gravity can propagate both on the brane and in the bulk \\cite{s1}\\cite{randall}\\cite{s3}. The extra dimensions need not be compact and in particular it was shown that it is possible to localize gravity on a 3-brane when there is one infinite extra dimension \\cite{randall}\\cite{s3}. One of the striking consequences of the theories with large extra dimensions is that the lowering of the fundamental gravity scale allows the production of mini black holes in the universe. Such mini black holes are centered on the brane and may have been created in the early universe due to density perturbations and phase transitions. Recently it was proposed that such mini black holes may also be produced in particle collision with the mass energy of TeV order or in the earth atmosphere due to the high energy cosmic ray showers \\cite{s4}\\cite{s5}\\cite{s6}. Once produced, these black holes will go through a number of stages in their lives \\cite{s5}\\cite{s6}\\cite{s7}, namely: i) the balding phase, where the black hole sheds the hair inherited from original ordinary object; ii) the spin-down phase, where the black hole loses its angular momentum; iii) the Schwarzschild quantum phase, where the black hole's mass will be decreased due to quantum process and finally iv) the Planck Phase. In the first phase, the black hole emits mainly gravitational radiation, while in the second and the third phases, the black hole will lose its energy mainly through the emission of the Hawking radiation. It was argued that during a certain time interval the evolution of the perturbation around black hole is dominated by damped single-frequency oscillation, called quasinormal modes (QNM), which carries unique fingerprint of the black hole and is expected to be detected through gravitational wave observations in the near future (see reviews on this topic and references therein \\cite{s8}). Recently the QNM of a brane-world black hole has been studied in [13]. The gravitational radiation of the mini black hole would be a characteristic sound and can tell us the existence of such black hole. Another possibility of observing signatures of this kind of mini black hole exists in particle accelerator experiments, where the spectrum of Hawking radiation emitted by these small black holes can be detected \\cite{s9}\\cite{s10}\\cite{s11}\\cite{s12}\\cite{park}. Since these small black holes carry information of extra dimensions and have different properties compared to ordinary 4-dimensional black holes, these two tools of detecting mini black holes can help to read the extra dimensions. The motivation of the present paper is to study signals of tiny black holes in the modern brane world scenarios[12]. Under the assumptions of the theory, most standard matter fields are brane-localized, therefore from the observational point of view it is much more interesting to study the brane-localized modes in the QNM and Hawking radiation. In the Hawking radiation study, it was found that the emission on the brane is dominated compared to that off the brane \\cite{s10}. We will investigate brane world black holes in the 4-dimensional background and study the QNM and the emission of brane-localized Hawking radiation. By comparing the properties of QNM and Hawking radiation to those in ordinary 4-dimensional black holes, we will argue that the dependence of these properties on the extra dimension of spacetime could be read if the spectrum of QNM and Hawking radiation are detected. In our following study, we will employ the exact black hole solution to the effective field equations on the brane, which is of the Reissner-Nordstrom type given as \\cite{Tdl} \\begin{equation} \\label{e0} ds_4^2 = - f dt^2+f^{-1}dr^2+r^2 d\\Omega ^2 \\end{equation} where $f=1-\\frac{2M}{M_{4}^{2}r}+\\frac{q}{\\tilde{M}_P^2 r^2}$, and $q$ is not the electric charge of the conventional Reissner-Nordstrom (RN) metric, but the 'tidal charge' arising from the projection onto the brane of the gravitational field in the bulk. Thus $q$ contains the information of the extra dimension. For $q>0$, this metric is a direct analogy to the Reissner-Nordstrom solution with two horizons and both horizons lie inside the Schwarzschild horizon $2M/M_4^2$. For $q<0$, the metric has only one horizon given by $ r_ + = \\frac{M}{{M_4^2 }}(1 + \\sqrt {1 - \\frac{{qM_4^4 }}{{M^2 \\tilde M^4_p }}} )$, which is larger than the Schwarzschild horizon. We will concentrate our attention on $q<0$, since it was argued that this negative tidal charge is the physically more natural case \\cite{Tdl}. Astrophysics limits the appearance of the Reissner-Nordstrom black hole with macroscopic electric charge, while there is no exact constraint on the emergence of the effective Reissner-Nordstrom black hole with tidal charge, especially negative tidal charge. However the tidal charge affects the geodesics and the gravitational potential, so that its value should receive at least indirect limit from observations, which needs careful study. On the other hand, tidal charge contains the information of extra dimensions. It is of great interest to investigate its influence on the gravitational wave and Hawking radiation observations and whether it can leave us the signature of the extra dimension. This paper is organized as following. In Sec.II, we will investigate the QNM of the braneworld black hole with tidal charge. We will study the scalar as well as the gravitational perturbations in the background of this black hole. In Sec.III, we will examine the absorption and emission problems for the scalar and the brane-localized graviton in the background of the braneworld black hole. Our conclusions are summarized in Sec.IV. For simplicity we will adopt the nature units in the following discussion. ", "conclusions": "In this paper, we have studied the perturbations around the brane-world black hole whose projected bulk effects are represented by the tidal charge. The positive tidal charged black hole is of the RN type and the negative one has only one horizon, which is larger than that of the 4D Schwarzschild case surrounding the central singularity. The effect of the negative tidal charge strengthens the gravitational field and lowers the temperature of the black hole \\cite{Tdl}. We have calculated the QNMs for both massless scalar field perturbation and gravitational perturbations. We have compared them to the case of the 4D Schwarzschild and RN black holes with the same mass. For the negative tidal charge case, unlike those observed in RN black hole, there is no critical charge where QNM differs below and above this critical charge \\cite{star}\\cite{wangb1}\\cite{wangb2}. When the negative tidal charge becomes more negative, the perturbations stay longer with less oscillations regardless of what types of perturbations. The QNMs of the negative tidal charge and positive tidal charge black holes are separated by that of the 4D Schwarzschild black hole. However, the influence of the tidal charge on the Hawking radiation are not as simple as that on the QNMs. In the scalar field, the negative tidal charge $Q$ suppresses the effective potential barrier and the absorption cross section increases correspondingly as a consequence of its effect strengthening the gravitational field. It is interesting that the emission spectrum does not simply decrease with the temperature as that in the background of RN black hole \\cite{s11}. There exist the critical values $Q_h$. When $Q > Q_h$, the peak of the emission spectrum is smaller than that of Schwarzschild black hole of the same mass. But if we take $Q < Q_h$, despite that the temperature is lower, the emission becomes stronger compared to the Schwarzschild black hole. This is the result of the enhancement of the absorption cross section and the decrease of the temperature according to Eq.(\\ref{e24}) caused by the negative tidal charge, which changes our usual picture about the Hawking radiation in 4D black holes. We emphasis that it is a new behavior which has not been observed before. As for the graviton, we find that when $ Q > \\frac{{9M^2 }}{{8n}}\\quad (l\\ne 0)$, the conservation of the graviton flux on the brane is broken and the leakage of the gravitons requires the exact solution of the brane world black hole in the bulk which has not been known yet. We have investigated only the Hawking radiation of the brane-localized gravitons under the condition $ Q < \\frac{{9M^2 }}{{8n}}\\quad (l\\ne 0)$. As pointed out before, there are two kinds of gravitational perturbations: $V_1$(or $V_3$) referring to those generated mainly by the fluctuation of the bulk effects and $V_2$(or $V_4$) to those created by the factors on the brane. All these potential barriers are suppressed by the negative tidal charge. As a result, the absorption cross sections increase. For the potential $V_2$(or $V_4$), the behaviors of the partial emission spectrums are much similar to those in the scalar field and there is a critical value $Q_h$, beyond which the peak of the emission spectrum is smaller and below this $Q_h$, the emission spectrum is stronger compared to that of the 4D Schwarzschild black hole, although with the negative tidal charge we always have lower black hole temperature. But for the potential $V_1$(or $V_3$), no matter what is the value of the tidal charge, the emission spectrum of the brane world black hole is always stronger than that of the Schwarzschild black hole for the brane-localized gravitons. The phenomenon of increasing emission spectrum with decreasing temperature in the brane world black hole has not been observed before. This is an effect of the negative tidal charge caused by the bulk effect. We conclude that both the QNMs and the Hawking radiation have given signatures of negative tidal charge due to the bulk effects in the brane-world black hole, which differs from that of the 4D Schwarzschild black hole. We expect that these signatures can be observed in the future experiments, which could help us learn the properties of the extra dimensions." }, "0607/hep-ph0607136_arXiv.txt": { "abstract": "We investigate Kolmogorov wave turbulence in QCD or, in other words, we calculate the spectrum of gluons as a function of time, $f_k(t)$, in the presence of a source which feeds in energy density in the infrared region at a constant rate. We find an early, an intermediate and a late time form for the gluon spectrum. Wave turbulence in QCD turns out to be somewhat different than the turbulence in the case of $\\phi^4$-type theories studied by Zakharov, L'vov and Falkovich. The hope is that a good understanding of QCD wave turbulence might lead to a better understanding of the instability problem in the early stages of the evolution after a heavy ion collision. \\vspace{1.cm} \\noindent {\\it Keywords}: Kolmogorov wave turbulence, Kolmogorov spectra, Instability, Thermalization, Heavy Ion Collisions ", "introduction": "One of the key questions in heavy ion physics is how rapidly and by what mechanism equilibrium is reached after a collision. Phenomenological analyses of experimental data at RHIC suggest a rapid thermalization~\\cite{Kolb:2001qz+X,Teaney:2003kp+X} but, so far, no convincing theoretical picture has emerged which naturally gives such a quick approach to equilibrium. The ``bottom-up'' picture of equilibration~\\cite{Baier:2000sb} is a detailed picture of how the initially produced hard gluons (at or near the saturation scale of the colliding ions) lose energy by radiating softer gluons. After they become sufficiently numerous these softer gluons equilibrate amongst themselves and continue to absorb energy from the initially produced hard gluons. After the hard gluons have lost all their energy the system has reached complete equilibration. Although difficult to evaluate precisely, the time for this equilibration was estimated to be on the order of 3 fm~\\cite{Baier:2002bt}. While it is an attractive picture the bottom-up scenario has a serious flaw in that the initially produced hard particles quickly acquire an instability as they spread out along the axis of the heavy ion collision as it was recently worked out by Arnold, Lenaghan and Moore~\\cite{Arnold:2003rq} and was long ago advocated by Mrowczynski~\\cite{Mrowczynski:1988dz+X} (For other early disscusions, see Refs.~\\cite{Weibel:1959,Buneman:1958,Heinz:1985vf+X}). This instability then becomes the dominant mechanism for the creation of softer gluons~\\cite{Arnold:2003rq}, more important than the Bethe-Heitler radiation used in the bottom-up picture~\\cite{Baier:2000sb,Wong:1996va}. The instability potentially speeds up equilibration, but so far it has been difficult to follow analytically how the QCD gluonic system evolves toward equilibration in the presence of the instability~\\cite{Mueller:2005un+X,Bodeker:2005nv}. Numerical studies~\\cite{Arnold:2005ef,Romatschke:2003ms+X,Rebhan:2004ur+X,Romatschke:2005pm+X,Dumitru:2005gp+X,Mrowczynski:2005ki} seem to indicate that the instability is effective at early times, but becomes less prominent in non-Abelian theories when the occupation number of low momentum modes becomes large. So far there is no analytic understanding of the numerical results. The phase space spectrum found in numerical studies by Arnold and Moore~\\cite{Arnold:2005ef} has a resemblance to that of the Kolmogorov spectrum in turbulence~\\cite{Arnold:2005qs}. Indeed, the problem of wave turbulence with an infrared source of energy discussed by Zakharov, L'vov and Falkovich (ZLF)~\\cite{ZLF} (for a nice overview, see~\\cite{Micha:2004bv}) would seem to have much in common with the instability problem currently under discussion for a non-expanding QCD medium and where the source for the instability is a collection of hard particles having a fixed asymmetry in their momentum distribution. Wave turbulence is a somewhat easier problem to deal with since one can take a spherically symmetric source of inflow of energy in low momentum modes. One can hope that a good understanding of QCD wave turbulence will lead to a better understanding of the early stages of evolution after a high-energy heavy ion collision. Wave turbulence in QCD appears to be somewhat different than the situation studied by Zakharov, L'vov and Falkovich~\\cite{ZLF}. The essence of the ZLF discussion is that waves, or particles, interact with each other locally in momentum. This is certainly the case, say, in a $\\phi^4$ theory where the scattering cross section for a high momentum particle on a low momentum particle is very small. Soft quanta in $\\phi^4$ theory interact very weakly with hard quanta because of the small overlap in their fields due to Lorentz contraction of the harder quanta. Soft gluons, on the other hand, interact with the currents of the harder gluons, and these currents do not decrease with the momentum of the harder gluons. In addition a gauge theory has a lower cutoff in frequency, the plasma frequencly, which has no analog in a theory of the $\\phi^4$ variety. Because of the strong interaction between soft and hard modes along with the plasma frequency infrared cutoff QCD turbulence seems to be quite different than the ZLF problem. In the problems considered by ZLF a source, say, of energy inserts energy at an infrared scale and a sink extracts that energy at a higher scale. The flow of energy then proceeds from a low momentum scale, through intermediate energy scales and exits at some high energy scale. A closely related process omits the high energy sink and has the energy flowing from low to high momentum through all intermediate scales. What is different in the QCD case is that energy can be absorbed by very high energy particles from a low energy source without the necessity of passing through intermediate scales. Particle number, however, cannot be directly transferred from low to high energy but must flow through intermediate scales, or be created at high energy through inelastic reactions. Thus in QCD there are important, even dominant, interactions which are not local in momentum. In this paper we calculate the analytic form of the spectrum of gluons in the presence of a source which feeds in energy density at a constant rate $\\dot{\\epsilon}_0=m_0^5/\\alpha .$ We suppose the energy comes into our system uniformly in space in the form of gluons which have a spherically symmetric momentum distribution and which are inserted uniformly in phase space just above the plasma frequency cutoff in a range $m < \\omega <\\bar{m}$ with $m$ the plasma frequency and $\\bar{m}$ on the order of $m.\\ m_0$ is the single dimensionful parameter in the discussion given in Sec.2 while in Sec. 3 we allow the incoming energy to be spread uniformly in phase space in a region $0 \\le k \\le k_0$ of momenta in which case $k_0$ is a separate dimensional parameter if we choose $k_0$ to be a scale larger than and independent of $m_0.$ We always suppose $\\alpha,$ the gluonic coupling, to be small. At late times after the source has become active, $m_0t>(1/\\alpha)^{9/5}$ in case the source energy is deposited in $m<\\omega <\\bar{m},$ the system of gluons is very close to thermal equilibrium. The incoming energy is transferred from the scale $m$ to the scale given by the temperature $T$ by direct absorption of soft gluons in an inelastic $3\\to 2$ process and, parametrically equally as important, by elastic scattering of soft gluons (scale $m$) on hard gluons (scale $T$) with the soft gluons losing energy to the hard gluons. Both $m$ and $T$ are slowly increasing with time. In the time domain $(1/\\alpha)^{7/5}2$ covering $\\sim 10$~Gpc$^3$ volume would be required for the LSS data to detect $|f_{\\rm NL}|\\simeq 100$. Minkowski Functionals are nicely complementary to the bispectrum because the Minkowski Functionals are defined in real space and the bispectrum is defined in Fourier space. This property makes the Minkowski Functionals a useful tool in the presence of real-world issues such as anisotropic noise, foreground and survey masks. Our formalism can be extended to scale-dependent $f_{\\rm NL}$ easily. ", "introduction": "\\label{sec:intro} Recent observations of cosmological fluctuations from the Cosmic Microwave Background (CMB) and Large Scale Structure (LSS) strongly support basic predictions of inflationary scenarios: primordial fluctuations are nearly scale-invariant \\citep{Spergel2003,Tegmark2004,Seljak2005,Spergel2006}, adiabatic \\citep{peiris2003,Bucher2004,bean/etal:2006}, and Gaussian \\citep[][and references therein]{Komatsu2002,Komatsu2003,Creminelli2005,Spergel2006}. In order to discriminate between more-than-100 candidate inflationary models, however, one needs to look for {\\it deviations} from the scale-invariance, adiabaticity as well as Gaussianity, for which different inflationary models make specific predictions. Inflationary models based upon a slowly rolling single-field scalar field generally predict very small deviations from Gaussianity; however, the post-inflationary evolution of non-linear metric perturbations inevitably generates ubiquitous non-Gaussian fluctuations. On the other hand, a broad class of inflationary models based upon different assumptions about the nature of scalar field(s) can generate significant primordial non-Gaussianity \\citep{lyth/etal:2003,dvali/etal:2004,arkani-hamed/etal:2004,alishahiha/etal:2004,BKMR2004}. Therefore, Gaussianity of the primordial fluctuations offers a direct test of inflationary models. It is customary to adopt the following simple form of primordial non-Gaussianity in Bardeen's curvature perturbations during the matter era \\citep[e.g., ][]{KS2001}: \\begin{equation} \\label{eq:ngpotential2} \\Phi=\\phi+f_{\\rm NL}(\\phi^2-\\langle\\phi^2\\rangle), \\end{equation} where $\\phi$ is an auxiliary random-Gaussian field and $f_{\\rm NL}$ characterizes the amplitude of a quadratic correction to the curvature perturbations. Note that $\\Phi$ is related to the primordial comoving curvature perturbations generated during inflation, ${\\cal R}$, by $\\Phi=(3/5){\\cal R}$. While this quadratic form is motivated by inflationary models based upon a single slowly-rolling scalar field, the actual predictions usually include momentum dependence in $f_{\\rm NL}$. (That is to say, $f_{\\rm NL}$ is not a constant.) Therefore, when precision is required, one should use the actual formula given by specific processes, either from primordial non-Gaussianity from inflation or the post-inflationary evolution of non-linear perturbations, to calculate a more accurate form of statistical quantities such as the angular bispectrum of CMB \\citep{Babich2004,Liguori2006}. Nevertheless, a constant $f_{\\rm NL}$ is still a useful parameterization of non-Gaussianity which enables us to obtain simple analytical formulae for the statistical quantities to compare with observations. The use of a constant $f_{\\rm NL}$ is also justified by the fact that the current observations are not sensitive enough to detect momentum-dependence of $f_{\\rm NL}$. Therefore, we adopt the constant $f_{\\rm NL}$ for our analysis throughout this paper. Note that it is actually straightforward to extend our formalism to a momentum-dependent $f_{\\rm NL}$. So far, analytical formulae for the statistical quantities of the CMB from primordial non-Gaussianity are known only for the angular bispectrum \\citep{KS2001,BZ2004,Liguori2006} and trispectrum \\citep{OH2002,KK2006}. The analytical formulae are extremely valuable especially when one tries to measure non-Gaussian signals from the data. Fast, nearly optimal estimators for $f_{\\rm NL}$ have been derived on the basis of these analytical formulae \\citep{KSW2005,Creminelli2005}, and have been successfully applied to the CMB data from the Wilkinson Microwave Anisotropy Probe (WMAP): the current constraint on $f_{\\rm NL}$ from the angular bispectrum is $-54$ to $114$ at the 95\\% confidence level \\citep{Komatsu2003,Spergel2006}. \\citep[See][for an alternative parameterization of $f_{\\rm NL}$.]{Creminelli2005} As for the LSS, the analytical formula is known only for the 3-d bispectrum \\citep{Verde2000,Scocci2004}. The LSS bispectrum contains not only the primordial non-Gaussianity, but also the late-time non-Gaussianity from gravitational instability and galaxy biasing, which potentially obscure the primordial signatures. In this paper, we derive analytical formulae for another statistical tool, namely the Minkowski Functionals (MFs), which describe morphological properties of fluctuating fields \\citep{MBW1994,SB1997,SG1998,WK98}. In $d$-dimensional space ($d=2$ for CMB and $d=3$ for LSS), $d+1$ MFs are defined, as listed in Table~\\ref{tab:MFs}. The ``Euler characteristic'' measures topology of the fields, and is essentially given by the number of hot spots minus the number of cold spots when $d=2$. This quantity is sometimes called the ``genus statistics'', which was independently re-discovered by \\citet{GMD1986} in search of a topological measure of non-Gaussianity in the cosmic density fields. (The Euler characteristic and genus are different only by a numerical coefficient, $-1/2$.) \\begin{table*}[t] \\caption{Minkowski Functionals defined in $d-$dimensional space: $d=2$ for CMB and $d=3$ for LSS.} \\begin{center} \\begin{tabular}{cccccc} \\hline\\hline observations & $d$ & $V_0$ & $V_1$ & $V_2$ & $V_3$ \\\\ \\hline CMB& 2 & Area & Total Circumference & Euler Characteristic & -- \\\\ LSS& 3 & Volume & Surface Area & Total Mean Curvature & Euler Characteristic \\\\ \\hline \\end{tabular} \\end{center} \\label{tab:MFs} \\end{table*} Why study MFs? Since different statistical methods are sensitive to different aspects of non-Gaussianity, one should study as many statistical methods as possible. Most importantly, the MFs and bispectrum are very different in that MFs are defined in real space, whereas the bispectrum is defined in Fourier (or harmonic) space. Therefore, these statistical methods are nicely complementary to each other. Previously there are several attempts to give constraints on the primordial non-Gaussianity using MFs \\citep[e.g,][]{NSM2000}. Although we shall show in this paper that the MFs do not contain information more than the bispectrum in the limit that non-Gaussianity is weak, the complementarity is still powerful in the presence of complicated real-world issues such as inhomogeneous noise, survey mask, foreground contamination, etc. The MFs have also been used to constrain $f_{\\rm NL}$. \\citet{Komatsu2003} and \\citet{Spergel2006} have used numerical simulations of non-Gaussian CMB sky maps to calculate the predicted form of MFs as a function of $f_{\\rm NL}$, and compared the predictions with the WMAP data to constrain $f_{\\rm NL}$, obtaining similar constraints to the bispectrum ones. This method (calculating the form of MFs from non-Gaussian simulations) is, however, a painstaking process: it takes about three hours to simulate one non-Gaussian map on one processor of SGI Origin 300. When cosmological parameters are varied, one needs to re-simulate a whole batch of non-Gaussian maps from the beginning --- this is a highly inefficient approach. Once we have the {\\it analytical} formula for the MFs as a function of $f_{\\rm NL}$, however, we no longer need to simulate non-Gaussian maps, greatly speeding up the measurement of $f_{\\rm NL}$ from the data. We use the perturbative formula for MFs originally derived by \\citet{Matsubara1994,Matsubara2003}: assuming that non-Gaussianity is weak, which has been justified by the current constraints on $f_{\\rm NL}$, we consider the lowest-order corrections to the MFs using the multi-dimensional Edgeworth expansion around a Gaussian distribution function. The organization of paper is as follows; In \\S~\\ref{sec:pb_general} we review the generic perturbative formula for the Minkowski Functionals. In \\S~\\ref{sec:mf_cmb} we derive the analytical formula for MFs of the CMB from primordial non-Gaussian fluctuations parameterized by $f_{\\rm NL}$. We also estimate projected statistical errors on $f_{\\rm NL}$ expected from the WMAP data from multi-year observations as well as from the Planck data. In \\S~\\ref{sec:mf_lss} we derive the analytical formula for MFs of the LSS from primordial non-Gaussianity, non-linear gravitational evolution, and galaxy biasing in a perturbative manner. \\S~\\ref{sec:summary} is devoted to summary and conclusions. In Appendix~\\ref{app:sim} we outline our method for computing the MFs from the CMB and LSS data. We also describe our simulations of CMB and LSS. In Appendix~\\ref{app:Bg} we derive the analytical formula for the galaxy bispectrum. In Appendix~\\ref{app:cmbsim} we compare the analytical MFs of CMB with non-Gaussian simulations in the Sachs--Wolfe limit. In Appendix~\\ref{app:poten}, we extend the corrections of primordial potential to $n$-th order, in order to examine more carefully validity of our perturbative expansion. Throughout this paper, we assume $\\Lambda$CDM cosmology with the cosmological parameters at the maximum likelihood peak from the WMAP first-year data only fit \\citep{Spergel2003}. Specifically, we adopt $\\Omega_b=0.049$, $\\Omega_{cdm}=0.271$, $\\Omega_\\Lambda=0.68$, $H_0=68.2~{\\rm km~s^{-1}~Mpc^{-1}}$, $\\tau=0.0987$, and $n_s=0.967$. The amplitude of primordial fluctuations has been normalized by the first acoustic peak of the temperature power spectrum, $l(l+1)C_l/(2\\pi)=(74.7~{\\mu}{\\rm K})^2$ at $l=220$ \\citep{WMAPpeak}. ", "conclusions": "\\label{sec:summary} We have derived analytical formulae of the MFs for CMB and LSS using a perturbation approach. The analytical formula is useful for studying the behavior of MFs and estimating the observational constraints on $f_{\\rm NL}$ without relying on non-Gaussian numerical simulations. The perturbation approach works when the skewness parameters multiplied by variance, $S^{a}\\sigma_0$, is much smaller than unity, i.e., $|f_{\\rm NL}|\\ll 3300$ for CMB and $|f_{\\rm NL}|\\ll 5000$ for LSS, both of which are satisfied by the current observational constraints from WMAP \\citep{Komatsu2003,Creminelli2005,Spergel2006}. We have shown that the perturbation predictions agree with non-Gaussian numerical realizations very well. We have used the Fisher matrix analysis to estimate the projected constraints on $f_{\\rm NL}$ expected from the observations of CMB and LSS. We have found that the projected 1-$\\sigma$ error on $f_{\\rm NL}$ from the WMAP should reach 50, which is consistent with the MF analysis given in \\cite{Komatsu2003,Spergel2006}, and is comparable to the current constraints from the bispectrum analysis given in \\cite{Komatsu2003,Creminelli2005,Spergel2006}. The MFs from the WMAP 8-year and Planck observations should be sensitive to $|f_{\\rm NL}|\\sim 40$ and 20, respectively, at the 68\\% confidence level. As the MFs are solely determined by the weighted sum of the bispectrum for $|f_{\\rm NL}|\\ll 3300$ for CMB and $|f_{\\rm NL}|\\ll 5000$ for LSS, the MFs do not contain information more than the bispectrum. However, this does not imply that the MFs are useless for measuring primordial non-Gaussianity by any means. The important distinction between the MFs and bispectrum is that the MFs are intrinsically defined in real space, while the bispectrum is defined in Fourier space. The systematics in the data are most easily dealt with in real space, and thus the MFs should be quite useful in this regard. Therefore, in the presence of real-world issues such as inhomogeneous noise, foreground, masks, etc., these two approaches should be used to check for consistency of the results. In this paper we have calculated the MFs from primordial non-Gaussianity with a scale-independent $f_{\\rm NL}$. It is easy to extend our calculations to a scale-dependent $f_{\\rm NL}$. All one needs to do is to calculate the form of the bispectrum with a scale-dependent $f_{\\rm NL}$ \\citep[e.g.,][]{Babich2004,Liguori2006}, and use it to obtain the skewness parameters, $S^{(a)}$ (Eqs.~[\\ref{eq:scmb0}--\\ref{eq:scmb2}] for CMB and Eqs.~[\\ref{eq:s0_lss}--\\ref{eq:s2_lss}] for LSS). The MFs are then given by equation~(\\ref{eq:MFs_perturb}) in terms of the skewness parameters. Also, we have not included non-Gaussianity from secondary anisotropy such as the Sunyaev-Zel'dovich effect, Rees-Sciama effect, patchy reionization, weak lensing effect, extragalactic radio sources, etc. It is again straightforward to calculate the MFs from these sources using our formalism, as long as the form of the bispectrum is known \\citep[e.g.,][]{SG1999,GS1999,CH2000,KS2001,VS2002}. The constraints on primordial non-Gaussianity from the MFs of LSS in a galaxy survey covering 1~$h^{-3}~{\\rm Gpc}^3$ volume are about 5 times weaker than those from the MFs of CMB in the WMAP data. One would therefore need the survey volume as large as 25~$h^{-3}~{\\rm Gpc}^3$ to make the LSS constraints comparable to the WMAP constraints (using the MFs only). This could be done by a survey covering $\\sim 2000$~deg$^2$ at $3.5$ = 5.55 $\\times$ 10$^{21}$~cm$^{-2}$~mag$^{-1}$), and the value of the extinction per unit distance in the direction of the sources, which is of $E_{B-V}/d\\sim0.3$~mag~kpc$^{-1}$ \\citep{lucke78}. For this purpose we use the most accurate values of $N_{\\rm H}$, obtained with the {\\it XMM-Newton} data. For 1WGA~J1346.5$-$6255 we consider the power-law model fit with $N_{\\rm H}\\sim(2.2\\pm0.4)\\times10^{21}$~atoms~cm$^{-2}$ (3$\\sigma$ uncertainties) and obtain a distance of $d=1.3\\pm0.3$~kpc, fully consistent with the distance to HD~119682 and NGC~5281. (We note that the column density derived for the best fit power-law+MEKAL or power-law+Gaussian is close enough to that derived using the power-law and thus yields a similar distance estimate within error). For the diffuse emission of SNR G309.2$-$00.6 we use the {\\it vnei} model fit with $N_{\\rm H}=6.5^{+4.5}_{-2.5}\\times10^{21}$~atoms~cm$^{-2}$ (3$\\sigma$ uncertainties) and obtain a distance of $d\\la3.9^{+2.4}_{-1.5}$~kpc, consistent with the $4\\pm2$~kpc value obtained by \\cite{rakowski01} from their {\\it ASCA} fit to the SNR. This distance is clearly incompatible with the distance to HD~119682/1WGA~J1346.5$-$6255, but consistent within errors with the distance of $5.4\\pm1.6$~kpc to the RCW~80 \\ion{H}{2} region proposed by \\cite{gaensler98} to be associated with SNR G309.2$-$00.6. \\subsection{Circumstellar envelope} \\label{envelope} The intrinsic $(J-K_s)_0$ infrared color of a normal B0,1\\,V star is $-0.22$ \\citep{ruelas91}. The 2MASS photometry provides a value of $J-K_s=0.40\\pm0.04$, leading to a total infrared color excess of $E_{J-K_s}=0.62\\pm0.04$. On the other hand, from the visual absorption of interstellar origin found above, $A_{V,~\\rm IS}=1.26\\pm0.23$~mag, and using the relationships by \\cite{rieke85} we found interstellar infrared absorptions of $A_{J,~\\rm IS}=0.36\\pm0.06$ and $A_{K_s,~\\rm IS}=0.14\\pm0.03$, and finally $E_{J-K_s,~\\rm IS}=0.22\\pm0.07$. Therefore, in addition to the interstellar color excess, there is an extra contribution to $E_{J-K_s}$ of $0.40\\pm0.08$. This is typical of Be stars with strong emission lines, usually adscribed to their circumstellar envelope. \\subsection{Age of HD~119682} \\label{age} \\cite{levenhagen04} report an age of $4\\pm1$~Myr and a mass of $M=18\\pm1~M_\\odot$ for HD~119682 based on the evolutionary tracks by \\cite{schaller92}. In fact, according to the new evolutionary tracks by \\cite{meynet03} (see their Fig.~5), the luminosity and effective temperature reported by \\cite{levenhagen04} are in agreement with a 18~$M_\\odot$ star at the beginning of the main sequence or with a slightly evolved 20~$M_\\odot$ star. However, since new calibrations such as the one of \\citealt{martins05} favor lower masses, we will consider as a strict upper limits for the age of HD~119682 the main sequence lifetime of a 15~$M_\\odot$ star. On the other hand, since HD~119682 has a high rotational velocity of $v \\sin i = 220\\pm20$~km~s$^{-1}$, we will use the evolutionary tracks by \\cite{meynet03} for an initial rotational velocity of 300~km~s$^{-1}$, leading to a total main sequence lifetime of 14.5~Myr (see their Table~1). Therefore, even if assuming the lowest possible mass and a high rotational velocity, the lifetime in the main sequence is less than 15~Myr, i.e., 3 times lower than the estimated age of NGC~5281 by \\cite{sanner01}. This clearly confirms the blue straggler nature of HD~119682, which is not clear in the $V$ vs. $(B-V)$ CMD of \\cite{sanner01}, slightly affected by the reddening of the envelope, but clearly evident when looking at the $V$ vs. $(U-B)$ CMD of \\cite{moffat73}. \\subsection{On the nature of HD~119682/1WGA~J1346.5$-$6255} \\label{nature} The unabsorbed X-ray luminosity of 1WGA~J1346.5$-$6255 from our {\\it Chandra} observations is $L_{(0.5-7.5~{\\rm keV})}=2.5^{+0.6}_{-0.7}\\,(d/1.3~{\\rm kpc})^2\\times10^{32}$~ergs~s$^{-1}$ (using the range of fluxes tabulated in Table~\\ref{table:chandra}). From the {\\it XMM-Newton} observations and using the range of fluxes obtained for the various models (Table~\\ref{table:xmm}), we obtain a luminosity of $L_{(0.5-8.5~{\\rm keV})}=4.5^{+0.1}_{-0.7}\\,(d/1.3~{\\rm kpc})^2\\times10^{32}$~ergs~s$^{-1}$, which corresponds to $L_{(0.5-7.5~{\\rm keV})}=4.1^{+0.2}_{-0.6}\\,(d/1.3~{\\rm kpc})^2\\times10^{32}$~ergs~s$^{-1}$, about a factor of 1.5--2 higher than the {\\it Chandra} value, regardless of the model used. In any case, this is clearly lower than that of Be/X-ray binaries containing neutron stars, even the faint persistent ones (which have $L_{\\rm X}\\sim10^{34}$~ergs~s$^{-1}$, see \\citealt{reig99}). Moreover, no X-ray pulsations have ever been detected, although \\cite{rakowski01} could only constrain the pulsed fraction to be less than $\\sim$85\\% in the frequency range 0.002--32~Hz (0.03--500~s) at a 99.99\\% confidence level. Therefore, the currently available information does not favor the hypothesis of a binary system with an accreting neutron star as the compact object. Another possibility would be to have a relatively quiescent accreting black hole in a binary system. Indeed, the photon index of 1WGA~J1346.5$-$6255 is similar to those found in accreting black holes while in the low/hard state (see \\citealt{fender04} and references therein). \\cite{gallo03} found a correlation between the 2--11~keV X-ray flux and the cm radio flux density for black holes in the low/hard state. >From our power-law model fit to the {\\it XMM-Newton} data, we obtain an unabsorbed flux of $F_{\\rm 2-11~keV}= 1.7\\times10^{-12}$~ergs~cm$^{-2}$~s$^{-1}$ and compute an expected radio flux density of $0.29\\pm0.20$~mJy at a distance of $1.3$~kpc. Given the uncertainties, such a source could be above or below the 3$\\sigma$ level in the ATCA image of \\cite{gaensler98}. Deeper radio observations, with a noise smaller than 0.02~mJy, and with long baselines to avoid contamination from the SNR, would be needed to properly search for a radio counterpart. Only simultaneous X-ray observations could be useful to exclude the black hole scenario, although this would also be uncertain due to the scatter of the radio/X-ray correlation. However, we note that there is no clear evidence that a Be+black hole binary system has been found up to now. In addition, the X-ray spectra are better fitted including one or several MEKAL models, which is not the case for black holes in the low/hard state. Therefore, the scenario of an accreting black hole orbiting the Be star is not favored by the available data. OB stars are known to be X-ray sources, presumably due to shocks arising in their stellar winds (see \\citealt{guedel04} for a review). Based on {\\it Einstein} data, \\cite{pallavicini81} found that the X-ray luminosities of $\\sim$30 OB stars followed approximately the empirical law $L_{\\rm X}\\simeq(1.4\\pm0.3)\\times10^{-7}L_{\\rm bol}$. A more complete study based on {\\it ROSAT} data for more than 200 OB stars allowed \\cite{berghoefer97} to propose the existence of a correlation with two different power laws above and below $L_{\\rm bol}=10^{38}$~ergs~s$^{-1}$. The $\\log L/L_\\odot=4.64\\pm0.10$ value given by \\cite{levenhagen04} provides a star bolometric luminosity of $L_{\\rm bol}=(1.7\\pm0.4)\\times10^{38}$~ergs~s$^{-1}$, leading, according to \\cite{berghoefer97}, to an X-ray luminosity of $L_{\\rm X}=(2^{+3}_{-1})\\times10^{31}$~ergs~s$^{-1}$. In contrast, the unabsorbed power-law X-ray luminosity from our {\\it Chandra} observations, $L_{(0.5-7.5~{\\rm keV})}=2.5^{+0.6}_{-0.7}\\,(d/1.3~{\\rm kpc})^2\\times10^{32}$~ergs~s$^{-1}$ is around one order of magnitude higher, while the {\\it XMM-Newton} luminosity is even a factor of about 2 higher. In fact, the X-ray emission of HD~119682 is not only stronger than expected, but obviously harder than that of isolated stars: our {\\it XMM-Newton} two-component MEKAL fit to the data requires a high temperature of $kT=13.0^{+2.6}_{-2.4}$~keV (14.3$^{+2.9}_{-2.6}\\times10^7$~K), more than an order of magnitude higher than usual in main sequence B stars, and a factor of at least four higher than the most extreme cases \\citep{cohen97}. As a matter of fact, the luminosity and spectral shape of HD~119682/1WGA~J1346.5$-$6255 are similar to those displayed by the notoriously peculiar Be star $\\gamma$~Cas: $L_{(2-10~{\\rm keV})}=6\\times10^{32}$~ergs~s$^{-1}$ and two-component MEKAL spectrum with temperatures of $0.05\\pm0.01$~keV and $12.3\\pm0.6$~keV \\citep{owens99}. Recently, \\cite{smith06} have proposed that the star HD~110432, which shows $L_{(2-10~{\\rm keV})}\\simeq5\\times10^{32}$~ergs~s$^{-1}$ and a MEKAL spectrum with $kT=10.6\\pm1.9$~kev \\citep{torrejon01}, is an almost perfect twin of $\\gamma$~Cas and have mentioned the possibility of the existence of a class of objects displaying similar characteristics, which they dub ``$\\gamma$~Cas analogs''. In addition, \\cite{motch05} have very recently presented a summary of X-ray and optical properties of their proposed currently existing five ``$\\gamma$~Cas-like objects'' (plus $\\gamma$~Cas itself). The properties of HD~119682/1WGA~J1346.5$-$6255, to be compared with their Table~1, are: spectral type B0.5\\,Ve, H$\\alpha$ EW of $-$24~\\AA, $L_{\\rm X~(0.2-12~keV)}=3.6$--$6.6\\times10^{32}$~ergs~s$^{-1}$ (from the {\\it Chandra} and {\\it XMM-Newton} power-law model fits, respectively), $kT_{\\rm soft}=1.7$~keV and $kT_{\\rm hard}=13$~keV, and $\\Gamma\\sim$1.7. These properties are practically identical to those of SS~397, with a slightly different equivalent width of H$\\alpha$. Furthermore, one of the objects in the list of \\cite{motch05} is a blue straggler in the 50~Myr old open cluster NGC~6649, closely matching the observed properties of HD~119682, also a blue straggler and located in the $\\sim$45~Myr old open cluster NGC~5281. In the previous spectroscopic study, we showed that the X-ray flux of the 1WGA~J1346.5$-$6255 source is a factor of $\\sim$2 brighter with the {\\it XMM} observation than with the {\\it Chandra} one, obtained 3 years later. This could be attributed to intrinsic variability of the source. The short-term variability detected on timescales of a few hundred seconds is similar to that seen in other $\\gamma$~Cas-like objects \\citep{motch05}. Moreover, the periodicity of $\\sim$1500~s we have detected in HD~119682/1WGA~J1346.5$-$6255, with a pulse fraction of 13--16\\%, is reminiscent to, although not so clear than, the oscillation found in the $\\gamma$~Cas analog HD~161103, with a period of $3250\\pm350$~s and a pulse fraction of 24\\% \\citep{lopes06}. However, in the case of HD~119682/1WGA~J1346.5$-$6255, this possible period has been detected in two different observations performed with two different satellites, indicating that this signal is probably stable along time (although other superimposed signals are detected in the {\\it XMM-Newton} data set). In summary, HD~119682/1WGA~J1346.5$-$6255 shares the following properties with the other $\\gamma$~Cas analogs: spectral fits including thermal components are better than simple power-law fits (see Tables~\\ref{table:chandra} and \\ref{table:xmm}); in the two-temperature MEKAL model fits, the cooler component is much fainter than the hard one (see Tables~\\ref{table:chandra} and \\ref{table:xmm}); the total photoelectric absorption in X-rays is not very different from that due to the interstellar medium; the X-ray luminosity is in the range $L_{\\rm X~(0.2-12~keV)}=10^{32}$--$10^{33}$~ergs~s$^{-1}$; there are no reported X-ray outbursts; the flux varies by a factor of a few from one observation to the other; variability on short timescales is detected, showing different quasi-periodic signals; the presence of an iron line near $\\sim$6.7~keV (see \\S4.1 and Fig.~\\ref{figure:XMM_powerlaw_fit}), difficult to interpret in the black hole scenario; a moderate N enhancement; and a circumstellar envelope that contributes to the NIR reddening. Therefore, we conclude that the overall properties of HD~119682 indicate that it is most likely a new $\\gamma$~Cas analog. The nature of $\\gamma$~Cas analogs is still a matter of debate (see \\citealt{motch05} and \\citealt{lopes06} for recent discussions). Possible scenarios to explain the properties of the detected X-ray emission are the magnetic interaction between the star and its circumstellar decretion disk, or accretion on to a compact object (neutron star, white dwarf, or even a black hole). We note that the possible $\\sim$1500~s period we have detected in HD~119682/1WGA~J1346.5$-$6255 could be the rotational period in the neutron star or white dwarf scenario (although a long observation with {\\it XMM-Newton} is needed to confirm it). In principle, optical spectroscopic observations of HD~119682 could unveil the presence of a compact object by means of a radial velocity curve, as has been done in the the case of $\\gamma$~Cas itself \\citep{harmanec00,miroshnichenko02}. However, Be stars in X-ray binaries typically have orbital periods in the range $\\sim$10--200~d, so a long-lasting observing program would be needed. In addition, the number and intensity of emission lines along the whole spectrum would render difficult to obtain accurate radial velocity measurements. Certainly, to obtain a radial velocity curve is not a straightforward task in the case of HD~119682. Whatever the nature of HD~119682/1WGA~J1346.5$-$6255, we stress that it is not related to the background SNR G309.2$-$00.6." }, "0607/astro-ph0607598_arXiv.txt": { "abstract": "We have used the instruments on the \\spitzer\\ {\\it Space Telescope} to study the Large Magellanic Cloud supernova remnants (SNRs) N11L, N44, N49, N206, N63A, and N157B. The two large SNRs N44 and N206 were not detected in any IRAC or MIPS wavebands; the remainder were detected at one or more wavelengths. In particular, the SNRs N49 and N63A each had features that were evident in all available IRAC and MIPS bands. Each of these two also displayed faint limb emission in the MIPS 24 \\micron\\ band only. IRS spectra obtained for the N49 SNR showed a number of prominent lines, with little continuum contribution. We therefore suggest that N49, and possibly N63A, are dominated by line emission, with thermal emission from hot dust being at most a secondary component. ", "introduction": "Supernova remnants (SNRs) are expected to produce infrared (IR) emission from such mechanisms as atomic and molecular line emission (dominated by ground-state fine structure lines), free-free emission from hot gas, synchrotron emission, and thermal continuum emission from dust heated by collisions with the post-shock gas. Of these, we expect only negligible contributions from free-free emission \\citep[e.g.,][]{FTV02} or from synchrotron emission \\citep{G+87a}, except where a pulsar-wind nebula (PWN) may be present. Thus, the IR emission is thought to be dominated by line emission and/or thermal emission from dust. The dust emission is of particular interest to the study of SNRs, as supernovae are thought to produce a significant fraction of dust in the ISM \\citep[e.g.,][]{D98}. Dust emission, in turn, is an important factor in cooling the hot plasma within SNRs through inelastic collisions with electrons and ions \\citep{OS73,SC01}. Infrared emission was detected with the {\\it Infrared Astronomy Satellite} (\\iras) from roughly a third of the known Galactic SNRs, demonstrating the difficulty in detecting many of these objects in the confusion with other IR sources in the ISM. Because of this confusion, most of the SNRs to be observed were comparatively bright, nearby objects \\citep{A89}. Similarly, \\iras\\ observations of the Large Magellanic Cloud (LMC), including pointed observations toward nine known remnants, only detected four SNRs \\citep{G+87a}. Using the more recent GLIMPSE survey from the {\\it Spitzer Space Telescope}, \\citet{R+06} clearly detected 18 of 95 Galactic SNRs known to be within the fields. \\spitzer\\ observations are sensitive to many forms of infrared emission from SNRs, depending on the instrument and wavelength band. The Infrared Spectrograph \\citep[IRS;][]{Ho+04} is ideal for the study of line emission. The Infrared Array Camera \\citep[IRAC;][]{F+04} is expected to include emission from polycyclic aromatic hydrocarbons (PAHs) and very small grains (VSGs); H$_2$ lines; and a number of atomic lines including Br$\\alpha$ 4.1 \\micron, [\\ion{Fe}{2}] 5.3 \\micron\\ and [\\ion{Ar}{2}] 7.0 \\micron. The Multiband Imaging Photometer for \\spitzer\\ \\citep[MIPS;][]{R+04} bands are also expected to include some atomic line emission, including [\\ion{Fe}{2}] 24.5 \\micron, [\\ion{O}{1}] 63.2 \\micron, and [\\ion{C}{2}] 157 \\micron; additionally, we expect significant thermal emission from collisionally heated dust. Dust emission from SNRs is of particular interest. It has been suggested that supernovae (SNe) are one of the principal sources of dust production \\citep[e.g.,][]{D98,T+01,D+03}. Theoretically, models have predicted as much as 0.2-4 M$_{\\sun}$ of dust production in a typical Type II supernova \\citep[e.g.,][]{T+01,D+03}. Certainly spectroscopic and photometric observations have suggested dust formation in at least five supernovae, including SN1987A \\citep{Da05}. Finding this freshly produced dust in SNRs, however, has proven problematic. Infrared fluxes attributed to warm ($\\sim$ 100-200 K) dust have been observed in young remnants such as Cas A, Tycho, and Kepler \\citep[e.g.,][]{B87,A+99,H+04}, but the inferred dust mass is generally several orders of magnitude less than that predicted. Furthermore, it is often unclear whether this dust was largely produced in the SN or has since been swept up from the surrounding medium. Dust at cooler temperatures ($\\le$ 25 K) has also been detected toward remnants such as Cas A \\citep{D+03} and Kepler \\citep{M+03} using sub-millimeter continuum observations. Substantially greater dust masses ($>$ 1 M$_{\\sun}$) have been deduced from these observations, but again, how much of this dust is ejecta from the recent SN is unknown. Worse, not all of the observed dust may be physically associated with the SNRs themselves. In fact, \\citet{WB05} suggest that ``at least one-half\" of the ``cold\" dust emission from the direction of Cas A is actually foreground emission. \\begin{deluxetable*}{lccrrclc} \\tablecaption{Spitzer Datasets used} \\tablehead{ \\colhead{SNR} & \\colhead{Inst.} & \\colhead{Mode} & \\colhead{Prgrm} & \\colhead{AOR} & \\colhead{PI} & \\colhead{Date} } \\startdata N11L & IRAC & 5$\\times$30s Map & 3565 & 11171840 & Chu & 2004 Nov. 30 \\\\ N11L & MIPS & Scan Map & 3565 & 11177728 & Chu & 2005 Mar. 7 \\\\ N44-SNR & IRAC & 5$\\times$30s Map & 3565 & 11172352 & Chu & 2005 Mar. 28 \\\\ N44-SNR & MIPS & Scan Map & 3565 & 11177984 & Chu & 2005 Apr. 7 \\\\ N49 & IRAC & 10$\\times$2s Map & 124 & 8152064 & Gehrz & 2004 May 5 \\\\ N49 & MIPS & 24\\micron\\ Phot & 124 & 8151808 & Gehrz & 2004 May 26 \\\\ N49 & MIPS & 70\\micron\\ Phot & 124 & 8791040 & Gehrz & 2005 May 26 \\\\ N49 & IRS & Stare & 124 & 6586112 & Gehrz & 2004 May 31 \\\\ N206-SNR & IRAC & 3$\\times$12s Map & 1061 & 6063104 & Gorjian & 2003 Nov. 21 \\\\ N206-SNR & MIPS & Scan Map & 717 & 7864320 & Rieke & 2003 Nov. 24 \\\\ N63A & IRAC & 5$\\times$30s Map & 3565 & 11173888 & Chu & 2004 Dec. 16 \\\\ N63A & MIPS & Scan Map & 3565 & 11178496 & Chu & 2005 Mar. 8\\\\ N157B & IRAC & 3$\\times$12s Map & 63 & 4379904 & Houck & 2004 Jan. 12\\\\ N157B & IRAC & 3$\\times$12s Map & 1032 & 6056960 & Brandl & 2003 Nov. 06 \\enddata \\label{tab:snrobs} \\end{deluxetable*} ", "conclusions": "The point that line emission is often a substantial contributor to the infrared emission from SNRs has been raised for some time \\citep[e.g.,][]{G+87b,A+92,O+97,O+99}. For instance, \\citet[][]{A+92} found that infrared line emission from shocked gas contributes ``a significant fraction\" to one component of the infrared emission - the component generally traced by the presence of optical emission. (Note that dust emission was also thought to contribute to this component, and that a second IR component, traced by an excess of IR to X-ray emission, was inferred to be generated only by collisionally heated dust.) More recently, \\citet{S+05} find that for the Small Magellanic Cloud SNR 1E 0102.2-7219, emission from the [\\ion{O}{4}] line may contribute up to 60\\% of the 24 \\micron\\ emission. Our observations of these remnants have primarily shown infrared emission, where present, arising from optically-bright areas of the SNRs. The implication is that line emission is a significant contributor to the infrared emission of these remnants. The IRAC-band ratios of \\citet{R+06} for the optical/IR-bright areas in N63 and N49 also indicate substantial contributions from ionic and molecular lines. This suggestion is greatly strengthened by the infrared spectroscopy of the optically-bright limb of N49, where the line contributions appear to overwhelmingly dominate the infrared flux density. Perhaps the most interesting (and mysterious) result is that for our two clearest detections, N49 and N63A, the outer limb of the SNR (in both cases, extremely faint in optical emission lines) is detected clearly at 24 \\micron\\ but not in the other wavebands. \\citet{S+05} find a similar result for 1E 0102.2-7219; the SNR is detected clearly at 24 \\micron, but not at 8 or 70 \\micron. Similarly, \\citet{Bo+06}, in a study of Type Ia SNRs in the Large Magellanic Cloud, find that none of their four SNRs are detected in the IRAC bands; two are detected at 24 and 70 \\micron, and two at 24 \\micron\\ alone. The comparative brightness of emission from these SNRs in the 24 \\micron\\ band, compared to the other IR bands, may be due in part to the line emission from [\\ion{O}{4}] and [\\ion{Fe}{2}], particularly in regions of bright optical emission. However, we are far from able to rule out a secondary dust component even in these regions. The areas of strong 24 \\micron\\ emission, weak optical emission, and moderate X-ray emission may point to a very different physical case, where IR emission from collisionally heated dust may yet predominate. Our future spectral mapping of N49 and N63A with \\spitzer's IRS are expected to elucidate this situation. The potentially significant contributions of infrared line emission from at least some SNRs has serious implications for the calculations of the dust content within those SNRs, and the use of such calculations to constrain models of dust production and destruction. For example, dust mass estimates based on the 24 \\micron\\ \\spitzer\\ flux or the 25 \\micron\\ \\iras\\ flux may overestimate the amount of dust due to overestimation of the flux resulting from inclusion of line emission. For some remnants, e.g. where nonradiative shocks dominate, this difference may be only slight, but for highly radiative SNRs, particularly those interacting with dense ambient material, the effect may be significant. Likewise, fits to SEDs may be unphysical if line contributions in one or more bands are not properly taken into account." }, "0607/astro-ph0607621_arXiv.txt": { "abstract": "In the last year many collaborations searching for the dark matter constituent have published results from their experiments. Here I give a review of direct detection searches reported by the DAMA, KIMS, CDMS-II, EDELWEISS-I, CRESST-II and ZEPLIN-I collaborations. I also outline the future plans of each collaboration as well as the XENON10 collaboration. ", "introduction": "The first evidence for the existence of dark matter came from the study of galaxy clusters by Zwicky\\cite{zwicky} in 1933. He observed that the motion of galaxies within the clusters was not consistent with the amount of matter. The idea that unseen dark matter is responsible for large, peculiar galaxy velocities as well as other phenomena is now generally accepted. Today most research in the field is focused on the nature of dark matter. The most stringent constraints on model parameters have been achieved by combining recent WMAP results with other experiments\\cite{WMAP}. These constraints indicate that about a quarter of the energy density of the universe consists of nonbaryonic dark matter. In order to accommodate the relic densities of dark matter needed to explain the observed universe, the dark matter constituents must have an interaction cross-section consistent with the weak interaction. These weakly interacting massive particles (WIMPs) are one class of dark matter particles that could be accommodated in big bang cosmology\\cite{jungman}. Many models of new physics outside the Standard Model offer a candidate for the dark matter constitutent. In supersymmetry this candidate is the neutralino ($\\chi$)\\cite{neutralino}, a linear combination of binos (\\~{B}), winos (\\~{W}$_{3}$) and higgsinos (\\~{H}$_{1}^{0}$ and \\~{H}$_{2}^{0}$). Neutralinos scatter elastically with atomic nuclei at a rate given by \\begin{equation} Rate \\sim Nn_{\\chi}<\\sigma_{\\chi}>, \\label{eqrate} \\end{equation} where N is the number of targets in the detector, $n_{\\chi}$ is the local neutralino density, and $<\\sigma_{\\chi}>$ is the WIMP-nucleon scattering cross-section. There are two different techniques for detecting dark matter. One technique, known as {\\em indirect detection}, is to look for products of neutralino annihilations in the sun or earth. Experiments such as AMANDA\\cite{amanda}, Super-K\\cite{superk} and EGRET\\cite{egret} rely on this technique. The other technique is called {\\em direct detection}. Experimenters using this technique look for evidence of neutralinos interacting with nuclei in their detector's target medium. This paper will review experiments using the latter technique. ", "conclusions": "The field of direct detection dark matter physics is in a very active and exciting time. Final results from many experiments world-wide including CDMS-II, Edelweiss-I, CRESST-II and Zeplin-I have been released in the last year and are shown in figure \\ref{limitsFig}. DAMA claims to have observed a dark matter signal at 6.3 $\\sigma$. This claim has not yet been confirmed by another experiment. All of these experiments are in the processes of commissioning the next stage of their detectors. Several new detectors, including XENON10, are also coming online this year. Plans for a ton-scale dark matter experiment are also being made. Proposals are in progress by the CDMS, Zeplin and XENON collaborations. Although not mentioned in this paper, the LHC will be coming online next year. It too will search for dark matter and will be a great compliment to the ton-scale experiments. \\begin{figure}[t] \\vspace{9.0cm} \\special{psfile=figs/graph_22278_2.eps voffset=0 hoffset=20 hscale=45 vscale=45 angle=0} \\caption{\\it Current experimental results for the case of spin-independent WIMP-nucleon interactions. DAMA is the only experiment to claim a signal (solid region). The best limit is given by the CDMS-II experiment (solid black). Limits from Zeplin I (dashed black), Edelweiss (dotted black), CRESST (solid grey), CDMS-II Si (dashed grey), and KIMS (dotted grey) are also shown. Theoretical SUSY models taken from A. Bottino {\\it et al} (light pink) and J. Ellis {\\it et al} (light green). \\label{limitsFig} } \\end{figure}" }, "0607/astro-ph0607417_arXiv.txt": { "abstract": "The kinematical properties of the Galactic Thick Disk are studied using absolute proper motions from the SPM3 Catalog and 2MASS near-infrared photometry for a sample of $\\sim$1200 red giants in the direction of the South Galactic Pole. The photometrically-selected sample is dominated by Thick Disk stars, as indicated by the number-density distribution that varies with distance from the Galactic plane as a single-valued exponential over the range $1